Datasets:

Hendrik-a

/

dancer-data

like 0

1005.3619

c

we have derived fe abundances via an ew analysis of and lines in high - resolution and moderate - s / n spectra of 16 ms dwarfs in the pleiades open cluster . the [ /h ] abundances increase dramatically relative to [ /h ] at @xmath0 below 5400 k , with the difference reaching over 0.8 dex in the coolest stars . this behavior is akin to what is seen in m34 , the hyades , and the uma moving group . comparison of the @xmath2 abundance patterns in the pleiades and hyades , as well as the [ o / h]@xmath41 abundances in the pleiades , hyades , and uma moving group , suggests that the trends may relax with age , though metallicity may yet prove to be a factor . abundances of cool dwarfs in additional open clusters or other stellar associations , especially those older than the hyades , are needed to determine if either age or metallicity are related to these anomalous abundances . the [ /h ] abundances are also higher in pleiads below 5400 k , but they show no evidence of an increase with decreasing @xmath0 . the inability to attribute the high [ /h ] abundances of the cool stars to the over - excitation effects illustrates the difficulty of quantifying this phenomenon . with lines of exceptionally high excitation potential such as the triplet , the over - excitation effect is clearly seen ( e.g. , * ? ? ? * ; * ? ? ? * ) , but for lines with excitation potentials @xmath70 ev , the effect is more difficult to pinpoint . our linelist includes transitions ranging in excitation potential from 2.18 to 5.10 ev , but for the individual pleiades stars , no increase in the line - by - line abundances as a function of excitation potential , like that seen in abundances of cool hyades dwarfs ( 1.83 ev @xmath71 4.42 ev ; * ? ? ? * ) , is evident ( figure 4 ) . line - to - line sensitivities to the over - excitation / ionization effects have yet to be clearly delineated , and it needs to be determined if there is an excitation potential threshold above which the abundances derived from these lines become enhanced by these effects . similarly , it needs to be determined if there is an excitation potential threshold _ below _ which the opposite occurs , the abundances derived from the low - excitation lines are lower due to these effects . such behavior would be expected if the overabundances of high - excitation ( singly ionized ) lines are due to the overpopulation of high - excitation ( singly ionized ) electronic states at the expense of depopulating low - excitation states . whether or not the over - excitation effects impact the spectroscopic derivation of stellar parameters ( @xmath0 , @xmath27 , and @xmath28 ) , an approach not adopted here , also needs to be determined . future investigations of these effects will require high - quality high - resolution spectroscopy so that accurate line - by - line abundances can be derived , even from features of just a few m in strength . a strong correlation between the @xmath2 and [ o / h]@xmath41 abundances of the pleiades dwarfs is evident in figure 2 , suggesting that the over - excitation / ionization effects share a common cause or origin . chromospheric emission and photospheric spots have been shown to be promising culprits , but to this point , the data are inconclusive . whereas strong correlations between the pleiades @xmath2 and chromospheric emission indicators and their residuals exist , they do not exist between [ o / h]@xmath41 and chromospheric emission . these contradictory results complicate the interpretation of the observed over - excitation / ionization effects and will have to be addressed by future studies . also , comparing abundances and chromospheric emission indicators measured using different spectra may not provide an accurate test of a true correlation because of potential temporal changes in chromospheric activity levels . future investigations into the over - excitation / ionization effects in cool open cluster dwarfs should make every effort to derive chromospheric activity levels from the same spectra so that any possible relation between the two can be more definitively delineated . determining the influence of photospheric spots on abundance derivations is more arduous . multicomponent model atmospheres- simulating photospheres with different areal coverages of hot , cool , and quiescent spots- have been shown to able to reproduce the measured ews of the triplet in a sample of hyades stars , but , while such exercises are useful and demonstrate the plausibility of the photospheric spot hypothesis , the results are only suggestive . observationally , a simultaneous photometric and spectroscopic monitoring program could be used to identify any correlated changes in spot coverage and spectral line strengths . such observational constraints would be helpful to determine if spotted photospheres affect high - resolution abundance derivations . despite the challenges , the possible connection between spots , over - excitation / ionization effects , and pre - ms li depletion should provide sufficient motivation for future efforts . a final conclusion that can be drawn from this study is that those carrying out spectroscopic abundance analyses of open clusters should heed caution when their samples include cool dwarfs , particularly those with @xmath0 @xmath1 k. including the abundances of these stars may skew cluster mean abundances . a similar caution may be needed for those studying cool ms dwarfs in the disk field , as well . further investigations into the sensitivity of the over - excitation / ionization effects to excitation potential , first ionization potential , stellar age , and stellar metallicity are needed in order to identify the extent and ubiquity of these effects . gratefully acknowledges her support for summer research from the national science foundation through their research experience for undergraduates program at cerro tololo inter - american observatory ( grant ast-0647604 ) . a.l.p . also thanks the c.v . starr - middlebury school in latin america for additional support . s.c.s . is grateful for support provided by the noao leo goldberg fellowship ; noao is operated by the association of universities for research astronomy , inc . , under a cooperative agreement with the national science foundation . j.r.k . gratefully acknowledges support for this work by grants ast 00 - 86576 and ast 02 - 39518 to j.r.k . from the national science foundation and by a generous grant from the charles curry foundation to clemson university . allende prieto , c. , barklem , p. s. , lambert , d. l. , & cunha , k. 2004 , , 420 , 183 an , d. , terndrup , d. m. , pinsonneault , m. h. , paulson , d. b. , hanson , r. b. , & stauffer , j. r. 2007 , , 655 , 233 asplund , m. 2005 , , 43 , 481 boesgaard , a. m. 1989 , , 336 , 798 boesgaard , a. m. , budge , k. g. , & ramsay , m. e. 1988 , , 327 , 389 boesgaard , a. m. & friel , e. d. 1990 , , 351 , 467 castelli , f. & kurucz , r. l. 2004 , arxiv : astro - ph/0405087 cayrel , r. , cayrel de strobel , g. , & campbell , b. 1988 , the impact of very high s / n spectroscopy on stellar physics , 132 , 449 chen , y. q. , zhao , g. , izumiura , h. , zhao , j. k. , liu , y. j. , honda , s. , & ohkubo , m. 2008 , , 135 , 618 dorazi , v. & randich , s. 2009 , , 501 , 553 fabbian , d. , asplund , m. , barklem , p. s. , carlsson , m. , & kiselman , d. 2009 , , 500 , 1221 feltzing , s. & gustafsson , b. 1998 , , 129 , 237 fitzpatrick , m. j. & sneden , c. 1987 , , 19 , 1129 funayama , h. , itoh , y. , oasa , y. , toyota , e. , hashimoto , o. , & mukai , t. 2009 , , 61 , 930 gebran , m. & monier , r. 2008 , , 483 , 567 king , j. r. , krishnamurthi , a. , & pinsonneault , m. h. 2000b , , 119 , 859 king , j. r. & schuler , s. c. 2005 , , 117 , 911 king , j.r . , schuler , s.c . , hobbs , l.m . , & pinsonneault , m.h . 2010 , , 710 , 1610 king , j. r. , soderblom , d. r. , fischer , d. , & jones , b. f. 2000a , , 533 , 944 kroupa , p. 2002 , science , 295 , 82 kupka , f. , piskunov , n. , ryabchikova , t. a. , stempels , h. c. , & weiss , w. w. 1999 , , 138 , 119 morel , t. & micela , g. 2004 , , 423 , 677 noyes , r. w. , hartmann , l. w. , baliunas , s. l. , duncan , d. k. , & vaughan , a. h. 1984 , , 279 , 763 oinas , v. 1974 , , 27 , 405 pinsonneault , m. h. , terndrup , d. m. , hanson , r. b. , & stauffer , j. r. 2004 , , 600 , 946 piskunov , n. e. , kupka , f. , ryabchikova , t. a. , weiss , w. w. , & jeffery , c. s. 1995 , , 112 , 525 ramrez , i. , allende prieto , c. , & lambert , d. l. 2007 , , 465 , 271 ryabchikova , t. a. , piskunov , n. e. , stempels , h. c. , kupka , f. , & weiss , w.w . 1999 , phys . t83 , 162 schuler , s. c. , hatzes , a. p. , king , j. r. , krster , m. , & the , l .- s . 2006a , , 131 , 1057 schuler , s. c. , king , j. r. , fischer , d. a. , soderblom , d. r. , & jones , b. f. 2003 , , 125 , 2085 schuler , s. c. , king , j. r. , hobbs , l. m. , & pinsonneault , m. h. 2004 , , 602 , l117 schuler , s. c. , king , j. r. , terndrup , d. m. , pinsonneault , m. h. , murray , n. , & hobbs , l. m. 2006b , , 636 , 432 shen , z .- x . , jones , b. , lin , d. n. c. , liu , x .- w . , & li , s .- l . 2005 , , 635 , 608 shen , z .- x . , liu , x .- w . , zhang , h .- w . , jones , b. , & lin , d. n. c. 2007 , , 660 , 712 sneden , c. 1973 , , 184 , 839 soderblom , d. r. , jones , b. f. , balachandran , s. , stauffer , j. r. , duncan , d. k. , fedele , s. b. , & hudon , j. d. 1993 , , 106 , 1059 takeda , y. 2003 , , 402 , 343 yong , d. , lambert , d. l. , allende prieto , c. , & paulson , d. b. 2004 , , 603 , 697 @xmath72 & & 0.65 & & 5715 & & 4.55 & & 0.96 + @xmath73 & & 0.75 & & 5339 & & 4.61 & & 0.58 + @xmath74 & & 0.65 & & 5715 & & 4.55 & & 0.96 + @xmath75 & & 0.84 & & 5048 & & 4.64 & & 0.30 + @xmath76 & & 0.89 & & 4899 & & 4.66 & & 0.16 + @xmath77 & & 0.75 & & 5373 & & 4.60 & & 0.61 + @xmath78 & & 0.73 & & 5407 & & 4.60 & & 0.65 + @xmath79 & & 0.82 & & 5098 & & 4.64 & & 0.35 + @xmath80 & & 0.73 & & 5407 & & 4.60 & & 0.65 + @xmath81 & & 0.81 & & 5142 & & 4.63 & & 0.39 + @xmath82 & & 0.74 & & 5363 & & 4.61 & & 0.60 + @xmath83 & & 0.78 & & 5239 & & 4.62 & & 0.48 + @xmath84 & & 0.78 & & 5239 & & 4.62 & & 0.48 + @xmath85 & & 0.72 & & 5442 & & 4.59 & & 0.68 + @xmath86 & & 0.80 & & 5174 & & 4.63 & & 0.42 + @xmath87 & & 0.82 & & 5117 & & 4.64 & & 0.37 + @xmath88 & & 0.53 & & 6172 & & 4.48 & & 1.42 + @xmath89 & & & & 5777 & & 4.44 & & 1.38 + & & 5793.92 & & 4.22 & & -1.70 & & 33.9 & 7.52 & & 46.0 & 7.63 & & 37.0 & 7.59 & & 47.5 & 7.61 & & 45.4 & 7.56 & & 39.3 & 7.51 + & & 5856.10 & & 4.29 & & -1.64 & & 36.6 & 7.57 & & 42.1 & 7.56 & & 30.8 & 7.46 & & 47.8 & 7.62 & & 51.1 & 7.67 & & 43.0 & 7.59 + & & 5927.80 & & 4.65 & & -1.09 & & 45.7 & 7.52 & & 54.0 & 7.58 & & 41.5 & 7.46 & & 60.3 & 7.63 & & 63.9 & 7.67 & & 56.5 & 7.63 + & & 6089.57 & & 5.02 & & -0.94 & & 41.8 & 7.64 & & 47.1 & 7.63 & & 35.3 & 7.52 & & 49.5 & 7.64 & & 56.6 & 7.74 & & 43.0 & 7.57 + & & 6093.65 & & 4.61 & & -1.50 & & 32.0 & 7.63 & & 39.6 & 7.67 & & 30.4 & 7.60 & & 45.9 & 7.76 & & 42.5 & 7.69 & & 39.4 & 7.68 + & & 6096.67 & & 3.98 & & -1.93 & & 41.2 & 7.64 & & 52.8 & 7.75 & & 42.4 & 7.68 & & 52.9 & 7.70 & & 55.1 & 7.72 & & 49.2 & 7.69 + & & 6151.62 & & 2.18 & & -3.29 & & 52.5 & 7.45 & & 66.5 & 7.58 & & 53.6 & 7.50 & & 76.0 & 7.64 & & & & & 61.4 & 7.48 + & & 6165.36 & & 4.14 & & -1.47 & & 48.3 & 7.46 & & 58.3 & 7.56 & & 47.2 & 7.47 & & 58.1 & 7.49 & & 57.6 & 7.47 & & 53.0 & 7.46 + & & 6270.23 & & 2.86 & & -2.71 & & 56.2 & 7.59 & & 70.8 & 7.77 & & 57.6 & 7.67 & & 68.8 & 7.62 & & 75.3 & 7.69 & & & + & & 6627.56 & & 4.55 & & -1.68 & & 30.9 & 7.71 & & 30.6 & 7.59 & & 30.0 & 7.69 & & 37.0 & 7.69 & & 39.1 & 7.74 & & 36.6 & 7.73 + & & 6806.86 & & 2.73 & & -3.21 & & 36.4 & 7.57 & & 53.1 & 7.72 & & 37.8 & 7.60 & & & & & & & & 47.1 & 7.61 + & & 6839.84 & & 2.56 & & -3.45 & & 37.2 & 7.65 & & & & & 35.8 & 7.62 & & & & & 58.4 & 7.76 & & 38.5 & 7.48 + & & 6842.69 & & 4.64 & & -1.32 & & 41.8 & 7.64 & & & & & 43.5 & 7.68 & & & & & & & & 44.1 & 7.60 + & & 6857.25 & & 4.07 & & -2.15 & & 23.4 & 7.54 & & 35.0 & 7.66 & & 23.8 & 7.54 & & & & & 37.1 & 7.66 & & 28.5 & 7.52 + & & 6861.94 & & 2.42 & & -3.89 & & 21.5 & 7.61 & & & & & & & & & & & 42.5 & 7.69 & & 32.5 & 7.63 + & & 6862.50 & & 4.56 & & -1.57 & & 29.9 & 7.58 & & & & & & & & & & & 45.9 & 7.76 & & 40.2 & 7.69 + & & 7284.84 & & 4.14 & & -1.75 & & 38.8 & 7.51 & & & & & 46.0 & 7.66 & & & & & 57.9 & 7.64 & & 51.0 & 7.65 + & & 7461.53 & & 2.56 & & -3.58 & & 29.2 & 7.59 & & 37.4 & 7.53 & & & & & 48.5 & 7.67 & & 50.4 & 7.68 & & 42.0 & 7.64 + & & 7547.90 & & 5.10 & & -1.35 & & 22.6 & 7.67 & & 27.5 & 7.70 & & & & & 33.3 & 7.81 & & & & & 29.4 & 7.74 + & & 5264.81 & & 3.23 & & -3.13 & & 43.8 & 7.40 & & 42.1 & 7.82 & & 42.9 & 7.54 & & 49.9 & 8.31 & & & & & 44.0 & 7.84 + & & 5414.05 & & 3.22 & & -3.65 & & 30.7 & 7.61 & & & & & & & & & & & & & & & + & & 5425.25 & & 3.20 & & -3.39 & & 39.4 & 7.53 & & 39.7 & 7.99 & & 40.8 & 7.71 & & 41.3 & 8.32 & & 39.6 & 8.45 & & 40.5 & 7.97 + & & 6084.10 & & 3.20 & & -3.88 & & 22.0 & 7.59 & & & & & & & & & & & & & & & + & & 6247.56 & & 3.89 & & -2.44 & & 52.6 & 7.49 & & 47.4 & 7.92 & & 56.8 & 7.77 & & 52.5 & 8.37 & & 47.2 & 8.43 & & 50.5 & 7.96 + & & 6432.68 & & 2.89 & & -3.69 & & 42.0 & 7.55 & & 39.8 & 7.96 & & 41.8 & 7.70 & & 48.1 & 8.48 & & 41.5 & 8.50 & & 37.2 & 7.86 + & & 6456.39 & & 3.90 & & -2.19 & & 63.4 & 7.45 & & 55.7 & 7.86 & & 69.6 & 7.79 & & 67.1 & 8.41 & & 53.0 & 8.32 & & 63.0 & 7.98 + & & 5793.92 & & 4.22 & & -1.70 & & 40.4 & 7.54 & & 48.0 & 7.62 & & 40.5 & 7.54 & & 49.5 & 7.66 & & 48.2 & 7.69 & & 46.5 & 7.62 + & & 5856.10 & & 4.29 & & -1.64 & & 40.9 & 7.56 & & 46.7 & 7.61 & & 40.9 & 7.56 & & 49.4 & 7.66 & & 41.0 & 7.23 & & 44.3 & 7.58 + & & 5927.80 & & 4.65 & & -1.09 & & 53.0 & 7.57 & & 58.0 & 7.60 & & 51.2 & 7.54 & & 60.5 & 7.64 & & 53.2 & 7.56 & & 55.5 & 7.58 + & & 6089.57 & & 5.02 & & -0.94 & & 42.4 & 7.56 & & 49.5 & 7.64 & & 41.6 & 7.55 & & 47.0 & 7.60 & & 46.7 & 7.63 & & 54.4 & 7.73 + & & 6093.65 & & 4.61 & & -1.50 & & 38.5 & 7.67 & & 41.0 & 7.67 & & 39.0 & 7.68 & & 40.7 & 7.66 & & 40.1 & 7.69 & & & + & & 6096.67 & & 3.98 & & -1.93 & & 51.0 & 7.74 & & 54.0 & 7.72 & & 46.8 & 7.65 & & 55.2 & 7.76 & & 56.1 & 7.83 & & 51.3 & 7.70 + & & 6151.62 & & 2.18 & & -3.29 & & 63.9 & 7.55 & & 69.5 & 7.53 & & 61.7 & 7.50 & & 80.0 & 7.75 & & 70.9 & 7.68 & & 72.9 & 7.66 + & & 6165.36 & & 4.14 & & -1.47 & & 53.5 & 7.48 & & 62.3 & 7.57 & & 49.9 & 7.41 & & 56.6 & 7.48 & & 51.6 & 7.43 & & 57.0 & 7.51 + & & 6270.23 & & 2.86 & & -2.71 & & 63.5 & 7.64 & & 67.0 & 7.60 & & 64.2 & 7.66 & & 73.1 & 7.73 & & 71.2 & 7.54 & & 69.5 & 7.70 + & & 6627.56 & & 4.55 & & -1.68 & & 35.1 & 7.70 & & 41.1 & 7.78 & & 35.1 & 7.70 & & 38.6 & 7.73 & & 42.2 & 7.84 & & 39.6 & 7.76 + & & 6806.86 & & 2.73 & & -3.21 & & 46.3 & 7.60 & & 55.1 & 7.68 & & 44.6 & 7.57 & & 58.7 & 7.77 & & 49.8 & 7.66 & & 53.7 & 7.70 + & & 6839.84 & & 2.56 & & -3.45 & & 42.5 & 7.58 & & 52.0 & 7.67 & & 39.1 & 7.51 & & 46.9 & 7.57 & & 43.8 & 7.59 & & 49.6 & 7.66 + & & 6842.69 & & 4.64 & & -1.32 & & 51.0 & 7.73 & & 48.5 & 7.64 & & 48.6 & 7.69 & & & & & 49.9 & 7.70 & & 48.1 & 7.65 + & & 6857.25 & & 4.07 & & -2.15 & & 33.2 & 7.64 & & 33.0 & 7.57 & & 31.1 & 7.59 & & 35.3 & 7.63 & & 30.2 & 7.56 & & 37.6 & 7.69 + & & 6861.94 & & 2.42 & & -3.89 & & 26.6 & 7.51 & & 33.4 & 7.53 & & 26.3 & 7.50 & & & & & 32.0 & 7.61 & & & + & & 6862.50 & & 4.56 & & -1.57 & & 36.4 & 7.62 & & 42.5 & 7.70 & & 35.0 & 7.59 & & & & & 43.0 & 7.74 & & 40.3 & 7.66 + & & 7284.84 & & 4.14 & & -1.75 & & 50.8 & 7.65 & & 55.6 & 7.69 & & 48.8 & 7.61 & & & & & 48.8 & 7.60 & & & + & & 7461.53 & & 2.56 & & -3.58 & & 45.0 & 7.72 & & 48.4 & 7.68 & & 36.6 & 7.54 & & 42.2 & 7.55 & & 45.9 & 7.72 & & 47.6 & 7.71 + & & 7547.90 & & 5.10 & & -1.35 & & & & & & & & & & & 31.7 & 7.78 & & 27.0 & 7.69 & & 31.7 & 7.78 + & & 5264.81 & & 3.23 & & -3.13 & & 37.7 & 7.64 & & 36.6 & 7.91 & & 41.9 & 7.76 & & 32.7 & 7.75 & & 34.7 & 7.60 & & 37.7 & 7.79 + & & 5414.05 & & 3.22 & & -3.65 & & & & & & & & & & & & & & & & & 27.0 & 7.98 + & & 5425.25 & & 3.20 & & -3.39 & & 37.3 & 7.86 & & & & & 38.0 & 7.88 & & 32.1 & 7.96 & & 39.6 & 7.96 & & 37.3 & 8.01 + & & 6084.10 & & 3.20 & & -3.88 & & & & & & & & 21.3 & 7.88 & & & & & 14.5 & 7.66 & & & + & & 6247.56 & & 3.89 & & -2.44 & & 51.0 & 7.93 & & 38.4 & 7.96 & & 48.5 & 7.87 & & 40.7 & 7.97 & & 44.0 & 7.81 & & 43.8 & 7.93 + & & 6432.68 & & 2.89 & & -3.69 & & 36.8 & 7.82 & & 30.5 & 7.95 & & 38.6 & 7.86 & & 29.9 & 7.88 & & 35.2 & 7.82 & & & + & & 6456.39 & & 3.90 & & -2.19 & & 62.4 & 7.94 & & 51.2 & 8.03 & & 58.7 & 7.86 & & 50.6 & 7.97 & & 54.5 & 7.82 & & 57.8 & 8.02 + & & 5793.92 & & 4.22 & & -1.70 & & 47.8 & 7.65 & & 43.1 & 7.61 & & 48.0 & 7.64 & & 43.5 & 7.54 & & 26.3 & 7.58 + & & 5856.10 & & 4.29 & & -1.64 & & 46.8 & 7.63 & & 38.9 & 7.52 & & 46.6 & 7.62 & & 46.8 & 7.61 & & 24.8 & 7.55 + & & 5927.80 & & 4.65 & & -1.09 & & 57.3 & 7.61 & & 50.5 & 7.54 & & 58.3 & 7.61 & & 55.2 & 7.55 & & 34.2 & 7.52 + & & 6089.57 & & 5.02 & & -0.94 & & 51.4 & 7.69 & & 40.2 & 7.53 & & 49.5 & 7.65 & & 48.6 & 7.63 & & 30.4 & 7.60 + & & 6093.65 & & 4.61 & & -1.50 & & 47.2 & 7.80 & & 32.6 & 7.55 & & 41.2 & 7.68 & & 43.5 & 7.72 & & 24.1 & 7.66 + & & 6096.67 & & 3.98 & & -1.93 & & 49.8 & 7.67 & & 45.5 & 7.64 & & 56.1 & 7.78 & & 55.1 & 7.75 & & 29.5 & 7.65 + & & 6151.62 & & 2.18 & & -3.29 & & 72.5 & 7.66 & & 55.4 & 7.38 & & 70.5 & 7.59 & & 73.2 & 7.61 & & 36.8 & 7.51 + & & 6165.36 & & 4.14 & & -1.47 & & 57.3 & 7.51 & & 51.3 & 7.45 & & 56.5 & 7.48 & & 55.9 & 7.46 & & 37.3 & 7.50 + & & 6270.23 & & 2.86 & & -2.71 & & 71.3 & 7.74 & & 61.2 & 7.61 & & 72.9 & 7.74 & & 65.9 & 7.59 & & 44.8 & 7.70 + & & 6627.56 & & 4.55 & & -1.68 & & 34.8 & 7.66 & & 36.0 & 7.73 & & 39.6 & 7.75 & & 30.2 & 7.55 & & 21.5 & 7.70 + & & 6806.86 & & 2.73 & & -3.21 & & 56.8 & 7.76 & & 43.5 & 7.56 & & 56.3 & 7.73 & & 57.1 & 7.73 & & 23.7 & 7.62 + & & 6839.84 & & 2.56 & & -3.45 & & 51.6 & 7.71 & & 42.5 & 7.60 & & 47.9 & 7.60 & & 49.3 & 7.62 & & & + & & 6842.69 & & 4.64 & & -1.32 & & 48.5 & 7.65 & & 44.5 & 7.62 & & 49.5 & 7.67 & & 52.8 & 7.72 & & 32.5 & 7.67 + & & 6857.25 & & 4.07 & & -2.15 & & 38.2 & 7.71 & & & & & 33.3 & 7.59 & & 35.5 & 7.63 & & 17.2 & 7.60 + & & 6861.94 & & 2.42 & & -3.89 & & 41.1 & 7.76 & & & & & 36.9 & 7.64 & & 36.6 & 7.61 & & 11.0 & 7.60 + & & 6862.50 & & 4.56 & & -1.57 & & & & & & & & 44.5 & 7.74 & & 38.4 & 7.62 & & 25.3 & 7.68 + & & 7284.84 & & 4.14 & & -1.75 & & & & & & & & 57.5 & 7.73 & & & & & & + & & 7461.53 & & 2.56 & & -3.58 & & 42.4 & 7.59 & & 34.2 & 7.51 & & 44.8 & 7.62 & & 42.0 & 7.54 & & 18.9 & 7.67 + & & 7547.90 & & 5.10 & & -1.35 & & & & & 24.3 & 7.64 & & & & & 30.1 & 7.74 & & & + & & 5264.81 & & 3.23 & & -3.13 & & 34.0 & 7.69 & & 37.8 & 7.61 & & 36.9 & 7.84 & & 38.9 & 7.95 & & 51.4 & 7.48 + & & 5414.05 & & 3.22 & & -3.65 & & 23.5 & 7.86 & & & & & 21.6 & 7.86 & & & & & 30.3 & 7.53 + & & 5425.25 & & 3.20 & & -3.39 & & 33.8 & 7.91 & & 36.1 & 7.79 & & 30.8 & 7.88 & & 34.8 & 8.07 & & 48.1 & 7.63 + & & 6084.10 & & 3.20 & & -3.88 & & & & & & & & & & & 39.4 & 7.97 & & 26.5 & 7.64 + & & 6247.56 & & 3.89 & & -2.44 & & 41.0 & 7.86 & & 51.3 & 7.90 & & 38.0 & 7.86 & & 34.5 & 8.05 & & 66.6 & 7.64 + & & 6432.68 & & 2.89 & & -3.69 & & & & & 42.4 & 7.93 & & 31.5 & 7.90 & & 52.0 & 8.03 & & 48.9 & 7.61 + & & 6456.39 & & 3.90 & & -2.19 & & 51.2 & 7.87 & & 60.8 & 7.86 & & 50.8 & 7.93 & & & & & 74.1 & 7.53 + h ii 193 & & 0.06 & & 0.02 & & 0.43 & & 0.01 & & + 0.37 + h ii 250 & & 0.01 & & 0.02 & & 0.22 & & 0.05 & & + 0.21 + h ii 263 & & 0.07 & & 0.02 & & 0.89 & & 0.04 & & + 0.82 + h ii 298 & & 0.09 & & 0.01 & & 0.92 & & 0.02 & & + 0.83 + h ii 571 & & 0.02 & & 0.02 & & 0.44 & & 0.04 & & + 0.42 + h ii 746 & & 0.04 & & 0.02 & & 0.35 & & 0.06 & & + 0.31 + h ii 916 & & 0.06 & & 0.01 & & 0.49 & & 0.05 & & + 0.43 + h ii 1593 & & 0.00 & & 0.01 & & 0.35 & & 0.02 & & + 0.35 + h ii 2126 & & 0.08 & & 0.03 & & 0.42 & & 0.04 & & + 0.34 + h ii 2284 & & 0.05 & & 0.03 & & 0.28 & & 0.06 & & + 0.23 + h ii 2311 & & 0.08 & & 0.01 & & 0.45 & & 0.04 & & + 0.37 + h ii 2366 & & 0.09 & & 0.02 & & 0.34 & & 0.04 & & + 0.25 + h ii 2406 & & -0.02 & & 0.01 & & 0.33 & & 0.05 & & + 0.35 + h ii 2462 & & 0.08 & & 0.02 & & 0.37 & & 0.04 & & + 0.29 + h ii 2880 & & 0.04 & & 0.02 & & 0.53 & & 0.02 & & + 0.49 + h ii 3179 & & 0.03 & & 0.01 & & 0.06 & & 0.03 & & + 0.03 + h ii 263 @xmath90 & & @xmath91 & & @xmath92 & & @xmath93 + & & @xmath94 & & @xmath95 & & @xmath96 + & & @xmath97 & & @xmath98 & & @xmath99 + h ii 2284@xmath100 & & @xmath31 & & @xmath101 & & @xmath102 + & & @xmath103 & & @xmath104 & & @xmath105 + & & @xmath33 & & @xmath106 & & @xmath107 + h ii 3179@xmath108 & & @xmath31 & & @xmath109 & & @xmath106 + & & @xmath103 & & @xmath110 & & @xmath111 + & & @xmath33 & & @xmath98 & & @xmath112 + @xmath113 & & 0.06 & & @xcite + @xmath114 & & 0.02 & & @xcite + @xmath115 & & 0.03 & & @xcite + @xmath116 & & 0.02 & & @xcite + @xmath117 & & 0.03 & & @xcite + @xmath117 & & 0.01 & & @xcite + @xmath118 & & 0.05 & & @xcite + @xmath119 & & 0.02 & & this work + & & & & + @xmath119 & & 0.01 & & mean +

we have derived fe abundances of 16 solar - type pleiades dwarfs by means of an equivalent width analysis of and lines in high - resolution spectra obtained with the hobby - eberly telescope and high resolution spectrograph . curiously , a correlation between the pleiades abundances and chromospheric emission indicators does not exist .

we have derived fe abundances of 16 solar - type pleiades dwarfs by means of an equivalent width analysis of and lines in high - resolution spectra obtained with the hobby - eberly telescope and high resolution spectrograph . abundances derived from lines are larger than those derived from lines ( herein referred to as over - ionization ) for stars with k , and the discrepancy ( = [ /h ] - [ /h ] ) increases dramatically with decreasing , reaching over 0.8 dex for the coolest stars of our sample . the pleiades joins the open clusters m34 , the hyades , ic2602 , and ic2391 , and the ursa major moving group , demonstrating ostensible over - ionization trends . the pleiades abundances are correlated with infrared triplet and h chromospheric emission indicators and relative differences therein . oxygen abundances of our pleiades sample derived from the high - excitation triplet have been previously shown to increase with decreasing , and a comparison with the abundances suggests that the over - excitation ( larger abundances derived from high excitation lines relative to low excitation lines ) and over - ionization effects that have been observed in cool open cluster and disk field main sequence ( ms ) dwarfs share a common origin . curiously , a correlation between the pleiades abundances and chromospheric emission indicators does not exist . star - to - star abundances have low internal scatter ( dex ) , but the abundances of stars with k are systematically higher compared to the warmer stars . the cool star [ /h ] abundances can not be connected directly to over - excitation effects , but similarities with the and triplet trends suggest the abundances are dubious . using the [ /h ] abundances of five stars with k , we derive a mean pleiades cluster metallicity of [ fe / h ] .

astro-ph0502575

c

it is conventional to believe ( e.g. leahy 1993 ) that the jets of wat sources are physically similar to those in classical double ( frii ) radio galaxies and quasars ; evidence for this comes from their narrow and sometimes one - sided appearance ( as in 3c465 ) , the polarization structure they exhibit and their frequent termination in compact , shock - like features similar to the hotspots in friis ( hardcastle & sakelliou 2004 ) . x - ray emission from frii - like ( ` type ii ' ) jets , particularly in quasars and other sources where the jet is _ a priori _ thought to make a small angle to the line of sight , is often attributed in the literature to inverse - compton scattering of the cosmic microwave background radiation by a jet with a high bulk lorentz factor ( @xmath72 ) , based on the observations that a one - zone synchrotron model can not be fitted to the data and that the inverse - compton process without beaming requires a large departure from the equipartition condition ( e.g. tavecchio et al . 2000 ) . however , the generally accepted picture of x - ray emission from fri - like ( ` type i ' ) jets is that they are due to synchrotron emission related to the strong deceleration of the jet flow expected on other grounds in the inner few kpc of these systems ( e.g. hardcastle et al.2001 , 2002 ) ; in these cases a one - zone synchrotron model is often a good fit to the overall jet spectrum ( although the details of jet structure may in fact be complex on scales that can not be resolved in most sources : hardcastle et al . 2003 ) while a beamed inverse - compton model requires implausible jet parameters . the detection of 3c465 s inner jet in x - rays , at a level similar to what would be predicted using a simple synchrotron model based on observations of type i jets , lends support to the idea that some type ii jets can be synchrotron sources . there is already evidence for this from the comparatively strong detections of the pre - flaring jets , thought to be type ii , in some fri sources , notably 3c66b and cen a ( hardcastle et al . 2001 , 2003 ) as well as x - ray detections of components of the jets of some frii sources ( e.g. 3c219 , comastri et al . 2003 ; 3c452 and 3c321 , hardcastle et al . 2004 ; 3c403 , kraft et al . 2005 ) for which it is highly implausible that a beamed inverse - compton model can be an explanation . 3c465 is particularly interesting in this context for three reasons : ( 1 ) it is generally accepted that wat jets are type ii , ( 2 ) the emission from the jet is not confused by inverse - compton emission from any surrounding lobes , and ( 3 ) there is no confusion between jet knots and hotspots , which can also be synchrotron x - ray sources . 3c465 s detection thus gives us a reason to reiterate a point already made by hardcastle et al . ( 2004 ) and kraft et al . ( 2005 ) : there are clearly x - ray - emitting type ii jets for which the beamed inverse - compton model is _ not _ the explanation . since there is no compelling reason to expect a one - zone synchrotron model to be a good fit in an frii jet in which the loss spatial scale may be two to three orders of magnitude smaller than the region of the jet being considered , observers should be wary of ruling out the possibility that synchrotron emission is responsible for the x - rays from a given type ii jet . on the other hand , the non - detection of the bases of the plumes of 3c465 in x - rays , at a level sufficient to show that their radio to x - ray spectral index is significantly greater than that seen in all fri jets , strongly suggests that these regions are _ not _ analogous to the bases of type i jets . this is as expected in the picture of these objects proposed by hardcastle et al . ( 1998 ) and hardcastle & sakelliou ( 2004 ) , in which the jet decelerates rapidly at a single point close to the base of the plume , rather than there being uniform distributed deceleration over a finite - sized region of the base of the plume . we conclude that the non - detection of the plumes in x - rays in the _ chandra _ observations is evidence that the bases of wat plumes are physically different from the bases of fri type i jets ; wat plumes are probably more similar to the large - scale jets of fris , which typically do not show either strong deceleration or x - ray emission . it is not clear why we do not see x - ray emission associated with the assumed rapid deceleration of the jet at the base of the plume , by analogy with the x - ray emission detected from the hotspots of many low - luminosity frii radio galaxies , thought to be synchrotron in origin ( hardcastle et al . 2004 and references therein ) . features as faint as the brightest knot at the end of the n jet ( nj2 in the notation of hardcastle & sakelliou : see also fig . [ pbase ] ) have been detected in x - ray observations of frii hotspots . however , as we noted above , there is no particularly good candidate feature analogous to a hotspot in the radio images of 3c465 ; the jet flow clearly continues after nj2 . _ chandra _ images of wats that _ do _ exhibit compact bright jet termination structures would be of considerable interest . we begin by noting that the conclusions of jetha et al . ( 2005 ) , based on less sensitive observations , apply to 3c465 as well . there is no evidence in the data either for discrete features associated with the jet termination points or plume bases or for discontinuities in the properties of the external medium at the distances of the jet - plume transition , such as the shock structures in the models of loken et al.(1995 ) ; in particular , our observations show that the apparent temperature discontinuity seen in the inner regions of the sources studied by jetha et al . was probably an unresolved temperature gradient . models that require a discontinuity in the external medium are thus ruled out by observation . the jet termination is on a similar scale to the transition between the central , cool x - ray core and the hotter cluster medium , again as observed by jetha et al . , but as yet we have no clear idea of why this should be the case . as jetha et al.suggest , numerical modelling of realistic ( fast , low - density ) jets and plumes in realistic cluster environments may be the way forward here . it is also interesting to note that the transition occurs close to the point where the jet internal ( minimum ) pressure starts to exceed the external thermal pressure ( section [ cluster ] and fig.[pprofile ] ) . x - ray and radio data of similar quality for a larger number of sources would be required to say whether this is more than a coincidence . our detailed x - ray and radio observations have given a very clear indication that the large - scale plumes of 3c465 behave similarly to the lobes or plumes of normal twin - jet fri sources : the minimum pressures in the lobes are about an order of magnitude below the external thermal pressures . this implies some sort of departure from the minimum - energy condition , either a departure from equipartition between the radiating electrons and magnetic fields , a population of non - radiating relativistic particles , or a contribution to pressure from particles that do not participate in equipartition , such as thermal protons . as usual ( e.g. hardcastle et al . 1998 , croston et al.2003 ) we can not tell which of these is the case . in the inner parts of normal fri jets the jet deceleration is thought to be by entrainment of external material , which has suggested that the entrained material might supply the missing pressure if it could be heated efficiently so as to avoid thermal x - rays from the lobes or plumes ( see e.g. croston et al . 2003 ) . in wat plumes , the jets are not expected to have entrained significant amounts of external material , since they do not show the deceleration - related structures of fri jets , and so this explanation seems less likely ( see also below , section [ bending ] ) . on the other hand , if we assume that equipartition does not hold , it is not clear why such a large departure from equipartition should be required in wat plumes but not in frii lobes , where inverse - compton evidence ( e.g. hardcastle et al . 2002 , belsole et al . 2004 , croston et al . 2004 ) suggests that equipartition is roughly maintained and that the electrons and magnetic fields alone are sufficient to provide pressure balance with the external medium . the departure from equipartition would correspond to around a factor 5 difference between the actual magnetic field strength and the equipartition value ; this is considerably larger than is typically seen in the x - ray - detected lobes of frii radio galaxies ( croston et al . in addition , a departure from equipartition sufficient to make the internal pressure equal to the external pressure in the direction of electron domination would give rise to strong inverse - compton emission from the plumes , at the level of a few thousand _ xmm _ counts , which is not observed the deficits in x - ray count rate seen at the bases of the plumes seem to rule out this model , so that the lobes would have to be magnetically dominated if electrons and magnetic field alone were to supply the required pressure . the third possibility is that the plumes contain an energetically dominant population of relativistic protons : this can not be ruled out directly by the data , but again arguments based on inverse - compton detections suggest that it is not the case in the hotspots ( hardcastle et al . 2004 ) or lobes ( croston et al . 2005 ) of frii sources . we return to the particle content of the plumes in the next section . it is instructive to revisit the arguments of e84 that led to the conclusion that the bending of 3c465 in particular , and of wats in general , was difficult to explain in any model . they assumed that the kinetic energy of the plume must supply the overall observed radio luminosity : that is , the bulk flow is cold and there is _ in situ _ particle acceleration at all points in the plume . this leads to a condition on the density and speed of a non - relativistic plume : @xmath73 ( their eq . 2 ) where @xmath74 is a factor giving the efficiency of conversion between kinetic energy and energy of relativistic electrons , @xmath75 is the plume radius , @xmath76 is its density and @xmath77 is its speed . they assumed _ in situ _ particle acceleration on the basis of estimates of the synchrotron loss timescale at 5 ghz , and so we note to begin with that it may not be necessary to satisfy this condition in its original form . the synchrotron age of material at the end of the plumes is @xmath78 years , based on 1.4 ghz and 330 mhz data taken from the vla archives , if we assume a jaffe & perola ( 1973 ) aged synchrotron spectrum , an injection index of 2.0 and a loss magnetic field strength around the minimum - energy value of 0.4 nt ( taking into account cmb losses ) , and consequently material can be transported out from the jets or plume bases without _ in situ _ reacceleration so long as @xmath79 , or 2500 km s@xmath2 . e84 used the lack of observed depolarization to place a limit ( @xmath80 m@xmath49 ) on the density of thermal material ( effectively , protons ) in the tails . this method of setting a density limit has fallen out of favour in recent years , partly because it does not provide a strong upper limit ( laing 1984 ) and partly because it seems unlikely that the limits it does provide are close to the true values , which may be much lower . it may thus not be appropriate to use a proton density around @xmath81@xmath82 m@xmath49 as a ` typical ' value in the plumes . observationally , there are few other constraints on the internal density available : as discussed above , we know from the x - ray deficits seen in both ordinary fris and in cluster - centre objects ( e.g. brzan et al . 2004 ) that the missing pressure in these sources is _ not _ supplied by thermal protons at the temperature of the external medium , but this only really gives the constraint that the internal density is likely to be much lower than the external one . a true limit comes from considering the minimum possible internal density . if we assume that the wat plumes are in pressure balance and that the internal pressure is supplied by relativistic particles and magnetic field , presumably with some departure from the minimum energy condition , the plume material remains very light compared to the external medium : we require @xmath83 , where @xmath84 , and the effective density of the internal material is @xmath85 , while the external density is @xmath86 , where @xmath87 is the mass per particle in atomic mass units ( about 0.6 ) . the density contrast is thus @xmath88 , or about @xmath89 for the temperature of the 3c465 cluster . the internal density ( in terms of an equivalent proton density ) in this model , which represents a minimum possible internal density if we think that the plumes are not underpressured , would be of the order of 0.04 m@xmath49 , much lower than the ` typical ' values used by e84 . such a light plume is comparatively easy perhaps even too easy to bend by ram pressure ; if the flow speed up the plume , @xmath77 , were comparable to the local external sound speed ( @xmath90 km s@xmath2 ) , bulk motions with speeds of @xmath91 would have a significant effect , while if @xmath77 were comparable to the speed estimated above from spectral ageing arguments , we would require external bulk motions with speeds around @xmath92 . if instead we adopt the luminosity condition ( eq . [ lum ] ) to estimate the flow speed , retaining this low internal density , then we would obtain a bulk speed around @xmath93 [ for a radiative efficiency of 0.1 , which may be generous , in view of the jet luminosity inferred for the weaker source 3c31 by laing & bridle ( 2002 ) ] and , because of the low internal density , bulk motions with speeds around @xmath94 would still be sufficient to bend the plume : as e84 comment , low jet densities are possible , but require high speeds if the luminosity condition is to be satisfied . if , on the other hand , the pressure deficit at minimum energy is supplied by hot , entrained , internal thermal material , as is possible for fri lobes , then we know only that the density contrast is @xmath95 ( since the thermal protons then dominate the density of the flow ) and in that case external bulk motions at the sound speed can bend the plume significantly if @xmath96 , which for moderate heating of the internal protons requires plume speeds that are at most mildly supersonic : for example , if the plume speed is 2500 km s@xmath2 , then the internal protons must be heated to around 20 kev ( which would also give a density for internal protons that was close to the ` limit ' on internal density inferred by e84 ) . bending by subsonic motions requires a correspondingly higher temperature , increasing as @xmath97 , but this is not impossible , as we have little information on the temperature of the internal protons . in the case of a proton - dominated plume the luminosity condition , eq . [ lum ] , requires @xmath98 , and it can be shown that very high internal temperatures ( @xmath99 k ) , correspondingly low densities , and high plume flow speeds ( @xmath100 ) are required for the plume to be bent by sub - sonic bulk motions in the external medium . to summarize , the plumes can be bent by bulk motions in the external medium if they are light , but they will only also meet the luminosity condition of e84 if they are also fast . if we drop the assumption of _ in situ _ acceleration ( and instead assume passive transport of radiating electrons ) then the plumes need only be mildly supersonic , but must still be light . this conclusion revives the possibility that the plumes are in fact bent by the motion of the galaxy with respect to the host cluster . e84 suggest that the systematic motion of ngc 7720 could be of order 100300 km s@xmath2 , or @xmath101@xmath102 . their estimate does not necessarily rule out higher galaxy speeds , because they base it on the radial velocity of the galaxy with respect to the cluster mean ( which does not tell us about motions in the plane of the sky ) and on studies of the radial velocities of cds in general ( which may not be representative of the systems hosting wats ) . the possible detection of an x - ray wake behind 3c465 ( section [ wake ] ) may indicate a higher speed . nevertheless , it is interesting that the speed estimate of e84 is within the range estimated above for plume bending either if the plume s internal pressure is supplied entirely by relativistic particles and the plume is fast , or if a slower plume s internal pressure is supplied by very hot thermal material . is this picture consistent with the fact that the inner jets are not bent ? this question is more difficult to answer because we do not have a reliable estimate for the inner jet speed or density , and we can not assume pressure balance in the jet . if we assume @xmath103 ( from the beaming analysis ) and take the minimum equivalent internal density to be given by @xmath104 , using the minimum pressure plotted in fig . [ pprofile ] , then we can use euler s equation in the form @xmath105 where @xmath106 is the radius of curvature of the jet , @xmath107 is the galaxy speed , and @xmath108 is the width of the jet . ( the non - relativistic approximation is not badly wrong in this case . ) from the radio maps we can see that the radio jet bends , if it bends at all , by less than one jet width over its observed length @xmath109 . taking @xmath110 arcsec ( which includes the effects of projection with @xmath111 ) , @xmath112 arcsec , and an external density comparable to that at the transition between the hot and cool components in the x - ray ( 1500 protons m@xmath49 ) we find that speeds of 120 km s@xmath2 could produce a bend of one jet width . as this is very consistent with the possible values for systematic motion of the galaxy , and as our estimate of @xmath113 is a lower limit , we are satisfied that an explanation for the plume bending in terms of the systematic motion of the galaxy through the cluster is not inconsistent with the observations of essentially straight inner jets , provided that the jets are fast . we emphasise , however , that this model is only viable if the plume is light . if there is substantial entrained thermal material , then the analysis above suggests that sub - sonic motion of the galaxy through the cluster material is too slow to have a significant effect on the plumes . it is interesting in this context that there is a population of wat sources with bent inner jets ( exemplified by 0647 + 693 , hardcastle & sakelliou 2004 ) : if the jets in these sources _ are _ bent by ram pressure due to the host galaxy s motion , then the model we have discussed suggests either that the jets in these sources are weaker ( transport less energy ) or slower , or that the systematic motions of the host galaxies are larger . we will investigate this possibility in more detail elsewhere . [ fri - wat ] these results , combined with those of the previous subsection , mean that there are some interesting possible differences between wats and classical fri sources . on the one hand , we have shown ( section [ pcontent ] ) that the large - scale plumes of 3c465 are very similar to the extended structure of other fri sources in terms of the difference between the minimum internal pressure and the external pressure : in fris one of the best explanations for this observation is that the pressure deficit is supplied by heated , entrained material . on the other hand , as we have seen , bending of the 3c465 plumes by galactic motions puts limits on the amount of entrained material that is present : for example , if we adopt a plume speed of 2500 km s@xmath2 , then significant bending by bulk speeds of @xmath114 km s@xmath2 require thermal densities several orders of magnitude below the external density ( with correspondingly high temperatures if we want this material to supply the missing pressure ) , while satisfying the e84 luminosity condition requires even lower densities . these densities are considerably lower than the densities estimated by odea ( 1985 ) from the bending of narrow - angle tail ( nat ) sources . although it would be interesting to revisit the calculations of odea with modern estimates of the external density , and to compare them to estimates of the proton density required to provide pressure in the lobes , this may be an indication that type ii jets such as those in 3c465 are less efficient at entraining thermal material than the type i jets of nats .

chandra _ and _ xmm - newton_. x - ray emission is detected from the active nucleus and the inner radio jet , as well as a small - scale , cool component of thermal emission , a number of the individual galaxies of the host cluster ( abell 2634 ) , and the hotter thermal emission from the cluster itself . the x - ray detection of the jet allows us to argue that synchrotron emission may be an important mechanism in other well - collimated , fast jets , including those of classical double radio sources . the bases of the radio plumes are not detected in the x - ray , which supports the model in which these plumes are physically different from the twin jets of lower - power radio galaxies . the plumes are in fact spatially coincident with deficits of x - ray emission on large scales , which argues that they contain little thermal material at the cluster temperature , although the minimum pressures throughout the source are lower than the external pressures estimated from the observed thermal emission . our observations confirm both spatially and spectrally that a component of dense , cool gas with a short cooling time is associated with the central galaxy . however , there is no evidence for the kind of discontinuity in external properties that would be required in many models of the jet - plume transition in wats . galaxies : active x - rays : galaxies galaxies : individual : 3c465 galaxies : jets radiation mechanisms : non - thermal

we have observed the prototypical wide - angle tail radio galaxy 3c465 with _ chandra _ and _ xmm - newton_. x - ray emission is detected from the active nucleus and the inner radio jet , as well as a small - scale , cool component of thermal emission , a number of the individual galaxies of the host cluster ( abell 2634 ) , and the hotter thermal emission from the cluster itself . the x - ray detection of the jet allows us to argue that synchrotron emission may be an important mechanism in other well - collimated , fast jets , including those of classical double radio sources . the bases of the radio plumes are not detected in the x - ray , which supports the model in which these plumes are physically different from the twin jets of lower - power radio galaxies . the plumes are in fact spatially coincident with deficits of x - ray emission on large scales , which argues that they contain little thermal material at the cluster temperature , although the minimum pressures throughout the source are lower than the external pressures estimated from the observed thermal emission . our observations confirm both spatially and spectrally that a component of dense , cool gas with a short cooling time is associated with the central galaxy . however , there is no evidence for the kind of discontinuity in external properties that would be required in many models of the jet - plume transition in wats . although the wat jet - plume transition appears likely to be related to the interface between this central cool component and the hotter intra - cluster medium , the mechanism for wat formation remains unclear . we revisit the question of the bending of wat plumes , and show that the plumes can be bent by plausible bulk motions of the intra - cluster medium , or by motion of the host galaxy with respect to the cluster , as long as the plumes are light . galaxies : active x - rays : galaxies galaxies : individual : 3c465 galaxies : jets radiation mechanisms : non - thermal

astro-ph0502575

c

our principal results can be summarized as follows : 1 . the detection of x - ray emission from the inner part of the jet of 3c465 adds significantly to the evidence that fast ( at least mildly relativistic ) , efficient , frii - like ( ` type 2 ' ) jets can be sources of x - ray synchrotron radiation . 2 . the non - detection of x - ray emission from the jet termination point and the bases of the plumes rules out a model in which the plume bases resemble the bright bases of fri radio jets , which are often x - ray synchrotron sources : this is consistent with our favoured model for the wat jet - plume transition , in which the deceleration of the fast inner jets takes place rapidly when the jet interacts with the external medium at the edge of the plume , rather than slowly by entrainment as in fri jets . however , it is interesting that there is no detection of synchrotron x - rays from any of the plausible candidates for the jet termination ` hotspot ' in 3c465 , unlike the situation in the ( supposedly physically similar ) hotspots in low - power frii sources . more x - ray observations of wats with bright jet termination hotspots are required . we have made detailed measurements of the temperature and density structure of 3c465 s host cluster , a2634 , and in particular have resolved a central dense , cool thermal component with a relatively short cooling time , which is similar to the components seen in other wats and to the central components of fri sources in groups . there is no evidence for any discontinuity in the properties of the external medium , and so we conclude , following jetha et al . ( 2005 ) , that models of wat formation that require such discontinuities are ruled out . the most likely scenario based on our observations and those of jetha et al . is that the location of the jet - plume transition is related to the emergence of the jet from the small - scale cool component associated with the host galaxy , but the physical mechanism that sets the location of the plume base remains unclear . our observations show that the minimum pressures of the plumes fall below the pressures estimated for the external thermal material : this implies either some departure from the minimum - energy ( equipartition ) condition or an additional contribution to pressure from low - density , hot thermal material . the plumes of wats are thus similar to the large - scale components of more normal fri sources . 5 . revisiting the arguments related to wat plume bending , we find that the plumes can be bent by sub - sonic bulk motions of the icm , or by plausible motion of the host galaxy through the icm , provided that they are extremely light with respect to the external medium ( as would be the case if their internal pressures were dominated by relativistic particles and/or magnetic field ) .

although the wat jet - plume transition appears likely to be related to the interface between this central cool component and the hotter intra - cluster medium , the mechanism for wat formation remains unclear . we revisit the question of the bending of wat plumes , and show that the plumes can be bent by plausible bulk motions of the intra - cluster medium , or by motion of the host galaxy with respect to the cluster , as long as the plumes are light .

we have observed the prototypical wide - angle tail radio galaxy 3c465 with _ chandra _ and _ xmm - newton_. x - ray emission is detected from the active nucleus and the inner radio jet , as well as a small - scale , cool component of thermal emission , a number of the individual galaxies of the host cluster ( abell 2634 ) , and the hotter thermal emission from the cluster itself . the x - ray detection of the jet allows us to argue that synchrotron emission may be an important mechanism in other well - collimated , fast jets , including those of classical double radio sources . the bases of the radio plumes are not detected in the x - ray , which supports the model in which these plumes are physically different from the twin jets of lower - power radio galaxies . the plumes are in fact spatially coincident with deficits of x - ray emission on large scales , which argues that they contain little thermal material at the cluster temperature , although the minimum pressures throughout the source are lower than the external pressures estimated from the observed thermal emission . our observations confirm both spatially and spectrally that a component of dense , cool gas with a short cooling time is associated with the central galaxy . however , there is no evidence for the kind of discontinuity in external properties that would be required in many models of the jet - plume transition in wats . although the wat jet - plume transition appears likely to be related to the interface between this central cool component and the hotter intra - cluster medium , the mechanism for wat formation remains unclear . we revisit the question of the bending of wat plumes , and show that the plumes can be bent by plausible bulk motions of the intra - cluster medium , or by motion of the host galaxy with respect to the cluster , as long as the plumes are light . galaxies : active x - rays : galaxies galaxies : individual : 3c465 galaxies : jets radiation mechanisms : non - thermal

astro-ph0412501

i

cataclysmic variables ( cvs ) are interacting binary systems in which a white dwarf ( wd ) accretes matter from a low - mass companion . in non - magnetic cvs , the in - falling matter forms a disk around the wd . dwarf novae ( dne ) are disk - accreting cvs that undergo sharp rises in luminosity , during which the optical brightness typically increases by 2 to 5mag . the outbursts occur when a sudden increase in the gas viscosity ( due to a change in the ionization state of the gas ) allows matter to flow more rapidly through the disk , causing an increase in the accretion luminosity . in most dne , outbursts last for a day to a few weeks . in the subclass of dne known as u gem systems , outbursts end with a steady decline to minimum that has a typical timescale of a day however , in z cam systems , the decline from outburst occasionally stalls about 1mag below the peak . such ` standstills ' can last for days to , in a few cases , years . other non - magnetic cvs ( nova - like variables ) appear to remain persistently in the outburst state . the z cam standstill state is comparable to that of a nova - like variable , i.e. a steady , high state , but somewhat lower in luminosity than the peak of a dwarf nova outburst . the difference in outburst behavior between u gem , z cam and nova - like systems is thought to reflect a difference in the rate of mass flow from the secondary ( ) . in dne , is below the upper stability boundary ( ) for which dn outbursts can occur , whereas in nova - like systems @xmath1 . in z cam - type stars @xmath2 lies very near , but below @xmath3 most of the time . standstill phases are triggered when an increase in @xmath2 @xcite is accompanied by a decrease in @xmath3 as a response to heating of the outer edge of the accretion disk by the increased mass transfer @xcite . thus , @xmath2 becomes greater than @xmath3 and the disk assumes a stable high state ( like a nova - like variable ) . it has been difficult to measure mass transfer rates in cvs observationally , and as a result reliable empirical determinations of and are lacking . on theoretical grounds , is thought to be about when heating of the mass transfer stream is taken into account @xcite . z cam is the brightest ( in outburst ) and most extensively studied object of its class . it has an orbital period @xmath4d , which implies that it , like other standstill systems , is above the period gap . ( a complete list of system parameters is presented in table [ t : param ] . ) the difference in the magnitude at outburst maximum ( @xmath5=10.4 ) and quiescence ( @xmath613 ) is small for dne in general , but typical of standstill systems . the standstill @xmath5 averages 11.5 @xcite . far ultraviolet ( fuv ) spectra of z cam in outburst were obtained with the hopkins ultraviolet telescope ( hut ) @xcite . the outburst spectrum resembles that of a steady - state accretion disk with a mass accretion rate = , modified by resonant scattering by material in a wind @xcite . here , we attempt to complete the picture of z cam by describing fuv spectra of z cam in the quiescent and standstill states . we utilize spectra of z cam in quiescence that we have obtained using _ far ultraviolet spectroscopic explorer _ ( ) , and a spectrum of z cam in standstill obtained with telescope that we retrieved from the mast archive . our discussion of the effort to understand the two datasets is organized as follows : in section [ s : obs ] , we describe the data reduction and the mean spectrum of z cam in quiescence . in section [ s : cont ] , we attempt to fit the spectrum of z cam in quiescence , while in section [ s : var ] we analyze the time variability of the spectral lines . in section [ s : stand ] , we describe the spectrum of z cam in standstill as well as our attempts to model that spectrum . finally , in section [ s : discus ] , we evaluate the available spectral models and compare the properties of z cam in its different luminosity states .

a second component could be due to a rotating accretion belt , a remnant of an outburst that ended eleven days prior to the observation , or ongoing accretion . we have compared the quiescent spectrum with an archival spectrum of z cam in standstill , obtained with .

we have obtained _ far ultraviolet spectroscopic explorer _ ( 9051187 ) spectra of the non - magnetic cataclysmic variable z cam during a period of quiescence . the spectrum resembles that of a hot ( 57,000k ) metal - enriched white dwarf . a high effective temperature is consistent with the expectation that dwarf nova systems that show standstills , as z cam does , have higher than normal time - averaged mass accretion rates . it is also consistent with current estimates for the mass and distance to z cam . a white dwarf model in which 29% of the surface has a temperature of 72,000 k and the rest of the surface is at 26,000k also reproduces the spectral shape and the continuum flux level at the nominal distance and wd mass , and is a somewhat better statistical fit to the data . a second component could be due to a rotating accretion belt , a remnant of an outburst that ended eleven days prior to the observation , or ongoing accretion . we favor the uniform temperature model for z cam , largely because the data do not require anything more complicated . we have compared the quiescent spectrum with an archival spectrum of z cam in standstill , obtained with . the standstill spectrum is described well by a disk accreting at = , where the errors depend on the assumed inclination . the quiescent spectra cover a full orbital period and are time - resolved tos . no variability was observed in the continuum during the observation , but the depth of many of the absorption lines increased between orbital phases 0.65 and 0.81 . we attribute this effect to absorption by material associated with the accretion stream , which is easier to understand if the inclination is near the maximum allowed value of 68 .

astro-ph0412501

c

in quiescence , we find that the fuv continuum of z cam as observed by is qualitatively well - described in terms of emission from a hot metal - enriched wd atmosphere . for log g=8.5 , the best fit single temperature models has a of 57,200 k and of . the mass of the wd in z cam is estimated to be [email protected] , which suggests of . thus the two measurements of can be reconciled if the distance to z cam is @xmath63 pc , where the errors here are calculated from the uncertainty in the mass . this is slightly larger than , but wholly consistent with the astrometric distance of @xmath64 pc . if this interpretation is correct , then the wd in z cam is hot when compared to the wds in other cvs . a recent compilation of reliably determined wd temperatures in cvs by @xcite contains 34 objects ( other than z cam ) ranging in temperature from 9,500 k in the polar ef eri to 50,000 k in the nova - like variable and sw sex star dw uma . a temperature of 57,200 k would make z cam the hottest wd in a cv , excepting only v1500 cyg , which is still cooling from its nova outburst in 1975 and for which there is a crude blackbody temperature estimate of 70,000 to 120,000 k for the wd @xcite . is this reasonable ? the surface temperatures of wds in cvs are thought to be primarily determined by the accretion history of the wd @xcite . this is because although the wds in cvs emerge from the common envelope phase as very hot wds , the wds cool to temperatures of 4,500 - 6000 k in the 3 - 4 gyr required to reach the mass - exchanging stage . accretion reheats the surface layers of the wd on the time scales of outbursts ( weeks ) , as is evidenced from observations of temperature changes in wds in cvs such as u gem @xcite and vw hyi @xcite . more importantly , accretion reheats the entire wd over longer time scales ( ) , as the envelope and core adjust to the increase in mass of the wd , an effect known as compressional heating , and through nuclear burning of the accreted material @xcite . recently , @xcite have determined the relationship between time - averaged accretion rate and wd temperature . for a temperature of 57,000 k , the required rate is @xmath7 ( with uncertainties of perhaps a factor of two ) . ( a somewhat lower value of is obtained for 45,000 k , the temperature derived from the combined + spectrum . ) thus the accretion rate required is high , although it may be consistent with the fact that z cam is commonly in ( or near ) the high state . indeed , of the four systems with wd temperatures greater than 40,000 k , mv lyr , tt ari , ru peg , and dw uma , all but ru peg are nova - like variables and hence , like z cam , are systems where there are other indications of high time - averaged accretion rates . given this , we believe a temperature of 57,000 k is reasonable . an alternative interpretation of the quiescent spectrum is that the surface temperature of the wd in z cam is not uniform . our analysis shows that a somewhat better , but still not statistically acceptable fit to the data can be obtained if 2 t models are considered . our best fit had a low temperature component with @xmath36 of 26,300 k covering 71% of the surface and a high temperature component with @xmath37 of 71,700 k covering 29% . if the low temperature component represents energy loss from the interior of the wd while the high temperature component is due to recent or ongoing accretion , then 94% of the luminosity of the wd is due to recent accretion . is this reasonable ? if the excess emission is due to ongoing accretion , then the observed luminosity is @xmath65 , where @xmath66 is the fraction of gravitational energy available that is radiated . for the best 2 t parameters , the excess luminosity is , and the implied value of @xmath67 is . this is a fairly large value for a disk in quiescence , especially since @xmath68 , but is hard to rule out because of our lack of good models for emission from quiescent disks and the brightness of the wd itself . on the other hand , the accretion belt hypothesis is that some of the kinetic energy of material that reaches the boundary layer between the disk and wd is stored in a rapidly rotating belt and released slowly during the interoutburst period . in standard accretion disk theory , half of the gravitational energy of material accreted by the wd is radiated by the disk and half remains as kinetic energy when it enters the boundary layer . for z cam , with a disk accretion rate of , the disk luminosity is . for the outburst that preceded the observation , z cam remained within 1 mag . of outburst maximum for about 4 days ( see fig.[f : aavso ] ) , and hence the total energy release was about . the observation occurred about 11 days from the return to quiescence and , as noted previously , the belt luminosity in the context of this hypothesis was still . we do not have the coverage to determine the time decay of the second component , but a conservative lower limit is obtained by assuming constant luminosity until the observation and zero from then on . by this argument the energy stored in the rotating layer was at least , or 8% of the energy of the outburst . if one assumes an exponential decay , then the energy release is @xmath69 , where t is the time from outburst , l(t ) the luminosity at the time of our observation , and @xmath70 the decay time constant . if @xmath70 was 11 days , then the total energy released would be @xmath71 greater , or 20% of the total outburst energy . in either case , the energy that must be stored in the belt is large , though not energetically disallowed . if the hot component is produced by short - term effects of accretion , then the low t component reflects the long - term average of accretion . for a temperature of 26,300 , this requires an accretion rate of . this is lower than the number inferred from 1 t fits , but still substantial . so how does one decide between the two models we have explored for z cam in quiescence ? the data are not definitive . the model parameters that result from both models are substantially in agreement with other known parameters of the system . is somewhat better for the 2 t model than the uniform temperature model . but is much greater than 1 per degree of freedom in the 2 t model , showing that it is at best a step toward a correct model of z cam in quiescence , and the systematic errors in the spectrum are difficult to quantify . the differences in the two model spectra are quite small ; there is no characteristic of the spectrum that one can point to that requires a second component . evidence that for the high temperature component is high compared to the low t component is lacking . in u gem @xcite , single temperature models lead to a physically implausible result , namely that the wd grows in radius during quiescence . since we do not have multiple observations of z cam through a quiescent period , we do not know whether single t models lead to a similar problem in z cam . there are other well - observed systems , e.g. vw hyi @xcite and wz sge @xcite , which do not exhibit `` radius growth '' in the context of one temperature models , and so one can not use that as a guide . there is also no consensus theoretically . while @xcite have argued that mass transfer from the disk is higher immediately after outburst in order to explain uv delays in dne , the expected x - ray evidence of a gradual decline in hard x - ray emission following outbursts is lacking . furthermore the x - ray data that exist indicates that the x - ray plasma that exists in quiescence is not rotating rapidly @xcite , and therefore it is hard to understand how it could produce a heated annular region around the surface . while @xcite postulated the existence of accretion belts and described an instability that might release energy in an accretion belt , there are to our knowledge no detailed calculations of expected properties of the emitting region and no attempt to calculate the luminosity of the rotating component with time . given these considerations , our answer to the question of which model is most likely correct is to appeal to occam s razor and assert that the simpler 1 t model provide a good physical description of the data and is therefore is the best - bet physical model to pursue . in outburst , @xcite showed that the shape of the fuv ( 9001800 ) continuum of z cam could be modeled as an accretion disk with @xmath72 at an inclination of 57 and distance of 170 pc ( very close to the best astrometric distance of 163 pc ) . in the high state , the fits are self - consistent . in standstill we find @xmath73 . approximately , then , there appears to be at least an order of magnitude increase in from quiescence to standstill , while is an order of magnitude higher still . our derived gives an upper limit on , as it is a reasonable assumption that is also the rate of mass flow from the secondary at that time , which must be greater than . @xcite calculate @xmath74 , using the formula of @xcite and taking into account heating by the stream disk impact ( or if tidal dissipation is also included ) . our upper limit of @xmath75 is just consistent with this and with the upper limit of @xmath76 derived by @xcite . the comparison of the befs standstill spectrum and the spectrum of z cam in quiescence illustrates the problem of separating a wd spectrum from that of a disk . except for the flux level the qualitative differences between the befs spectrum and the spectrum are relatively small . in both states , the continuum slope is almost the same and , other than a slight increase in ionization temperature , their line spectra are very similar . yet , we find that in quiescence a wd alone can make up the continuum flux , whereas in standstill an accretion disk is the best model for the continuum flux . ultimately , the key model parameter that has enabled us to fit one model and reject others is the normalization required to match the observed and the modeled flux levels , which in turn depends on the wd mass and the distance from the sun . without a reasonably accurate knowledge of these parameters it would be impossible to judge between a wide range of models that reproduce the continuum slope and line spectrum equally well . in the time - resolved spectra of z cam in quiescence , we have seen enhanced line absorption between phase 0.65 and 0.81 , which covers the full range of atomic species and excitation levels . the enhanced absorption is very well centered on the mean line profile , which , in turn , is centered 100200 to the red of the transition rest velocity . the change in absorption width is greatest in those lines that are weak or absent from the mean spectrum , but also encompasses lines which are thought to be optically thick in the mean spectrum , leaving only the n@xmath13924 , o@xmath13@xmath321032,1038 and c@xmath13977 and 1176 lines unchanged . flux variations at around phase 0.50.8 are seen in a selection of disk - accreting systems , in phenomena such as dips in the uv @xcite and x - ray @xcite light curve of low - mass x - ray binaries and cv , and humps in the optical light curves of cv ( e.g. * ? ? ? this is typically attributed the interaction of the accretion stream from the secondary with the outer edge of the accretion disk . this can result in a luminous ` bright spot ' at the outer edge of the disk , which is thought to be responsible for the optical light curve modulation . also , the shock of the in - falling stream can cause the disk to bulge outwards at smaller radii or cause the stream to be deflected upward and arc over the surface of the accretion disk ( e.g. * ? ? ? this can cause phase dependent absorption as colder material in the accretion stream moves in front of the continuum source . the additional absorption seen in z cam is qualitatively compatible with some sort of stream - disk interaction . the enhanced absorption is narrow and centered on the mean line profile , indicative of matter that is moving across , rather than parallel to the line of sight . there is no corresponding change in the ionization state and the increase in absorption affects those transitions that appear to be optically thick in the mean spectrum . this is consistent with a disk bulge or stream temporarily moving to occult a larger area of the continuum source , or to place a greater column density of absorbing ions into the line of sight . also , the 0.650.81 phasing is in the phase range that stream - overflow effects are thought to be seen . in these respects z cam compares well with the spectrum of u gem in outburst @xcite , which was observed to show strong enhanced absorption at phase 0.53 to 0.79 . the line variability in u gem is of a similar strength to that seen in z cam , i.e. residual line core flux reducing to as little as 40% of its mean value in some lines , with the o lines being the least affected , although the n@xmath13924 , c@xmath131176 and other weak lines participate in the variation in u gem . smooth particle hydrodynamical models of the disk - stream interaction are now being produced ( e.g. * ? ? ? * ; * ? ? ? * ; * ? ? ? these models make predictions of the observational qualities which we are now able to set against the evidence for stream - related variability seen in the spectrum of z cam . the models are generally in agreement that stream overflow can occur at low mass - accretion rates , i.e. quiescent disks . @xcite predict that efficient cooling of the hot - spot region can occur in systems with low mass - accretion rates ( @xmath77 , or @xmath78 ) , allowing material to overflow the inner disk in a coherent stream , whereas , in higher- systems the stream impact region tends to push the disk into a bulging shape . using a different method , @xcite concurs with armitage & livio that up to @xmath59% of the matter in the accretion stream from the secondary is able to overflow the disk . if the phase - dependent absorption is a result of stream - overflow , it may have implications for the estimated inclination of the system . logically , for absorption by the stream , the absorption strength should have a strong correlation with inclination angle . at higher inclinations , material above the disk plane presents a higher column density to the bright continuum regions , i.e. the wd and inner disk . @xcite estimate that in outburst , enhanced absorption can be seen in systems of inclinations down to 65 , based on the maximum scale height above the disk attained by material in the accretion stream . however , they predict that , at low , material is not as strongly deflected upward as for high and that an inclination angle of at least 75is needed before the overflowing stream can cause accretion dips . at an inclination of only @xmath79(as taken from * ? ? ? * ) , z cam would fall below this lower limit . furthermore , an inclination above the current upper limit of 68 for z cam would produce observable eclipses and therefore can be ruled out . however , given the similarity of the time - dependent absorption to that seen in u gem ( @xmath80 ) , an inclination angle near 68seems far more plausible that the opposite extreme of 48 .

the spectrum resembles that of a hot ( 57,000k ) metal - enriched white dwarf . a high effective temperature is consistent with the expectation that dwarf nova systems that show standstills , as z cam does , have higher than normal time - averaged mass accretion rates . it is also consistent with current estimates for the mass and distance to z cam . a white dwarf model in which 29% of the surface has a temperature of 72,000 k and the rest of the surface is at 26,000k also reproduces the spectral shape and the continuum flux level at the nominal distance and wd mass , and is a somewhat better statistical fit to the data .

we have obtained _ far ultraviolet spectroscopic explorer _ ( 9051187 ) spectra of the non - magnetic cataclysmic variable z cam during a period of quiescence . the spectrum resembles that of a hot ( 57,000k ) metal - enriched white dwarf . a high effective temperature is consistent with the expectation that dwarf nova systems that show standstills , as z cam does , have higher than normal time - averaged mass accretion rates . it is also consistent with current estimates for the mass and distance to z cam . a white dwarf model in which 29% of the surface has a temperature of 72,000 k and the rest of the surface is at 26,000k also reproduces the spectral shape and the continuum flux level at the nominal distance and wd mass , and is a somewhat better statistical fit to the data . a second component could be due to a rotating accretion belt , a remnant of an outburst that ended eleven days prior to the observation , or ongoing accretion . we favor the uniform temperature model for z cam , largely because the data do not require anything more complicated . we have compared the quiescent spectrum with an archival spectrum of z cam in standstill , obtained with . the standstill spectrum is described well by a disk accreting at = , where the errors depend on the assumed inclination . the quiescent spectra cover a full orbital period and are time - resolved tos . no variability was observed in the continuum during the observation , but the depth of many of the absorption lines increased between orbital phases 0.65 and 0.81 . we attribute this effect to absorption by material associated with the accretion stream , which is easier to understand if the inclination is near the maximum allowed value of 68 .

astro-ph0412501

c

we have obtained observations with the satellite of the dwarf nova , z cam , during a period of quiescence . the spectrum is characterized by a fairly flat continuum , with a peak in the flux at 1000 . the line spectrum covers a broad range of atomic species and ionization stages and is dominated by absorption lines of fwhm @xmath12 . qualitatively , the spectrum can be self - consistently described in terms of emission from the wd with a temperature of 57,000 k. modest improvements in the fits are obtained with 2 t models in which most of the emission comes from a small fraction ( 29% ) of the surface heated to @xmath81k . the remainder is at 26,300k . we favor the uniform temperature model for z cam , largely because the data do not require anything more complicated . there are emission lines ( e.g. c@xmath131176 ) that do not come from the wd . the widths of these lines are consistent with a disk origin . in the phase - resolved quiescent spectrum we have observed transient enhanced line absorption . from the 0.650.81 phasing of this absorption and its effect on a wide range of ionization species , we attribute it to material raised from the disk plane into the los , due to interaction of the accretion stream with the disk , or the stream itself moving over the disk . such effects may be easier to understand if inclination of z cam is close to the upper limit of 68 . the standstill continuum is best described by an optically - thick disk accreting at a rate of , depending on the actual inclination of z cam . the continuum slope and the line absorption spectra differ little between quiescence and standstill , despite that , during standstill , an accretion disk creates ( most of ) the continuum flux . the qualitative similarity of the two spectra highlights the difficulty of separating wd from disk signatures without detailed model fitting and , also , underlines the reliance of the model fits on an secure knowledge of the distance and the wd mass . the high temperature of the wd in z cam and the estimates of mass accretion rate in standstill and outburst are all consistent with a mass transfer rate from the secondary star that is higher than in normal dne . additional fuv observations of z cam in the period following a normal outburst are highly desirable both to better understand the physics of standstill systems and to cleanly distinguish between single and multi - temperature models of the wd in z cam . in this research , we have used , and acknowledge with thanks , data from the aavso international database , based on observations submitted to the aavso by variable star observers worldwide . we also gratefully acknowledge the financial support from nasa through grant nag 5 - 11885 . all of the data presented in this paper was obtained from the multimission archive at the space telescope science institute ( mast ) . stsci is operated by the association of universities for research in astronomy , inc . , under nasa contract nas5 - 26555 . support for mast for non - hst data is provided by the nasa office of space science via grant nag5 - 7584 and by other grants and contracts . we thank the anonymous referee for their helpful comments and suggestions . anderson , n. 1988 , , 325 , 266 araujo - betancor , s. a. , g " ansicke , b. t. , long , k. s. , beuermann , k. , de martino , d. , sion , e. m. & szkody , p. 2005 , , submitted armitage , p. j. & livio , m. 1996 , , 470 , 1024 armitage , p. j. & livio , m. 1998 , , 493 , 898 baraffe , i. , kolb , u. 2000 , mnras , 318 , 354 baskill , d. , wheatly , p. j. , osbourne , j. p. 2001 , , 328 , 71 berriman , g. , szkody , p. , and capps , r. w. 1985 , , 217 , 327 buat - m ' enard , v. , hameury , j .- m . , & lasota , j .- 2001 , , 369 , 925 froning , c. s. , long , k. s. , drew , j. e. , knigge , c. , proga , d. 2001 , , 963 g " ansicke , b. t. & beuermann , k. 1996 , , 309 , l47 g " ansicke , b. t. , hoard , d. w. , beuermann , k. , sion , e. m. , & szkody , p. 1998 , , 338 , 933 hameury , j. , menou , k. , dubus , g. , lasota , j. , & hure , j. 1998 , , 298 , 1048 hellier , c. , garlick , m. a. , & mason , k. o. 1993 , , 260 , 299 hessman , f. v. 1999 , , 510 , 867 hubeny , i. 1988 , comput . commun . , 52 , 103 hubeny , i. , & lanz , t. 1995 , , 439 , 875 hubeny , i. , lanz , t. , & jeffery , c. s. 1994 , in newsletter on analysis of astronomical spectra ( st . andrews : st . andrews univ . ) , 20 , 30 hurwitz , m. et al.1998 , , 500 , l1 king , a. r. 1997 , , 288 , l16 king , a. r. & cannizzo , j. k. 1998 , , 499 , 348 kiplinger , a. l. 1980 , , 236 , 839 kippenhahn , r. & thomas , h .- c . 1978 , , 63 , 265 knigge , k. , long , k. s. , blair , w. p. , and wade , r. a. 1997 , , 476 , 291 kunze , s. , speith , r. , & hessman , f. v. 2001 , , 322 , 499 lasota , j .- p . 2001 , new astronomy review , 45 , 449 liu , f. k. , meyer , f. , & meyer - hofmeister , e. 1995 , , 300 , 823 long , k. s. et al . 1991 , , 381 , l25 long , k. s. , blair , w. p. , bowers , c. w. , davidsen , a. f. , kriss , g. a. , sion , e. m. , & hubeny , i.1993 , , 405 , 327 mauche , c. w. , wheatley , p. j. , long , k. s. , raymond , j. c. 2004 , & szkody , p. 2004 , in asp conf . the astrophysics of cataclysmic variables and related objects , ed . j. m. hameury & j. p. lasota , in press meyer , f. & meyer - hofmeister , e. 1994 , , 288 , 175 moos , h. w. et al . 2000 , , 538 , l1 oppenheimer , b. d. , kenyon , s. j. , & mattei , j. a. 1998 , , 115 , 1175 sahnow , d. j. et al.2000 , , 538 , l7 schreiber , m. r. & g " ansicke , b. t. 2002 , , 382 , 124 shafter , a. w.1983 , ph.d . thesis , ucla sion , e. m. 1999 , , 111 , 532 schmidt , g. d. , liebert , j. , & stockman , h. s. 1995 , , 441 , 414 sion , e. m. 1985 , , 297 , 538 sion , e. m. 1995 , , 438 , 876 sion , e. m. & urban , j. 2002 , , 572 , 456 sion , e. m. , szkody , p. , g " ansicke , b. , cheng , f. h. , ladous , c. , & hassall , b. 2001 , , 555 , 834 stehle , r. , king , a. , & rudge , c. 2001 , , 323 , 584 szkody , p. & wade , r. a. 1981 , , 251 , 201 szkody , p. & mateo , m. 1986 , , 301 , 286 szkody , p. , g " ansicke , b. t. , howell , s. b. , & sion , e. m. 2002 , , 575 , l79 lcccc n * + h ... & 922924 ; 923.2 & [email protected] & 1030@xmath6265 + h ... & 930.748 & [email protected] & 570@xmath6275 + s ... & 933.4 & [email protected] & 565@xmath6230 + h ... & 937.8 & [email protected] & 545@xmath6245 + s ... & 944.5 & [email protected] & 470@xmath6220 + h+ p ... & 949.7 ; 950.7 & [email protected] & 900@xmath6260 + h ... & 972.5 & 1.09@xmath93 & 600@xmath6250 + c ... & 977.0 & [email protected] & 590@xmath6240 + he * + n ... & 989.8 ; 992.3 & [email protected] & 1150@xmath6240 + si * ... & 997.4 & @xmath94 & @xmath95 + c * ... & 1010.0 & [email protected] & 420@xmath6290 + s ... & 1012.5 & [email protected] & 240@xmath6260 + s ... & 1015.5 & [email protected] & 410@xmath6230 + s ... & 1021.1 & [email protected] & 500@xmath6230 + ... & 1025.7 & [email protected] & 470@xmath6290 + o ... & 1031.9 & [email protected] & 610@xmath6220 + o ... & 1037.6 & [email protected] & 700@xmath6220 + s ... & 1062.7 & [email protected] & 480@xmath6230 + s ... & 1073.0 & [email protected] & 480@xmath6230 + s * ... & 1077.2 & [email protected] & 280@xmath6261 + he * + n ... & 1084.9 ; 1085.0 & [email protected] & 1020@xmath6280 + si * ... & 1108.4 ; 1110.0 & [email protected] & 890@xmath6220 + si * ... & 1113.2 & [email protected] & 500@xmath6220 + p ... & 1118.0 & @xmath96 & @xmath97 + si * ... & 1122.5 & [email protected] & 730@xmath6250 + p+ si * ... & 1128.0 ; 1128.3 & [email protected] & 850@xmath6280 + c * ... & 1175.3 & [email protected] & 940@xmath6230 + cccccccc + da & 4.3 & 58,100 & - & - & 163 & 6.7 & 672 + wd & 4.4 & 57,200 & 330 & - & 163 & 10.1 & 1146 + 2wd & 6.0 ( 71%/29% ) & 26,300/71,700 & 170/370 & - & 163 & 9.4 & 1146 + disk ( 46 ) & - & - & - & 6.6 & 163 & 32.1 & 1146 + disk ( 57 ) & - & - & - & 9.1 & 163 & 27.1 & 1146 + disk ( 68 ) & - & - & - & 12.9 & 163 & 22.0 & 1146 + disk ( 57 ) & - & - & - & 180.0 & 710 & 11.9 & 1146 + disk ( 46 ) & wd & 4.0 & 62,000 & 306 & 1.2 & 163 & 9.8 & 1146 + disk ( 57 ) & wd & 4.0 & 62,000 & 295 & 1.7 & 163 & 9.8 & 1146 + disk ( 68 ) & wd & 3.9 & 62,000 & 293 & 2.5 & 163 & 9.8 & 1146 + + wd & 6.0 & 44,600 & - & - & 163 & 8.2 & 471 + 2wd & 18.4(97%/3% ) & 17,600/82,400 & - & - & 163 & 4.9 & 471 + disk ( 46 ) & - & - & - & 6.0 & 163 & 11.1 & 471 + disk ( 57 ) & - & - & - & 8.7 & 163 & 9.6 & 471 + disk ( 68 ) & - & - & - & 12.9 & 163 & 8.1 & 471 + disk ( 57 ) & - & - & - & 182.0 & 710 & 5.8 & 471 + disk ( 46 ) & wd & 3.0 & 82,000 & - & 3.0 & 163 & 4.7 & 471 + disk ( 57 ) & wd & 3.0 & 81,000 & - & 4.0 & 163 & 5.0 & 471 + disk ( 68 ) & wd & 3.0 & 81,000 & - & 6.3 & 163 & 5.4 & 471 + + wd & 12.5 & 59,700 & 240 & - & 163 & 2.6 & 1505 + disk ( 46 ) & - & - & - & 43.0 & 163 & 2.7 & 1505 + disk ( 57 ) & - & - & - & 69.0 & 163 & 2.6 & 1505 + disk ( 68 ) & - & - & - & 112.0 & 163 & 2.5 & 1505 + disk ( 57 ) & - & - & - & 195.0 & 270 & 2.4 & 1505 + [ fits ] lcccc h ... & 930.748 & @xmath98 & @xmath99 + s ... & 933.4 & @xmath100 & @xmath101 + h ... & 937.8 & @xmath102 & @xmath103 + s ... & 944.5 & @xmath104 & @xmath105 + h + p ... & 949.7 & @xmath106 & @xmath107 + h ... & 972.5 & @xmath108 & @xmath109 + c ... & 977.0 & @xmath110 & @xmath111 + n + si * ... & 989.8 ; 993.5 & @xmath112 & @xmath113 + si * ... & 994.8 & @xmath114 & @xmath115 + si * ... & 997.4 & @xmath116 & @xmath117 + s ... & 1012.5 & @xmath96 & @xmath115 + s ... & 1015.5 & @xmath96 & @xmath115 + s ... & 1021.1 & @xmath118 & @xmath119 + ... & 1025.7 & @xmath120 & @xmath121 + o ... & 1031.9 & @xmath122 & @xmath123 + o ... & 1037.6 & @xmath124 & @xmath125 + s ... & 1062.7 & @xmath126 & @xmath127 + s ... & 1073.0 & @xmath128 & @xmath129 + he * + n ... & 1084.9 ; 1085.0 & @xmath130 & @xmath131 + si * ... & 1108.4 ; 1110.0 & @xmath132 & @xmath133 + si * ... & 1113.2 & @xmath134 & @xmath135 + p ... & 1118.0 & @xmath136 & @xmath137 + si * ... & 1122.5 & @xmath138 & @xmath139 + p + si * ... & 1128.0 ; 1128.3 & @xmath120 & @xmath140 +

we have obtained _ far ultraviolet spectroscopic explorer _ ( 9051187 ) spectra of the non - magnetic cataclysmic variable z cam during a period of quiescence . we favor the uniform temperature model for z cam , largely because the data do not require anything more complicated . the standstill spectrum is described well by a disk accreting at = , where the errors depend on the assumed inclination . we attribute this effect to absorption by material associated with the accretion stream , which is easier to understand if the inclination is near the maximum allowed value of 68 .

we have obtained _ far ultraviolet spectroscopic explorer _ ( 9051187 ) spectra of the non - magnetic cataclysmic variable z cam during a period of quiescence . the spectrum resembles that of a hot ( 57,000k ) metal - enriched white dwarf . a high effective temperature is consistent with the expectation that dwarf nova systems that show standstills , as z cam does , have higher than normal time - averaged mass accretion rates . it is also consistent with current estimates for the mass and distance to z cam . a white dwarf model in which 29% of the surface has a temperature of 72,000 k and the rest of the surface is at 26,000k also reproduces the spectral shape and the continuum flux level at the nominal distance and wd mass , and is a somewhat better statistical fit to the data . a second component could be due to a rotating accretion belt , a remnant of an outburst that ended eleven days prior to the observation , or ongoing accretion . we favor the uniform temperature model for z cam , largely because the data do not require anything more complicated . we have compared the quiescent spectrum with an archival spectrum of z cam in standstill , obtained with . the standstill spectrum is described well by a disk accreting at = , where the errors depend on the assumed inclination . the quiescent spectra cover a full orbital period and are time - resolved tos . no variability was observed in the continuum during the observation , but the depth of many of the absorption lines increased between orbital phases 0.65 and 0.81 . we attribute this effect to absorption by material associated with the accretion stream , which is easier to understand if the inclination is near the maximum allowed value of 68 .

astro-ph0304371

i

galaxies in dense clusters evolve mainly due to tidal interactions , mergers and ram pressure stripping with accretion contributing to the buildup of the cd galaxies that often reside at the bottom of the cluster potential ( dressler 1984 ) . these processes are expected to be most effective in the cluster cores where higher densities will lead to more efficient ram pressure stripping and tidal effects should be stronger . indeed , observations of in abell 2029 failed to detect emission - line galaxies projected within 600 kpc of the center of the cluster , with the exception of a background galaxy , whereas two - thirds of the galaxies located outside this radius were detected ( dale & uson 2000 ) . it is well established that galaxies in clusters tend to be deficient in their neutral hydrogen ( ) emission ( e.g. haynes , giovanelli & chincarini 1984 ) . more recently , solanes et al . ( 2001 ) analyzed a sample of 1900 galaxies in the fields of eighteen nearby clusters ( @xmath7 ) and found that deficient galaxies are more likely located near the cluster cores . this trend holds true in abell 2670 , which is one of the most distant clusters to be imaged in neutral hydrogen ( @xmath8 , van gorkom 1996 ) . a truly extreme case is that of abell 2029 which is at the same redshift as abell 2670 but in which a deeper search for by one of us ( jmu ) resulted in an order of magnitude lower detection rate . deficiency in cluster galaxies is usually attributed to ram pressure stripping : as galaxies approach the central regions of clusters at a speed @xmath9 , they experience a strong interstellar - gas intracluster - gas interaction which is proportional to @xmath10 . this interaction is likely to induce a burst of star formation and subsequently strip away most of the the remaining galactic interstellar gas ( fujita 1998 ; fujita & nagashima 1999 ; balogh et al . 1999 ; quilis , moore & bower 2000 ) . new results from the sloan digital sky survey show cluster effects that can not be reconciled by the morphology - density relation alone . using a volume - limited sample of 8598 galaxies with @xmath11 , gomez et al . ( 2003 ) find that the galaxy star formation rate decreases with increasing galaxy density . in a study of the rotation curves of 510 galaxies in 53 clusters spread throughout @xmath12 , dale et al . ( 2001 ) found support for this galaxy - cluster interaction scenario in the normalized radial extent of emission in cluster spiral galaxies : it increases 4@xmath132% per mpc of ( projected ) separation of the galaxy from the cluster - core . analysis of the kinematical asymmetries of the galaxies also shows a trend with cluster - centric distance rotation curve asymmetry is greater by a factor of two for inner cluster early - type spiral galaxies . these two results lead these authors to claim that _ `` such trends are consistent with spiral disk perturbations or even the stripping of the diffuse , outermost gaseous regions within the disks as galaxies pass through the dense cluster cores '' _ ( dale et al . 2001 ) . are such observables different in @xmath14 clusters ? are they different for richer , more x - ray luminous clusters , thus solidifying the interpretation that the intracluster environment can significantly impact galaxy evolution in clusters ? in an effort to explore the tully - fisher relation at @xmath5 , we have undertaken a new imaging and spectroscopic campaign of four rich abell clusters : abell 1413 , abell 2029 , abell 2218 , and abell 2670 . results from the first cluster surveyed , abell 2029 , were presented in dale & uson ( 2000 ) . here , we present a full analysis of the data for all four clusters . we also include results from abell 2295 , the only @xmath8 cluster observed in the dale et al . ( 1999 ) survey .

we have obtained new optical imaging and spectroscopic observations of 78 galaxies in the fields of the rich clusters abell 1413 ( ) , abell 2218 ( ) and abell 2670 ( ) . we have combined these data with our previous observations of abell 2029 and abell 2295 ( both at ) , which yields a sample of 156 galaxies .

we have obtained new optical imaging and spectroscopic observations of 78 galaxies in the fields of the rich clusters abell 1413 ( ) , abell 2218 ( ) and abell 2670 ( ) . we have detected line emission from 25 cluster galaxies plus an additional six galaxies in the foreground and background , a much lower success rate than what was found (% ) for a sample of 52 lower - richness abell clusters in the range . we have combined these data with our previous observations of abell 2029 and abell 2295 ( both at ) , which yields a sample of 156 galaxies . we evaluate several parameters as a function of cluster environment : tully - fisher residuals , equivalent width , and rotation curve asymmetry , shape and extent . although is more easily detectable in galaxies that are located further from the cluster cores , we fail to detect a correlation between extent and galaxy location in those where it is detected , again in contrast with what is found in the clusters of lesser richness . we fail to detect any statistically significant trends for the other parameters in this study . the zero - point in the tully - fisher relation is marginally fainter ( by 1.5 ) than that found in nearby clusters , but the scatter is essentially unchanged .

astro-ph0304371

c

the goal of this survey is to investigate the impact that the intracluster medium has on galaxies residing in a range of cluster environments at @xmath94 . a total of 156 galaxies were observed spectroscopically in five clusters ranging from rich , x - ray luminous systems to richness class 0 clusters undetected in x - rays . about half of the sample ( 45% ) was detected in . we observe no significant trend in the extent of emission as a function of cluster - centric distance . this is in sharp contrast to what is observed in the nearby cluster sample , particularly when the projected cluster - centric distance is expressed in units of @xmath66 . we also fail to detect significant trends in rotation curve asymmetry , @xmath43 band tully - fisher residuals , rotation curve shape , and equivalent width . the relatively small number of rotation curves obtained for the dense clusters might be the main reason for our inability to detect statistically significant trends with cluster richness or location within the clusters . however , it is also possible that the low detection rate is a consequence of ram pressure stripping of the galaxies as they traverse the cores of the clusters , so that the galaxies that we do detect in would have avoided the cluster cores irrespective of their projected cluster - centric separation . indeed , a vla study of abell 2670 found @xmath95 sources , but _ none _ within 0.5 mpc of the cluster center ( van gorkom 1996 ) . but a deep vla survey of abell 2029 , sensitive to a 5@xmath6 mass of @xmath96 m@xmath97 within 1.5 mpc of the cluster center , has detected only three sources ( uson , in preparation ) . there is a clear difference in the cluster - centric distribution of actively star - forming and quiescent galaxies , with the star - forming galaxies preferentially found more in the cluster peripheries . this result is perhaps not surprising , given that the well - known cluster morphology - density relation likely contributes to this effect . although ground - based imaging precluded perfect sample selection , we strived to only observe spiral galaxies . moreover , as can be seen in figure [ fig : env ] , our sample selection appears to have been skewed towards late - type spirals near the cluster cores . thus , if the morphology - density relation has contributed to our observed distribution of strong and weak emitters , our sample selection seems to have minimized its impact . the zero - point in the @xmath5 @xmath43 band tully - fisher relation ( equation [ eq : tfhighz ] ) is a bit fainter ( by 1.5@xmath6 ) than that observed for nearby clusters ( equation [ eq : tflowz ] ) . the data for abell 2218 , even more distant at @xmath19 , are an additional 1.8@xmath6 fainter than the @xmath5 template . are these discrepancies related to the comparatively high redshift of the clusters ? is this discrepancy evidence of evolution or observational biases ? it is unclear whether or not the difference is due to systematic effects in the observational program : though seeing effects push deprojected rotational velocity widths to artificially high values , `` slit smearing '' of the disk rotation profile conversely biases the widths low . in case the correction recipes applied are unfit at this relatively high redshift , the simulations and empirical prescriptions used to generate the correction recipes have been revisited . for example , dale et al . ( 1997 ) have calibrated the effects of seeing on the inferred inclinations . they found a linear relation between the circularization of isophotes and the ratio of the seeing full - width half - maximum to face - on disk scale length @xmath98 : _ true/_obs = 1 + 0.118 , where @xmath99 and @xmath100 are the `` intrinsic '' and `` observed '' disk ellipticities . however , they only tested this relation for galaxies observed out to @xmath36 . using a separate technique from that explored in dale et al . ( 1997 ) , we are able to both reproduce the seeing correction of dale et al . ( 1997 ) over their sample s @xmath101 redshift range , and find that it is applicable even for galaxies as distant as those in abell 2218 . in the simulations artificial , inclined disk galaxies ( created using mkobjects in the artdata package of iraf ) are progressively smoothed such that they exhibit increasingly circular isophotes . isophotal ellipses are fit to the simulated galaxies using the same techniques used for the observed galaxies . the relation first quantified by dale et al . ( 1997 ) tends to non - linearity only when @xmath98 exceeds 2.0 , and all of the galaxies observed in abell 2218 have this ratio less than 1.8 . both the total and intrinsic scatter of the @xmath43 band tully - fisher relation at @xmath14 are similar to that in nearby clusters . the total scatter is comparable to that found in giovanelli et al . ( 1997 ) and dale et al . ( 1999 ) : 0.35 and 0.38 mag , respectively . the intrinsic scatter in the tully - fisher relation is the portion that can not be reconciled by measurement uncertainties . this parameter presumably reflects the various feedback mechanisms and range of parameter space involved in galaxy formation . andersen et al . ( 2001 ) suggest that perhaps up to half of the intrinsic scatter can be accounted for if disk galaxies simply have an average inherent ellipticity of 0.05 , a value they find for their sample of seven nearly face - on spiral galaxies . compared to the @xmath102 mag intrinsic scatter found by dale et al . ( 1999 ) in clusters at lower redshifts , there is no significant difference in the intrinsic @xmath43 band tully - fisher scatter in these four clusters at @xmath5 . in short , no significant evolution in the @xmath43 band tully - fisher properties of cluster galaxies appears to have occurred since @xmath14 . we are grateful for the assistance richard cool provided in preliminary explorations for correlations with cluster x - ray data . we thank steve boughn for his permission to use the unpublished _ i_-band observations of abell 2029 and abell 2218 and heinz andernach for his insight on the redshifts of abell 2218 . the comments of the referee helped to improve the presentation of this work . the results presented in this paper are based on observations carried out at the palomar observatory ( po ) and at the kitt peak national observatory ( kpno ) . the hale telescope at the po is operated by the california institute of technology under a cooperative agreement with cornell university and the jet propulsion laboratory . kpno is operated by the association of universities for research in astronomy , inc . , under a cooperative agreement with the national science foundation . this research has made use of the nasa / ipac extragalactic database ( ned ) which is operated by the jet propulsion laboratory , california institute of technology , under contract with nasa , and the image reduction and analysis facility ( iraf ) which is distributed by the national optical astronomy observatories , which are operated by the association of universities for research in astronomy , inc . , under a cooperative agreement with the national science foundation . the nrao is a facility of the national science foundation which is operated under cooperative agreement by associated universities , inc . the digitized sky surveys were produced at the space telescope institute under u.s . government grant nag w-2166 . the images of these surveys are based on photographic data obtained using the oschin schmidt telescope on palomar mountain and the uk schmidt telescope . dum abell , g. o. 1958 , , 3 , 211 andersen , d. r. , bershady , m. a. , sparke , l. s. , gallagher iii , j. s. & wilcots , e.m . 2001 , , 551 , l131 balogh , m. l. , schade , d. , morris , s. l. , yee , h. k. c. , carlberg , r. g. , & ellingson , e. 1998 , , 504 , l75 balogh , m. l. , morris , s. l. , yee , h. k. c. , carlberg , r. g. & ellingson , e. 1999 , , 527 , 54 balogh m. l. , couch , w. j. , smail , i. , bower , r. g. & glazebrook , k. 2002 , , 335 , 10 bird , c. 1994 , , 422 , 480 boughn , s. p. , uson , j. m. , blount , c. d. & gupta , g. 2000 , , 32 , 1499 carlberg , r. g. , yee , h. k. c. & ellingson , e. 1997 , , 478 , 462 conselice , c. j. & gallagher iii , j. s. 1999 , , 117 , 75 dale , d. a. , giovanelli , r. , haynes , m. p. , scodeggio , m. , hardy , e. & campusano , l. 1997 , , 114 , 455 dale , d. a. , giovanelli , r. , haynes , m. p. , scodeggio , m. , hardy , e. & campusano , l. 1998 , , 115 , 418 dale , d. a. 1998 , ph.d . thesis , cornell university dale , d. a. , giovanelli , r. , haynes , m. p. , hardy , e. & campusano , l. 1999 , , 118 , 1489 dale , d. a. & uson , j. m. 2000 , , 120 , 552 dale , d. a. , giovanelli , r. , haynes , m. p. , hardy , e. & campusano , l. 2001 , , 121 , 1886 dressler , a. 1978 , , 226 , 55 dressler , a. 1984 , , 22 , 185 dressler , a. , oemler jr . , a. , couch , w. j. , smail , i. , ellis , r. s. , barger , a. j. , butcher , h. , poggianti , b. m. & sharples , r. m. 1997 , , 490 , 577 ebeling , h. , edge , a. c. , fabian , a. c. , allen , s. w. , crawford , c. s. & bhringer h. 1997 , , 479 , l101 fujita , y. 1998 , , 509 , 587 fujita , y. & nagashima , m. 1999 , , 516 , 619 gavazzi , g. , boselli , a. , pedotti , p. , gallazzi , a. & carrasco , l. 2002 , , 396 , 449 giovanelli , r. , haynes , m. p. , herter , t. , vogt , n. p. , da costa , l. n. , freudling , w. , salzer , j. j. & wegner , g. 1997 , , 113 , 53 gomez , p.l . , nichol , r.c . , miller , c.j . et al . ( 2003 ) , 584 , 210 haynes , m. p. , giovanelli , r. & chincarini , g. l. 1984 , , 22 , 445 haynes , m. p. , hogg , d. e. , maddalena , r. j. , roberts , m. s. & van zee , l. 1998 , , 115 , 62 jones , c. & forman , w. 1999 , , 511 , 65 kennicutt , r.c . 1998 , , 36 , 189 kneib , j .- p . , mellier , y. , pell , r. miralda - escud , j. , le borgne , j .- f . , bhringer , h. & picat , j .- 1995 , , 303 , 27 kneib , j .- ellis , r. s. , smail , i. , couch , w. j. & sharples , r. m. 1996 , , 471 , 643 kogut , a. et al . 1993 , , 419 , 1 koopmann , r.a . & kenney , j.d.p . 2002 , , submitted , astro - ph/0209547 kornreich , d. a. , haynes , m. p. , lovelace , r. v. e. & van zee , l. 2000 , , 120 , 139 landolt , a. u. 1983 , , 88 , 439 lewis , i. , et al . 2002 , , 334 , 673 magri , c. , haynes , m. p. , forman , w. , jones , c. & giovanelli , r. 1988 , , 333 , 136 oke , j. b. & gunn , j. e. 1982 , , 94 , 586 persic , m. & salucci , m. 1991 , , 368 , 60 poggianti , b. m. 1997 , , 122 , 399 quilis , v. , moore , b. & bower , r. 2000 , science , 288 , 1617 richter , o .- g . & sancisi , r. 1994 , , 290 , l9 rubin , v. c. , waterman , a. h. & kenney , j. d. p. 1999 , , 118 , 236 sarazin , c. 1986 , rev . , 58 , 1 solanes , j. m. , manrique , a. , garca - gmez , c. , gonzlez - casado , g. , giovanelli , r. & haynes , m. p. 2001 , , 548 , 97 swaters , r. a. , schoenmakers , r. h. m. , sancisi , r. & van albada , t. s. 1999 , , 304 , 330 uson , j. m. , boughn , s. p. & kuhn , j. r. 1991 , , 369 , 46 van gorkom , j. h. 1996 , in _ the minnesota lectures on extragalactic hydrogen _ , ed . e. d. skillman , asp conference series , vol . 106 , p. 293 whitmore , b. c. , forbes , d. a. & rubin , v. c. 1988 , , 333 , 542 ziegler , b. l. , bower , r. g. , smail , i. , davies , r. l. & lee , d. 2001 , , 325 , 1571 ccccccclccc a1413 & 115248 & @xmath103233900 & 226,@xmath10377 & 42453570 & 42767 & 1490 & 9 & 10.30.5 & 4.5 + a2029 & 150830 & @xmath103055700 & 7,@xmath10351 & 23220170 & 23400 & 1471 & 92 & 16.40.6 & 4.4 + a2218 & 163542 & @xmath103661900 & 98,@xmath10338 & 52514170 & 52497 & 1437 & 93 & 9.30.4 & 4.3 + a2295 & 180018 & @xmath103691300 & 99,@xmath10330 & 24623200 & 24555 & 563 & 9 & @xmath1040.1 & 0 + a2670 & 235136 & @xmath105104100 & 81,@xmath10569 & 22841 60 & 22494 & 895 & 235 & 2.00.1 & 1.4 + [ tab : cluster_sample ] llcccccc 1999 apr 20 - 21 & a2218 & 0.7 & 1 & 1/ 8 & 1.52 & 72507920 & 30,30062,100 + 1999 may 14 & a2218 & 1 & 3 & 4/11 & 1.52.5 & 73308000 & 33,50065,700 + 1999 oct 31-nov 1 & a2670 & 2 & 16 & 15/28 & 1.52 & 67807450 & 9,80040,300 + 2001 may 20 - 21 & a1413,a2218 & 2 & 7 & 12/31 & 1.52.5 & 73308000 & 33,50065,700 + [ tab : spec_runs ] a1413 - 211894 & 115208.1 & @xmath103233828 & ... & 0 & a2218 - 261717 & 163612.6 & @xmath103663031 & 52378 & 1 + a1413 - 211903 & 115216.8 & @xmath103235549 & ... & 0 & a2218 - 261274 & 163618.1 & @xmath103662436 & ... & 0 + a1413 - 211916 & 115228.0 & @xmath103235115 & ... & 0 & a2218 - 261718 & 163644.7 & @xmath103661913 & 51601 & 1 + a1413 - 211924 & 115238.6 & @xmath103235007 & 42900 & 2 & a2218 - 261719 & 163647.4 & @xmath103661425 & ... & 0 + a1413 - 211934 & 115243.6 & @xmath103234059 & 42844 & 0 & a2218 - 261720 & 163647.5 & @xmath103662135 & 54557 & 2 + a1413 - 211938 & 115244.8 & @xmath103234332 & ... & 0 & a2218 - 261286 & 163650.7 & @xmath103662455 & ... & 0 + a1413 - 211956 & 115301.1 & @xmath103234748 & ... & 0 & a2218 - 261721 & 163658.6 & @xmath103660743 & ... & 0 + a1413 - 211957 & 115301.6 & @xmath103234725 & 41400 & 2 & a2218 - 261289 & 163700.4 & @xmath103661548 & ... & 0 + a1413 - 211968 & 115331.3 & @xmath103235648 & ... & 0 & a2218 - 261290 & 163705.2 & @xmath103661137 & ... & 0 + a2218 - 261222 & 163237.5 & @xmath103662428 & ... & 0 & a2218 - 261723 & 163752.9 & @xmath103661412 & ... & 0 + a2218 - 261223 & 163248.0 & @xmath103661744 & 56018 & 1 & a2218 - 261291 & 163755.8 & @xmath103662815 & 51553 & 1 + a2218 - 261702 & 163250.8 & @xmath103660732 & ... & 0 & a2670 - 630560 & 235008.6 & @xmath105110012 & 22883 & 0 + a2218 - 261224 & 163255.6 & @xmath103663400 & ... & 0 & a2670 - 630563 & 235024.5 & @xmath105105752 & 21810 & 1 + a2218 - 261703 & 163302.7 & @xmath103661240 & ... & 0 & a2670 - 630565 & 235031.5 & @xmath105105736 & 21948 & 1 + a2218 - 261704 & 163314.2 & @xmath103662632 & ... & 0 & a2670 - 630567 & 235034.8 & @xmath105110058 & 22561 & 1 + a2218 - 261705 & 163318.4 & @xmath103662500 & 55706 & 1 & a2670 - 630568 & 235034.9 & @xmath105104620 & 22892 & 1 + a2218 - 261706 & 163340.0 & @xmath103660322 & ... & 0 & a2670 - 630570 & 235044.7 & @xmath105103449 & 22607 & 0 + a2218 - 261226 & 163421.0 & @xmath103662518 & 52793 & 1 & a2670 - 630571 & 235044.9 & @xmath105103944 & 22291 & 2 + a2218 - 261708 & 163456.9 & @xmath103663036 & ... & 0 & a2670 - 630574 & 235053.7 & @xmath105104718 & 21843 & 0 + a2218 - 261709 & 163502.8 & @xmath103661355 & ... & 0 & a2670 - 630577 & 235111.6 & @xmath105105034 & 23501 & 1 + a2218 - 261244 & 163518.8 & @xmath103662121 & 42855 & 1 & a2670 - 630584 & 235124.6 & @xmath105105820 & 23662 & 0 + a2218 - 261245 & 163523.0 & @xmath103661938 & ... & 0 & a2670 - 630585 & 235126.2 & @xmath105103826 & 23757 & 0 + a2218 - 261247 & 163524.0 & @xmath103662003 & ... & 0 & a2670 - 630318 & 235129.7 & @xmath105103642 & 20968 & 0 + a2218 - 261710 & 163524.6 & @xmath103661204 & ... & 0 & a2670 - 630317 & 235130.4 & @xmath105104231 & 21597 & 1 + a2218 - 261248 & 163527.3 & @xmath103662630 & ... & 0 & a2670 - 630319 & 235132.2 & @xmath105104208 & 22380 & 0 + a2218 - 261711 & 163529.1 & @xmath103662719 & 52572 & 1 & a2670 - 630588 & 235134.1 & @xmath105105129 & 22562 & 1 + a2218 - 261251 & 163530.2 & @xmath103661213 & ... & 0 & a2670 - 630322 & 235139.2 & @xmath105103951 & 23398 & 0 + a2218 - 261712 & 163531.1 & @xmath103660345 & 37100 & 2 & a2670 - 630593 & 235139.5 & @xmath105104006 & 23924 & 1 + a2218 - 261252 & 163532.9 & @xmath103662302 & ... & 0 & a2670 - 630324 & 235142.6 & @xmath105103136 & 22674 & 0 + a2218 - 261713 & 163538.9 & @xmath103661645 & 53500 & 2 & a2670 - 630597 & 235143.0 & @xmath105104138 & 24556 & 0 + a2218 - 261259 & 163539.5 & @xmath103661912 & ... & 0 & a2670 - 630599 & 235144.0 & @xmath105104136 & 23128 & 0 + a2218 - 261714 & 163545.7 & @xmath103661000 & ... & 0 & a2670 - 630600 & 235144.2 & @xmath105103031 & 23229 & 1 + a2218 - 261265 & 163550.3 & @xmath103661751 & 54862 & 0 & a2670 - 630328 & 235153.3 & @xmath105103855 & 22918 & 0 + a2218 - 261715 & 163554.1 & @xmath103661404 & 39100 & 2 & a2670 - 630607 & 235153.3 & @xmath105104040 & 27928 & 1 + a2218 - 261268 & 163556.3 & @xmath103662422 & 53852 & 1 & a2670 - 630615 & 235210.2 & @xmath105103337 & 23141 & 1 + a2218 - 261269 & 163600.5 & @xmath103661742 & ... & 0 & a2670 - 630621 & 235226.9 & @xmath105102654 & 23605 & 0 + a2218 - 261271b & 163602.1 & @xmath103662726 & 55300 & 2 & a2670 - 630624 & 235242.7 & @xmath105110249 & 13010 & 1 + a2218 - 261271 & 163604.4 & @xmath103662730 & ... & 0 & a2670 - 630625 & 235246.9 & @xmath105103855 & 22361 & 1 + a2218 - 261272 & 163611.9 & @xmath103661725 & ... & 0 & a2670 - 630627 & 235255.7 & @xmath105102030 & 12640 & 2 + [ tab : galaxy_sample ] crcccc a2218 - 261223 & 3900 & 5600307 & 397 & 70 & 37021 + a2218 - 261705 & 4800 & 5569007 & 538 & 60 & 53027 + a2218 - 261226 & 4500 & 5277710 & 300 & 58 & 35428 + a2218 - 261244 & 900 & 4283814 & 427 & 69 & 41324 + a2218 - 261711 & 4800 & 5255507 & 426 & 69 & 40922 + a2218 - 261268 & 5700 & 5383510 & 402 & 64 & 38226 + a2218 - 261717 & 4500 & 5236007 & 367 & 65 & 35023 + a2218 - 261718 & 6300 & 5158310 & 305 & 66 & 29923 + a2218 - 261291 & 3300 & 5153408 & 398 & 57 & 44125 + a2670 - 630563 & 3600 & 2146310 & 372 & 53 & 45033 + a2670 - 630565 & 900 & 2160110 & 342 & 66 & 37223 + a2670 - 630567 & 1800 & 2221510 & 528 & 70 & 52125 + a2670 - 630568 & 2400 & 2254510 & 453 & 79 & 45422 + a2670 - 630571 & 1200 & 2194410 & 131 & 54 & 15431 + a2670 - 630577 & 2400 & 2315410 & 390 & 61 & 42223 + a2670 - 630317 & 1200 & 2125010 & 323 & 63 & 35227 + a2670 - 630588 & 2400 & 2221510 & 292 & 79 & 35443 + a2670 - 630593 & 900 & 2357710 & 338 & 71 & 41345 + a2670 - 630600 & 3600 & 2288210 & 443 & 90 & 43221 + a2670 - 630607 & 600 & 2758110 & 427 & 72 & 48626 + a2670 - 630615 & 3000 & 2279410 & 316 & 64 & 36345 + a2670 - 630624 & 300 & 1266410 & 123 & 42 & 19737 + a2670 - 630625 & 4200 & 2201410 & 451 & 70 & 49533 + a2670 - 630627 & 600 & 1229310 & 204 & 67 & 22223 + [ tab : spec ] crrrccccc a2218 - 261223 & 3 : & 18 & 7 & 0.6130.049 & 2.2 & 6.0 & 17.35 & -21.390.07 + a2218 - 261705 & 2 : & 16 & 100 & 0.4640.033 & 1.4 & 4.4 & 17.06 & -21.670.09 + a2218 - 261226 & 1 & 10 & 17 & 0.4320.043 & 1.3 & 4.0 & 17.75 & -20.860.10 + a2218 - 261244 & 2 & 3 & 18 & 0.5960.064 & 4.2 & 8.5 & 16.32 & -21.840.10 + a2218 - 261711 & 3 : & 8 & 0 & 0.5810.040 & 1.2 & 3.7 & 17.72 & -20.880.06 + a2218 - 261268 & 5 & 6 & 114 & 0.5400.063 & 1.6 & 5.8 & 17.25 & -21.350.07 + a2218 - 261717 & 3 & 12 & 45 & 0.5350.040 & 1.8 & 5.0 & 17.50 & -21.100.06 + a2218 - 261718 & 6 & 6 & 175 & 0.5660.058 & 1.0 & 3.8 & 17.75 & -20.850.09 + a2218 - 261291 & 2 : & 16 & 67 & 0.4270.021 & 1.1 & 3.6 & 17.59 & -20.970.09 + a2670 - 630563 & 4 & 24 & 130 & 0.3970.042 & 4.3 & 13.0 & 14.66 & -22.000.05 + a2670 - 630565 & 5 & 23 & 83 & 0.5830.036 & 4.0 & 12.0 & 14.82 & -21.850.05 + a2670 - 630567 & 3 & 25 & 135 & 0.5650.048 & 4.3 & 14.3 & 14.95 & -21.780.07 + a2670 - 630568 & 2 & 16 & 159 & 0.7230.045 & 3.2 & 9.8 & 15.18 & -21.580.14 + a2670 - 630571 & 4 & 13 & 43 & 0.2060.112 & 3.3 & 10.8 & 15.12 & -21.640.07 + a2670 - 630577 & 5 & 11 & 51 & 0.4910.025 & 2.6 & 7.7 & 15.02 & -21.740.04 + a2670 - 630317 & 5 & 2 & 105 & 0.6220.063 & 2.0 & 7.2 & 15.39 & -21.370.05 + a2670 - 630588 & 5 & 11 & 97 & 0.7600.009 & 2.9 & 9.6 & 15.25 & -21.510.10 + a2670 - 630593 & 7 & 1 & 110 & 0.6430.037 & 2.3 & 8.1 & 14.88 & -21.880.05 + a2670 - 630600 & 2 : & 11 & 138 & 0.8040.015 & 4.7 & 14.6 & 15.05 & -21.710.14 + a2670 - 630607 & 5 & 4 & 145 & 0.6640.018 & 3.2 & 9.2 & 15.50 & -21.700.05 + a2670 - 630615 & 3 : & 11 & 151 & 0.5430.057 & 5.4 & 15.0 & 15.31 & -21.450.07 + a2670 - 630624 & 7 & 27 & 115 & 0.2550.040 & 2.1 & 7.0 & 15.86 & -19.650.09 + a2670 - 630625 & 3 & 18 & 86 & 0.5400.099 & 2.7 & 8.9 & 14.84 & -21.920.08 + a2670 - 630627 & 1 : & 28 & 80 & 0.5720.043 & 5.7 & 14.5 & 14.65 & -20.800.10 + [ tab : phot ] cccccccc @xmath5 clusters & 70 of 156 & 0.47 & 0.36 & 12.61.2&11.11.2&1.180.07&15.92.7 + @xmath106 sbc & & & & 10.91.3&10.42.1&1.290.09&12.72.8 + @xmath107 sb & & & & 13.81.7&11.61.5&1.100.09&17.94.0 + @xmath108 clusters & 582 of 897 & 0.75 & & 14.00.4&10.10.4&1.110.02 & + field & & & & 12.51.0&10.61.1&1.180.04&22.81.3 + [ tab : results ] ccccccc a2029 & 14 & @xmath1050.100 & @xmath1050.0520.157 & 0.53 ( 0.34 ) & 23400 & @xmath1035501700 + a2218 & 8 & @xmath1030.237 & @xmath1030.2510.138 & 0.34 ( 0.29 ) & 52497 & @xmath10564003300 + a2295 & 10 & @xmath1050.042 & @xmath1050.0260.137 & 0.38 ( 0.32 ) & 24555 & @xmath1032901500 + a2670 & 11 & @xmath1030.003 & @xmath1030.0120.116 & 0.32 ( 0.17 ) & 22494 & @xmath1051201200 + [ tab : tf_results ]

we evaluate several parameters as a function of cluster environment : tully - fisher residuals , equivalent width , and rotation curve asymmetry , shape and extent . although is more easily detectable in galaxies that are located further from the cluster cores , we fail to detect a correlation between extent and galaxy location in those where it is detected , again in contrast with what is found in the clusters of lesser richness . we fail to detect any statistically significant trends for the other parameters in this study . the zero - point in the tully - fisher relation is marginally fainter ( by 1.5 ) than that found in nearby clusters , but the scatter is essentially unchanged .

we have obtained new optical imaging and spectroscopic observations of 78 galaxies in the fields of the rich clusters abell 1413 ( ) , abell 2218 ( ) and abell 2670 ( ) . we have detected line emission from 25 cluster galaxies plus an additional six galaxies in the foreground and background , a much lower success rate than what was found (% ) for a sample of 52 lower - richness abell clusters in the range . we have combined these data with our previous observations of abell 2029 and abell 2295 ( both at ) , which yields a sample of 156 galaxies . we evaluate several parameters as a function of cluster environment : tully - fisher residuals , equivalent width , and rotation curve asymmetry , shape and extent . although is more easily detectable in galaxies that are located further from the cluster cores , we fail to detect a correlation between extent and galaxy location in those where it is detected , again in contrast with what is found in the clusters of lesser richness . we fail to detect any statistically significant trends for the other parameters in this study . the zero - point in the tully - fisher relation is marginally fainter ( by 1.5 ) than that found in nearby clusters , but the scatter is essentially unchanged .

astro-ph0303275

i

we have presented high - sensitivity images of the isolated , edge - on , ` superthin ' spiral galaxy ugc 7321 with a resolution of 24.4 khz and a total bandwidth of 3.052 mhz ( 125 channels ) . in the optical , ugc 7321 exhibits a diffuse stellar disk with little dust obscuration , an extremely small global scale height , and no discernible spheroid component . these features indicate that ugc 7321 is a late - type , lsb spiral seen at high inclination . ugc 7321 thus allows us to explore both the radial and vertical structure of the ism of an ordinary lsb galaxy whose disk has not been recently perturbed . ugc 7321 shows similar global characteristics to other late - type , lsb spirals : it has a gas - rich disk with = @xmath77 d@xmath175 , and an inferred ratio of @xmath6 = 1.0(in solar units ) , which is at the high end of the observed values for spiral galaxies . assuming a stellar mass - to - light ratio similar to the sun ( based on its global @xmath103 color ) ugc 7321 contains a similar fraction of its visible mass in gas as in stars , suggesting that it has been a relatively inefficent star former . ugc 7321 is undetected in the 21-cm radio continuum ( @xmath9mjy ) , which is consistent with its status as a low star formation rate system . assuming a spherical mass distribution , the ratio of the dynamical mass to the total visible mass ( total gas+stars ) is @xmath176 , implying that ugc 7321 is a highly dark - matter dominated galaxy . to first order , the total- distribution of ugc 7321 is regular and symmetric with a slight depression near its center , and extends to approximately 1.5 times the stellar radius . the deprojected surface density is low ( peak value @xmath177 pc @xmath57 ) , which is consistent with other lsb spirals . the outer disk of ugc 7321 displays mild flaring in addition to an `` integral sign '' warp that commences near the edge of the stellar disk . the warped material bends back toward the equatorial plane in the outermost regions . the origin of this warp is still not understood , as ugc 7321 appears to be an isolated galaxy . the nearest known companions to ugc 7321 are two dwarf galaxies , both with projected distances of 0.34 d@xmath8 mpc and @xmath178 which would have had their last encounter with ugc 7321 @xmath93 years ago . in addition , the observations described in this paper rule out the presence of any close gas - rich companions with masses greater than 2.2@xmath179 ( @xmath180 ) within @xmath181 ( 36 d@xmath8 kpc ) . the disk of ugc 7321 appears thicker and more complex than what a single , cold layer confined to the midplane would display . however , the projection effects expected from warping and flaring of the gas layer have to be modelled in order to untangle the three - dimensional structure of the galaxy , which is complicated by its near - edge - on orientation . in addition , we detect filaments of gas extending to distances of up to @xmath1060@xmath0 ( 2.9 d@xmath8 kpc ) from the mid - plane in the total intensity images and in the position - velocity cuts along the minor axis of the disk , which suggests the presence of high - latitude gas . the position - velocity plot along the major axis of ugc 7321 displays a `` figure-8 '' pattern that we interpret as a manifestation of the warping of the disk . in addition , the diagram reveals the presence of a small bar or inner arm . overall ugc 7321 has a rather regular velocity field with only minor perturbations . similarly , the rotation curve of ugc 7321 exhibits only a mild asymmetry about the two sides of the disk . the rotation curve rises rather slowly and linearly throughout the stellar disk , indicating a low central matter density ; this effect can not be attributed to the finite spatial resolution of our observations . a peak rotational velocity of @xmath10110 is reached near 0.9 optical radii , and a slight decline in the rotation curve occurs in the last few measured points . two possible interpretations of this decline are that the dark matter halo of ugc 7321 is truncated near the edge of the disk or that we are seeing a projection effect due to the warping of the layer . @xmath132 ( j2000.0 ) & 12 17 33.8 @xmath182 ( j2000.0 ) & + 22 32 25 hubble type & sd iv distance ( mpc ) & 7 13 p. a. ( deg ) & @xmath183 inclination ( deg ) & @xmath184 @xmath185 & 10.3 @xmath186 ( arcmin ) & 5.6 @xmath186 ( kpc ) & 16.3 d@xmath8 @xmath187 ( kpc ) & @xmath188 d@xmath8 @xmath189 ( pc ) & @xmath190 d@xmath8

we have used the very large array to image the isolated `` superthin '' galaxy ugc 7321 in the line with a spatial resolution of 16 and a spectral resolution of 24 khz ( 5.2 ) . ugc 7321 has a gas - rich disk with d and ( d is the distance to ugc 7321 in units of 10 mpc , the value adopted in this paper ) , and no detectable radio continuum emission ( mjy ) . the global distribution of ugc 7321 is rather symmetric and extends to.5 times the optical radius ( at atoms ) . the ratio of the inferred dynamical mass to the mass in gas and stars is d , implying that ugc 7321 is a highly dark - matter dominated galaxy . 0.6truein 0.2truein

we have used the very large array to image the isolated `` superthin '' galaxy ugc 7321 in the line with a spatial resolution of 16 and a spectral resolution of 24 khz ( 5.2 ) . we have reached a sensitivity of ( 0.36 0.40 ) mjy / beam per channel , which correspond to a column density of ( 8 9) atoms ( ) . ugc 7321 has a gas - rich disk with d and ( d is the distance to ugc 7321 in units of 10 mpc , the value adopted in this paper ) , and no detectable radio continuum emission ( mjy ) . the global distribution of ugc 7321 is rather symmetric and extends to.5 times the optical radius ( at atoms ) . an `` integral sign '' warp is observed in the disk , commencing near the edge of the stellar distribution , and twisting back toward the equatorial plane in the outermost regions . in addition , the position - velocity diagram suggests the presence of a bar or inner arm within from the center . the rotation curve of ugc 7321 is slowly rising ; it reaches its asymptotic velocity of at from the center ( about 0.9 optical radii ) and declines near the edge of the disk . the ratio of the inferred dynamical mass to the mass in gas and stars is d , implying that ugc 7321 is a highly dark - matter dominated galaxy . 0.6truein 0.2truein

1409.8576

i

in many applications from a wide variety of fields , the data to be processed can partially ( or even almost completely ) be affected by severe noise in several phases , e.g. , occlusions during a visual recording or packet losses during transmission in a communication channel . the partial , i.e. , localized , corruptions resulted in the data due to such problems often severely degrade the performance of the target application ; for instance , face recognition or pedestrian detection under occlusion @xcite . in order to reduce the impact of this adverse effect , we develop a complete and novel framework , which efficiently detects , localizes and imputes corruptions by identifying the local anomalies in a given suspicious data instance . we emphasize that neither the existence nor , if exists , the location of a corruption is known in our framework . moreover , the proposed algorithms do not assume a model but operate in a data driven manner . we consider the local corruptions as statistical deviations from the nominal distribution of the uncorrupted ( clean ) observations . to detect and localize corruptions , i.e. , such statistical deviations , we model a corruption as an anomaly due to an external factor ( communication failure in a channel or occluder object in an image ) , which locally overwrites a data instance and moves it outside the support of the nominal distribution . however , corruptions that we consider as examples of anomalies have further specific properties such that ( a ) the corruptions in an instance are confined to unknown intervals along the data attributes , i.e. , localized , and ( b ) not only a corrupted part but also all of its subparts are anomalous . thus , a corruption does not provide an anomaly due to an incompatible combination of normal subparts . based on these properties that accurately model a wide variety of real life applications , we characterize the event of corruption and formulate the corresponding detection / localization as an anomaly detection problem , cf . @xcite . the introduced algorithm applies a series of statistical tests with a pre - specified false alarm rate to the parts of the suspicious instance after extracting the nominal statistics from a reference ( training ) data set of uncorrupted ( clean ) observations . as a result , each part is labeled as anomalous / normal and the local anomalies are identified . these parts are generated and organized through a binary tree partitioning of the data attributes , each node of which corresponds to a part of the suspicious instance . once the nodes ( or parts ) are labeled as anomalous / normal on this tree , the patterns of corruption are identified using the aforementioned characterization to detect and localize corruptions . we point out that this localization procedure transforms the nominal distribution into a multivariate bernoulli distribution with a success probability that precisely coincides with the constant false alarm rate of the local anomaly tests . considering the hierarchy among the binary labels implied by the tree as a directed acyclic graph , the resulting multivariate bernoulli distribution achieves a certain dependency structure . under this condition , we derive the false alarm rate of the proposed framework in detecting the corruptions and show that it is a constant rate , namely , no parameter tuning is required even if the data change . if a corruption is localized , then we impute / replace the affected attributes with the estimates of the underlying unknown true attributes . for this purpose , we additionally develop a novel _ maximum a posteriori _ ( map ) estimator using the score function " defined in @xcite . our estimator exploits the local dependencies among the data attributes , where the locality is encoded in the binary partitioning tree . we point out that the implementation of this map estimator does not load extra computational cost since it utilizes the outputs of our anomaly detection approach , which are computed prior to the imputation phase . furthermore , we also propose a novel distance measure named ranked euclidian distance " as a generalization to the standard euclidean distance , which is used in the course of the labeling of each part as anomalous / normal . the proposed distance measure is compared with the standard euclidean distance in the experiments and shown to be superior in terms of detecting and localizing corruptions . we conduct tests over several well - known machine learning data sets @xcite , which are exposed to severe data corruptions . our experiments indicate that the proposed framework achieves significant improvements after imputation up to @xmath0 in terms of the classification purposes and outperforms the typical approaches . the proposed algorithms are also empirically shown to be robust to varying training phase conditions with strong corruption separation capabilities . in this study , the corrupted attributes are considered to be statistically independent with the underlying unobserved true data , i.e. , corrupted attributes are of no use in estimation of the uncorrupted counterparts . hence , if one knows which attributes are corrupted in an instance , then those attributes can readily be treated as missing data , cf . @xcite . for example , classification and clustering with missing data is a well - studied problem in the machine learning literature . the corresponding studies such as @xcite are related to inference with incomplete data @xcite and generative models @xcite , where bayesian frameworks @xcite are used for inference under missing data conditions . alternatively , pseudo - likelihood @xcite and dependency network @xcite approaches solve data completion problem by learning conditional distributions . in @xcite , the probability density of the missing data is modeled conditioned on a set of introduced latent variables and thereafter , a map based inference is used . however , all of the studies @xcite either assume the knowledge of the location of the missing attributes or impose strong modeling constraints ; as opposed to the model free solutions in this paper . on the other hand , imputation is commonly used as a pre - processing tool @xcite . the mixture of factor analyzers @xcite approach replaces the missing attributes with samples drawn from a parametric density , which models the distribution of the underlying true data . the proposed imputation techniques in @xcite are , whereas , both non - parametric and based on the inference of the posterior densities via certain kernel expansions . on the contrary , the map estimator in this study does not even attempt to estimate the posterior density either in a parametric or non - parametric manner . instead , the introduced method is only based on the sufficient rank statistics . we emphasize that unlike our approach , the incomplete data approaches generally assume the knowledge of the missing attributes , i.e. , they are precisely localized and provided beforehand . for example , the occluded pixels in the event of occlusion of a target object in an image can not be known a priori , which requires a detection and localization step . since the existing studies do not have such a step , an exhaustive list of the occluded pixels as the result of a manual inspection of the missing attributes is required as an input to the algorithms proposed in the corresponding literature . in this regard , our study is the first to jointly handle the issues of detecting / localizing missing attributes , i.e. , corruptions , as well as their imputation in one complete framework . hence , the generic local corruption detection and imputation algorithm of our framework complements the missing data imputation approaches as an additional merit . data imputation and completion is also essential in image processing for handling corrupted images , e.g. , @xcite . generally , a corrupted image is restored by explicitly learning the image statistics @xcite or by using neural networks @xcite . these denoising studies do not attempt to localize corruptions in an image , but treat them as a noise and filter it out using statistical approaches applied to the image globally . even though this is a valid approach for image enhancement , an attempt to correct / enhance an image globally in case of only a localized corruption might be even detrimental since the uncorrupted parts are also affected by global operations . additionally , it is not usually possible to locally impute corrupted portions using denoising approaches . there exist several studies that aim localization as well . studies such as @xcite indicate that occlusion , as an example of corruption , is a common phenomenon and detrimental in pedestrian detection as well as face recognition applications . in this regard , detection of occluded , i.e. , corrupted , visual objects had been previously investigated in a number of studies @xcite . in these studies , occlusion detection is performed using domain specific knowledge ( visual cues ) or external information ( object geometry ) , which , however , are not always available in general data imputation setting . from the machine learning perspective , descriptors are extracted from various parts of the occluded object in @xcite and similarly ; part - based descriptors are weighted with the occlusion measure in @xcite to relieve the corresponding degrading effects . since these approaches do not directly target handling occlusions , i.e. , corruptions , they only provide partial or limited solutions . several other studies propose solutions via extracting occlusion maps , e.g. , @xcite . in @xcite , hog based classification errors ; and in @xcite , template based reconstruction errors are used to generate such an occlusion map . however , both studies assume rigid models and significantly rely on domain specific knowledge ; and in general fail to remain applicable if the data source belongs to another domain . in this study , we assume that data is generic and no domain information is available , yet detection and imputation of corruption is necessary for improving the subsequent processing stages , such as classification . 1 . this study is the first to jointly handle localized data corruptions in one statistical framework that is designed completely model free for the goal of separating a corruption and imputing the affected data attributes . we also provide a false alarm rate ( in detecting corruptions ) analysis of the framework via directed acyclic graphs . 2 . a novel map estimator for imputation and a novel distance measure for corruption localization purposes is proposed . the proposed framework is computationally efficient in the sense that ( i ) it effectively utilizes a binary search for corruption separation , and ( ii ) the computational load due to our map based imputation is insignificant . we propose a characterization for anomalies , e.g. , rarities , incompatible combinations and corruptions , which is a novel notion . + in section [ sec : pd ] , we provide the problem description . we then present our algorithm in section [ sec : thethemethod ] and the associated computational complexity in section [ sec : complexity ] . we report the corruption detection / localization performance of the proposed algorithm as well as the improvement in classification tasks achieved by the imputation in section [ sec : experiments ] . the paper concludes with a discussion in section [ sec : conclusion ] .

once each part is labeled as anomalous vs normal , the corresponding binary patterns over this tree that characterize corruptions are identified and the affected attributes are imputed . under a certain conditional independency structure assumed for the binary patterns , localized corruption , occlusion , map based imputation , anomaly detection .

we introduce a comprehensive and statistical framework in a model free setting for a complete treatment of localized data corruptions due to severe noise sources , e.g. , an occluder in the case of a visual recording . within this framework , we propose i ) a novel algorithm to efficiently separate , i.e. , detect and localize , possible corruptions from a given suspicious data instance and ii ) a _ maximum a posteriori _ ( map ) estimator to impute the corrupted data . as a generalization to euclidean distance , we also propose a novel distance measure , which is based on the ranked deviations among the data attributes and empirically shown to be superior in separating the corruptions . our algorithm first splits the suspicious instance into parts through a binary partitioning tree in the space of data attributes and iteratively tests those parts to detect local anomalies using the nominal statistics extracted from an uncorrupted ( clean ) reference data set . once each part is labeled as anomalous vs normal , the corresponding binary patterns over this tree that characterize corruptions are identified and the affected attributes are imputed . under a certain conditional independency structure assumed for the binary patterns , we analytically show that the false alarm rate of the introduced algorithm in detecting the corruptions is independent of the data and can be directly set without any parameter tuning . the proposed framework is tested over several well - known machine learning data sets with synthetically generated corruptions ; and experimentally shown to produce remarkable improvements in terms of classification purposes with strong corruption separation capabilities . our experiments also indicate that the proposed algorithms outperform the typical approaches and are robust to varying training phase conditions . localized corruption , occlusion , map based imputation , anomaly detection .

1409.8576

c

in this paper , we proposed a comprehensive framework for handling localized and severe data corruptions . the novel contributions of the proposed framework includes ( i ) a first algorithm to jointly detect and localize such corruptions by identifying the local anomalies , ( ii ) a _ maximum a posteriori _ based estimator for imputation ; and a distance measure for corruption separation purposes , ( iii ) computational efficiency via the binary searches and the fast imputations , and ( iv ) a characterization for anomalous observations , e.g. , rarities , incompatible combinations and corruptions . we point out that our algorithm does not assume prior information or a model for the input data and instead , works in a completely data driven way . furthermore , we conducted a false alarm rate analysis and showed that the desired false alarm rate in detecting corruptions can be set independently with the input data . our algorithm is tested against the synthetically generated corruptions in several well - known machine learning data sets and experimentally shown to provide signicant improvements in terms of classication purposes with strong corruption separation capabilities . the proposed algorithms outperform the typical approaches and are robust to varying training phase conditions . stefanos zafeiriou , georgios tzimiropoulos , maria petrou , and tania stathaki , `` regularized kernel discriminant analysis with a robust kernel for face recognition and verification , '' , vol . 3 , pp . 526534 , 2012 . roberto perdisci , guofei gu , and wenke lee , `` using an ensemble of one - class svm classifiers to harden payload - based anomaly detection systems , '' in _ sixth international conference on data mining . ieee icdm06 . _ , 2006 , pp . 488498 . david heckerman , david m. chickering , christopher meek , robert rounthwaite , and carl kadie , `` dependency networks for inference , collaborative filtering , and data visualization , '' , vol . 1 , pp . 4975 , 2001 . nelson h. yung and andrew h. lai , `` detection of vehicle occlusion using a generalized deformable model , '' in _ proceedings of the ieee international symposium on circuits and systems _ , 1998 , . 4 , pp .

, we propose i ) a novel algorithm to efficiently separate , i.e. , detect and localize , possible corruptions from a given suspicious data instance and ii ) a _ maximum a posteriori _ ( map ) estimator to impute the corrupted data . as a generalization to euclidean distance the proposed framework is tested over several well - known machine learning data sets with synthetically generated corruptions ; and experimentally shown to produce remarkable improvements in terms of classification purposes with strong corruption separation capabilities . our experiments also indicate that the proposed algorithms outperform the typical approaches and are robust to varying training phase conditions .

we introduce a comprehensive and statistical framework in a model free setting for a complete treatment of localized data corruptions due to severe noise sources , e.g. , an occluder in the case of a visual recording . within this framework , we propose i ) a novel algorithm to efficiently separate , i.e. , detect and localize , possible corruptions from a given suspicious data instance and ii ) a _ maximum a posteriori _ ( map ) estimator to impute the corrupted data . as a generalization to euclidean distance , we also propose a novel distance measure , which is based on the ranked deviations among the data attributes and empirically shown to be superior in separating the corruptions . our algorithm first splits the suspicious instance into parts through a binary partitioning tree in the space of data attributes and iteratively tests those parts to detect local anomalies using the nominal statistics extracted from an uncorrupted ( clean ) reference data set . once each part is labeled as anomalous vs normal , the corresponding binary patterns over this tree that characterize corruptions are identified and the affected attributes are imputed . under a certain conditional independency structure assumed for the binary patterns , we analytically show that the false alarm rate of the introduced algorithm in detecting the corruptions is independent of the data and can be directly set without any parameter tuning . the proposed framework is tested over several well - known machine learning data sets with synthetically generated corruptions ; and experimentally shown to produce remarkable improvements in terms of classification purposes with strong corruption separation capabilities . our experiments also indicate that the proposed algorithms outperform the typical approaches and are robust to varying training phase conditions . localized corruption , occlusion , map based imputation , anomaly detection .

cond-mat0001235

i

the microscopic understanding of the effect of charge carrier doping on spin@xmath0 antiferromagnetic ( afm ) mott insulators is the central issue of the high - temperature superconducting cuprates . @xcite many puzzling experimental features of these systems @xcite suggest that a fundamental law of nature remains to be recognized . extremely low doping ( @xmath1 charge carriers per site ) leads to a complete destruction of the long - range afm order , and a transition to an unusual non - fermi - liquid metal . this unusual metal becomes superconducting , with the transition temperature @xmath2 strongly dependent on the doping @xmath3 . the maximum @xmath2 is reached for dopings around @xmath4 . for higher dopings the critical temperature decreases to zero , and in the overdoped region a crossover towards a ( non - superconducting ) fermi - liquid takes place . two central questions require resolution . the first one concerns the nature of the charge carriers responsible for this non - fermi - liquid metallic behavior . this is a fundamental issue , since it lies outside the framework of landau s fermi - liquid theory and it necessitates understanding the appearance of non - quasiparticle - like charge carriers in a system of interacting electrons . the second question concerns the nature of strong attractive pairing between these charge carriers , given the purely repulsive interaction between the constituent electrons . in conventional superconductors , the pairing attraction is due to overscreening of the electron - electron coulomb repulsion by the ionic lattice . in the case of the high - temperature superconducting cuprates , it has been suggested @xcite that pairing is an intrinsic property of the electron gas itself mediated by afm spin - fluctuations of the doped system . accordingly , the challenge is to identify a strong attractive force based purely on repulsive coulomb interactions . in this paper , we derive such a force and demonstrate that it leads to d - wave pairing of charge carrying holes introduced by doping a quantum , spin-@xmath5 , mott - hubbard antiferromagnet . the simplest model hamiltonians used to investigate the cuprate physics are the hubbard model and the closely related t - j model . unlike the 1d problem , an exact solution for the 2d hubbard hamiltonian is not known . as a result , various approximations are necessary . although the application of the mean - field theory has been severely criticized in this context , it provides a valuable reference point for incorporating fluctuation effects . moreover , even for the 1d hubbard model , essential features of the exact solution may be recaptured by judiciously incorporating fluctuation and tunneling effects into mean - field theory . @xcite the most straightforward mean - field theory is the hartree - fock approximation ( hfa ) . at half - filling ( @xmath6 ) the hfa predicts an afm mott insulator ground - state . as the system is doped , hfa suggests that charge carrier holes in the afm background assemble in charged stripes , which are quasi - one - dimensional structures . @xcite a large effort has been devoted to studying these charged stripes and relating them to certain features of the cuprates . @xcite recently , a more fundamental investigation of the many - electron problem has suggested the possibility of an alternative model hamiltonian for the cuprate physics . this model hamiltonian , called the spin - flux model , @xcite suggests that the long - range coulomb interaction between spin-@xmath5 electrons leads to qualitative new physics , not apparent in the conventional hubbard model ( see section 2 ) . the results of the hartree - fock study @xcite of this spin - flux model are summarized in section 3 . they suggests that the undoped parent compound is also an afm mott insulator . however , unlike the conventional hubbard model , the one - electron dispersion relations of the afm mean - field in the spin - flux model match those measured experimentally through angle - resolved photo - emission for undoped cuprates . a proper description of the highest occupied electronic states ( as provided by the spin - flux model ) , is crucial to considerations of doping . the spin - flux model and the conventional hubbard model differ dramatically in this regard . at the hf level , the doping holes added to the afm background of the spin - flux model are trapped in the core of antiferromagnetic spin vortices . this composite object ( meron - vortex ) is a bosonic charged collective mode of the many - electron system ( the total spin of the magnetic vortex is zero ) . the reversal of the spin - charge connection provides a microscopic basis for non - fermi - liquid behavior . a magnetic vortex is strongly attracted to an antivortex . this attraction increases logarithmically with the distance between the vortex cores , and is stronger than the unscreened coulomb repulsion between the charge meron - vortex cores . in effect , the increase in coulomb energy between a given pair of holes is more than offset by the lowering in exchange energy between the background electrons as their vortices approach each other from far away . as the inter - vortex distance increases , more and more spins are rotated out of their afm background orientation and the total energy of the system increases . thus , even at the hf level , the spin - flux model provides a fundamental underpinning for the origin of both non - fermi - liquid behavior , and strong pairing between the charge carriers . while providing a good starting point , the hartree - fock approximation also has serious shortcomings . for instance , the ground - state wavefunction in the presence of doping is non - homogeneous ( the static meron - vortices of the spin - flux model , or the charged stripes of the conventional hubbard model , break translational symmetry ) . physically , one expects that these charge carriers can move along the planes , resulting in a wavefunction which preserves the translational symmetry of the original hamiltonian . quantum dynamics of the charge carriers also determines whether the doped ground - state is really a metal . charge carriers in the optimally doped cuprates are quite mobile excitations , although their scattering rates are radically different from electrons in a conventional fermi liquid . a consistent way of treating the quantum dynamics of the charge carriers is provided by the configuration interaction ( ci ) method,@xcite described in section 4 . here , a linear combination of hf wavefunctions is used in order to restore the various broken symmetries . for instance , in a doped system the ci wavefunction is chosen to be a linear combination of hf wavefunctions with the charge carrier localized at different sites . certain types of charge carriers can lower their total energy substantially by quantum mechanically hopping from one site to the next . we tested the accuracy of the ci method against the exact solution @xcite of the one - dimensional hubbard model in reference 3 . in the 1d hubbard model the ci method describes the quantum dynamics of charged domain wall solitons in the afm background . by including these effects as fluctuation corrections to the hartree - fock mean - field theory , the ci method provides excellent agreement with the exact bethe ansatz solution for the ground - state energy of the doped 1d hubbard chain , over the entire @xmath7 range . the ci method also leads to a clear demonstration of the spin - charge separation in 1d . addition of one doping hole to the half - filled antiferromagnetic chain results in the appearance of two different carriers : a charged bosonic domain - wall ( which carries the charge but no spin ) and a neutral spin-1/2 domain wall ( which carries the spin but no charge ) . this study @xcite demonstrates the effectiveness of the ci method . in this paper we use the ci method to investigate dynamics of the charged meron - vortices in the spin - flux model . throughout this paper we exploit and refer to the analogy between the charge excitations of the 1d hubbard model and the 2d spin - flux model , @xcite apparent in the ci approach . @xcite the ci results for the spin - flux model ( presented in section 4 ) confirm that the meron - vortices are very mobile , suggesting that a collection of such mobile bosonic charge carriers is a non - fermi - liquid metal . the ci method also allows us to identify the rotational symmetry of the meron - antimeron pair wavefunction to be d - wave for the most stable pairs . an energetically more expensive metastable s - wave pairing is also possible . the possibility of spin - charge separation in 2d is elucidated . a summary of the results , their interpretation and conclusions is provided in section 5 .

the magnetic exchange energy of the distorted afm background leads to a logarithmic vortex - antivortex attraction which overcomes the direct coulomb repulsion between holes localized on the vortex cores . the ci method systematically describes fluctuation and quantum tunneling corrections to the hartree - fock approximation ( hfa ) .

we investigate a microscopic model for strongly correlated electrons with both on - site and nearest neighbor coulomb repulsion on a 2d square lattice . this exhibits a state in which electrons undergo a `` somersault '' in their internal spin - space ( spin - flux ) as they traverse a closed loop in external coordinate space . when this spin-1/2 antiferromagnetic ( afm ) insulator is doped , the ground state is a liquid of charged , bosonic meron - vortices , which for topological reasons are created in vortex - antivortex pairs . the magnetic exchange energy of the distorted afm background leads to a logarithmic vortex - antivortex attraction which overcomes the direct coulomb repulsion between holes localized on the vortex cores . this leads to the appearance of pre - formed charged pairs . we use the configuration interaction ( ci ) method to study the quantum translational and rotational motion of various charged magnetic solitons and soliton pairs . the ci method systematically describes fluctuation and quantum tunneling corrections to the hartree - fock approximation ( hfa ) . we find that the lowest energy charged meron - antimeron pairs exhibit d - wave rotational symmetry , consistent with the symmetry of the cuprate superconducting order parameter . for a single hole in the 2d afm plane , we find a precursor to spin - charge separation in which a conventional charged spin - polaron dissociates into a singly charged meron - antimeron pair . this model provides a unified microscopic basis for ( i ) non - fermi - liquid transport properties , ( ii ) d - wave preformed charged carrier pairs , ( iii ) mid - infrared optical absorption , ( iv ) destruction of afm long range order with doping and other magnetic properties , and ( v ) certain aspects of angled resolved photo - emission spectroscopy ( arpes ) .

cond-mat0001235

r

the configuration interaction method utilizes a linear combination of judiciously chosen hartree - fock wavefunctions.@xcite in this section , we provide a short review of the relevant hartree - fock results for the spin - flux model . a full comparison between the hfa for the spin - flux model and the conventional hubbard model has been published elsewhere . @xcite one of the most widely used approximations for the many - electron problem is the static hartree - fock approximation ( hfa ) . in this approximation the many - body problem is reduced to one - electron problems in which each electron moves in a self - consistent manner depending on the mean - field potential of the other electrons in the system . while this method is insufficient , by itself , to capture all of the physics of low dimensional electronic systems with strong correlations , it provides a valuable starting point from which essential fluctuation corrections can be included . in particular , we use the hartree - fock method to establish the electronic structure and the static energies of various magnetic soliton structures . in the more general configuration interaction ( ci ) variational wavefunction , the solitons acquire quantum dynamics and describe large amplitude tunneling and fluctuation effects that go beyond mean field theory . in the hf approximation , the many - body wavefunction @xmath58 is decomposed into a slater determinant of effective one - electron orbitals . the one - electron orbitals are found from the condition that the total energy of the system is minimized @xmath59 in order to approximate the ground state of the spin - flux hamiltonian ( [ 1.104 ] ) , we consider a slater determinant trial - wavefunction of the form @xmath60 is the vacuum state , @xmath61 is the total number of electrons in the system and the one - electron states are given by @xmath62 here , the one - particle wave - functions @xmath63 form a complete and orthonormal system . using the wavefunction ( [ 2.3 ] ) in equation ( [ 2.1 ] ) , and minimizing with respect to the one - particle wavefunctions @xmath63 , we obtain the hartree - fock eigen - equations : where @xmath64 are the pauli spin matrices and the charge density , @xmath65 and the spin density , @xmath66 must be computed self - consistently . the notation @xmath67 appearing in ( [ 2.5 ] ) means that the sum is performed over the sites @xmath15 which are nearest - neighbors of the site @xmath12 . the self - consistent hartree - fock equations ( [ 2.5],[2.6],[2.7 ] ) must be satisfied by the occupied orbitals @xmath68 , but can also be used to compute the empty ( hole ) orbitals . the ground - state energy of the system in the hfa is given by @xmath69 where the single particle energies are obtained from ( [ 2.5 ] ) . the approximation scheme described above is called the unrestricted hartree - fock approximation , because we did not impose constraints on the wavefunction @xmath58 which would require it to be an eigenfunction of various symmetry operations which commute with the hamiltonian ( [ 1.104 ] ) . if these symmetries are enforced , the method is called the restricted hartree - fock approximation . we use the unrestricted hfa since it leads to lower energies . the breaking of symmetries in our case implies that electronic correlations are more effectively taken into account . @xcite the restoration of these symmetries is deferred until the ci wavefunction is introduced . in the undoped ( half - filled ) case , the self - consistent hartree - fock equations can be solved analytically for the infinite system , using plane - wave one - particle wave - functions . in the unrestricted hartree - fock approach , doping the system leads to the appearance of inhomogeneous solutions , which break the translational invariance . in this case , we solve the unrestricted self - consistent hartree - fock equations numerically on a finite lattice . starting with an initial spin and charge distribution @xmath70 and @xmath71 , we numerically solve the eigenproblem ( [ 2.5 ] ) and find the hf eigenenergies @xmath72 and wavefunctions @xmath73 . these are used in eqs . ( [ 2.6 ] ) and ( [ 2.7 ] ) to calculate the new spin and charge distribution , and the procedure is repeated until self - consistency is reached . numerically , we define self - consistency by the condition that the largest variation of any of the charge or spin components on any of the sites of the lattice is less that @xmath74 between successive iterations . for the undoped system , the hartree - fock equations ( [ 2.5 ] ) for an infinite system are easily solved . in the cuprates , long - range afm order is experimentally observed . accordingly , we choose a spin distribution at the site @xmath75 of the form @xmath76 , where @xmath77 is the unit vector of some arbitrary direction , while the charge distribution is @xmath78 . in the spin - flux phase , it is convenient to choose a square unit cell , in order to simplify the description of the @xmath48 phase - factors . we make the simplest gauge choice compatible with the spin - flux condition for the @xmath79-matrices , namely that @xmath80 ( see fig . ( [ fig1.55 ] ) ) . this leads to the reduced square brillouin zone @xmath81 . from the hartree - fock equations we find two electronic bands , characterized by the dispersion relations : @xmath82 where each level is four - fold degenerate and @xmath83 are the noninteracting electron dispersion relations in the presence of spin - flux . the hf ground - state energy of the spin - flux afm background is given by ( see eq . ( [ 2.8 ] ) ) : @xmath84 where the afm local moment amplitude is determined by the self - consistency condition ( [ 2.7 ] ) @xmath85 at half - filling the valence band ( @xmath86 ) is completely filled , the conduction band ( @xmath87 ) is completely empty , and a mott - hubbard gap of magnitude @xmath88 opens between the valence and the conduction bands . the ground - state of the undoped spin - flux model is an afm mott insulator . it is interesting to note that the quasi - particle dispersion relation obtained in the presence of the spin - flux ( eq . ( [ 2.27 ] ) ) closely resembles the dispersion as measured through angle - resolved photo - emission spectroscopy ( arpes ) in a compound such as sr@xmath9cuo@xmath9cl@xmath9 @xcite ( see fig . [ fig2.13 ] ) . there is a a large peak centered at @xmath89 with an isotropic dispersion relation around it , observed on both the @xmath90 to @xmath91 and @xmath92 to @xmath93 lines . the spin - flux model in hfa exhibits another smaller peak at @xmath94 which is not resolvable in existing experimental data . this minor discrepancy may be due to next nearest neighbor hopping or other aspects of the electron - electron interaction which we have not yet been included in our model . @xcite the quasi - particle dispersion relation of the conventional hubbard model ( @xmath95 ) has a large peak at @xmath96 on the @xmath90 to @xmath91 line ( see fig . [ fig2.13 ] ) , but it is perfectly flat on the @xmath92 to @xmath93 line ( which is part of the large nested fermi surface of the conventional 2d hubbard model ) . also , it has a large crossing from the upper to the lower band - edge on the @xmath90 to @xmath92 line . this dispersion relation is very similar to that of the @xmath97 model ( see ref . 17 ) . while both the conventional and spin - flux model predict afm insulators at half - filling ( at least at the hf level ) , the spin - flux also provides a much better agreement with the dispersion relations , as measured by arpes . as in the 1d case , @xcite the effect of doping is the appearance of discrete levels deep inside the mott - hubbard gap . these levels are drawn into the gap from the top ( bottom ) of the undoped valence ( conduction ) bands . accordingly , the type of excitations created by doping depends strongly on topology of the electronic structure near the band edges . if we introduce just one hole in the plane , the self - consistent hfa solution is a conventional spin - polaron or `` spin - bag '' ( see fig . this type of excitation is the 2d analog of the 1d spin - polaron . @xcite the doping hole is localized around a particular site , leading to the appearance of a small ferromagnetic core around that site . the spin and charge distribution at the other sites are only slightly affected . in fact , the localization length of the charge depends on @xmath7 , and becomes very large as @xmath98 , since in this limit the mott - hubbard gap closes . for intermediate and large @xmath7 , the doping hole is almost completely localized on the five sites of the ferromagnetic core . the spin - bag is a charged fermion , as can be seen by direct inspection of its charge and spin distributions . this is also confirmed by its electronic structure ( see fig.3b ) . thus , the 2d spin - bag is indeed the analog of the 1d spin - polaron . @xcite the 2d analog of the 1d charged domain wall is the meron - vortex ( see fig . [ fig2.22 ] ) . like the 1d domain - wall , the meron - vortex is also a topological excitation , characterized by a topological ( winding ) number @xmath99 ( the spins on each sublattice rotate by @xmath100 on any closed contour surrounding the center of the meron ) . as such , a single meron - vortex can not be created in an extended afm background with cyclic boundary conditions by the introduction of a single hole ( just as a single isolated charged domain wall can not be created on an afm ( even ) chain with cyclic boundary conditions , by the introduction of a single hole ) . from a topological point of view , this is so because the afm background has a winding number 0 , and the winding number must be conserved , unless topological excitations migrate over the boundary into the considered region . however , excitations can be created in pairs of total topological number 0 . in the 1d case , this means creation of pairs of domain walls , while in 2d this means the creation of vortex - antivortex pairs . from figs . 4(a ) and 4(b ) we can see that the meron - vortex is a charged boson . the total spin of such a configuration is zero , while it carries the doping charge trapped in the vortex core . moreover , from its electronic spectrum ( fig . 4b ) we can see that only the extended states of the valence band are occupied . they are the only ones contributing to the total spin . since only one state is drawn from the valence band into the gap , to become a discrete bound level , it appears that an odd ( unpaired ) number of states remains in the valence band . however , one must remember that for topological reasons , merons must appear in vortex - antivortex pairs . therefore , the valence band has an even number of ( paired ) levels , and the total spin is zero . this argument for the bosonic character of the meron - vortex is identical to that for the charged domain wall in polyacetylene . @xcite unlike in the 1d case,@xcite we can not directly compare the excitation energy of the spin - bag with the excitation energy of the meron - vortex . the reason is that the excitation energy of the latter increases logarithmically with the size of the sample , and therefore an isolated meron - vortex is always energetically more expensive than a spin - bag . however , topology requires that merons and antimerons are created in pairs . the excitation energy of such a meron - antimeron pair is finite , allowing a meaningful comparison between excitation energies of a pair of spin - bags and a meron - antimeron pair . in figs . 5(a ) and 5(b ) we show the self - consistent spin and charge distributions for the lowest energy self - consistent hf configuration found when we add 2 holes to the afm background , in the spin - flux model , for @xmath101 . this configuration consists of a meron and an antimeron centered on second nearest - neighbor sites . as a result of interactions , the cores of the vortices are somewhat distorted . if the vortices were uncharged , vortex - antivortex pair annihilation would be possible . however , for charged vortices , the fermionic nature of the underlying electrons prevents two holes from being localized at the same site , in spite of the bosonic character of the collective excitation . a very interesting feature of this configuration is the strong topological attraction between the vortex and the antivortex . the closer the two cores are to each other , the smaller is the region in which the spins are rotated out of their background afm orientation by the vortices , and therefore the smaller is the excitation energy of the pair . since the holes are localized in the cores of the vortices , this topological attraction between vortices is an effective attraction between holes in the purely repulsive 2d electron system . this effect is unique to the spin - flux phase . in the conventional hubbard model , vortices are not stable excitations . the vortex - antivortex attraction increases as the logarithm of the distance between the cores . therefore , the pair of vortices should remain bound even if full unscreened @xmath102 coulomb repulsion exists between the charged cores , providing a compelling scenario for the existence of strongly bound pre - formed pairs in the underdoped regime . there is another possible self - consistent state for the system with two holes , consisting of two spin - bags far from each other ( such that their localized wave functions do not overlap ) . the excitation energy of such a pair of spin - bags is simply twice the excitation energy of a single spin - bag . when this excitation energy is compared to the excitation energy of the tightly bound meron - antimeron pair , we find that it is higher by @xmath103 ( for @xmath101 ) . in fact , for @xmath104 the hfa predicts that the meron - antimeron pair is the low - energy charged excitation , while for @xmath105 , the spin - bag is the low - energy charge carrier . this is analogous with the situation in 1d , where the spin - bag was predicted to be the low energy excitation for @xmath106 , in the hfa . @xcite as in 1d , however , we expect that this conclusion will be drastically modified once the charged solitons are allowed to move along the planes and the lowering of kinetic energy through translations is also taken into consideration . we complete this review of the hf results by pointing out that the strong analogy between the 1d hubbard model and the 2d spin - flux model is due to the similarity between the electronic structures at zero doping . as seen from fig.[fig2.13 ] , the 2d spin - flux model has isotropic dispersion relations about the @xmath96 point . this acts as a fermi point for the noninteracting system as it does in the 1d system . the two empty discrete levels drawn deep inside the mott - hubbard gap in the presence of the meron - vortex split from the @xmath96 peaks of the electron dispersion relation . the different topology of the large nested fermi surface of the conventional hubbard model leads to instability of the meron - antimeron pair . in fact , in the conventional hubbard model doping holes assemble in charged stripes , as opposed to the liquid of meron - antimeron pairs which is the low - energy state of the doped spin - flux model .

we use the configuration interaction ( ci ) method to study the quantum translational and rotational motion of various charged magnetic solitons and soliton pairs . we find that the lowest energy charged meron - antimeron pairs exhibit d - wave rotational symmetry , consistent with the symmetry of the cuprate superconducting order parameter . for a single hole in the 2d afm plane , we find a precursor to spin - charge separation in which a conventional charged spin - polaron dissociates into a singly charged meron - antimeron pair .

we investigate a microscopic model for strongly correlated electrons with both on - site and nearest neighbor coulomb repulsion on a 2d square lattice . this exhibits a state in which electrons undergo a `` somersault '' in their internal spin - space ( spin - flux ) as they traverse a closed loop in external coordinate space . when this spin-1/2 antiferromagnetic ( afm ) insulator is doped , the ground state is a liquid of charged , bosonic meron - vortices , which for topological reasons are created in vortex - antivortex pairs . the magnetic exchange energy of the distorted afm background leads to a logarithmic vortex - antivortex attraction which overcomes the direct coulomb repulsion between holes localized on the vortex cores . this leads to the appearance of pre - formed charged pairs . we use the configuration interaction ( ci ) method to study the quantum translational and rotational motion of various charged magnetic solitons and soliton pairs . the ci method systematically describes fluctuation and quantum tunneling corrections to the hartree - fock approximation ( hfa ) . we find that the lowest energy charged meron - antimeron pairs exhibit d - wave rotational symmetry , consistent with the symmetry of the cuprate superconducting order parameter . for a single hole in the 2d afm plane , we find a precursor to spin - charge separation in which a conventional charged spin - polaron dissociates into a singly charged meron - antimeron pair . this model provides a unified microscopic basis for ( i ) non - fermi - liquid transport properties , ( ii ) d - wave preformed charged carrier pairs , ( iii ) mid - infrared optical absorption , ( iv ) destruction of afm long range order with doping and other magnetic properties , and ( v ) certain aspects of angled resolved photo - emission spectroscopy ( arpes ) .

cond-mat0001235

m

the essence of the configuration interaction ( ci ) method is that the ground - state wavefunction , for a system with @xmath61 electrons , is not represented by just a single @xmath107 slater determinant ( as in the hfa ) , but a judiciously chosen linear combination of such slater determinants . @xcite given the fact that the set of all possible slater determinants ( with all possible occupation numbers ) generated from a complete set of one - electron orbitals constitute a complete basis of the @xmath61-particle hilbert space , our aim is to pick out a subset of slater determinants which captures the essential physics of the exact solution . consider the ci ground - state wavefunction given by @xmath108 is a distinct @xmath107 slater determinant and the coefficients @xmath109 are chosen to satisfy the minimization principle : @xmath110 this leads to the system of ci equations @xmath111 where @xmath112 is the energy of the system in the @xmath113 state , @xmath114 are the matrix elements of the hamiltonian in the basis of slater determinants @xmath115 , and @xmath116 are the overlap matrix elements of the slater determinants ( which are not necessarily orthogonal ) . the ci solution is easily found by solving the linear system of equations ( [ 3.3 ] ) , once the basis of slater determinants @xmath117 is chosen . if we denote by @xmath118 the @xmath119 one - electron occupied orbitals of the slater determinant @xmath120 , these matrix elements are given by : @xmath121{ccc } \beta_{1,1}^{nm } & ... & \beta_{1,n_e}^{nm } \\ \vdots & & \vdots\\ \beta_{n_e,1}^{nm } & ... & \beta_{n_e , n_e}^{nm } \\ \end{array } \right|\ ] ] the matrix elements of the hamiltonian ( [ 1.104 ] ) can be written as : @xmath122 where the expectation values of the hopping and on - site interaction terms are : @xmath123{ccccc } \beta_{1,1}^{nm}& ... &t_{1,p}^{nm}& ... &\beta_{1,n_e}^{nm}\\ \vdots & & \vdots & & \vdots\\ \beta_{n_e,1}^{nm}& ... &t_{n_e , p}^{nm}& ... &\beta_{n_e , n_e}^{nm } \\ \end{array } \right|\ ] ] and @xmath124{ccccccc } \beta_{1,1}^{nm } \!\!& ... \!\!&u_{1,p_1}^{nm}(i)\!\ ! \!\!&d_{1,p_2}^{nm}(i)\!\!& ... \!\ ! & \beta_{1,n_e}^{nm } \\ \vdots & & \vdots & & \vdots & & \vdots\\ \beta_{n_e,1}^{nm}\!\!& ... \!\!&u_{n_e , p_1}^{nm}(i)\!\ ! & ... \!\!&d_{n_e , p_2}^{nm}(i)\!\!& ... \!\ ! & \beta_{n_e , n_e}^{nm}\\ \end{array } \right|_{.}\ ] ] here , @xmath125 @xmath126 @xmath127 and @xmath128 we now consider the specific choice of the slater determinant basis @xmath115 . strictly speaking , one may choose an optimized basis of slater determinants from the general variational principle : @xmath129 however , implementation of this full trial - function minimization scheme ( also known as a multi - reference self - consistent mean - field approach @xcite ) is numerically cumbersome even for medium - sized systems . instead , we select the slater determinant basis @xmath115 from the set of broken symmetry , unrestricted hartree - fock wavefunctions ( [ 2.3 ] ) , their symmetry related partners and their excitations . clearly , ( [ 2.3 ] ) satisfies ( [ 3.4 ] ) by itself , provided that the @xmath130 coefficients corresponding to the other slater determinants in eq.([3.1 ] ) are set to zero ( see eq . ( [ 2.1 ] ) ) . since this unrestricted hf wavefunction is not translationally invariant ( the doping hole is always localized somewhere on the lattice ) , we can restore the translational invariance of the ci ground - state wavefunction by also including in the basis of slater determinants all the possible lattice translations of this unrestricted hf wavefunction . furthermore , if the self - consistent configuration is not rotationally - invariant ( e.g. a meron - antimeron pair ) , all possible rotations must be performed as well . by rotation we mean changing the relative position of the meron and antimeron while keeping their center of mass fixed . clearly , all the translated hf slater determinants lead to the same hf ground - state energy @xmath131 as defined by eq . ( [ 2.8 ] ) . the ci method lifts the degeneracy between states with the hole - induced configuration localized at different sites ( see eq . ( [ 3.3 ] ) ) , thereby restoring translational invariance . we may identify the lowering in the total energy due to the lifting of this degeneracy as quantum mechanical kinetic energy of deconfinement which the doping - induced configuration saves through hopping along the lattice . in addition , quantum fluctuations in the internal structure of a magnetic soliton can be incorporated by including the lowest order excited state configurations of the static hartree - fock energy spectrum . such wavefunctions are given by @xmath132 , where @xmath133 labels an excited particle state and @xmath134 labels the hole which is left behind ( see eq . ( [ 2.3 ] ) ) . once again , all possible translations ( and non - trivial rotations ) of this `` excited '' configurations must be included in the full ci wavefunction . these additions can describe changes in the `` shape '' of the soliton as it undergoes quantum mechanical motion along the plane . the ci method is described in more detail in reference 3 , where it is used to study the 1d hubbard chain in order to gauge its accuracy by comparing its results with the exact bethe ansatz solution . we showed that the ci method recaptures the essential physical features of the exact solution of the 1d hubbard chain , such as spin - charge separation , as well as leading to remarkable agreement of ground state energies of doped chains for all values of @xmath7 . the main difference between the 1d case and the 2d case is the computation time required . the computation time for one matrix element @xmath135 scales roughly like @xmath136 , where @xmath137 is the number of sites . the number of configurations included in the ci set scales as @xmath138 when @xmath139 solitons are present . for both an @xmath137-site chain and an @xmath140 lattice , the hf `` bulk '' limit is reached for @xmath141 . in the 1d case @xcite we used chains with @xmath142 , and numerical calculations can be easily performed . however , in 2d the smallest acceptable system has 100 sites , leading to an enormous increase in the computation time . nevertheless , our sample of results in 2d suggest a simple and clear physical picture which we describe below . the charged spin - bag carries a spin of 1/2 . let @xmath143 be the hf determinants for the spin - bag centered at any two nearest neighbor sites , respectively , and let @xmath144 be the total spin operator in the @xmath145-direction . then , @xmath146 . since the hubbard hamiltonian commutes with @xmath147 , it follows that @xmath148 . from the ci equation ( [ 3.3 ] ) we conclude that there is no mixing between states with different total spin . as a result , it is enough to include in the ci set only those configurations with the spin bag localized on the same magnetic sublattice . let us denote by @xmath149 the initial static hartree - fock configuration , and by @xmath150 the configuration obtained through its translation by @xmath151 sites in the x - direction and @xmath152 sites in the y - direction ( cyclic boundary conditions are imposed ) . the condition that only configurations on the same sublattice are included means that @xmath153 must be an even number , and the cyclic boundary conditions mean that @xmath154 , for a @xmath155 lattice . as explained in detail in the 1d analysis , @xcite mixing configurations with the charged spin - bag localized at different sites and then subtracting out the contribution of the undoped afm background allows us to calculate the dispersion band of the spin - bag itself : @xmath156 here , the total energy of the lattice with the spin - bag @xmath157 and the ci wave - function @xmath158 ) . the finite size of the lattice and cyclic boundary conditions restricts the calculation to @xmath159-points of the form @xmath160 , where @xmath161 is any pair of integer numbers . as usual , only @xmath159-points inside the first brillouin zone need to be considered . an analysis of the dependence of the spin - bag dispersion relation @xmath162 on the size @xmath140 of the lattice is shown in fig . [ fig3.20 ] , for the conventional hubbard model ( upper panel ) and spin - flux model ( lower panel ) , and @xmath101 . we used 6x6 , 8x8 , 10x10 and ( only for the spin - flux model ) 12x12 lattices . the dispersion relation is plotted along lines of high symmetry of the full brillouin zone . for comparison , we also show the excitation energy @xmath163 obtained in the static hfa as a full line . for both models , we see that the spin - bag dispersion band is almost converged , even though we used quite small lattices . the convergence is somewhat slower in the spin - flux case , as seen most clearly at the ( 0,0 ) point . although the values obtained from the four lattices all differ at ( 0,0 ) , the extremum values correspond to the 6x6 and the 8x8 lattices , while the values for the 10x10 and 12x12 lattices are indistinguishable . we conclude that the fit ( [ 3.40 ] ) is legitimate . from fig . [ fig3.20 ] we also see that the dispersion relations for the spin - bag in the two different models are very different . the dispersion relations over the full 2d brillouin zone are shown in fig . [ fig3.21 ] , and they are seen to mimic the electronic dispersion relation of the underlying undoped afm background , shown in fig . [ fig2.13 ] . this is consistent with the quasi - particle nature of this charged spin-1/2 spin - bag . in the conventional hubbard model , the undoped afm background has a large nested fermi surface along the @xmath92 to @xmath93 line , and it is exactly along this line that the spin - bag dispersion band has a minimum . similarly , the lowest energy of the spin - bag of the spin - flux model is at @xmath89 , corresponding to the fermi points of the underlying undoped spin - flux afm background . the extra kinetic energy @xmath164 saved by the spin - bag through quantum hopping is @xmath165 in the conventional model and @xmath166 in the spin - flux model ( for @xmath101 ) . since the spin - bag is confined to one magnetic sublattice , it must tunnel two lattice constants to the next allowed site . consequently , the energy gained through hopping ( of order @xmath167 ) is small . this is displayed , for the spin - bag of the spin - flux model , in fig . [ fig3.24 ] , where we plot the lowering in kinetic energy of the deconfined spin - bag @xmath168 as a function of @xmath169 . a similar dependence for the spin - bag of the conventional hubbard model is presented elsewhere . @xcite as in the 1d case , we conclude that the spin - bag in 2d is a rather immobile quasiparticle - like excitation . in the 1d model it is energetically favorable for the immobile spin - bag to decay into a charged bosonic domain wall and a neutral fermionic domain wall , resulting in spin - charge separation . @xcite the analog of the 1d charged bosonic domain wall is the 2d charged bosonic meron - vortex of the spin - flux model . if the spin - bag decays into a charged meron - vortex , a magnetic antivortex must also be created for topological reasons . unlike the pair of domain walls in the 1d case , the vortex - antivortex pair is tightly bound by a topological binding potential that increases as the logarithm of the vortex - antivortex separation . therefore , we expect that the doping charge is shared between the two magnetic vortices . one technical problem for testing this hypothesis is that such a configuration ( a vortex - antivortex pair sharing one doping hole ) is not self - consistent at the static hartree - fock level . in the static approximation we require two doping holes to stabilize two vortex cores and create a meron - antimeron pair . we can , however , construct a trial wavefunction to describe the singly - charged vortex - antivortex pair , by adding one electron in the first empty state of the self - consistent doubly - charged meron - antimeron configuration . the first empty levels of the meron - antimeron pair are the localized levels bound in the vortex cores , two for each vortex ( see fig . because of degeneracy between the two lower discrete levels of the pair , we have in fact two distinct trial wave - functions , obtained by adding one electron in either of these two lower localized gap electronic states of a self - consistent meron - antimeron pair . these wavefunctions are not invariant to rotations ( see fig . [ fig2.25 ] ) . therefore we must include in the ci set of slater determinants the configurations obtained through @xmath170 rotations of the vortex - antivortex pair about its fixed center of mass in addition to translated configurations . as a result , we have a total of @xmath171 configurations describing the singly charged vortex - antivortex pair localized at all possible sites with all possible orientations about the center of mass . we performed this ci analysis for a @xmath172 lattice and @xmath101 . the hf energy of a simple static spin - bag is @xmath173 ( measured with respect to the hf energy of the undoped afm background , equal to @xmath174 ) . the energy of the static singly - charged vortex - antivortex pair is @xmath175 . thus , we see that because this singly - charged pair trial wavefunction is not self - consistent , in the static case this configuration is energetically much more costly than the self - consistent spin - bag configuration . however , if we allow for quantum motion of these configurations , the situation changes dramatically . performing the ci analysis for the set of all possible translated spin - bag configurations , we find that the energy of the spin - bag is lowered to @xmath176 . performing the ci analysis for the set of all translated and rotated singly - charged vortex - antivortex pairs we find that this configuration s energy is lowered to @xmath177 . this shows that the vortex - antivortex pair has lowered its translational and rotational kinetic energy by almost @xmath178 , thereby becoming the low - energy charge carrier . this large number is not surprising , since unlike the spin - bag , the vortex - antivortex pair is not constrained to motion on one magnetic sublattice . as a result , such configurations lower their kinetic energy by an amount on the scale of @xmath24 , as opposed to @xmath167 for the spin - bag configuration . for larger @xmath7 values this effect is even more pronounced . we conclude that these results strongly support the hypothesis of spin - bag dissociation into a much more mobile singly - charged vortex - antivortex pair , analogous to the 1d spin - bag dissociation into a pair of a charged bosonic domain - wall and a neutral fermionic domain - wall . @xcite unlike in the 1d case , however , we do not have distinct charge and spin carriers for the composite excitation . instead the spin and the charge are shared equally between the vortex and the antivortex . if , on the other hand , there was a mechanism whereby the vortices became unbound , complete spin - charge separation could occur , in which one vortex traps the hole ( and is therefore a charged meron ) and the other vortex carries the spin , in a lotus - flower@xcite ( or undoped magnetic meron ) configurations . in the absence of the corresponding self - consistent static hf configuration we are not able to settle this question . at very low doping , the strong vortex - antivortex topological attraction binds the spin and charge together . this is different from the 1d case , where the absence of long - range interactions between the domain - walls allow for a complete spin - charge separation at any doping and even at zero temperature . this scenario opens a new avenue for research into how the system evolves with doping . if each hole is dressed into a singly - charged vortex - antivortex pair , when two such pairs overlap it is possible that both doping charges move to the same pair , creating a meron - antimeron pair of charged bosons . such pre - formed charge pairs may condense into a superconducting state at low temperatures . the other uncharged vortex - antivortex pair may either collapse and disappear ( this is likely to happen at low - temperatures ) or remain as a magnetic excitation of the system ( at higher temperatures ) , mediating the destruction of the long - range afm order , the renormalization of the spin - wave spectrum and the opening of the spin pseudogap . from the static hf analysis we found that the most stable static self - consistent configuration with two doping holes added to the afm background of the spin - flux model is the meron - antimeron pair , for @xmath179 . at larger @xmath7 , two charged spin - bags become more stable , in the static hf approximation . this is in close analogy to the prediction that the spin - bag is energetically more favorable than the static charged domain - wall for @xmath180 , in the hfa of the 1d hubbard model . @xcite however , in the 1d case the charged domain - wall is considerably more mobile than the charged spin bag , gaining a kinetic energy of the order of @xmath24 as opposed to @xmath167 energy gained by the spin - bag . as a result , when this kinetic energy of deconfinement is taken into account within the ci method , the charged domain - wall is found to be the relevant charged excitation for all values of @xmath7 . a similar picture emerges in the 2d case , because the meron - vortices are much more mobile than the spin - bags . for the 2d system , we have shown that the charged spin - bag has very similar behavior to the 1d charged spin - bag . the analog of the 1d charged bosonic domain - wall is the 2d charged bosonic meron - vortex . we now consider the properties of the doubly - charged meron - antimeron pair . all the numerical results quoted in the rest of this section refer to a meron - antimeron pair on a 10x10 lattice , in the spin - flux model with @xmath101 . as already discussed , the meron - antimeron pair is not rotationally invariant . we can find the rotational kinetic energy saved by the pair as it rotates about its center of mass . in the present case , only 4 configurations need to be included , corresponding to the four possible self - consistent arrangements of the meron and antimeron about their fixed center of mass ( see fig . [ fig3.28 ] ) . simple rotation by @xmath170 of the one - particle orbitals @xmath181 about the center of mass is not , however , sufficient to generate the rotated configurations . first of all , the @xmath170 rotation also changes the spin - flux parameterization . if the spin - flux of the initial configuration is @xmath57 , a @xmath170 rotation leads to a state corresponding to the rotated configuration @xmath182 . thus , following the @xmath170 rotation , a unitary transformation must be performed in order to restore the initial spin - flux parameterization . for the case cited above , this simply implies the change in the one - particle orbitals @xmath183 for all sites @xmath184 which are a type 2 site of the unit cell , in other words sites with even @xmath185 and odd @xmath186 ( also , see fig . [ fig1.55 ] ) . the second observation is that the rotation by @xmath170 also changes ( flips ) all the spins of the afm background surrounding the pair . thus , an extra @xmath187 rotation about an axis perpendicular to the lattice plane is necessary to restore the alignment of the afm background . following these transformations it is straightforward to generate the slater determinants @xmath188 and @xmath189 from the initial self - consistent hf meron - antimeron pair described by @xmath190 . the ci method can be used to find the rotational energy saved by superposing these rotated meron - antimeron configurations . the lowest ci energy found is @xmath191 below the energy of the static pair , and corresponds to d - wave symmetry . by this we mean that the coefficients @xmath109 multiplying the 4 rotated states in the ci wave - function @xmath192 . translation of a pair over the whole lattice can also be investigated . since the pair does not carry any spin , all possible translations must be included ( there is no restriction to same magnetic sublattice configurations ) . this leads to a total of @xmath193 possible configurations for a @xmath140 lattice . again , when various configurations are generated from the initial self - consistent hf meron - antimeron state @xmath194 , care must be taken to preserve the same spin - flux parameterization and the same afm background orientation . this can be achieved performing transformations similar to the ones described above . as a result of performing the ci method on the set of translated states , we find the dispersion relation of the ( un - rotated ) meron - antimeron pair . this is shown in fig . [ fig3.25 ] . in this plot we show the total energy of the lattice with the moving meron - antimeron pair , as a function of the total momentum of the pair . quantum hopping of the meron - antimeron pair lowers its total energy by an extra @xmath195 . two other interesting features are observed in fig . [ fig3.25 ] . the first one is that the dispersion relation of this rigidly polarized pair is not invariant to rotations by @xmath170 , as expected . more important is that the minima of the dispersion relation occur at the @xmath91 points . since the momentum of the pair is twice the momentum of either the meron or the antimeron , this is consistent with the fact that in their lowest energy state , both the meron and the antimeron have momenta of @xmath89 , in the spin - flux model . the doubling of the size of the brillouin zone is also a direct consequence of this doubling of total momentum ( for comparison with undoped dispersion relation , see fig . [ fig2.13 ] ) . however , to obtain the true energy of the charged pair , we must mix all @xmath196 rotated and translated meron - antimeron configurations . all have the same static hf energy and are equally important in the ci method . let @xmath197 denote the initial self - consistent static hartree - fock meron - antimeron configuration , and @xmath198 denote the configuration obtained through translation by @xmath151 sites in the x - direction and @xmath152 sites in the y - direction , as well as a rotation by an angle of @xmath199 of the pair about its center of mass . here @xmath200 and @xmath201 ( cyclic boundary conditions are imposed ) . the ci wavefunctions are then given by : @xmath202 the dispersion relation @xmath203 obtained from this complete set is shown in fig . [ fig3.26 ] . the reference point is the hf energy of the self - consistent meron - antimeron pair @xmath204 . thus , we see that the total kinetic energy saved by the freely moving and rotating meron - antimeron pair is @xmath205 , when the total momentum of the pair is @xmath91 . this energy equals the sum @xmath206 of rotational and translational kinetic energies found before ( the number of significant figures indicates the estimated accuracy of the computational method ) . the rotational invariance of the dispersion band is also restored . besides the absolute minima about the @xmath91 points , there is a more shallow minimum region about the @xmath90 point . the rotational symmetry of the meron - antimeron pair wavefunction , defined by the coefficients @xmath207 , is a function of the total momentum carried by the pair , as shown in fig . [ fig3.27 ] . the absolute minima points @xmath91 and the area around them correspond to pairs with d - wave symmetry . by this , we mean that the coefficients @xmath208 have the form @xmath209 with @xmath210 , i.e. @xmath211 . the core area , about the local minimum @xmath90 point , corresponds to s - wave symmetry . in this region the coefficients @xmath212 again satisfy eq . ( [ 3.101 ] ) , but for @xmath213 , i.e. @xmath214 . the intermediary area appears to be a mixture of different @xmath215 values . a simple decomposition of the form ( [ 3.101 ] ) is no longer possible . instead , a sum of such terms corresponding to different @xmath215 values is required . since we only have rotations by @xmath170 , a unique identification of the composite symmetry is not possible . moreover , the energy of the states in this intermediary area is at the top of the dispersion band . in order to find the correct ci states for energies well above the static hf value ( i.e. larger than zero , in this case ) we must add to the ci set the first set of excited hf states . for a meron - antimeron pair , excitation of an electron from the valence band onto the empty localized levels inside the mott - hubbard gap costs about 1.5 t of energy , for @xmath101 , so such states should contribute significantly in the ci states with positive energies and modify their dispersion and symmetry ( for this reason , we do not show the upper three high - energy bands in fig . [ fig3.26 ] ) . consequently , both the energy and the symmetry of the states in the intermediary area may be modified from the ones shown in fig . [ fig3.26 ] . however , the minima at @xmath91 and @xmath90 are at energies well below zero . their energies and rotational symmetry are unaffected by additions of higher energy configurations to the ci set . the fact that we obtain two distinct minima is not very surprising . as argued before , we expect that individual merons and antimerons are created with momenta of @xmath216 . as a result , two different couplings are possible . a @xmath217 ) meron can pair with a @xmath217 ) antimeron , creating a pair of total momentum @xmath91 . this is the most stable coupling , leading to the lowest possible energy of @xmath218 below the static hf energy . this pair has d - wave symmetry . the second possible coupling is between a @xmath217 ) meron and a @xmath219 ) antimeron . this pair has a total momentum of @xmath90 , and s - wave symmetry . however , this coupling is less strong . for the @xmath101 case considered , the energy of the s - wave @xmath90 pair is @xmath220 above the energy of the d - wave @xmath91 pair . the existence of both d - wave and s - wave pairing , and the dominance of the d - wave pairing , have been established experimentally for the high - t@xmath23 cuprates . @xcite we are not aware of any other microscopic theory that predicts the two types of pairing to appear in different regions of the brillouin zone . the total kinetic energy saved by the meron - antimeron pair through quantum hopping and rotation is of order @xmath24 , as expected , since the pair is not restricted to one magnetic sublattice and tunneling is not required for motion . consequently , we expect that the energy saved by the meron - antimeron pair for larger values of @xmath7 is comparably large . on the other hand , the energy saved by the spin bag through tunneling motion scales like @xmath167 . in fact , we argued that a spin - bag may dissociate into a singly - charged vortex - antivortex pair in order to enhance its mobility . however , even if dissociation does not occur , the kinetic energy saved by a pair of spin - bags is significantly smaller than the kinetic energy saved by the meron - antimeron pair . this shows that for @xmath101 the meron - antimeron pair is even more favorable energetically than the hfa predicts and suggests that the @xmath7 range where meron - antimeron pair formation occurs may extend well beyond the @xmath221 limit found within the hfa . @xcite in the 1d case , the range of stability of the charged domain wall versus the charged spin - bag is extended ( from the hf prediction of @xmath222 ) to all @xmath7 range.@xcite a numerical analysis is needed to determine if the limit is extended to infinity in the 2d case as well .

when this spin-1/2 antiferromagnetic ( afm ) insulator is doped , the ground state is a liquid of charged , bosonic meron - vortices , which for topological reasons are created in vortex - antivortex pairs . this model provides a unified microscopic basis for ( i ) non - fermi - liquid transport properties , ( ii ) d - wave preformed charged carrier pairs , ( iii ) mid - infrared optical absorption , ( iv ) destruction of afm long range order with doping and other magnetic properties , and ( v ) certain aspects of angled resolved photo - emission spectroscopy ( arpes ) .

we investigate a microscopic model for strongly correlated electrons with both on - site and nearest neighbor coulomb repulsion on a 2d square lattice . this exhibits a state in which electrons undergo a `` somersault '' in their internal spin - space ( spin - flux ) as they traverse a closed loop in external coordinate space . when this spin-1/2 antiferromagnetic ( afm ) insulator is doped , the ground state is a liquid of charged , bosonic meron - vortices , which for topological reasons are created in vortex - antivortex pairs . the magnetic exchange energy of the distorted afm background leads to a logarithmic vortex - antivortex attraction which overcomes the direct coulomb repulsion between holes localized on the vortex cores . this leads to the appearance of pre - formed charged pairs . we use the configuration interaction ( ci ) method to study the quantum translational and rotational motion of various charged magnetic solitons and soliton pairs . the ci method systematically describes fluctuation and quantum tunneling corrections to the hartree - fock approximation ( hfa ) . we find that the lowest energy charged meron - antimeron pairs exhibit d - wave rotational symmetry , consistent with the symmetry of the cuprate superconducting order parameter . for a single hole in the 2d afm plane , we find a precursor to spin - charge separation in which a conventional charged spin - polaron dissociates into a singly charged meron - antimeron pair . this model provides a unified microscopic basis for ( i ) non - fermi - liquid transport properties , ( ii ) d - wave preformed charged carrier pairs , ( iii ) mid - infrared optical absorption , ( iv ) destruction of afm long range order with doping and other magnetic properties , and ( v ) certain aspects of angled resolved photo - emission spectroscopy ( arpes ) .

1108.2281

i

in this paper we begin the systematic study of legendrian graphs in @xmath6 with the standard contact structure . these are embedded graphs that are everywhere tangent to the contact planes . legendrian graphs have appeared naturally in several important contexts in the study of contact manifolds . they are used in giroux s proof of existence of open book decompositions compatible with a given contact structure , see @xcite . legendrian graphs also appear in the proof of eliashberg s and fraser s result which says that in a tight contact structure the unknot is determined up to legendrian isotopy by the invariants @xmath7 and @xmath8 , see @xcite . yet no study of legendrian graphs , until now , has been undertaken . we remedy this by establishing the foundations for what we expect will be a very rich field . a _ spatial graph _ is an embedding of a graph into @xmath9 an _ abstract graph _ is a set of vertices together with a set of edges between them , without any specified embedding . we will throughout this paper refer to an abstract graph as simply a graph . in section [ graphs ] , we show there is no obstruction to having a legendrian realization of any spatial graph . we extend the classical invariants thurston - bennequin number , @xmath7 , and rotation number , @xmath8 , from legendrian knots to legendrian graphs . in @xcite , mohnke proved that the borromean rings and the whithead link can not be represented by legendrian links of _ trivial unknots _ , that is unknots with @xmath1 and @xmath2 . the trivial unknot is the one unknot among all unknots that attains the maximal thurston - bennequin number of its topological class . as an application of our invariants , we ask which graphs admit legendrian embeddings with all cycles trivial unknots . in section [ maxtb ] , we give a full characterization of these graphs , in the form of the following : a graph @xmath10 admits a legendrian embedding in @xmath0 with all its cycles trivial unknots if and only if @xmath10 does not contain @xmath3 as a minor . the proof of this theorem relies partly on the fact that the trivial unknot has an odd thurston - bennequin number , that is @xmath1 . we prove the reverse implication in more generality , for @xmath11 which represents the set of topological knot classes with odd maximal thurston - bennequin number . let @xmath10 be a graph that contains @xmath3 as a minor . there does not exist a legendrian realization of @xmath10 such that all its cycles are knots in @xmath12 realizing their maximal thurston - bennequin number . it is known that certain legendrian knots and links are determined by the invariants @xmath7 and @xmath8 : the unknot ( see @xcite ) , the torus knots and the figure eight knot ( see @xcite ) , the links consisting of an unknot and a cable of that unknot ( see @xcite ) . in section [ classification ] , we ask what types of spatial graphs are classified up to legendrian isotopy by the pair @xmath4 . we show that the pair @xmath4 does not classify graphs which contain both cycles and either cut vertices or cut edges , independent of the chosen topological class . this means that even uncomplicated graphs carry more information than the set of knots represented by their cycles . in order to have a classification by the pair @xmath4 , we must first narrow to a specific topological class . we investigate _ topological planar graphs _ , that is an embedded graph that is ambient isotopic to a plannar embedding . not to be confused with _ planar graphs _ , which refers to an abstract graph that has a planar embedding . in the case of the handcuff graph , for topological planar graphs , we prove there are exactly four legendrian realizations for each pair @xmath4 . for the lollipop graph , for topological planar graphs , we prove there are exactly two legendrian realizations for each pair @xmath4 . the authors would like to thank tim cochran for his support and interest in the project . they would also like to thank john etnyre for helpful conversations .

we investigate legendrian graphs in . we extend the classical invariants , thurston - bennequin number and rotation number to legendrian graphs . we prove that a graph can be legendrian realized with all its cycles legendrian unknots with and if and only if it does not contain as a minor .

we investigate legendrian graphs in . we extend the classical invariants , thurston - bennequin number and rotation number to legendrian graphs . we prove that a graph can be legendrian realized with all its cycles legendrian unknots with and if and only if it does not contain as a minor . we show that the pair does not characterize a legendrian graph up to legendrian isotopy if the graph contains a cut edge or a cut vertex . for the lollipop graph the pair determines two legendrian classes and for the handcuff graph it determines four legendrian classes .

1106.4320

c

we have investigated signals of nucleon destruction that can arise in specific theories where dark matter is antibaryonic . antibaryonic dm is motivated by unified mechanisms for dm and baryon generation ( referred to as hylogenesis ) which address the cosmic coincidence between the energy densities of dark and visible matter in the universe . in the hylogenesis scenario considered here , the dm consists of an asymmetric density of fermions @xmath11 and scalars @xmath12 . these particles can scatter inelastically with nucleons via the reactions @xmath321 and @xmath322 , where @xmath30 is a nucleon and @xmath27 is a meson . these induced nucleon decay ( ind ) processes lead to distinctive signatures in nucleon decay searches , stellar evolution , and hadron colliders . ind is a novel signature of dm that can be searched for in terrestrial nucleon decay searches . the effective nucleon lifetime is expected to be @xmath323 years , if baryon transfer between dark and visible sectors is mediated by new physics at a scale @xmath324 . since the dm states are unobserved , ind events mimic standard nucleon decay with neutrino final states , but typically with greater final state meson energy @xmath325 . due to these different kinematics , existing searches do not directly apply in general . our study therefore motivates new searches in these experiments . we expect that the resulting sensitivities should be comparable to those in standard nucleon decay searches . the coupling of hylogenic dm to quarks responsible for ind can also give rise to observable signals in hadron colliders . in particular , such couplings can potentially lead to monojet signals at the tevatron and the lhc . existing monojet searches at the tevatron place a lower bound on the coupling of hylogenic dm to quarks , and this bound will be significantly improved at the lhc . we find that the coupling strengths that can be probed at the lhc are of the same size as those that can produce an observable effect in nucleon decay search experiments . this correlation may permit the characterization of hylogenic dm through a very diverse set of experimental probes . if hylogenic dm also scatters elastically with nucleons , it can be captured in stars . once captured , it will thermalize with the baryons in the star and sink to the stellar core . in compact stars , the local baryon density in the stellar core may be large enough for ind to occur at a significant rate , destroying both a nuclon and the hdm particle , and producing an anti - hdm particle and a meson . this process deposits energy in the star , an effect that may be probed by observing old and cool neutron stars or white dwarfs . however , current observations do not provide a definitive bound on or evidence for these processes .

we investigate new and unusual signals that arise in theories where dark matter is asymmetric and carries a net antibaryon number , as may occur when the dark matter abundance is linked to the baryon abundance . these processes lead to an effective nucleon lifetime of years in terrestrial nucleon decay experiments , if baryon number transfer between visible and dark sectors arises through new physics at the weak scale .

we investigate new and unusual signals that arise in theories where dark matter is asymmetric and carries a net antibaryon number , as may occur when the dark matter abundance is linked to the baryon abundance . antibaryonic dark matter can cause _ induced nucleon decay _ by annihilating visible baryons through inelastic scattering . these processes lead to an effective nucleon lifetime of years in terrestrial nucleon decay experiments , if baryon number transfer between visible and dark sectors arises through new physics at the weak scale . the possibility of induced nucleon decay motivates a novel approach for direct detection of cosmic dark matter in nucleon decay experiments . monojet searches ( and related signatures ) at hadron colliders also provide a complementary probe of weak - scale dark - matter induced baryon number violation . finally , we discuss the effects of baryon - destroying dark matter on stellar systems and show that it can be consistent with existing observations . baryon destruction by asymmetric dark matter hooman davoudiasl , david e. morrissey , kris sigurdson , sean tulin + department of physics , brookhaven national laboratory , + upton , ny 11973 , usa + theory group , triumf , + 4004 wesbrook mall , vancouver , bc v6 t 2a3 , canada + department of physics and astronomy , university of british columbia , + vancouver , bc v6 t 1z1 , canada + email : [email protected] , [email protected] , [email protected] , [email protected]

0801.3305

r

figure [ fig:4 ] shows the gas structure around the hill sphere for model m04 at @xmath186 . figure [ fig:4 ] left panel shows the structure on the cross section in the @xmath187 plane , in which red lines indicate the streamlines . in this panel , the gas enters @xmath188 grid from upper @xmath83 boundary for @xmath189 ( from lower @xmath83 boundary for @xmath190 ) and goes downward ( upward for @xmath190 ) according to the keplerian shear motion . the shocks ( crowded contours near the hill radius ) are seen in the upper right and lower left region from the protoplanet . in this model , the shock front almost corresponds to the hill radius ( fig . [ fig:4 ] left panel ) . when the gas approaches the protoplanet , the streamlines are bent by the gravity of the protoplanet . according to @xcite , the gas flow is divided into three region : the pass - by region ( @xmath191 ) , the horseshoe region ( @xmath192 , and @xmath193 ) , and the planet atmosphere region ( @xmath194 ) . note that although @xcite classified the flow pattern in their two dimensional calculation , their classification is useful for global flow pattern in three dimensions . in the pass - by region , the flow is first attracted toward the protoplanet , and then causes a shock after passing by the protoplanet . at the shock front , the density reaches a local peak and the streamlines bend suddenly . on the other hand , the gas entering the horseshoe region turns round by the coriolis force and goes back . the outermost streamlines in the horseshoe region ( i.e. , the streamlines passing very close to the protoplanet ) pass through the shock front , while the gas on the streamlines far from the protoplanet does not experience the shock . in the atmospheric region , the gas is bound by the protoplanet and forms a circumplanetary disk that revolves circularly around the protoplanet in the prograde ( counterclockwise ) direction . although the streamlines on @xmath187 plane ( fig . [ fig:4 ] left panel ) are similar to those in recent two dimensional calculations ( e.g. , * ? ? ? * ; * ? ? ? * ) , there are important differences . in two - dimensional calculations , a part of the gas near the hill sphere can accrete onto the protoplanet . @xcite showed that gas only in a narrow band that distributed from the lower left to the upper right region against the protoplanet for @xmath128 spirals inward toward the protoplanet passing through the shocked region and finally accretes onto the protoplanet ( for details , see figs . 4 , and 8 of * ? ? ? * ) . on the other hand , in our three - dimensional calculation , gas only flows out from the hill sphere and thus does not accrete onto the protoplanet on the midplane . figure [ fig:4 ] left panel shows that although gas flows into the hill sphere , a part of the gas flows out from the central region . figure [ fig:4 ] right panel is three - dimensional view at the same epoch as the left panel . in this panel , only the streamlines flowing into the high - density region of @xmath195 are drawn for @xmath196 which are inversely integrated from the high - density region . this panel clearly shows the gas flowing into the protoplanet in the vertical direction . to investigate gas flowing into the protoplanet system in detail , in figure [ fig:5 ] , we plot the streamlines at the same epoch as figure [ fig:4 ] with different grid levels ( @xmath93=3 , 5 , and 7 ) . in this figure , each upper panel shows three - dimensional view , while each lower panel shows the structure on the cross section in the @xmath197 plane . note that , in lower panels , the streamlines are projected onto the @xmath197 plane . figure [ fig:5]_a _ shows only the streamlines in a narrow bundle flowing into the protoplanet system . this feature is similar to that shown in two - dimensional calculations @xcite . however , the streamlines in figure [ fig:5]_a _ indicate that gas rises upward near the shock front and then falls into the central region in the vertical direction . gas flowing in the vertical direction spirals into the inner region ( fig . [ fig:5]_c _ ) . in this process , vortices appear as shown in figure [ fig:5]_d_. as shown in figure [ fig:4 ] left panel , also in figure [ fig:5]_d _ , gas is flowing out from the central region on @xmath187 plane . gas in the proximity of the protoplanet rotates circularly in the prograde direction as shown in figure [ fig:5]_e_. figure [ fig:5]_f _ shows that a part of the gas flowing into the upper boundary of @xmath198 grid level contributes to the disk formation around the protoplanet , while a remainder is bent and flows out from the central region . when we look down the protoplanetary disk from the above along the @xmath84-axis , streamlines may seem to be almost the same as those in two - dimensional calculations . however , gas moves also in the vertical direction : streamlines go upward near the shock front ( @xmath199 ) , and vertically falls into the central region at @xmath195 . this feature of streamlines is also seen in models m04s01 and m04s03 , in which the sink cell is adopted . thus , different features of streamlines in two- and three - dimensional calculations are not caused by the effect of the pressure gradient force , but caused by the dimensions ( because the same feature appears in both models with and without the sink ) . this flow pattern is also seen in other three - dimensional calculation @xcite . in two - dimensional calculation , since the vertical motion is restricted , the flow pattern is different from that in three dimensions . this difference affects the accretion rate onto the protoplanet and migration rate . @xcite showed that the migration rate is different between two- and three - dimensional calculations . in the present study , however , since we focus on the angular momentum of a protoplanet system , we do not discuss them any more . we will discuss the accretion and migration rate in the subsequent papers . finally , we comment on the circumplanetary disk . in figure [ fig:5 ] lower panels , the green surface ( i.e. , iso - density surface ) indicates the high - density structure around the protoplanet . these panels show the disk - like structure in the proximity of the protoplanet , and the disk becomes thinner as it approaches the protoplanet ( i.e. , the origin ) . we discuss the circumplanetary disk in [ sec : dis - disk ] . we have shown the evolution of the protoplanet system for model with 0.4@xmath2 in [ sec : typical ] . in this subsection , we investigate the evolution of the protoplanet system with different protoplanet masses . figure [ fig:6 ] upper panels show the accumulated masses ( eq . [ eq : mr ] ) within @xmath200 ( @xmath201 , fig . [ fig:6]_a _ ) and @xmath202 ( @xmath203 , fig . [ fig:6]_b _ ) against the elapsed time for different models , while figure [ fig:6 ] middle panels show the corresponding angular momenta ( eq . [ eq : jr ] ) in the same regions ( fig . [ fig:6]_c _ for @xmath204 , and fig . [ fig:6]_d _ for @xmath205 ) . in figure [ fig:6 ] , both masses and angular momenta for all models increase rapidly for @xmath206 . this is because the protoplanet with mass of @xmath207 suddenly appears in the protoplanetary disk at @xmath208 . however , this rapid growth phase ( @xmath206 ) is not real , because the gas is considered to begin to accrete onto the protoplanet when the mass of the solid core exceeds @xmath6 @xcite . the growth rates of the mass and angular momentum begin to decrease at @xmath209 in all models , then both masses and angular momenta increase with an almost constant rate until the end of the calculation ( @xmath210 ) . @xcite calculated the mass accretion rate onto the protoplanet as @xmath211 thus , the growth time @xmath212 is @xmath213yr ( i.e. , @xmath214 ) . in our calculation , we continue to calculate the evolution of the protoplanet system for @xmath215yr ( @xmath216 ) by fixing the planet mass . we think that this treatment is not problematic , since the growth timescale is longer than our calculation time . note that since @xcite calculated the evolution of the protoplanet system in two dimensions , the growth rate might be different from that in three - dimensional calculations . figure [ fig:6 ] lower panels show the evolution of the average specific angular momentum in the region of @xmath200 ( fig . [ fig:6]_e _ for @xmath217 ) and @xmath202 ( fig . [ fig:6]_f _ for @xmath218 ) . the masses and angular momenta flowing into the protoplanet system increase with a constant rate for @xmath206 , while the average specific angular momenta are saturated at certain values for @xmath219 . this saturation means that flow around the protoplanet is in the steady state . figures [ fig:6]_e _ and _ f _ also indicate that the average specific angular momentum brought into the protoplanet system increases with the mass of the protoplanet . in figures [ fig:6]_e _ and _ f _ , the average specific angular momenta in the region of @xmath200 is larger than those in the region of @xmath202 indicating that the protoplanet system has a larger average specific angular momentum in the more distant place from the protoplanet . we will investigate the angular momentum acquired in the protoplanet system in [ sec : dis - evo ] . figures [ fig:7 ] and [ fig:8 ] show the density distributions ( upper panels ) and jacobi energy contours ( lower panels ) on the cross section in the @xmath187 ( fig . [ fig:7 ] ) and @xmath197 plane ( fig . [ fig:8 ] ) around the hill sphere after the steady state is achieved ( @xmath220 ) for models m02 ( left panels ) , m04 ( middle panels ) , and m06 ( right panels ) . the white - dotted line in each panel indicates the hill radius @xmath71 . in each upper panel , the shock appears from the upper left to lower right near the hill radius . these shocks are frequently seen in similar calculations @xcite . in figure [ fig:7 ] upper panels , the round structures are seen in the proximity to the protoplanet ( @xmath195 ) , while the ellipsoidal structures are seen in the regions of @xmath221 . this is because gas distributed near the protoplanet is more strongly bound by the protoplanet . figure [ fig:8 ] upper panels show that contours of the central region sags in the center of a concave structure , and thin disks are formed around the protoplanet ( @xmath195 ) . in addition , the butterfly - like structure is also seen inside the hill radius in figure [ fig:8 ] upper panels . these structures are considered to be formed by the rapid rotation of the central circumplanetary disk : similar structure is seen in a rapidly rotating protostar ( e.g. , fig.1 of * ? ? ? we will discuss the disk structure in [ sec : dis - disk ] . in contrast to celestial mechanics , it is difficult to find fluid elements bound by the protoplanet because thermal energy is important in addition to the gravitational and kinetic energies . to discern gas bound by the gravity of the protoplanet , we use the jacobi energy as an indicator @xcite . in our unit , the jacobi energy is given by @xmath222 in equation ( [ eq : jacobi ] ) , the first term is the kinetic energy , the second and third terms are the tidal energy of the central star , and the forth term is the gravitational energy of the protoplanet . the jacobi energy is the conserved quantity in a rotating system . in the case of fluid , the jacobi energy is not strictly appropriate because the thermal energy is ignored . however , we can use this for rough estimation of gas bound by the protoplanet . we determine fluid elements bound by the protoplanet from the contour of the jacobi energy in the lower panels in figures [ fig:7 ] and [ fig:8 ] . fluid elements with lower jacobi energy are strongly bound by the protoplanet as described in equation ( [ eq : jacobi ] ) . lower panels of figures [ fig:7 ] and [ fig:8 ] show that inside the region of @xmath223 , each contour has a closed ellipse . since fluid elements move on this closed orbit , it is considered that gas distributed in the region of @xmath223 is bound by the protoplanet when the thermal effect can be ignored . we discuss the thermal effect in [ sec : dis - evo ] . on the other hand , although fluid elements exist inside the hill radius , outside @xmath224 , they are not bound by the protoplanet , because the contours of the jacobi energy straddle the hill radius . for example , in lower left panel of figure [ fig:7 ] , the contour of @xmath225 straddle the hill radius , and thus fluid elements with this jacobi energy freely move on this contour . thus , although these elements transiently stay inside the hill radius , they flow out from the hill sphere . in summary , figures [ fig:7 ] and [ fig:8 ] suggest that fluid elements inside @xmath226 are bound by the protoplanet . figure [ fig:9 ] upper panel shows the accumulated mass @xmath227 for different models after the steady state is achieved ( @xmath174 ) . in this panel , a thin solid line indicates the initial value ( or the value of the protoplanetary disk ) , and the circle is the hill radius @xmath228 for each model . the accumulated mass in any model is larger than the initial value , because gas flows into the hill sphere . this panel indicates that the massive protoplanet has a massive envelope . since these mass distributions are in a steady state at the fixed mass of the protoplanet , it can be considered that different curves correspond to snapshots at different evolution phases . namely , a gas envelope increases its mass with time ( or the protoplanet mass ) . outside the hill radius , the accumulated mass in each model converges to the initial value indicating that the mass distribution for @xmath229 does not change from the initial state because the influence of the protoplanet is small . figure [ fig:9 ] lower panel shows the absolute value of the angular momentum @xmath230 . there are spikes in all models , at which the sign of the angular momentum is reversed . as shown in [ sec : sink ] , the angular momentum has a positive sign around the protoplanet ( @xmath231 ) , while it becomes negative outside the hill radius ( @xmath229 ) . the sign of the angular momentum is reversed outside the hill radius for models m01 , m02 , m04 , and m06 , while it is reversed inside the hill radius for model m005 . this is because the protoplanet system for model m005 does not acquire a sufficient mass and angular momentum owing to the shallow gravitational potential and relatively large thermal pressure ( for details , see [ sec : dis - evo ] ) . figure [ fig:9 ] upper panel shows that the protoplanet system has the envelope mass of only @xmath232 for model m005 , in which the protoplanet mass is @xmath233 . except for model m005 , the angular momenta gradually decrease after they reach their peak around @xmath234 , then it becomes negative at @xmath235 . thus , the angular momenta bound by the protoplanet system are limited in the region of @xmath236 at the maximum . figure [ fig:9 ] lower panel shows that the angular momentum keeps an almost constant value around the hill radius ( @xmath237 ) , until they are reversed . this means that the angular momentum with plus sign and that with minus sign are mixed in this region , as shown in figure [ fig:4 ] left panel . as a result , wherever we estimate the angular momentum in the range of @xmath238 , we can obtain almost the same values of the angular momentum . the distributions of the average specific angular momentum @xmath175 for different models are shown in figure [ fig:10 ] . the circles in this figure mean the hill radii @xmath228 . the crosses indicate the jacobi radii @xmath239 inside which gas is considered to be bound by the gravity of the protoplanet . we determine the jacobi radii from the contour of the jacobi energies as in figures [ fig:7 ] and [ fig:8 ] lower panels . the jacobi radii in all models are distributed in the range of @xmath240 . figure [ fig:10 ] indicates that a more massive protoplanet has an envelope with larger amount of the specific angular momentum . thus , the specific angular momentum accreting onto the protoplanet system increases as the protoplanet mass increases . the rapid drops at large radii indicate the reverse of the rotation axis as shown in figures [ fig:3 ] and [ fig:9 ] . in figure [ fig:10 ] , in model m01 ( @xmath241 ) , @xmath175 at @xmath180 is twice of that at @xmath242 , while , in models m02 , m04 , m06 , m1 , and m3 ( @xmath243 ) , there are little differences between the average specific angular momentum derived from the hill radius @xmath244 and that derived from jacobi radius @xmath245 . thus , we can safely estimate the average specific angular momentum using either the hill radius or jacobi radius for models with @xmath246 . to properly calculate the angular momentum of the protoplanet system , we have to calculate the planetary growth from the solid core with @xmath6 to the present mass . however , it takes huge computation time to calculate all evolution phases . thus , we estimate the angular momentum of the protoplanet system according to the following procedure : ( i ) we calculate the average specific angular momenta of the gas flowing into the protoplanet system at the fixed masses of the protoplanet ( i.e. , under the same parameter @xmath71 ) in models changing the mass of the protoplanet , then ( ii ) derive the relation between the average specific angular momentum and mass of the protoplanet , and describe it as a function of the protoplanet mass , and ( iii ) estimate the angular momentum of the protoplanet system integrating the average specific angular momentum by mass up to the present value of gas giant planets . at first , we analytically estimate the specific angular momentum of the protoplanet system , then compare it with numerical results . we assume that the gas that overcomes the hill potential flows into the hill sphere ( or a protoplanet system ) with the kepler velocity of the protoplanet at the hill radius @xmath228 . when the protoplanet mass is @xmath0 , the kepler velocity at @xmath228 is given by @xmath247 note that @xmath248 is constant when @xmath33 is fixed . the specific angular momentum can be written as @xmath249 thus , the specific angular momentum is proportional to @xmath250 . in addition , when the mass of the central star is fixed , the specific angular momentum is proportional to @xmath251 thus , the specific angular momentum increases with 2/3 power of the protoplanet mass . figure [ fig:11 ] shows the average specific angular momentum derived from all models against the protoplanet mass . the average specific angular momenta are estimated in the region of @xmath252 ( @xmath253 ) , @xmath254 ( @xmath255 ) , @xmath256 ( @xmath257 ) , and @xmath258 ( @xmath259 ) . although we fixed the protoplanet mass in each model , we can consider the horizontal axis in figure [ fig:11 ] as time sequence of the protoplanet system . figure [ fig:11 ] clearly shows that more massive protoplanet can acquire the envelope with larger average specific angular momentum . although the average specific angular momenta for @xmath260 differ from those for @xmath261 for models with @xmath262 , there are little difference for models with @xmath263 . in addition , when we adopt the average specific angular momenta in the region of @xmath264 , we underestimate them for any model as shown in figure [ fig:11 ] . therefore , we can properly estimate the average specific angular momentum of the protoplanet system in any region of @xmath265 for models with @xmath266 . in figure [ fig:11 ] , the red and blue lines are fitting formulae of the evolution of the average specific angular momentum as a function of the protoplanet mass for @xmath267 ( low mass , @xmath268 ; red line ) , and @xmath269 ( high mass , @xmath270 : blue line ) . they are given by @xmath271 and @xmath272 figure [ fig:11 ] shows that the evolution of the average specific angular momentum for @xmath273 is well described by @xmath274 which corresponds to the equation ( [ eq : ana - j ] ) . on the other hand , the growth rate of the average specific angular momentum for @xmath275 is larger than @xmath274 , and can be fitted by @xmath3 . we ignored the thermal effect when we derive equation ( [ eq : ana - j ] ) . when @xmath0 is small , the gas flowing into the protoplanet system is affected relatively strongly by the thermal pressure . we can estimate the mass at which the gravity dominates the thermal pressure force from the balance between the thermal pressure gradient and gravitational forces . near the hill radius , the thermal pressure gradient force is more dominant than the gravity when the kepler speed is slower than the sound speed ( i.e. , @xmath276 ) . on the other hand , when @xmath277 , the gas flow is controlled mainly by the gravity of the protoplanet even near the hill radius . using equation ( [ eq : kepler ] ) , @xmath278 is realized when @xmath279 which corresponds to @xmath280 of the protoplanet mass at @xmath79au . thus , the gas flow is largely affected by the thermal pressure for @xmath281 , while it is not so affected by the thermal pressure for @xmath282 . in figure [ fig:11 ] , the evolution of the angular momentum for @xmath283 well corresponds to the analytical solution ( @xmath4 ) indicating that the thermal effect is negligible for @xmath284 . however , for @xmath285 , the average specific angular momentum is smaller than @xmath270 ( blue line ) . thus , when the protoplanet is @xmath275 , the thermal effect is not negligible for the acquisition process of the angular momentum because the gravitational potential is relatively shallow . thus , the thermal pressure seems to remain important for @xmath286 , although @xmath278 realized at @xmath287 . to verify the relation of the thermal and gravitational effects , in figure [ fig:12 ] , we plot the ratio of the azimuthal to kepler velocity ( @xmath288 ) around the protoplanet on the midplane for models m02 , m04 , and m06 . in this figure , closed contours inside the hill radius indicate that gas revolves around the protoplanet . for example , the closed contour of @xmath289 in figure [ fig:12 ] means that gas rotates along the contour with 50% of the kepler velocity . the black circles represent contours of @xmath290 , inside which gas revolves around the protoplanet with supersonic velocity . thus , the gravitational force of the protoplanet is dominant inside the black circles , while the thermal pressure gradient force is dominant outside the black circles . figure [ fig:12 ] shows that the radius of the @xmath290 contour increases with the protoplanet mass indicating that the region that dominated by gravity of the protoplanet extends with the protoplanet mass . as the protoplanet mass increases , the flow speed inside the hill radius approaches to the kepler velocity , and the region dominated by the gravity of the protoplanet spreads outward . in this way , since the thermal effect decreases as the protoplanet mass increases , the growth rate of the average specific angular momentum approaches to the analytical solution ( eq . [ eq : ana - j ] ) .

we find that the gas flows onto the protoplanet system in the vertical direction crossing the shock front near the hill radius of the protoplanet , which is qualitatively different from the picture established by two - dimensional simulations . the specific angular momentum of the gas accreted by the protoplanet system increases with the protoplanet mass . at jovian orbit , when the protoplanet mass is , where is jovian mass , the specific angular momentum increases as . on the other hand , it increases as when the protoplanet mass is . the stronger dependence of the specific angular momentum on the protoplanet mass for is due to thermal pressure of the gas . the estimated total angular momentum of a system of a gas giant planet and a circumplanetary disk is two - orders of magnitude larger than those of the present gas giant planets in the solar system . we also discuss the satellite formation from the circumplanetary disk .

we investigate the accretion of angular momentum onto a protoplanet system using three - dimensional hydrodynamical simulations . we consider a local region around a protoplanet in a protoplanetary disk with sufficient spatial resolution . we describe the structure of the gas flow onto and around the protoplanet in detail . we find that the gas flows onto the protoplanet system in the vertical direction crossing the shock front near the hill radius of the protoplanet , which is qualitatively different from the picture established by two - dimensional simulations . the specific angular momentum of the gas accreted by the protoplanet system increases with the protoplanet mass . at jovian orbit , when the protoplanet mass is , where is jovian mass , the specific angular momentum increases as . on the other hand , it increases as when the protoplanet mass is . the stronger dependence of the specific angular momentum on the protoplanet mass for is due to thermal pressure of the gas . the estimated total angular momentum of a system of a gas giant planet and a circumplanetary disk is two - orders of magnitude larger than those of the present gas giant planets in the solar system . a large fraction of the total angular momentum contributes to the formation of the circumplanetary disk . we also discuss the satellite formation from the circumplanetary disk .

1601.07675

i

the complex behavior of mixtures of bosonic gases has been extensively investigated experimentally in particular , in cold - atom systems @xcite and theoretically @xcite . these systems exhibit a rich behavior , at zero and finite temperature , with several different phases separated by transition lines , along which one or more components of the system undergo bose - einstein condensation ( bec ) . in this paper we consider three - dimensional ( 3d ) mixtures of two different bosonic gases with short - range interactions that only depend on the local densities of the two gases . the hamiltonian of these systems is invariant under u(1 ) transformations acting independently on each species , so that the model is u(1)@xmath0u(1 ) symmetric . in particular , we consider the 3d two - component bose - hubbard model with an on - site inter - species density - density interaction . this is a realistic model for two bosonic species in optical lattices @xcite . we determine the finite - temperature phase diagram by using a variety of techniques . first , we consider the mean - field ( mf ) approximation , determining the qualitative phase behavior of the system as a function of the model parameters , such as the chemical potentials and the on - site inter- and intra - species couplings . we find several different phases , in which each species may be in the normal or superfluid state , and identify critical lines and multicritical points ( mcps ) where some transition lines meet . the 3d phase diagram is investigated in the hard - core ( hc ) limit of each species by a finite - size scaling ( fss ) analysis of quantum monte carlo ( qmc ) simulations . the numerical data allows us to identify the universality class of the transition lines that correspond to the bec of one of the two species . we show that , independently whether the noncritical component is in the normal or superfluid phase , the critical behavior of the condensing species belongs to the 3d xy universality class characterized by the breaking of a global u(1 ) symmetry and by short - range effective interactions . this is the same universality class associated with the bec of a single bosonic gas @xcite ( and also with the superfluid transition in @xmath3he @xcite , with transitions in some liquid crystals characterized by density or spin waves and in magnetic systems with easy - plane anisotropy , etc . this result implies an effective decoupling of the critical modes of the condensing species from those of the noncritical component , independently whether the latter is in the normal or superfluid phase . the phase diagram of mixtures of bosonic gases also presents particular points where some transition lines meet . multicritical behaviors develop at these mcps , arising from the competition of the two u(1 ) order parameters associated with the bec of the two species . to identify the possible universality classes of the multicritical behaviors , we use the field - theoretical approach , considering the effective landau - ginzburg - wilson ( lgw ) @xmath1 theory for two complex fields with global u(1)@xmath0u(1 ) symmetry . we study the renormalization - group ( rg ) flow in the quartic - parameter space , identifying the stable fixed points ( fps ) , which control the critical behavior , and their attraction domain . the paper is organized as follows . in sec . [ model ] we define the bose - hubbard model for two lattice bosonic gases . in sec . [ phmf ] we determine the phase diagram of the model in the mf approximation , showing that , by changing the model parameters , one can obtain qualitatively different behaviors . in sec . [ critbeh ] we present our numerical results and determine numerically the critical behavior along the transition lines where one species undergoes bec . in sec . [ multi ] we study the multicritical behaviors at mcps where some transition lines meet in the phase diagram . finally , in sec . [ conclu ] we draw our conclusions . [ 5loop ] reports the five - loop perturbative series of the minimal - subtraction scheme , which are used in the rg study of the multicritical behavior .

we study the phase diagram and the critical behaviors along the transition lines characterized by the bose - einstein condensation of one or both species . we present mean - field calculations and numerical finite - size scaling analyses of quantum monte carlo data . we determine the possible multicritical behaviors by using field - theoretical perturbative methods .

we investigate the bose - einstein condensation patterns , the critical and multicritical behaviors of three - dimensional mixtures of bosonic gases with short - range density - density interactions . these systems have a global u(1)u(1 ) symmetry , as the system hamiltonian is invariant under independent u(1 ) transformations acting on each species . in particular , we consider the three - dimensional bose - hubbard model for two lattice bosonic gases coupled by an on - site inter - species density - density interaction . we study the phase diagram and the critical behaviors along the transition lines characterized by the bose - einstein condensation of one or both species . we present mean - field calculations and numerical finite - size scaling analyses of quantum monte carlo data . we also consider multicritical points , close to which it is possible to observe the condensation of both gas components . we determine the possible multicritical behaviors by using field - theoretical perturbative methods . we consider the u(1)u(1)-symmetric landau - ginzburg - wilson theory and determine the corresponding stable fixed points of the renormalization - group flow . the analysis predicts that , in all cases , the multicritical behavior is analogous to the one that would be observed in systems of two identical gases , with an additional exchange symmetry .

1601.07675

c

in this paper we study the critical and multicritical behaviors that can be observed in 3d mixtures of bosonic gases interacting by short - range density - density interactions . these systems have a global u(1)@xmath0u(1 ) symmetry , related to independent u(1 ) transformations acting on each species . as a representative of this class of systems , we consider the 3d bose - hubbard model for two lattice bosonic gases coupled by an on - site inter - species density - density interaction , whose hamiltonian is given in eq . ( [ sbh ] ) . however , the qualitative features of the finite - temperature phase diagram and the results for the universality classes associated with the critical and multicritical behaviors apply to generic bosonic mixtures . the generic features of the phase diagram of the 2bh model have been determined in the mf approximation and additionally confirmed by qmc simulations . the qualititave behavior depends on the model parameters , such as the chemical potentials and the on - site inter- and intra - species couplings . by varying them , one can observe several transition lines , along which one of the two species undergoes a normal - to - superfluid transition , and different types of multicritical behavior . the transition lines separating the different phases generally correspond to the bec condensation of one of the two species . we show that , independently whether the other species is in the normal or superfluid phase , the critical behavior of the condensing species belongs to the 3d xy universality class , characterized by the breaking of a global u(1 ) symmetry and short - ranged effective interactions , which is the same universality class associated with the bec of a single bosonic gas . therefore , the critical modes of the condensing gas effectively decouple from those of the other species , independently whether the latter is in the normal or superfluid phase . the phase diagram of mixtures of bosonic gases also presents particular points where some transition lines meet . at these points multicritical behaviors develop , due to the competition of the u(1 ) order parameters related to the two bosonic gases . we investigate them by a field - theoretical approach based on the effective lgw @xmath1 theory for two complex scalar fields with global u(1)@xmath0u(1 ) symmetry . the possible universality classes that describe the multicritical behaviors are associated with the stable fps of the rg flow . they can be determined by studying the rg trajectories in the critical theory , starting from the unstable gaussian fp in the quartic - parameter space . for this purpose , we consider the so - called @xmath223 scheme without @xmath222 expansion @xcite . we start from the five - loop @xmath340 @xmath221 functions , resum them using the pad - borel technique , and solve the flow equations . we find two stable fps , that also belong to the @xmath1 theory ( [ phi4s ] ) , which has an additional @xmath2 symmetry related to the exchange of the two order parameters . this more symmetric model has already been discussed in the context of the critical behavior of a mixture of two identical gases @xcite . if the system has a tetracritical continuous transition , see fig . [ mcpd ] , the critical behavior is controlled by a decoupled fp . each component shows an xy critical behavior correspondingly , the rg dimensions of the two relevant operators are @xmath341 but with very slowly - decaying scaling corrections ( they decay as @xmath342 , where @xmath184 is the correlation length ) due to inter - species coupling . if , instead , the system undergoes a bicritical continuous transition , the critical behavior is associated with a different asymmetric fp , with @xmath343 and @xmath344 . recent experiments on atomic gas mixtures @xcite , either using two different atomic species or the same atomic species in two different states , have already obtained several interesting results on the properties of the low - temperature condensed phase and on the interplay of the different condensates . they have also demonstrated the possibility of a robust control of the model parameters , which may allow the observation of the different phases , such as those found in the present study , and the determination of the nature of the critical and multicritical behaviors . our results should provide a complete characterization of the possible bec patterns and of the critical behaviors that these systems may develop along their transition lines . most cold - atom experiments have been performed in inhomogeneous conditions , due to the presence of space - dependent trapping forces , which effectively confine the atomic gas within a limited space region @xcite . the trapping potential is effectively coupled to the particle density , which may be taken into account by adding a further hamiltonian term to the 2bh hamiltonian ( [ sbh ] ) , i.e. , @xmath345 where @xmath65 is the space - dependent potential associated with the external force . the inhomogeneity arising from the trapping potential introduces an additional length scale @xmath346 into the problem , which drastically changes the general features of the behavior at the phase transitions . experimental data for inhomogeneous trapped cold - atom systems are usually analyzed using the local - density approximation , see , e.g. , ref . . however , this approach fails to describe the emergence of large - scale correlations @xcite . this problem may be overcome experimentally by using ( almost ) flat traps , giving rise to a finite space region where the system is effectively homogenous @xcite . otherwise , one may infer the critical behavior by studying the scaling behavior with respect to the trap size @xmath346 , which is expected to be universal and controlled by the critical exponents of the universality class of the corresponding homogenous system , in the large trap - size limit @xcite .

we investigate the bose - einstein condensation patterns , the critical and multicritical behaviors of three - dimensional mixtures of bosonic gases with short - range density - density interactions . these systems have a global u(1)u(1 ) symmetry , as the system hamiltonian is invariant under independent u(1 ) transformations acting on each species . in particular , we consider the three - dimensional bose - hubbard model for two lattice bosonic gases coupled by an on - site inter - species density - density interaction . we also consider multicritical points , close to which it is possible to observe the condensation of both gas components . the analysis predicts that , in all cases , the multicritical behavior is analogous to the one that would be observed in systems of two identical gases , with an additional exchange symmetry .

we investigate the bose - einstein condensation patterns , the critical and multicritical behaviors of three - dimensional mixtures of bosonic gases with short - range density - density interactions . these systems have a global u(1)u(1 ) symmetry , as the system hamiltonian is invariant under independent u(1 ) transformations acting on each species . in particular , we consider the three - dimensional bose - hubbard model for two lattice bosonic gases coupled by an on - site inter - species density - density interaction . we study the phase diagram and the critical behaviors along the transition lines characterized by the bose - einstein condensation of one or both species . we present mean - field calculations and numerical finite - size scaling analyses of quantum monte carlo data . we also consider multicritical points , close to which it is possible to observe the condensation of both gas components . we determine the possible multicritical behaviors by using field - theoretical perturbative methods . we consider the u(1)u(1)-symmetric landau - ginzburg - wilson theory and determine the corresponding stable fixed points of the renormalization - group flow . the analysis predicts that , in all cases , the multicritical behavior is analogous to the one that would be observed in systems of two identical gases , with an additional exchange symmetry .

1308.1094

l

[ sec : studies ] in this section we develop collider analyses aimed at reconstructing the @xmath30 angle in ( [ eq : thetadef ] ) . from ( [ eq : matrix_element : full ] ) and ( [ eq : greatformula ] ) , the matrix element squared for the @xmath224 decay has a term proportional to @xmath225 : the @xmath30 distribution is thus sensitive to the @xmath0 phase @xmath31 as its minimum is located at @xmath226 . as before , we fix @xmath227 and therefore the only new parameter we introduce is @xmath31 . we implement the @xmath31 phase in ( [ eq : lpheno ] ) and the effective vertices in ( [ eq : m : tau_to_rho - nu ] ) and ( [ eq : m : rho_to_pi - pi ] ) into a feynrules v.1.6.0 @xcite model . we then generate monte carlo events in madgraph 5 @xcite for @xmath228 production at the lhc with @xmath43 tev as well as @xmath229 production at the ilc with @xmath230 gev : in either case , the higgs decays via @xmath231 . in order to retain quantum interference effects , the full @xmath232 body process is simulated . for the lhc study , we also generate a background sample of @xmath233 production with the subsequent decay @xmath234 . we will first study the effectiveness of the @xmath30 distribution at truth level , assuming the neutrino momenta are known : this facilitates a comparison to the @xmath235 variable @xcite , which was previously proposed for studying @xmath0 violation in the higgs coupling to taus . after demonstrating the superior qualities of the @xmath30 variable , we present a sensitivity study for reconstructing @xmath30 at the ilc , where the neutrino four - momentum can be reconstructed up to a two - fold ambiguity . finally , we turn to the lhc , where the neutrinos can not be reconstructed and the irreducible @xmath10 background is significant . in this case , we find that using a collinear approximation @xcite for the neutrino momenta in addition to the standard hard cuts for higgs events still allows the @xmath30 distribution to retain significant discrimination power between different underlying @xmath31 signal models . we do not include pileup or perform any detector simulation in this work , aside from implementing flat efficiencies for @xmath1-tagging for the lhc study . pileup effects are expected to complicate the primary vertex determination necessary for measuring charged pion tracks as well as contribute extra ambient radiation in the electromagnetic calorimeter ( ecal ) , making neutral pion momenta measurements more difficult . furthermore , finite tracking and calorimeter resolutions are expected to smear the @xmath30 distribution . in particular , the ability to distinguish between charged and neutral pion momenta when both pions are overlapping also could affect the @xmath30 measurement . note , however , that because of the magnetic field , the softer @xmath133 and @xmath134 could be separated at the ecal . even if the two pions overlap in the ecal , the @xmath134 momentum can be obtained by subtracting the track momentum from the total momentum measured in ecal , assuming negligible contamination from other sources of energy deposition . we also neglect the neutral pion combinatoric issue , which is justified if the respective parent rho mesons are boosted far apart as a result of the higgs decay . in general , the @xmath133 and @xmath134 coming from the same @xmath92 parent are mostly collinear . this fact has been exploited in the hadronic tau tagging algorithm . for example , the hps algorithm used by cms requires that the charged and neutral hadrons are contained in a cone of the size @xmath236 , where @xmath237 is the transverse momentum of the reconstructed tau @xcite . since the two tau candidates are usually required to be well separated , the combinatorics problem in determining the correct @xmath92 parents can be ignored . recall from ( [ eq : matrix_element : full ] ) and ( [ eq : greatformula ] ) that the minimum of the @xmath30 distribution is located at @xmath238 , and so constructing the @xmath30 distribution allows us to read off the @xmath31 phase of the underlying signal model . in figure [ fig : true_varydelta ] , we show the @xmath30 distribution in @xmath228 events where we have temporarily assumed the neutrinos are fully reconstructed . the various signal models with @xmath239 ( @xmath0-even ) , @xmath240 ( maximal @xmath0 admixture ) , and @xmath241 ( @xmath0-odd ) clearly show the large @xmath242 contribution of the matrix element as seen in ( [ eq : greatformula ] ) . we also superimpose the @xmath30 distribution from @xmath233 event . note that it is flat . clearly , observing the cosine oscillation in experimental data will require both a favorable signal to background ratio as well as a solution for the neutrino momenta that preserves the inherently large amplitude of the @xmath30 oscillation . the @xmath30 distributions ( compare with ( [ eq : greatformula ] ) ) for the higgs with @xmath239 ( @xmath0-even ) , @xmath240 ( maximal @xmath0 admixture ) , and @xmath243 ( @xmath0-odd ) , and the @xmath10 , assuming neutrinos are fully reconstructed . the relative normalization of the @xmath10 line is arbitrary . ] we now compare @xmath30 at truth level with the @xmath235 variable proposed in refs . @xcite : here , @xmath235 is the acoplanarity angle between the decay planes of @xmath244 and @xmath122 in the @xmath245 rest frame . the sign of @xmath235 is defined as the sign of the product of @xmath246 . following @xcite , the events are divided into two classes , @xmath247 and @xmath248 , where the two classes are differ by a @xmath249 phase shift . in order to make a direct comparison with our @xmath30 variable , we combine the @xmath235 distributions of the two classes with a @xmath249 phase shift so the phases of the two classes agree . note that while @xmath235 does not refer to the neutrinos , this classification into the two classes still requires the knowledge of the neutrino momenta ( see ( [ eq : var : y+- ] ) ) . assuming the neutrinos are fully reconstructed , the @xmath30 and @xmath235 distributions for @xmath250 events are shown in figure [ fig : thetavsworek ] with @xmath239 . we readily see that oscillation amplitude of the @xmath30 distribution is larger than that of the acoplanarity angle @xmath235 by about @xmath251 . compared to @xmath235 , the @xmath30 variable thus provides superior sensitivity to the @xmath0 phase @xmath31 . the distributions of our @xmath30 and the @xmath235 variable of ref . @xcite for @xmath239 . the @xmath235 distribution is aggregated from the two @xmath248 and @xmath252 classes as explained in the text to make the direct comparison clearer . ] having considered the case where the neutrinos from the tau decays are fully reconstructed , we next turn to the lepton collider environment , where we will find the neutrinos can be fully reconstructed up to a two - fold ambiguity . at a lepton collider running at @xmath230 gev , such as the ilc , the main production mode for the higgs is via associated production with a @xmath10 boson . our prescribed decay mode for the higgs , @xmath253 , has two neutrinos that escape the detector . we use the known initial four momenta , two tau mass and two neutrino mass constraints to solve for each neutrino momentum component . note we will assume the @xmath10 decays to visible states , which will reduce our event yield by 20% . solving the system of equations for the neutrino momenta gives rise to a two - fold ambiguity , where one solution is equal to the truth input neutrino momenta while the other gives a set of wrong neutrino momenta . note both solutions are consistent with four - momentum conservation and therefore correctly reconstruct the higgs mass . since these solutions are indistinguishable in the analysis , we assign each solution half an event weight . the resulting distribution of @xmath30 for @xmath239 is given in figure [ fig : ilctheta ] , where we superimpose the truth level @xmath30 distribution for @xmath254 events for easy comparison . we can see that the oscillation amplitude at the ilc is degraded from the truth level result by @xmath255 . we also show the reconstructed distribution for @xmath239 , @xmath240 , and @xmath256 in figure [ fig : ilcvarydelta ] . while the two - fold ambiguity for the neutrino momenta solution set does degrade the truth level result , the _ @xmath30 distribution in figure [ fig : ilcvarydelta ] shows significant discrimination power between various @xmath31 signal models . note the amplitude of pseudoscalar distribution ( @xmath256 ) is slightly higher than the scalar amplitude : here , the `` wrong solution '' approximates the correct neutrino momenta on average better than the other @xmath239 or @xmath240 cases . this small effect can be traced back to equation ( [ eq : m : h_to_tautau : higgs_frame ] ) where we derived that a pseudoscalar decays to two taus in the singlet spin state . as a result , in this case the two tau spins point in opposite directions , regardless of the spin quantization axis . in the pseudoscalar case the two tau decays thus tend to occur with opposite orientation and the two neutrinos are slightly more back - to - back and consequently the two solutions for their momenta are closer together . the truth and reconstructed @xmath30 distributions at the ilc for @xmath239 . ] the reconstructed @xmath30 distribution at the ilc for @xmath239 , @xmath257 , and @xmath258 . ] we now discuss the projected ilc sensitivity for measuring @xmath31 . at the ilc , the cross section for @xmath259 production at @xmath230 gev with polarized beams @xmath260 for @xmath261 gev is 0.30 pb @xcite . distribution is insensitive to the polarization of the @xmath262-@xmath263 beams . ] assuming a higgs branching fraction to tau pairs of 6.1% , a @xmath264 branching fraction of 26% , and a @xmath10-to - visible branching fraction of 80% , we calculate the ilc should have 990 events with 1 ab@xmath5 of luminosity . since the solved neutrino momenta correctly reconstruct the higgs mass , the @xmath265 backgrounds are negligible and will be ignored . [ tab : ilc ] to estimate the expected ilc accuracy for measuring @xmath31 , we perform a log likelihood ratio test for the sm hypothesis with @xmath266 against an alternative hypothesis with @xmath267 . in general , the likelihood ratio in @xmath268 bins is given by @xmath269 where @xmath270 , @xmath271 and @xmath272 are the number of background events , signal events assuming @xmath239 , and signal events assuming @xmath267 in bin @xmath273 of the @xmath30 distribution . in our ilc treatment , we neglect @xmath265 and @xmath274 continuum backgrounds and so we set @xmath275 . here , @xmath276 is the usual poisson distribution function , @xmath277 . we parametrize the signal @xmath30 distribution with a @xmath278 fit function , where the offset constant @xmath279 and oscillation amplitude @xmath280 are fixed by the fit of the standard model @xmath30 distribution with @xmath239 , giving @xmath281 and @xmath282 respectively . then , the resulting @xmath283 signal @xmath30 distribution is given by @xmath284 . we construct the binned likelihood bins , though we verified the number of bins is immaterial for our results . ] according to ( [ eq : lr ] ) for various @xmath285 hypotheses to test the discrimination against the sm hypothesis . with 1 ab@xmath5 of ilc luminosity , we find @xmath286 discrimination at @xmath287 , which is a highly promising degree of sensitivity for measuring the @xmath0 phase of the higgs coupling to taus . we summarize our rate estimate and accuracy result in table [ tab : ilc ] . we remark that this sensitivity estimate is only driven by statistical uncertainties , and systematic uncertainties are expected to reduce the efficacy of our result . also , detector resolution effects and sm backgrounds , while expected to be small , will also slightly degrade our projection . based on our results , which surpass earlier accuracy estimates of @xmath39 @xcite , a full experimental sensitivity study incorporating these subleading effects is certainly warranted . we now develop an lhc study for reconstructing the @xmath30 distribution in @xmath250 in the @xmath288 final state . we use the @xmath289 final state for a couple of reasons . first , since hadronic taus can be faked by jets , @xmath290 two hadronic taus faces an immense background from multijet qcd . by requiring another object in the final state , we gain handles to suppress the background . second , the collinear approximation gives ambiguous results if the two taus are back - to - back , so the requirement of an additional object in the event guarantees we are away from this configuration . one option is associated production of a higgs wit a @xmath291 . however this rate is quite small , especially once the branching ratios for @xmath291 into clean final states are taken into account . other possibilities include higgs production via vector boson fusion and in association with a jet . both of these options give promising signal - to - background ratios and both should be considered . for concreteness we will consider @xmath292 here as a demonstration of the feasibility of our technique . as mentioned before , the neutrinos are not reconstructible in the hadron collider environment , and so we will employ the collinear approximation @xcite for the neutrino momenta . in figure [ fig : truthvscollinear ] , we show a comparison between the truth level @xmath30 distribution and the @xmath30 distribution using the collinear approximation for neutrino momenta , for the @xmath239 benchmark . while the collinear approximation reduces the oscillation amplitude of the distribution , the location of the minimum of the distribution does not change . therefore , measuring @xmath31 is a viable possibility at the lhc using the collinear approximation for the neutrino momenta . we remark that in the collinear approximation , @xmath30 is equivalent to the acoplanarity angle @xmath235 @xcite . yet , we are the first feasibility study for measuring @xmath0 violation in the higgs coupling to taus at hadron colliders using prompt tau decays and kinematics . with a more sophisticated scheme than the collinear approximation , the @xmath30 variable will be superior to @xmath235 . the distributions of the truth-@xmath30 and @xmath30 from the collinear approximation for @xmath239 . ] at the lhc , the dominant background for the @xmath293 signal process is the irreducible @xmath294 background , where the @xmath10 decays to the same final state as the higgs . as shown earlier in figure [ fig : true_varydelta ] , the @xmath30 distribution from @xmath10 events is flat : importantly , this is true regardless of possible mass window cuts on the reconstructed @xmath295 resonance . we remark that the @xmath0 phase in the higgs coupling to taus does manifest in the @xmath10@xmath1@xmath1 vertex at one loop . since this effect is suppressed by @xmath296 , whereas the signal to background ratio will be @xmath297 , we can safely ignore the loop induced @xmath0 phase in the @xmath10@xmath1@xmath1 vertex . in addition , we will assume that the qcd background contribution also has a flat @xmath30 distribution , since the qcd contamination in the signal region is not expected to have any particular spin correlations . using our @xmath293 and @xmath294 event samples from madgraph 5 for a 14 tev lhc , we first isolate the signal region with a series of hard cuts . first , we apply a preselection requirement on the leading jet @xmath298 gev with @xmath299 . using mcfm v.6.6 @xcite with these preselection requirements on the leading jet , we obtain a @xmath293 nlo inclusive cross section of 2.0 pb with @xmath23 gev and a @xmath294 nlo inclusive cross section of 420 pb . after applying the appropriate higgs , @xmath10 , and tau branching fractions , we calculate a signal cross section of 8.2 fb and @xmath10 background cross section of 970 fb . , while for the background we use @xmath300 . these scale choices are motivated by agreement with higher order ( nnlo ) calculations ( where they exist ) . ] next , we impose hard kinematic cuts to isolate the signal . motivated by @xcite , we choose the signal region to be : * @xmath301 gev , * @xmath302 gev , * @xmath303 , * @xmath304 gev , where @xmath305 is the reconstructed higgs mass by using the collinear approximation . the hard @xmath305 cut strongly suppresses the @xmath306 background , but is less effective on multijet qcd . to reduce the multijet component and its accompanying uncertainty to less than 10% of the total background we impose a high @xmath307 cut . the net efficiencies for signal and @xmath10 background after these cuts are 18% and 0.24% , respectively . rather than simulate the qcd contribution , we account for qcd contamination in the signal region by increasing the @xmath10 background rate by 10% : a complete treatment of the expected qcd background is beyond the scope of this study . finally , for hadronic @xmath1 tagging efficiency , we consider a standard 50% efficiency and a more optimistic 70% efficiency @xcite . we therefore expect 1100 signal events and 1800 @xmath308 qcd background events with 3 ab@xmath5 of luminosity from the 14 tev lhc , assuming 50% @xmath1 tagging efficiency . these rates are summarized in table [ tab : input ] . [ tab : input ] we note that although we generated signal and background samples independently , there is a small interference between higgs and @xmath10 diagrams in the @xmath309 diagram . our checks of this interference on the @xmath30 distributions for combined signal and background events versus separate signal and background events showed a negligible effect : we thus ignore this interference effect . we now perform a likelihood analysis ( [ eq : lr ] ) to quantify how effectively the @xmath30 distribution distinguishes between signal hypotheses with different @xmath0 phases in the presence of @xmath308 qcd background . first , we test the discrimination between a pure scalar and a pure pseudoscalar @xmath49@xmath1@xmath1 coupling . we find that these two hypotheses can be distinguished at @xmath310 sensitivity with 550 ( 300 ) fb@xmath5 assuming 50% ( 70% ) @xmath1 tagging efficiency . we can attain @xmath311 sensitivity between pure scalar and pseudoscalar couplings with 1500 ( 700 ) fb@xmath5 luminosity assuming 50% ( 70% ) efficiency . we also estimate the possible accuracy for the lhc experiments to measure @xmath31 with an upgraded luminosity of @xmath312 . we adopt the same procedure as with the ilc accuracy estimate described in the previous section , modified to account for the @xmath308 qcd background , which is fixed to be flat in @xmath30 . we find that the accuracy in measuring @xmath31 is @xmath42 ( @xmath44 ) assuming 50% ( 70% ) hadronic @xmath1 tagging efficiency . the scalar versus pseudoscalar discrimination and the accuracy estimates are summarized in table [ tab : lhc ] . [ tab : lhc ] again , these estimates are based only on statistical uncertainties without performing a full detector simulation . the effects from pileup and detector resolution are expected to degrade these projections , but corresponding improvements in the analysis , such as a more precise approximation for the neutrino momenta , improved background understanding ( from other lhc measurements ) or multivariate techniques , could counterbalance the decrease in sensitivity . the promising results of our study strongly motivate a comprehensive analysis by the lhc experiments for the prospect of measuring the @xmath0 phase @xmath31 .

we investigate the lhc and higgs factory prospects for measuring the phase in the higgs-- coupling a new , ideal observable is identified from an analytic calculation for the channel . it is demonstrated to have promising sensitivity at the lhc and superior sensitivity at the ilc compared to previous proposals . it is the first proposal for such a measurement at the lhc . , we project that a 250 gev run with 1 ab luminosity can measure the phase to accuracy .

we investigate the lhc and higgs factory prospects for measuring the phase in the higgs-- coupling . currently this phase can be anywhere between ( even ) and ( odd ) . a new , ideal observable is identified from an analytic calculation for the channel . it is demonstrated to have promising sensitivity at the lhc and superior sensitivity at the ilc compared to previous proposals . our observable requires the reconstruction of the internal substructure of decaying taus but does not rely on measuring the impact parameter of tau decays . it is the first proposal for such a measurement at the lhc . for the 14 tev lhc , we estimate that about 1 ab data can discriminate-even versus-odd at the level . with 3 ab , the phase should be measurable to an accuracy of . at an higgs factory , we project that a 250 gev run with 1 ab luminosity can measure the phase to accuracy .

1503.04578

i

let @xmath4 be some closed orientable surface , bordering a compact inner @xmath5 and outer @xmath6 domains . by @xmath7 we denote a subsurface of @xmath8 , which has two faces @xmath9 and @xmath10 and inherits the orientation from @xmath8 : @xmath10 borders the inner domain @xmath5 and @xmath9 borders the outer domain @xmath11 . @xmath7 has the smooth boundary @xmath12 , which is decomposed into two closed parts @xmath13 , consisting each of finite number of smooth arcs , having in common only endpoints . let @xmath14 , @xmath15 be the unit normal vector field on the surface @xmath7 and @xmath16 be the normal derivative . let us consider the laplace - beltrami operator in @xmath0 written in terms of the gnter s tangent derivatives ( see @xcite for more details ) @xmath17 let @xmath18 , @xmath19 , be the unit normal vector field on the boundary @xmath20 , which is tangential to the surface @xmath7 and directed outside of the surface . and , finally , let @xmath21 be the normal derivative on the boundary of the surface , which is the outer tangential derivative on the surface . we study the following mixed boundary value problem for the laplace - beltrami equation @xmath22 u^+(\tau)=g(\tau ) , \qquad & \tau\in\gamma_d , \\[0.2 cm ] ( \partial_{{{\boldsymbol{\nu}}}_\gamma}u)^+(\tau)=h(\tau),\qquad & \tau\in\gamma_n . \end{array}\right.\end{aligned}\ ] ] where @xmath23 and @xmath24 denote respectively the dirichlet and the neumann traces on the boundary . we need the bessel potential @xmath25 , @xmath26 , @xmath27 and sobolev - slobodekii @xmath28 spaces , where @xmath8 is a closed smooth surface ( without boundary ) , which contains @xmath7 as a subsurface , @xmath29 . the bessel potential space @xmath30 is defined as a subset of the space of schwartz distributions @xmath31 endowedp with the norm ( see @xcite ) @xmath32 where @xmath33 is the bessel potential and @xmath34 , @xmath35 are the fourier transformations . for the definition of the sobolev - slobodekii space @xmath36 see @xcite . the space @xmath37 coincides with the trace space of @xmath38 on @xmath8 and is known that @xmath39 for @xmath40 , @xmath3 ( see @xcite ) . we use , as common , the notation @xmath41 and @xmath42 for the spaces @xmath43 and @xmath44 ( the case @xmath45 ) . the spaces @xmath46 and @xmath37 are defined by a partition of the unity @xmath47 subordinated to some covering @xmath48 of @xmath8 and local coordinate diffeomorphisms ( see @xcite for details ) @xmath49 the space @xmath50 is defined as the subspace of @xmath46 of those functions @xmath51 , which are supported in the closed sub - surface @xmath52 , whereas @xmath53 denotes the quotient space @xmath54 , and @xmath55 is the complemented sub - surface . for @xmath56 the space @xmath53 can be identified with the space of those distributions @xmath57 on @xmath58 which admit extensions @xmath59 , while @xmath53 is identified with the space @xmath60 , where @xmath61 denotes the restriction from @xmath8 to the sub - surface @xmath7 . it is worth noting that for an integer @xmath62 the sobolev spaces @xmath63 and @xmath64 coincide and the equivalent norm is defined with the help of the gnter s derivatives ( see @xcite ) : @xmath65^{\frac1p } , \quad\mbox { where } \quad { { \mathcal d}}^\alpha:={{\mathcal d}}^{\alpha_1}_1{{\mathcal d}}^{\alpha_2}_2{{\mathcal d}}^{\alpha_3}_3\ ] ] and the gnter s derivatives @xmath66 are defined in . let us also consider @xmath67 , a subspace of @xmath68 , orthogonal to @xmath69 @xmath70 consists of those distributions on @xmath8 , belonging to @xmath68 which have their supports just on @xmath20 and @xmath68 can be decomposed into the direct sum of subspaces : @xmath71 the space @xmath72 is non - empty ( see @xcite ) and excluding it from @xmath68 is needed to make bvps uniquelly solvable ( cf . @xcite and the next theorem [ t0.1 ] ) . the lax - milgram lemma applied to the bvp gives the following result . [ t0.1 ] the bvp has a unique solution in the classical weak setting : @xmath73 from theorem [ t0.1 ] we can not even conclude that a solution is continuous . if we can prove that there is a solution @xmath74 for some @xmath75 , we can enjoy even a hlder continuity of @xmath76 . it is very important to know maximal smoothness of a solution as , for example , while designing approximation methods . to this end we will investigate the solvability properties of the bvp in the following non - classical setting @xmath77 and find necessary and sufficient conditions of solvability . note , that the constraint @xmath2 is necessary to ensure the existence of the trace @xmath23 on the boundary . to formulate the main theorem of the present work we need the following definition . [ t0.2 ] the bvp , is fredholm if the homogeneous problem @xmath78 has a finite number of linearly independent solutions and only a finite number of orthogonality conditions on the data @xmath79 ensure the solvability of the bvp . we prove below the following theorem ( see the concluding part of 5 ) . [ t0.3 ] let @xmath3 , @xmath2 . the bvp is fredholm in the non - classical setting if and only if : @xmath80 in particular , the bvp has a unique solution @xmath76 in the non - classical setting if @xmath81 note , that conditions and are independent of the parameter @xmath82 . the proof of the foregoing theorem [ t0.3 ] in [ sect5 ] is based on the theorem [ t0.3a ] and theorem [ t0.4 ] . [ t0.3a ] let @xmath3 , @xmath2 . let @xmath83 and @xmath84 be some fixed extensions of the boundary data @xmath85 and @xmath86 ( non - classical formulation ) , initially defined on the parts of the boundary @xmath13 . a solution to the bvp is represented by the formula @xmath87 here @xmath88 , @xmath89 and @xmath90 are the newton s , double and single layer potentials , defined below @xmath91see @xmath92 and @xmath93 , @xmath94 in are solutions to the following system of pseudodifferential equations @xmath95 { \displaystyle}\frac12\psi_0+r_d{{\boldsymbol w}}^*_{\gamma,0 } \psi_0-r_d{{\boldsymbol v}}_{\gamma,+1}\varphi_0=h_0\qquad & \text{on}\quad\gamma_d , \end{array}\right . \end{array}\\[2 mm ] \label{e0.7 } \begin{array}{c } \varphi_0\in{\widetilde}{{\mathbb{w}}}^{s-1/p}_p(\gamma_n),\quad \psi_0\in{\widetilde}{{\mathbb{w}}}^{s-1 - 1/p}_p(\gamma_d),\\[3 mm ] g_0\in{\mathbb{w}}^{s-1/p}_p(\gamma_n),\qquad h_0\in{\mathbb{w}}^{s-1 - 1/p}_p(\gamma_d ) , \end{array}\end{aligned}\ ] ] where @xmath96 and @xmath97 are given functions and the participating pseudodifferential operators are defined in 1 below . vice versa : if @xmath76 is a solution to the bvp , @xmath98 , @xmath99 and @xmath100 , @xmath101 are some fixed extensions of @xmath102 anf @xmath103 to @xmath20 , then @xmath104 , @xmath105 are solutions to the system . the system of boundary pseudodifferential equations has a unique pair of solutions @xmath106 and @xmath107 in the classical setting @xmath45 , @xmath108 . the proof of theorem [ t0.3a ] is exposed in [ sect1 ] . for the system we can remove the constraint @xmath2 and prove the following result for arbitrary @xmath109 . [ t0.4 ] let @xmath3 , @xmath110 . the system of boundary pseudodifferential equations is fredholm in the sobolev - slobodekii space setting @xmath111 g_0\in{\mathbb{w}}^r_p(\gamma_n),\qquad h_0\in{\mathbb{w}}^{r-1}_p(\gamma_d ) \end{array}\end{aligned}\ ] ] and also in the bessel potential space setting @xmath112 g_0\in{\mathbb{h}}^r_p(\gamma_n),\qquad h_0\in{\mathbb{h}}^{r-1}_p(\gamma_d ) \end{array}\end{aligned}\ ] ] if the following condition holds : @xmath113 in particular , the system has a unique solution in both settings and if : @xmath114 the proof of the foregoing theorem [ t0.4 ] in [ sect5 ] is based on the auxiliary theorem [ t0.5 ] . to formulate the theorem consider the following model system of boundary integral equations ( bies ) @xmath115 \psi(t ) + { { \boldsymbol k}}^1_{-1}{\varphi}(t)=h(t),\qquad & t\in{\mathbb{r}}^+,\end{array}\right.\\[2 mm ] & & \varphi,\psi\in{\widetilde}{{\mathbb{w}}}^{s-1 - 1/p}_p({\mathbb{r}}^+),\qquad g,\ ; h\in{\mathbb{w}}^{s-1 - 1/p}_p({\mathbb{r}}^+ ) , \nonumber\end{aligned}\ ] ] where @xmath116 is a mellin convolution operator with the kernel homogeneous of order @xmath117 ( see @xcite ) . [ t0.5 ] let @xmath3 , @xmath110 . + the system of boundary pseudodifferential equations is fredholm in the sobolev - slobodekii and bessel potential space settings if the system of boundary integral equations is locally invertible at @xmath118 in the sobolev - slobodekii @xmath119 and the bessel potential space @xmath120 settings , respectively . [ r0.7 ] theorem [ t0.5 ] is proved at the end of 1 . for the proof we apply a quasi - localization of the bvp with some model bvps on the half space ( see lemma [ l1.4 ] and lemma [ l1.5 ] ) . the constraint @xmath110 is due to this approach , since the boundary value problems are involved . in a forthcoming paper will be proved directly the local quasi - equivalence of the equation and the system at the points where the dirichlet and neumann boundary conditions collide and some simpler equations , which are uniquely solvable , at all other points . then the constraint @xmath110 can be dropped and replaced by @xmath109 . correspondingly , theorem [ t0.4 ] is also valid for all @xmath109 and the condition acquires the form @xmath121 a quasi - localization means `` freezing coefficients '' and `` rectifying '' underling contours and surfaces . for details of a quasi - localization we refer the reader to the papers @xcite and @xcite , where the quasi - localization is well described for singular integral operators and for bvps , respectively . we also refer to @xcite , where is exposed a short introduction to quasi - localization . in the present case under consideration we get 3 different model problems by localizing the mixed bvp to : * inner points of @xmath7 . * inner points on the boundary @xmath122 and @xmath123 . * points of the boundary @xmath20 where different boundary conditions collide ( endpoints of @xmath123 and @xmath122 ) . the model bvps obtained by a quasi - localization , are well investigated in the first two cases and such model problems have unique solutions without additional constraints . in the third case we get a mixed bvp on the half plane for the laplace equation ( cf . the system is related to this model mixed problem just as bvp is related to the system ( cf . lemma [ l1.5 ] below ) . the investigation of the boundary integral equation system is based on recent results on mellin convolution equations with meromorphic kernels in bessel potential spaces ( see r. duduchava @xcite , r. duduchava and v. didenko @xcite ) . the symbol @xmath124 of the system is a continuous function on some infinite rectangle @xmath125 and is responsible for the fredholm property and the index of the system . this provides necessary and sufficient conditions for the fredholm property of which is then used to prove the solvability of the original bvp in the non - classical setting . a rigorous analysis of solvability of the above and similar problems with dirichlet , neumann , mixed and impedance boundary condition for the helmholtz and other other elliptic equations are very helpful for a general understanding of elliptic boundary value problems in conical domains ( see @xcite ) . in @xcite the authors suggest another approach to the investigation of the model mixed problem for the helmholtz equation : they write explicit formulae for a solution with two different methods . but the setting is classical only ( the case @xmath45 ) and the approach can not be applied to the non - classical setting . other known results are either very limited to special situations such as the rectangular case @xcite or apply rather sophisticated analytical methods @xcite , or are missing a precise setting of appropriate functional spaces ( see , e.g. , @xcite ) . for the historical survey and for further references we recommend @xcite . there is another approach , which can also be applied is the limiting absorption principle , which is based on variational formulation and lax - milgram lemma and its generalizations . such approach is presented , e.g. , in @xcite . but again , these results are for the classical setting . in 1960 s there was suggested to solve canonical diffraction problems in sobolev spaces , based on the recent development in pseudodifferential equations in domains with corners and , more generally , with a lipschitz boundary . it was popularized by e. meister @xcite , e. meister and f .- o . speck @xcite , w.l . wendland @xcite , a. ferreira dos santos @xcite and their collaborators in the 1980 s . also see the book of vasilev @xcite with a considerable list of references . the results are also restricted to the classical setting .

we investigate the mixed dirichlet - neumann boundary value problems for the laplace - beltrami equation on a smooth bounded surface with a smooth boundary in non - classical setting in the bessel potential space for , . to the initial bvp we apply a quasi - localization and obtain a model bvp for the laplacian . mce is ivestigated in both bessel potential and sobolev - slobodekii spaces . the symbol of the obtained system is written explicitly and is responsible for the fredholm properties and the index of the system . an explicit criterion for the unique solvability of the initial bvp in the non - classical setting is derived as well . * key words : * boundary value problem , mixed boundary conditions , potential method , fredholm criteria , symbol , banach algebra of operators , mellin convolution equation , meromorphic kernel . bessel potential space , besov space . * ams subject classifications : * primary 35j57 , secondary 45e10 , 47b35 * funding : * the research was supported by shota rustaveli national science foundation grants no . 13/14 and 31/39 . -5 mm

we investigate the mixed dirichlet - neumann boundary value problems for the laplace - beltrami equation on a smooth bounded surface with a smooth boundary in non - classical setting in the bessel potential space for , . to the initial bvp we apply a quasi - localization and obtain a model bvp for the laplacian . the model mixed bvp on the half plane is reduced to an equivalent system of mellin convolution equation ( mce ) in sobolev - slobodekii space ( potential method ) . mce is ivestigated in both bessel potential and sobolev - slobodekii spaces . the symbol of the obtained system is written explicitly and is responsible for the fredholm properties and the index of the system . an explicit criterion for the unique solvability of the initial bvp in the non - classical setting is derived as well . * key words : * boundary value problem , mixed boundary conditions , potential method , fredholm criteria , symbol , banach algebra of operators , mellin convolution equation , meromorphic kernel . bessel potential space , besov space . * ams subject classifications : * primary 35j57 , secondary 45e10 , 47b35 * funding : * the research was supported by shota rustaveli national science foundation grants no . 13/14 and 31/39 . -5 mm

hep-lat0103026

c

in this paper , we have studied the @xmath56 improvement of wilson quarks on anisotropic lattices . at the tree - level we find that a certain choice of the parameters , @xmath82 @xcite , is well - behaved in the region of practical interest for charmed hadrons , namely @xmath222 , while @xmath223 is small . on the other hand , with a different choice , @xmath85 @xcite , continuum behavior is reached only for @xmath224 . with this latter choice a non - relativistic interpretation @xcite is still possible , but a mass - independent renormalization , which was proposed in ref . @xcite , is obstructed . the choice @xmath82 also simplifies tree - level @xmath56 improvement . the action does not require separate temporal and spatial hopping parameters . the currents require mass - dependent matching factors , but no intrinsically dimension - four terms . we therefore have started to examine the behavior of this choice at the one - loop level . we have computed the one - loop contributions to the rest mass and to the matching factors of the vector and axial vector currents . the matching factors depend significantly on @xmath24 . a more critical observation is that they are well approximated by taylor expansions @xmath225 for @xmath226 and @xmath227@xmath228 . this region encompasses the one suitable for the charmed quark with currently available computer resources . there are several issues that remain to be studied . the first is to compute the one - loop corrections to the ratio of hopping parameters @xmath46 , the clover coefficients @xmath32 and @xmath47 , and dimension - four terms in the currents . the calculation of @xmath46 is especially difficult , because it requires the one - loop kinetic mass @xmath229}$ ] . as at the tree - level , it is crucial to compute the full mass dependence , so one can check whether low - order taylor expansions work well for @xmath222 . only with the full mass dependence can one check whether @xmath24 , which comes with the couplings in the action , and @xmath223 , which also comes from the on - shell condition , come together to form @xmath230 . if not , then one could proceed with a non - perturbative calculation of @xmath57 , @xmath58 , @xmath59 , @xmath60 , etc . a more practical problem is to define renormalized couplings . the scale - setting scheme of brodsky , lepage , and mackenzie ( blm ) is usually a good way to absorb the dominant part of two- and higher - order contributions @xcite . on an anisotropic lattice , it may make sense to define separate scales for temporal and spatial gluons . these results are of interest in any case : even if anisotropic lattice calculations require a non - relativistic interpretation for heavy quarks , anisotropy remains a useful tool for improving the signal - to - noise ratio . finally , after these problems are resolved , it will be important to combine the results with numerical simulation data to obtain the matrix elements relevant to experimental measurements of charmed hadrons .

we investigate the symanzik improvement of the wilson quark action on anisotropic lattices . taking first a general action with nearest - neighbor and clover interactions , we study the mass dependence of the ratio of the hopping parameters , the clover coefficients , and an improvement coefficient for heavy - light vector and axial vector currents . we show how tree - level improvement can be achieved . for

we investigate the symanzik improvement of the wilson quark action on anisotropic lattices . taking first a general action with nearest - neighbor and clover interactions , we study the mass dependence of the ratio of the hopping parameters , the clover coefficients , and an improvement coefficient for heavy - light vector and axial vector currents . we show how tree - level improvement can be achieved . for a particular choice of the spatial wilson coupling , the results simplify , and improvement is possible . ( here is the bare quark mass and the temporal lattice spacing . ) with this choice we calculate the renormalization factors of heavy - light bilinear operators at one - loop order of perturbation theory employing the standard plaquette gauge action .

1511.01450

i

the low - energy effective theory describing the dynamics of a collection of wrapped d - branes at small separations is given by quiver quantum mechanics @xcite . these d - branes form bound states and , depending on parameters , can either be well separated or on top of each other . when well separated , the theory is on the coulomb branch , parametrized by a macroscopic vev @xmath2 for the fields describing the separation between the branes . conversely , when the branes are on top of each other , the theory is on the higgs branch , parametrized by a large vev @xmath3 for the field representing the light stretched string mode between the branes . in this paper we will consider the simplest example of quiver quantum mechanics , describing the dynamics of a pair of d - branes with a single light stretched string mode between them . being a supersymmetric theory , the relative separation between the branes lives in a vector multiplet @xmath4 and the lightest stretched string mode lives in a chiral multiplet @xmath5 . these models have proven fruitful playgrounds for the study of supersymmetric black hole bound states at weak coupling @xcite . however , one obvious question remains unsolved : what is the structure of the bound state wavefunctions ? while supersymmetry allows us to make some statements about the structure of the ground states , the full ground state wavefunctions remain unknown @xcite . this is because , while simpler than the full schrdinger equation , the bps equations remain too difficult to solve analytically . the situation is worse for the excited states , where we have to abandon the crutch of supersymmetry altogether . it is our interest in this paper to construct the bps and excited states of this model . we will not be able to do so analytically , but there exist numerical methods to compute the eigenspectra of differential operators on a finite domain , see e.g. @xcite . using these techniques we numerically solve the schrdinger equation , plot bound state wavefunctions , and determine their energies . in our analysis , we find that the physics is governed by a dimensionless quantity @xmath1 defined in ( [ nu ] ) , which is inversely proportional to the fayet - iliopoulos parameter @xmath6 . the quantity @xmath1 dictates whether the higgs branch or the coulomb branch is the dominant description of the dynamics . in studying the wavefunctions and their dependence on @xmath1 , we uncover that the two branches never quite decouple and the wavefunction has nonzero support in both branches for any value of @xmath1 . since this model is supersymmetric , much work has gone into studying features of the ground states . for example , the ground state wavefunction was previously studied in a born - oppenheimer approximation @xcite , which involves splitting the hamiltonian @xmath7 into a vector part @xmath8 and a chiral part @xmath9 and putting the chiral degrees of freedom in the harmonic oscillator ground state of @xmath9 . one then ` integrates out ' the chiral multiplet and solves the effective supersymmetric quantum mechanics on the vector multiplet degrees of freedom . this approximation is self - consistent if we are on the coulomb branch , that is if the ground state is such that the branes are well separated , and is only valid for large values of @xmath10 . one easy way to see this is that integrating out the chiral degrees of freedom generates a nontrivial metric on the moduli space of the vector multiplet , which appears in the effective lagrangian as @xmath11 the moduli space metric is singular at the origin where the approximation breaks down . the structure of the ground state wavefunction near the origin has previously been out of reach , and herein we fill this gap . interestingly the effective quantum mechanics on the coulomb branch of this model has an enhanced symmetry , allowing for a full determination of the spectrum on the coulomb branch @xcite . in computing excited state wavefunctions we numerically compute the energy gap between the supersymmetric ground state and the first excited state as a function of @xmath1 . we verify that the gap approaches the analytically determined value on the coulomb branch as we increase @xmath1 . recently @xcite the witten index @xmath12 , twisted by a combination of global symmetries , has been computed for various quivers ; and for the model of interest in this paper , @xmath13 ( depending on the sign of the fayet - iliopoulos parameter @xmath6 ) . we indeed find a single bosonic ground state in the correct representation of the global symmetry , meaning our result can be thought of as the first _ in silico _ experimental confirmation of the mathematically predicted bps spectrum of the model . moreover , using the methods described in this note we have access to more than the count of bps states . this includes plots of the ground state and excited state wavefunctions , the gap with the first excited state as well as various field expectation values . the numerical methods we describe will hopefully make the study of the quantum nature of weakly coupled black hole bound states more tractable , and can potentially provide a useful experimental testing ground for other witten index calculations . the organization of this paper is as follows : in section [ setup ] we present the supercharges and symmetry generators of the theory . these fix the wavefunctions up to their dependence on radial variables . we also give a quick review of the stringy interpretation of this model . in section [ groundstatesec ] we provide the bps equations obeyed by the ground state wavefunction , suitably reduced via symmetries . we also plot numerical solutions ( bps and non - bps ) for the full schrdinger problem in the singlet sector of the bosonic symmetry group . we study the dependence of the bps and non - bps wavefunctions on @xmath1 . we conclude in section [ outlook ] . we have collected formulae in appendices [ sphereapp ] and [ wavefunceqs ] and review the born - oppenheimer approximate ground state wavefunction in appendix [ boapp ] .

we numerically construct the bps and non - bps wavefunctions of an quiver quantum mechanics with two abelian nodes and a single arrow . this model captures the dynamics of a pair of wrapped d - branes interacting via a single light string mode . we also numerically compute the energy gap between the ground state and the first excited states as a function of . an expression for the the gap , computed on the coulomb branch , matches nicely with our numerics at large but deviates at small where the higgs branch becomes the relevant description of the physics . in the appendix , we provide the schrdinger equations fully reduced via symmetries which , in principle , allow for the numerical determination of the entire spectrum at any point in moduli space . for the ground states , this numerical determination of the spectrum can be thought of as the first _ in silico _ check of various witten index calculations . mit - ctp-4731 susy in silico : numerical d - brane bound state spectroscopy tarek anous _ center for theoretical physics , massachusetts institute of technology , cambridge , ma 02139 , usa _

we numerically construct the bps and non - bps wavefunctions of an quiver quantum mechanics with two abelian nodes and a single arrow . this model captures the dynamics of a pair of wrapped d - branes interacting via a single light string mode . a dimensionless parameter , which is inversely proportional to the the fayet - iliopoulos parameter , controls whether the bulk of the wavefunctions are supported on the higgs branch or the coulomb branch . we demonstrate how the bps and excited states morph as is tuned . we also numerically compute the energy gap between the ground state and the first excited states as a function of . an expression for the the gap , computed on the coulomb branch , matches nicely with our numerics at large but deviates at small where the higgs branch becomes the relevant description of the physics . in the appendix , we provide the schrdinger equations fully reduced via symmetries which , in principle , allow for the numerical determination of the entire spectrum at any point in moduli space . for the ground states , this numerical determination of the spectrum can be thought of as the first _ in silico _ check of various witten index calculations . mit - ctp-4731 susy in silico : numerical d - brane bound state spectroscopy tarek anous _ center for theoretical physics , massachusetts institute of technology , cambridge , ma 02139 , usa _

0705.4319

i

hercules x-1 ( discovered by tananbaum et al . 1972 ) is one of the most frequently observed x - ray binary systems . the intermediate mass of the donor star , @xmath5 m@xmath6 , leads to a wealth of behavior seen in both low mass and high mass systems . although we suspect the mass flow is primarily through roche lobe overflow giving rise to an accretion disk , as in the low mass x - ray binaries ( lmxb ) , the variable p cygni profiles observed in uv lines suggest either a transient stellar wind or a stellar wind that is photoionized in some regions ( boroson , kallman , & vrtilek 2001 ) . there are theoretical grounds as well to expect that winds may arise in this system when x - rays from the neutron star heat the surface of the normal star or accretion disk ( arons 1973 ; davidson & ostriker 1973 ; basko et al . 1977 ; london , mccray , & auer 1981 ; begelman , mckee , & shields 1983 ; begelman & mckee 1983 ) . most lmxbs , including z - sources and atoll - sources , do not show persistent pulsations , perhaps because they have neutron stars with low magnetic fields . these sources thus lack an essential window into their kinematics . light travel - time delays in the 1.24 second her x-1 x - ray pulsation period determine that the neutron star is in a nearly circular 1.7 day orbit with semimajor axis a@xmath7 seconds . the orbital period is slowly lengthening from mass loss in the system ( deeter et al . 1991 ) . fortuitously , her x-1 has high orbital inclination ( @xmath8 degrees ) , allowing total x - ray eclipses to constrain the size of the donor star , and the high galactic latitude offers only small reddening , allowing the system to be observed at crucial uv and even euv wavelengths . the system s most mysterious variability is the 35-day x - ray high and low cycle ( giacconi et al . 1973 ) . for 811 days ( the `` main - on state '' ) of this cycle her x-1 emits x - rays with a @xmath9 erg s@xmath10 total x - ray luminosity . for a @xmath11 day cycle ( the `` short - on state '' ) , halfway through the 35 day cycle , her x-1 s output is several times lower . for the remaining portion of the 35 days , the off state , the observed x - ray flux is only several percent of that during the main - on state . the behavior of the optical flux and x - ray spectrum over the 35-day period shows that the 35-day variability is not isotropic . instead , the x - ray emission is merely obscured by a warped accretion disk whose shape precesses " globally throughout a 35-day period . the optical emission from the normal star varies with the 1.7 day binary orbit . its spectral type is variously classified as a through f , because x - rays from the neutron star heat the surface of the facing side of the roche lobe . the spectral type continues to change with the orbital period throughout the longer 35-day cycle , implying that x - rays continue to heat the surface of hz her . the transition between the off and main - on state is only a few hours , whereas the short - on state is entered into more gradually . this indicates that regions of the disk at different radii may obscure the central source , in which case the disk must precess globally with the same period . the variation in pulse shape with the long - term period is sometimes taken to imply that the innermost edge of the accretion disk precesses with the 35 day period , obscuring local regions on the neutron star ( scott , leahy , & wilson , 2000 ) . the physical cause of a 35-day global disk precessing disk warp , or a definitive relation between system parameters and disk precession period , remains unknown and is a major goal of investigations of the system . it has been proposed , but not confirmed , that reprocessed x - ray radiation pressure ( maloney & begelman 1997 ; wijers & pringle 1999 ) or an x - ray driven disk wind and corona ( meyer , meyer - hofmeister , 1984 ) could propagate the warped shape in the disk . similar " long - term periods have been observed in other x - ray binaries , including lmc x-4 . her x-1 serves as a prototype of this behavior . x - ray dips present another variability that has not been explained definitively . crosa & boynton ( 1980 ) showed that the dips recur every 1.65 days , near , but not at , the 1.62 day beat period between orbital and precessional periods . this was confirmed by long - term x - ray observations ( scott & leahy , 1999 ) . the dips are thought to be associated with the gas stream between the stars , but there is no consensus explanation for them or their period . as befits such a frequently observed , enigmatic , yet prototypical source , the methods of spectroscopy have been applied to the system at a wide range of wavelengths . these studies have borne fruit with determinations of the system s elemental abundances and the orbital motion of the normal star ( the neutron star s pulse delays indicate only the motion of one component of the binary ) . high resolution spectroscopy in the x - ray range using _ xmm _ ( jimenez - garate et al . 2002 ) shows a multitude of x - ray emission lines presumably from an accretion disk corona . the line ratios indicate that the gas is enhanced from cno processing from a massive progenitor . jimenez - garate et al . ( 2005 ) confirmed the cno enhancement by observations of more than two dozen emission lines originating in an accretion disk corona . model predictions of the disk corona s response to illumination by the central x - ray source are in reasonable agreement with the observed fluxes for low and moderate z elements ( o through s ) , but the fexxv , fexxvi lines are several times brighter than predicted ( jimenez - garate et al . 2005 ) . from observations of optical absorption lines ( reynolds et al . 1997 ) , and from optical pulsations that result from the x - rays from the neutron star periodically striking the surface of the normal star ( middleditch & nelson 1976 ) , we know that the normal star has a mass @xmath12 m@xmath6 and the neutron star has a mass 1.5@xmath13 m@xmath6 . the optical spectrum shows absorption lines from hz her , emission at the bowen blend near 4640 ( see schachter , filippenko , & kahn 1989 for a discussion of the formation of these lines ) , and emission from heii@xmath14 . still et al . ( 1997 ) performed doppler tomography on these optical emission lines but could not conclude that they arose in a symmetric accretion disk . the uv spectrum is crucial for an understanding of the accretion disk and the x - ray illuminated face of hz her . the continuum emission from both the illuminated star and the disk peaks in the uv and the disk contributes a greater fraction than at optical wavelengths ( cheng , vrtilek , & raymond 1995 ) . from models of the variable uv continuum as observed with the international ultraviolet explorer ( _ iue _ , vrtilek & cheng 1996 ) showed that a change in the accretion disk precession could explain an anomalously low period of x - ray emission . to study global accretion one should observe gas at ionization stages and temperatures resulting from x - ray illumination of the disk , and one would need to resolve doppler - shifted velocities that correspond to the orbital motion of the neutron star ( 160 km s@xmath10 ) , motion in the accretion disk ( expected to be @xmath15 km s@xmath10 at the edge of the disk ) , and the velocities expected in a disk or stellar wind ( @xmath16 km s@xmath10 ) . the strong uv resonance lines from nv , siiv , and civ , first seen with _ iue _ ( howarth & wilson 1983b ) , presumably result from photoionization of the accretion disk and hz her ( raymond 1993 , ko & kallman 1994 ) . these lines are much stronger than the strongest optical high ionization line , heii@xmath14 , both in absolute flux and equivalent width . observations with the faint object spectrograph ( fos ) aboard the hubble space telescope ( _ hst _ ) showed that these lines are still present at a few percent of maximum brightness during mid - eclipse when the disk and heated star should be entirely obscured ( anderson et al . the source of this emission may be an expanding wind . the goddard high resolution spectrograph ( ghrs ) on _ hst _ first resolved these emission lines at @xmath17 km s@xmath10 resolution ( boroson et al . , 1996 ) to discern variable broad and narrow emission components . the _ hst _ space telescope imaging spectrograph ( stis ) confirmed that the resonance lines have at least two components ( vrtilek et al . a broad component arises on the accretion disk while a narrow line component may be associated with hz her . prior to the stis observations , the accretion disk had only been observed as it contributed to the continuum light curve or obscured the central x - rays . during eclipse ingress and egress , the broad lines seen with stis behaved as expected for lines from an accretion disk rotating prograde with the orbital direction . the blue edge of the line was obscured first in eclipse ingress and appeared first in egress . observations with the far ultraviolet spectroscopic explorer ( _ fuse _ ) are complementary to the existing _ hst _ observations . long - term variability in the system makes it impossible to combine rigorously data from different epochs . analyzed separately , however , the _ fuse _ wavelength range of 900 - 1200 offers similar advantages and powerful consistency checks to analysis of the _ hst _ bandpass of 1200 - 1700 . both wavelength ranges have strong resonance lines that respond to x - ray photoionization . observations with the hopkins ultraviolet telescope ( _ hut _ ) showed that the ovi doublet has flux comparable to the nv doublet , the brightest near uv line ( boroson et al . 1997 ) . the resonance line doublets offer optical depth information through their doublet ratios , but if the lines are as broad as the doublet separation , it may be impossible to determine the individual contribution of each line where they overlap . the separations of civ@xmath18 , siiv@xmath19 , nv@xmath20 , ovi@xmath21 , svi@xmath22 correspond to doppler shifts of 500 , 1900 , 960 , 1650 , 3600 km s@xmath10 , respectively . the far uv lines ovi and svi compare favorably with the near uv lines ( i.e. have greater separation and are less likely to overlap ) , except for siiv , which suffers from confusion with an oiv blend near 1400 . for the present observations , we observed an entire 1.7 day binary orbit with _ fuse_. a major goal of this program was to apply the method of doppler tomography ( 7 ) , which has the advantage of diagnosing the accretion flow of different systems without bringing to bear more than a few assumptions . table 1 shows the measured physical parameters of the system and parameters we adopt for our models .

we observed an entire 1.7 day orbit of the x - ray binary hercules x-1 with the far ultraviolet spectroscopic explorer ( _ fuse _ ) . changes in the ovi line profiles through eclipse ingress and egress indicate a keplerian accretion disk spinning prograde with the orbit . these observations may show the first double - peaked accretion disk line profile to be seen in the hercules x-1 system . simulations show that the bright spot is too far offset from the roche lobe to result from uneven x - ray heating of its surface . the absence of disk signatures in the tomogram can be reproduced in simulations which include absorption from a stellar wind . = 1 based on observations made with the nasa - cnes - csa far ultraviolet spectroscopic explorer .

we observed an entire 1.7 day orbit of the x - ray binary hercules x-1 with the far ultraviolet spectroscopic explorer ( _ fuse _ ) . changes in the ovi line profiles through eclipse ingress and egress indicate a keplerian accretion disk spinning prograde with the orbit . these observations may show the first double - peaked accretion disk line profile to be seen in the hercules x-1 system . doppler tomograms of the emission lines show a bright spot offset from the roche lobe of the companion star hz her , but no obvious signs of the accretion disk . simulations show that the bright spot is too far offset from the roche lobe to result from uneven x - ray heating of its surface . the absence of disk signatures in the tomogram can be reproduced in simulations which include absorption from a stellar wind . we attempt to diagnose the state of the emitting gas from the ciii , ciii , and niii emission lines . the latter may be enhanced through bowen fluorescence . = 1 based on observations made with the nasa - cnes - csa far ultraviolet spectroscopic explorer . fuse is operated for nasa by the johns hopkins university under nasa contract nas5 - 32985

astro-ph0406272

i

the galactic x - ray transient h 1743@xmath0322 was classified as a black hole candidate by white & marshall ( 1984 ) based on its x - ray spectral characteristics . the source was discovered during its 19771978 outburst with _ heao-1 _ and _ ariel v _ ( doxsey et al . 1977 , kaluzienski & holt 1977 ) . the next convincing detection of h 1743@xmath0322 was made on 2003 march 21 with _ integral _ ( revnivtsev et al . 2003 ) and later with _ rxte _ ( markwardt & swank 2003a ) , earning the source the additional designations igr j17464@xmath03213 and xte j1746@xmath0322 . markwardt & swank ( 2003b ) pointed out that although the original _ _ detection reported two possible locations , it was very likely that the potentially new transient was actually h 1743@xmath0322 . following its 2003 re - activation in x - rays , the source was extensively monitored in optical , ir , and radio bands ( see , e.g. , rupen , mioduszewski , & dhawan 2003 ; steeghs et al . 2003 and ref . therein ) . two discoveries made during the 2003 outburst of h 1743@xmath0322 have provided additional evidence that this source harbors a black hole primary . first , homan et al . ( 2003 ) reported the detection of quasi - periodic oscillations ( qpos ) at 240 hz in the _ rxte _ x - ray flux . this frequency is typical of black hole candidates , and is tantalizingly close to the keplerian orbital frequency expected at the innermost stable circular orbit ( isco ) around a schwarzschild black hole ( @xmath6 ) . moreover , there is evidence for qpos at 160 hz , and therefore in the 2:3 ratio which has recently emerged in a few galactic black hole systems ( gro j1655@xmath040 : strohmayer 2001a , grs 1915@xmath7105 : strohmayer 2001b , xte j1550@xmath0564 : miller et al . 2001 , and possibly xte j1650@xmath0500 : homan et al . . the apparent 2:3 pairing may be evidence of a parametric epicyclic resonance in black hole accretion disks ( see , e.g. , abramowicz & kluzniak 2003 ) . however , there is evidence that black hole qpo pairs may drift out of a 2:3 ratio ( miller et al.2001 , remillard et al . 2002 ) , and the 2:3 ratio described by the resonance model for so - called khz qpos in neutron star systems is at odds with data ( miller 2003 ) . second , radio observations have revealed a jet in h 1743@xmath0322 , moving at @xmath8 ( rupen , mioduszewski , & dhawan 2004 ) . on the basis of this relativistic jet , h 1743@xmath0322 may be termed a `` microquasar '' . indeed , the x - ray spectroscopic results presented in this paper will further establish close similarities between h 1743@xmath0322 and microquasars like gro j1655@xmath040 and grs 1915@xmath7105 , in addition to their radio jet and x - ray timing similarities . although highly relativistic jets are not a unique black hole primary signature ( fender et al . 2004 ) , such jets are are more commonly observed in black hole systems than in neutron star systems . we observed h 1743@xmath0322 simultaneously with _ chandra _ and _ rxte _ on four occasions during the bright phase of its 2003 outburst , as part of our larger multi - wavelength black hole observing program . herein , we report the results of an analysis of the x - ray spectra resulting from these observations . most importantly , the _ chandra_/hetgs spectra reveal variable he - like fe xxv and h - like fe xxvi resonance absorption lines , which may be due to absorption in an outflow . variable , seyfert - like `` warm - absorber '' outflows have recently been detected in the x - ray spectra of the galactic black holes xte j1650@xmath0500 and gx 339@xmath04 ( miller et al.2004a , miller et al . we examine the absorption in h 1743@xmath0322 within the context of these recent results . the results of an analysis of the x - ray timing properties of the _ rxte _ observations will be reported in a companion paper ( homan et al.2004 ) . the results of radio observations made in partial support of this program will be reported in a separate paper by rupen et al . ( 2004 , in prep . ) , and results from optical and ir observations will be reported in steeghs et al . ( 2004 , in prep . )

chandra _ and _ rxte _ on four occasions . the _ chandra_/hetgs spectra reveal narrow , variable ( he - like ) fe xxv and ( h - like ) fe xxvi resonance absorption lines . in the first observation this wind may be a hot precursor to the seyfert - like , outflowing `` warm absorber '' we discuss these findings in the context of ionized fe absorption lines found in the spectra of other galactic sources , and connections to warm absorbers , winds , and jets in other accreting systems . [ firstpage ]

we observed the bright phase of the 2003 outburst of the galactic black hole candidate h 1743 in x - rays simultaneously with _ chandra _ and _ rxte _ on four occasions . the _ chandra_/hetgs spectra reveal narrow , variable ( he - like ) fe xxv and ( h - like ) fe xxvi resonance absorption lines . in the first observation , the fe xxvi line has a fwhm of km / s and a blue - shift of km / s , suggesting that the highly ionized medium is an outflow . moreover , the fe xxv line is observed to vary significantly on a timescale of a few hundred seconds in the first observation , which corresponds to the keplerian orbital period at approximately ( where ) . our models for the absorption geometry suggest that a combination of geometric effects and changing ionizing flux are required to account for the large changes in line flux observed between observations , and that the absorption likely occurs at a radius less than for a black hole . viable models for the absorption geometry include cyclic absorption due to an accretion disk structure , absorption in a clumpy outflowing disk wind , or possibly a combination of these two . if the wind in h 1743 has unity filling factor , the highest implied mass outflow rate is 20% of the eddington mass accretion rate . this wind may be a hot precursor to the seyfert - like , outflowing `` warm absorber '' geometries recently found in the galactic black holes gx 339 - 4 and xte j1650 . we discuss these findings in the context of ionized fe absorption lines found in the spectra of other galactic sources , and connections to warm absorbers , winds , and jets in other accreting systems . [ firstpage ]

astro-ph0406272

c

the high equivalent neutral hydrogen column density along the line of sight is consistent with a central galactic distance to h 1743@xmath0322 . for distances greater than @xmath853 kpc , the highest 0.510.0 kev unabsorbed flux observed ( see table 2 ) exceeds the eddington limit for a neutron star . moreover , the parameters obtained from broad - band spectral fits to h 1743@xmath0322 with the phenomenological mcd plus power - law model are typical for galactic black hole systems in bright phases ( for a review , see mcclintock & remillard 2003 , see also homan et al . 2001 ) . these facts and the presence of 240 hz qpos in h 1743@xmath0322 make it a near certainty that h 1743@xmath0322 harbors a black hole primary . however , the apparent absence of reflection features in the spectrum , even in phases with 1 ) a strong hard component , 2 ) a small inferred inner disk radius , and 3 ) high - frequency qpos , is at odds with the majority of galactic black holes . it is likely that the low reflection fraction obtained is due to a combination of astrophysical effects ( ionization and disk inclination ) and the limitations of our reflection modeling . zdziarski et al.(2003 ) have shown that when the disk is fully ionized , the reflected spectrum is nearly featureless , and might therefore be confused with a zero reflection scenario . high disk ionization tends to smooth - out reflection features ; the dips we have detected in the lightcurve of _ chandra _ observation 3 and in additional public _ rxte _ observations of h 1743@xmath0322 point to a system viewed at high inclination . at high inclinations , doppler shifts become very important , and tend to smear - out reflection features . the cdid reflection model we fit to _ rxte _ observation 2b indicates that the disk in h 1743@xmath0322 is likely highly ionized and seen at a high inclination . while these constraints are likely robust , it must be acknowledged that even the cdid model has limitations . it is implemented into xspec as a library of angle averaged spectra . finally , we note that if any reflection spectrum also passes through the hard x - ray emitting corona , which is likely if a corona hugs a disc viewed at high inclination , then the reflection component will be comptonized and much less detectable ( petrucci et al 2001 ) . it is possible , then , that the total reflection model the cdid model smeared with a line function ( for which inclination is a fit parameter ) does not address high inclinations correctly . were the combination of ionization and inclination effects more accurately described by our model , it is likely that a high reflection fraction would have been within errors . we have clearly detected variable , narrow fe xxv and fe xxvi resonance absorption lines in the spectra of h 1743@xmath0322 . in the recent past , similar absorption lines have been detected in the microquasars gro j1655@xmath040 ( ueda et al . 1998 ) , and grs 1915@xmath7105 ( kotani et al.2000 ) with _ asca _ , and in grs 1915@xmath7105 with _ ( lee et al . 2002 ) . unfortunately , each of these prior spectra suffered heavily from photon pile - up . the extraordinary sensitivity of the h 1743@xmath0322 spectra has allowed us to improve upon previous implications based on fe xxv / xxvi absorption lines in microquasars in two important ways : first , we have measured blue - shifts in the lines detected in observation 1 , providing direct evidence that the highly ionized absorbing geometry may be an outflow ( previously , outflows were only asserted ) . second , we have resolved the lines in two cases , and in other cases the lines are likely narrower than the instrumental resolution . the fwhm of the lines is vital for constraining the nature of the absorption . viable descriptions of the absorption geometry in h 1743@xmath0322 must be able to account for a number of our key findings , including : ( 1 ) the blue - shifts observed in the first _ observation , ( 2 ) the large line flux changes observed between observations ( in particular , the absence of absorption lines in the second _ chandra _ observation ) , and the stronger absorption seen in the dip observed during the third _ chandra _ observation , and ( 3 ) the absence of fe emission lines in the spectra , ( 4 ) the variability of the absorption lines in observation 1 on a timescale of a few hundred seconds , and ( 5 ) the absorption must occur at a radius smaller than the maximum radius at which our photoionization models suggest the lines may be produced . two geometries satisfy or partially satisfy most of these requirements : cyclic absorption in a disk structure resulting from the high inclination , and absorption in an ( potentially clumpy ) outflowing disk wind . the absence of fe emission lines in the _ chandra _ and _ rxte _ spectra suggests that the absorbing geometry is not spherical ( else , emission line photons would have scattered into our line of sight ) . indeed , modeling of the absorbers in gro j1655@xmath040 ( ueda et al.2004 ) and grs 1915@xmath7105 ( kotani et al . 2000 ) also suggested a cylindrical , perhaps pancake - like absorbing geometry . absorption in a disk structure perhaps a disk atmosphere is consistent with this constraint . however , this picture relies on relatively large covering and filling factors . the lack of emission lines in the spectra is also consistent with the absorption occurring in relatively dense clumps in a disk wind , with large covering factors but small filling factors . a clumpy wind is particularly appealing when one considers that at high mass accetion rates ( such as in the first _ chandra _ observation ) , our models for the absorber indicate that the wind would already be optically - thick at approximately @xmath86 . yet , we clearly see continuum spectral components consistent with viewing the innermost flow without significant obscuration . moreover , if qpos are tied to keplerian frequencies , the qpos detected ( homan et al . 2004 , in prep . ) indicate that emission from the innermost flow is either seen directly or scattered into our line of sight . a significant fraction of the absorption needs to occur in clumps with a high covering factor but small filling factor to be consistent with the observed continuum spectral properties and variability properties . the absence of absorption lines due to a combination of geometric effects and ionizing flux ( as implied by our photoionization models ) in the second _ chandra _ observation is easily explained by the clumpy outflow model . simply , a clump may not have intersected our line of sight to the central accretion observation within this observation . it is more difficult to explain the absence of absorption lines if the absorption occurs via cyclic absorption in a disk structure . the disk structure would have to be dampened or disappear in the second _ chandra _ observation . however , the ionizing flux is higher in this observation , and it has been shown that irradiation can induce warps in accretion disks ( see , e.g. , maloney , begelman , & pringle 1996 ) . to explain the enhanced absorption line fluxes measured within the dip in the third _ chandra _ observation , it is easy to envision a cool clump absorbing fe xxv more strongly than the hotter surrounding wind . alternatively , it is also possible that the dip observed is due to an unrelated geometric effect , and that the enhanced absorption line flux seen within the dip occurs because the wind along our line of sight is temporarily shielded and allowed to cool . if the absorption lines are instead due to absorption in a disk structure , a second warp or disk structure may need to be invoked to explain the dip itself . the variability of the fe xxv absorption line flux on timescales of a few hundred seconds in the first _ chandra _ observation can be explained by both models . indeed , this short timescale line variability is one of the most compelling arguments in favor of cyclic absorption in a disk structure . the period of the line variability broadly agrees with the keplerian orbital period at @xmath3 for a @xmath5 black hole . if h 1743@xmath0322 harbors a black hole with @xmath87 , then @xmath3 corresponds to @xmath88 cm , which is less than @xmath77 in observations 1 and 3 ( see table 5 ) . however , the observed variability can not be termed a qpo , and it is therefore less easily tied to a cyclic disk structure than if it qualified as a true qpo . if the absorption occurs in a disk structure , it would suggest that the variability itself is due to changes in the absorption in our line of sight . this is at odds with the fact that similar variability is observed in the absence of absorption lines in the second _ chandra _ observation . if the variability were due to absorption , the variability should be more pronounced at low energies ; in fact , the fractional rms of the peak at @xmath57 hz increases with higher energy ( homan et al . 2004b ; see figures 2 and 5 ) . although it is less exciting , absorption in a clumpy outflowing wind can likely account for the rapid line flux variability at least as well as cyclic absorption in a disk structure . simply , the gas ionization and recombination timescale is much faster than the variability timescale observed . if the variability is not due to a disk structure , but rather just due to fluctuations in the intensity of the central engine , an absorbing clump may be able to rapidly respond to the changing flux . finally , the blue - shifts observed in the first _ observation are more easily explained through an outflowing disk wind , than by absorption in the disk . the modest blue - shifts of the lines are a small fraction of the speed of light , and similar to the blue - shifted lines observed in gx 339@xmath04 and xte j1650@xmath0500 ( miller et al . 2004a ) , which are consistent with a disk wind . if the absorption is instead due to a cyclic disk structure , it is rather hard to explain the blue - shifts without significant fine - tuning . if the structure were a warp , its motion should be tangential to our line of sight , not into our line of sight . instead , the absorption might need to occur in a disk wind flowing outward along the disk , and further postulate that the wind density is modulated by a cyclic disk structure underlying the atmosphere . it has been shown that accretion disks may drive outflows with a significant component along the disk ( see , e.g. , murray & chiang 1997 , proga 2003 ) . on balance , absorption in an outflowing , clumpy wind appears to be the more compelling than absorption in a disk structure . both possibilities require a degree of fine - tuning to satisfy the observational constraints . a combination of these simple models may also be a viable description of the absorption geometry . the highly ionized outflow revealed in these _ chandra _ observations of h 1743@xmath0322 is very likely connected to the variable , outflowing , seyfert - like warm absorbers recently discovered in gx 339@xmath04 and xte j1650@xmath0500 ( miller et al . 2004a , miller et al.2004b ) . in those systems , the outflowing absorbers are consistent with disk winds or shell ejections . the absorbers in gx 339@xmath04 and xte j1650@xmath0500 are blue - shifted , and so clearly outflowing . in gx 339@xmath04 , he - like ne ix , mg xi , and o vii absorption lines were detected simultaneously with more strongly blue - shifted ne ii and ne iii lines , suggesting a `` clumpy '' outflow with regions of differing temperature and density . the fe xxv and fe xxvi lines in h 1743@xmath0322 also need to come from distinct regions , though the contrast needed is much less than for the wind observed in gx 339@xmath04 . the outflow observed in gx 339@xmath04 was likely detected at a lower fraction of the eddingtion mass accretion rate in that source ; its lower ionization can therefore be attributed to a lower central source flux . moreover , the outflow in gx 339@xmath04 implied an order of magnitude lower mass outflow rate than than the highest rate implied in h 1743@xmath0322 ( perhaps as high as 22% of the eddington inflow rate ) . the outflow in h 1743@xmath0322 may be clumpy , however , and have a low filling factor ; in that case , the outflow rate in h 1743@xmath0322 could be lower by a factor of @xmath89 , or more . it is plausible that the outflow in h 1743@xmath0322 is a hot , high @xmath90 precursor to the cooler , lower @xmath90 outflows observed in gx 339@xmath04 and xte j1650@xmath0500 ( miller et al . 2004a , 2004b ) . the mass outflow rate in the fourth _ chandra _ observation of h 1743@xmath0322 is far less than that in the first observation ( see table 5 ) , further indicating an evolution along the lines we suggest here . the outflows in these black hole systems are likely disk winds , and resemble the warm absorbers revealed in some seyfert galaxies ( see , e.g. , reynolds 1997 , elvis 2000 , morales & fabian 2002 ) . such absorption regions could be tied to rather equatorial disk outflows ( murray & chiang 1997 , proga 2003 ) , rather than to a highly collimated flow perpendicular to the disk . focused studies of galactic black hole winds in the future may enable us to learn about the evolution seyfert warm absorbers ; their evolution is difficult to study directly given the inherently longer timescales . it is worth addressing the possibility that a putative disk structure and the outflow are not merely coincident , but potentially causally related . for instance , a family of models has been discussed wherein jets or a wind can be launched vertically from a disk through a coupling between a spiral density wave in the disk and a rossby vortex ( tagger & pellat 1999 ; varniere , rodriguez , & tagger 2002 ; rodriguez , varniere , tagger , & durouchoux 2002 ) . it is possible , then , that the @xmath52s variations seen in h 1743@xmath0322 are due to a density structure which acts to drive jets close to the black hole , and that the same wave modulates a disk wind far from the black hole . it is not yet clear if such a mechanism could be at work in driving both jets and a disk wind , but if so , then the appearance of ionized fe absorption lines in strong jet sources like gro j1655@xmath040 , grs 1915@xmath7105 , and now h 1743@xmath0322 would be explained naturally . we note that using reduction methods exactly like those detailed for h 1743@xmath0322 , we have detected fe xxv and fe xxvi lines in a 60 ksec _ chandra_/hetgs spectrum of cygnus x-1 obtained on 3 march 2004 , when the source was in a canonical low / hard state ( miller et al.2004c , in prep . ) . cygnus x-1 is known to produce a compact , steady radio jet in this state ( stirling et al . highly ionized fe absorption lines have also been detected in the neutron star system gx 13@xmath71 ( sidoli et al . interestingly , gx 13@xmath71 is also a radio emitter in which a delay between its x - ray spectral hardness and radio brightness has been observed which is very similar to that seen in grs 1915@xmath7105 ( homan et al . 2004a ) , strongly suggesting a compact jet is also at work in this system . however , simultaneous radio observations of h 1743@xmath0322 with the vla / vlba reveal that the radio source was quite strong during the first and second _ chandra _ observations nearly 10 mjy at ghz frequencies ) but only 2 mjy during the third observation , and less than 1 mjy during the fourth observation ( rupen et al . 2004 , in prep . ) . thus , strong radio activity ( indicative of a jet ) is seen during the second _ chandra _ observation ( where absorption lines are not observed ) , and absorption lines are clearly observed in the third and fourth _ chandra _ observations when the radio flux was low . at minimum , this suggests that if jets and ionized absorption are causally related , they are not related in a simple way . both may merely be the effect of a strongly active corona , for instance . moreover , the detection of highly ionized fe absorption lines in high inclination neutron star binaries may signify that it is unlikely that ionized absorption lines are tied to jet production through a mechanism like aei . similar lines have also been seen in the persistent emission of the ( edge - on ) dipping neutron star binaries x 1254@xmath0690 ( boirin & parmar 2003 ) , x 1624@xmath0490 ( parmar et al.2002 ) , mxb 1658@xmath029 ( sidoli et al . 2001 ) and 4u 1916@xmath0053 ( boirin et al . 2004 ) . in contrast to the absorption lines in the black hole systems h 1743@xmath0322 , gx 339@xmath04 , and xte j1650@xmath0500 , the absorption lines in these edge - on neutron star binaries are not significantly blue - shifted . although the appearance of lines in both black hole and neutron star systems suggests a related physical origin , the lack of blue - shifts in the lines seen in neutron stars may hint at a different physical picture . without significant blue - shifts , it is difficult to clearly tie the lines to an outflow . the lines may not be due to absorption in an outflow , but might be due to absorption in ambient gas within the system . the absorption lines in mxb 1658@xmath029 , for instance , are not observed to vary with orbital phase or even during dipping events ( sidoli et al . 2001 ) , which suggests that the absorbing gas is not as close to the compact object as the absorbing gas in black hole systems . + we wish to thank cxc director harvey tananbaum and the cxc staff , and jean swank and the _ rxte _ staff , for making these too observations . we thank john houck for assistance with isis . jmm acknowledges helpful discussions with mike nowak , david huenemoerder , aneta siemiginowska , rob fender , herman marshall , and patrick wojdowski . jmm gratefully acknowledges support from the nsf through its astronomy and astrophysics postdoctoral fellowship program . whgl and jh gratefully acknowledge support from nasa . ds acknowledges support from the sao clay fellowship . this research has made use of the data and resources obtained through the heasarc on - line service , provided by nasa - gsfc . houck , j. c. , & denicola , l. a. , 2000 , astronomical data analysis software and systems ix , in asp conference proceedings , vol . 216 , eds . n. manset , c. veillet , & d. crabtree ; isis is available at http://space.mit.edu/cxc/isis lll instrument & start time & net exposure ( ksec ) + cxo / acis@xmath7hetgs & 2003 - 05 - 01t21:41:03 & 48.8 + rxte / pca 1a & 2003 - 05 - 01t17:00:32 & 25.2 + rxte / pca 1b & 2003 - 05 - 02t00:02:24 & 0.8 + rxte / pca 1c & 2003 - 05 - 02t00:54:08 & 3.3 + cxo / acis@xmath7hetgs & 2003 - 05 - 28t04:09:21 & 45.5 + rxte / pca 2a & 2003 - 05 - 28t05:28:48 & 1.5 + rxte / pca 2b & 2003 - 05 - 28t06:44:16 & 16.1 + rxte / pca 2c & 2003 - 05 - 28t14:28:32 & 5.6 + cxo / acis@xmath7hetgs & 2003 - 06 - 23t15:56:10 & 50.0 + rxte / pca 3 & 2003 - 06 - 23t17:05:20 & 13.3 + cxo / acis@xmath7hetgs & 2003 - 07 - 30t15:57:58 & 50.2 + rxte / pca 4a & 2003 - 07 - 30t19:48:48 & 0.6 + rxte / pca 4b & 2003 - 07 - 30t21:22:40 & 0.6 + rxte / pca 4c & 2003 - 07:31t00:32:48 & 0.6 + lllllllll observation & 1a & 1b & 1c & 2a & 2b & 2c & 3 & 4 + @xmath91 & @xmath92 & @xmath92 & @xmath92 & @xmath92 & @xmath92 & @xmath92 & @xmath92 & @xmath92 + @xmath93 ( kev ) & @xmath94 & @xmath94 & @xmath95 & @xmath95 & @xmath96 & @xmath97 & @xmath98 & @xmath99 + @xmath100 & @xmath101 & @xmath102 & @xmath103 & @xmath104 & @xmath105 & @xmath106 & @xmath107 & @xmath108 + @xmath109 & @xmath110 & @xmath111 & @xmath112 & @xmath113 & @xmath114 & @xmath115 & @xmath116 & @xmath117 + @xmath118 & @xmath119 & @xmath120 & @xmath121 & @xmath122 & @xmath123 & @xmath124 & @xmath125 & @xmath126 + @xmath127 & @xmath128 & @xmath129 & @xmath130 & @xmath131 & @xmath132 & @xmath133 & @xmath134 & @xmath135 + @xmath136 ( 3 - 100 ) & @xmath137 & @xmath137 & @xmath137 & @xmath138 & @xmath139 & @xmath140 & @xmath141 & @xmath142 + @xmath143 ( 3 - 100 ) @xmath144 & @xmath145 & @xmath146 & @xmath147 & @xmath148 & @xmath149 & @xmath150 & @xmath151 & @xmath152 + @xmath153 & @xmath154 & @xmath155 & @xmath156 & @xmath157 & @xmath158 & @xmath159 & @xmath160 & @xmath161 + @xmath136 ( 0.5 - 10 ) & @xmath162 & @xmath163 & @xmath164 & @xmath165 & @xmath166 & @xmath167 & @xmath168 & @xmath169 + @xmath143 ( 0.5 - 10 ) @xmath144 & @xmath170 & @xmath171 & @xmath172 & @xmath156 & @xmath173 & @xmath174 & @xmath175 & @xmath161 + fe k@xmath9 ( fwhm@xmath1850 kev ) ( ev ) & @xmath186 & @xmath186 & @xmath187 & @xmath188 & @xmath189 & @xmath190 & @xmath190 & @xmath191 + fe k@xmath9 ( fwhm@xmath1852 kev ) ( ev ) & @xmath191 & @xmath188 & @xmath188 & @xmath192 & @xmath193 & @xmath194 & @xmath195 & @xmath189 + lllllllllllll obs . & & theor . & meas . & shift & & flux & w & n@xmath196 & n@xmath197 + & & ( ) & ( ) & ( km / s ) & ( @xmath198 ) & ( km / s ) & ( @xmath198 ph/@xmath199/s ) & ( m ) & ( @xmath200 ) & ( @xmath201 ) + 1 & & 1.850 & 1.848(1 ) & @xmath202 & @xmath203 & @xmath204 & @xmath205 & @xmath206 & @xmath207 & @xmath208 + 1 & & 1.780 & 1.776(1 ) & @xmath209 & @xmath210 & @xmath1 & @xmath211 & @xmath212 & @xmath213 & @xmath214 + 3 & & 1.850 & 1.851(2 ) & @xmath221 & @xmath222 & @xmath223 & @xmath224 & @xmath225 & @xmath226 & @xmath227 + 3 & & 1.780 & 1.778(1 ) & @xmath228 & @xmath229 & @xmath230 & @xmath231 & @xmath232 & @xmath233 & @xmath234 + 4 & & 1.850 & 1.850(1 ) & @xmath235 & @xmath236 & @xmath237 & @xmath238 & @xmath239 & @xmath240 & @xmath241 + 4 & & 1.780 & 1.781(1 ) & @xmath242 & @xmath243 & @xmath244 & @xmath245 & @xmath246 & @xmath247 & @xmath119 + lllllllllllll obs . & & theor . & meas . & shift & & flux & w & n@xmath196 & n@xmath197 + & & ( ) & ( ) & ( km / s ) & ( @xmath198 ) & ( km / s ) & ( @xmath198 ph/@xmath199/s ) & ( m ) & ( @xmath200 ) & ( @xmath201 ) + 1 ( @xmath248mean ) & & 1.850 & & & @xmath249 & @xmath250 & @xmath251 & @xmath252 & @xmath220 & @xmath217 + 1 ( @xmath253mean ) & & 1.850 & 1.848(1 ) & @xmath202 & @xmath249 & @xmath250 & @xmath254 & @xmath148 & @xmath255 & @xmath256 + 1 ( @xmath248mean ) & & 1.780 & 1.776(1 ) & @xmath209 & @xmath257 & @xmath258 & @xmath134 & @xmath259 & @xmath260 & @xmath151 + 1 ( @xmath253mean ) & & 1.780 & 1.776(1 ) & @xmath209 & @xmath261 & @xmath262 & @xmath128 & @xmath156 & @xmath172 & @xmath263 + 1 ( mean@xmath75 c / s ) & & 1.850 & & & @xmath249 & @xmath250 & @xmath220 & @xmath264 & @xmath265 & @xmath218 + 1 ( mean@xmath05 c / s ) & & 1.850 & 1.849(1 ) & @xmath266 & @xmath249 & @xmath250 & @xmath234 & @xmath133 & @xmath267 & @xmath268 + 1 ( mean@xmath75 c / s ) & & 1.780 & 1.775(1 ) & @xmath269 & @xmath257 & @xmath258 & @xmath128 & @xmath270 & @xmath271 & @xmath145 + 1 ( mean@xmath05 c / s ) & & 1.780 & 1.777(1 ) & @xmath272 & @xmath273 & @xmath274 & @xmath206 & @xmath156 & @xmath172 & @xmath263 + 3 ( norm . ) & & 1.850 & & & @xmath249 & @xmath250 & @xmath218 & @xmath275 & @xmath251 & @xmath217 + 3 ( dip ) & & 1.850 & 1.850(1 ) & @xmath235 & @xmath249 & @xmath250 & @xmath276 & @xmath277 & @xmath206 & @xmath278 + 3 ( norm . ) & & 1.780 & 1.777(1 ) & @xmath272 & @xmath249 & @xmath250 & @xmath279 & @xmath280 & @xmath281 & @xmath226 + 3 ( dip ) & & 1.780 & 1.778(1 ) & @xmath228 & @xmath249 & @xmath250 & @xmath254 & @xmath282 & @xmath283 & @xmath145 + c r r r r r r r r r observation & @xmath284 & @xmath285 & @xmath286 & @xmath287 & @xmath288 + 1 & @xmath289 & @xmath290 & @xmath291 & @xmath292 & @xmath293 + 2 & @xmath294 & & @xmath295 & @xmath296 & + 3 & @xmath297 & @xmath298 & @xmath299 & @xmath300 & @xmath301 + 4 & @xmath302 & @xmath303 & @xmath304 & @xmath305 & @xmath306 +

, the fe xxvi line has a fwhm of km / s and a blue - shift of km / s , suggesting that the highly ionized medium is an outflow . moreover , the fe xxv line is observed to vary significantly on a timescale of a few hundred seconds in the first observation , which corresponds to the keplerian orbital period at approximately ( where ) . our models for the absorption geometry suggest that a combination of geometric effects and changing ionizing flux are required to account for the large changes in line flux observed between observations , and that the absorption likely occurs at a radius less than for a black hole . viable models for the absorption geometry include cyclic absorption due to an accretion disk structure , absorption in a clumpy outflowing disk wind , or possibly a combination of these two . geometries recently found in the galactic black holes gx 339 - 4 and xte j1650 .

we observed the bright phase of the 2003 outburst of the galactic black hole candidate h 1743 in x - rays simultaneously with _ chandra _ and _ rxte _ on four occasions . the _ chandra_/hetgs spectra reveal narrow , variable ( he - like ) fe xxv and ( h - like ) fe xxvi resonance absorption lines . in the first observation , the fe xxvi line has a fwhm of km / s and a blue - shift of km / s , suggesting that the highly ionized medium is an outflow . moreover , the fe xxv line is observed to vary significantly on a timescale of a few hundred seconds in the first observation , which corresponds to the keplerian orbital period at approximately ( where ) . our models for the absorption geometry suggest that a combination of geometric effects and changing ionizing flux are required to account for the large changes in line flux observed between observations , and that the absorption likely occurs at a radius less than for a black hole . viable models for the absorption geometry include cyclic absorption due to an accretion disk structure , absorption in a clumpy outflowing disk wind , or possibly a combination of these two . if the wind in h 1743 has unity filling factor , the highest implied mass outflow rate is 20% of the eddington mass accretion rate . this wind may be a hot precursor to the seyfert - like , outflowing `` warm absorber '' geometries recently found in the galactic black holes gx 339 - 4 and xte j1650 . we discuss these findings in the context of ionized fe absorption lines found in the spectra of other galactic sources , and connections to warm absorbers , winds , and jets in other accreting systems . [ firstpage ]

1102.3360

i

in the 18th century , euler considered the problem of optimizing functionals depending not only on some unknown function @xmath0 and some derivative of @xmath0 , but also on an antiderivative of @xmath0 ( see @xcite ) . similar problems have been recently investigated in @xcite , where lagrangians containing higher - order derivatives and optimal control problems are considered . more generally , it has been shown that the results of @xcite hold on an arbitrary time scale @xcite . here we study such problems within the framework of fractional calculus . roughly speaking , a fractional calculus defines integrals and derivatives of non - integer order . let @xmath1 be a real number and @xmath2 be such that @xmath3 . here we follow @xcite and @xcite . let @xmath4\to\mathbb{r}$ ] be piecewise continuous on @xmath5 and integrable on @xmath6 $ ] . the left and right riemann liouville fractional integrals of @xmath7 of order @xmath8 are defined respectively by @xmath9 here @xmath10 is the well - known gamma function . then the left @xmath11 and right @xmath12 riemann liouville fractional derivatives of @xmath7 of order @xmath8 are defined ( if they exist ) as @xmath13 and @xmath14 the fractional derivatives and have one disadvantage when modeling real world phenomena : the fractional derivative of a constant is not zero . to eliminate this problem , one often considers fractional derivatives in the sense of caputo . let @xmath7 belong to the space @xmath15;\mathbb{r})$ ] of absolutely continuous functions . the left and right caputo fractional derivatives of @xmath7 of order @xmath8 are defined respectively by @xmath16 and @xmath17 these fractional integrals and derivatives define a rich calculus . for details see the books @xcite . here we just recall a useful property for our purposes : integration by parts . for fractional integrals , @xmath18 ( see , , ( * ? ? ? * lemma 2.7 ) ) , and for caputo fractional derivatives @xmath19_a^b\ ] ] ( see , , ( * ? ? ? * eq . ( 16 ) ) ) . in particular , for @xmath20 one has @xmath21_a^b.\ ] ] when @xmath22 , @xmath23 , @xmath24 , @xmath25 is the identity operator , and gives the classical formula of integration by parts . the fractional calculus of variations concerns finding extremizers for variational functionals depending on fractional derivatives instead of integer ones . the theory started in 1996 with the work of riewe , in order to better describe non - conservative systems in mechanics @xcite . the subject is now under strong development due to its many applications in physics and engineering , providing more accurate models of physical phenomena ( see , , @xcite ) . with respect to results on fractional variational calculus via caputo operators , we refer the reader to @xcite and references therein . our main contribution is an extension of the results presented in @xcite by considering lagrangians containing an antiderivative , that in turn depend on the unknown function , a fractional integral , and a caputo fractional derivative ( section [ sec : fundprob ] ) . transversality conditions are studied in section [ sec : natbound ] , where the variational functional @xmath26 depends also on the terminal time @xmath27 , @xmath28 , and where we obtain conditions for a pair @xmath29 to be an optimal solution to the problem . in section [ sec : isoprob ] we consider isoperimetric problems with integral constraints of the same type as the cost functionals considered in section [ sec : fundprob ] . fractional problems with holonomic constraints are considered in section [ sec : holonomic ] . the situation when the lagrangian depends on higher order caputo derivatives , , it depends on @xmath30 for @xmath31 , @xmath32 , is studied in section [ sec : higher ] , while section [ sec : fracopt ] considers fractional lagrange problems and the hamiltonian approach . in section [ sec : sufconditions ] we obtain sufficient conditions of optimization under suitable convexity assumptions on the lagrangian . we end with section [ sec : numsim ] , discussing a numerical scheme for solving the proposed fractional variational problems . the idea is to approximate fractional problems by classical ones . numerical results for two illustrative examples are described in detail .

isoperimetric problems , problems with holonomic constraints and depending on higher - order caputo derivatives , as well as fractional lagrange problems , are considered . 49k05 , 49s05 , 26a33 , 34a08 . calculus of variations , fractional calculus , caputo derivatives , fractional necessary optimality equations .

we obtain necessary optimality conditions for variational problems with a lagrangian depending on a caputo fractional derivative , a fractional and an indefinite integral . main results give fractional euler lagrange type equations and natural boundary conditions , which provide a generalization of previous results found in the literature . isoperimetric problems , problems with holonomic constraints and depending on higher - order caputo derivatives , as well as fractional lagrange problems , are considered . 49k05 , 49s05 , 26a33 , 34a08 . calculus of variations , fractional calculus , caputo derivatives , fractional necessary optimality equations .

1307.6388

i

we have examined the stability properties of a relativistic magnetized cylindrical flow in the approximation of zero thermal pressure , neglecting also the effects of rotation and focusing only on the @xmath0 mode . in this configuration we have two kinds of instability that may be present : kelvin - helmholtz and current driven . the instability behavior depends of course on the chosen equilibrium configuration and this is somewhat arbitrary since we have no direct information on the magnetic field structure , although some indications are provided by the acceleration models ( see e.g. * ? ? ? nevertheless , the general outcome and the properties of the solutions obtained for a particular configuration , such as the one we adopted here , can be considered valid for a more general class of equilibria . our results can then be considered representative of an equilibrium configuration characterized by a distribution of current concentrated in the jet , with the return current assumed to be mainly found at very large distances . we can summarize our results by considering the behavior of the system for different values of the ratio between matter and magnetic energy densities . for matter dominated flows , the dominant instability is kh and the wavenumber corresponding to the maximum growth rate as well as the growth rate itself scale both as @xmath170 . somewhat above equipartion khi reaches its maximum growth rate and then it becomes rapidly stabilized . below this stabilization limit the dominant instability becomes cdi the dependence of the khi on the value of the pitch is relatively weak and only for the smallest value of the axial pitch ( @xmath61 ) , in the relativistic case , we observe a displacement of the stability limit towards lower values of @xmath108 and a merging with cdi . cdi are therefore prevailing for flows in equipartition or magnetically dominated and the dependence on @xmath108 in these regimes is quite weak . the wavenumber corresponding to the maximum growth rate scale as @xmath117 , the growth rate itself increases with decreasing @xmath61 . and the modes have a resonant character with a peak in the eigenfunction at the radial position where @xmath188 . at low jet velocity , our equilibrium has no return current inside the domain , while at relativistic velocities we have a small portion of the return current corresponding to the velocity shear region . the corresponding steepening of the pitch profiles induces a stabilization of the modes for which the resonance condition corresponds to radial positions where the return current is found . we then observe a splitting of the cdi in two branches , one at high wavenumbers ( the inner mode ) characterized by an eigenfunction with a resonant peak in the inner radial part of the flow , and one at smaller wavenumbers ( the outer mode ) for which the resonant peak is outside the jet region . which of the two branches is dominant depends on the width of the current distribution and on the lorentz factor of the flow . an increase in the current concentration ( small @xmath45 ) favors the growth of the inner mode , while an increase in @xmath73 enhances the development of the outer mode . the different behavior in the explored parameter ranges may have crucial implications for the nonlinear stages as distinct types of instability may evolve differently . this study is therefore an essential first step for the interpretation of the results of numerical simulations that will be presented in a following paper and for their comparison with astrophysical data .

we perform a linear analysis of the stability of a magnetized relativistic non - rotating cylindrical flow in the aproximation of zero thermal pressure , considering only the mode . the current driven mode is splitted in two branches , the branch at high wavenumbers is characterized by the eigenfunction concentrated in the jet core , the branch at low wavenumbers is instead characterized by the eigenfunction that extends outside the jet velocity shear region .

we perform a linear analysis of the stability of a magnetized relativistic non - rotating cylindrical flow in the aproximation of zero thermal pressure , considering only the mode . we find that there are two modes of instability : kelvin - helmholtz and current driven . the kelvin - helmholtz mode is found at low magnetizations and its growth rate depends very weakly on the pitch parameter . the current driven modes are found at high magnetizations and the value of the growth rate and the wavenumber of the maximum increase as we decrease the pitch parameter . in the relativistic regime the current driven mode is splitted in two branches , the branch at high wavenumbers is characterized by the eigenfunction concentrated in the jet core , the branch at low wavenumbers is instead characterized by the eigenfunction that extends outside the jet velocity shear region . [ firstpage ] galaxies : jets , mhd , instabilities , relativistic processes

astro-ph0512272

i

one of the great unsolved problems of planet formation is how to form planetesimals , the kilometer - sized precursors of real planets @xcite . at this size solid bodies in a protoplanetary disk can attract each other through gravitational two - body encounters , whereas gravity is insignificant between smaller bodies . starting from micrometer - sized dust grains , the initial growth is caused by the random brownian motion of the grains ( e.g. * ? ? ? * ; * ? ? ? * see @xcite for a review ) . the vertical component of the gravity from the central object causes the gas in the disk to be stratified with a higher pressure around the mid - plane . even though the dust grains do not feel this pressure gradient , the strong frictional coupling with the gas prevents small grains from having any significant vertical motion relative to the gas however , once the grains have coagulated to form pebbles with sizes of a few centimeters , the solids are no longer completely coupled to the gas motion . they are thus free to fall , or sediment , towards the mid - plane of the disk . the increase in dust density opens a promising way of forming planetesimals by increasing the local dust density around the mid - plane of the disk to values high enough for gravitational fragmentation of the dust layer @xcite . there are however two major unresolved problems with the gravitational fragmentation scenario . any global turbulence in the disk causes the dust grains to diffuse away from the mid - plane , and thus the dust density is kept at values that are too low for fragmentation . a turbulent @xmath0-value of @xmath1 is generally enough to prevent efficient sedimentation towards the mid - plane @xcite , whereas the @xmath0-value due to magnetorotational turbulence @xcite is from a few times @xmath2 ( found in local box simulations with no imposed magnetic field ) to 0.1 and higher ( in global disk simulations ) . the presence of a magnetically dead zone around the disk mid - plane @xcite may not mean that there is no turbulence in the mid - plane , as other instabilities may set in and produce significant turbulent motion @xcite . the magnetically active surface layers of the disk can even induce enough turbulent motion in the mid - plane to possibly prevent efficient sedimentation of dust @xcite . the presence of a dead zone may actually _ in itself _ be a source of turbulence . the sudden fall of the accretion rate can lead to a pile up of mass in the dead zone , possibly igniting the magnetorotational instability in bursts @xcite or a rossby wave instability @xcite . the second major problem with the gravitational fragmentation scenario is that even in the absence of global disk turbulence , the dust sedimentation may in itself destabilize the disk . protoplanetary disks have a radial pressure gradient , because the temperature and the density fall with increasing radial distance from the central object , so the gas rotates at a speed that is slightly below the keplerian value . the dust grains feel only the gravity and want to rotate purely keplerian . close to the equatorial plane of the disk , where the sedimentation of dust has increased the dust - to - gas ratio to unity or higher , the gas is forced by the dust to orbit at a higher speed than far away from the mid - plane where the rotation is still sub - keplerian . thus there is a vertical dependence of the gas rotation velocity . such shear flow can be unstable to the kelvin - helmholtz instability ( khi ) , depending on the stabilizing effect of vertical gravity and density stratification . a necessary criterion for the khi is that the energy required to lift a fluid parcel of gas and dust vertically upwards by an infinitesimal distance is available in the relative vertical motion between infinitesimally close parcels @xcite . the turbulent motions resulting from the khi are strong enough to puff up the dust layer and prevent the formation of an infinitesimally thin dust sheet around the mid - plane of the disk @xcite . modifications to the gravitational fragmentation scenario have been suggested to overcome the problem of kelvin - helmholtz turbulence . @xcite ( 1998 , hereafter referred to as s98 ) found that if the mid - plane of the disk is in a state of constant richardson number , as expected for small grains whose settling time is long compared to the growth rate of the khi , then an increase in the global dust - to - gas ratio can lead to the formation of a high density dust cusp very close to the mid - plane of the disk , reaching potentially a dust - to - gas ratio of 100 already at a global dust - to - gas ratio that is 10 times the canonical interstellar value of @xmath3 . the appearance of a superdense dust cusp in the very mid - plane has been interpreted by @xcite as an inability of the gas ( or of the khi ) to move more mass than its own away from the mid - plane . as a source of an increased value of the global dust - to - gas ratio , @xcite suggest that the dust grains falling radially inwards through the disk pile up in the inner disk . a slowly growing radial self - gravity mode in the dust density has also been suggested as the source of an increased dust - to - gas ratio at certain radial locations @xcite . trapping dust boulders in a turbulent flow is a mechanism for avoiding the problem of self - induced kelvin - helmholtz turbulence altogether @xcite . if the dust can undergo a gravitational fragmentation locally , because the boulders are trapped in features of the turbulent gas flow such as vortices or high - pressure regions , then there is no need for an extremely dense dust layer around the mid - plane . @xcite found that meter - sized dust boulders are temporarily trapped in regions of slight gas overdensity in magnetorotational turbulence , increasing the dust - to - gas ratio locally by up to two orders of magnitude . they estimate that the dust in such regions should have time to undergo gravitational fragmentation before the high - pressure regions dissolve again . @xcite , on the other hand , find that vortices can even form in magnetorotationally turbulent disks , keeping dust boulders trapped for hundreds of disk rotation periods . the khi can not operate inside a vortex because there is no radial pressure gradient , and thus no vertical shear , in the center of the vortex @xcite . from a numerical side it has been shown many times that a pure shear flow , i.e. one that is not explicitly supported by any forces , is unstable , both with magnetic fields @xcite and without @xcite . but the key point here is that the vertical shear formed in a protoplanetary disk is due to the sedimentation of dust , and that the shear is able to regenerate as the dust falls down again , thus keeping the flow unstable to khi . the description of the full non - linear outcome of such a system requires numerical simulations that include dust that can move relative to the gas . linear stability analysis of dust - induced shear flows in protoplanetary disks have been performed for simplified physical conditions @xcite , but also with coriolis forces and keplerian shear included @xcite . recently @xcite ( 2005 , hereafter referred to as go05 ) took an approach to include the dust into their numerical simulations of the kelvin - helmholtz instability by having the dust grains so extremely well - coupled to the gas that they always move with the instantaneous velocity of the gas . this is indeed a valid description of the dynamics of tiny dust grains . however , the strong coupling to the gas does not allow the dust grains to fall back towards the mid - plane . thus the saturated state of the kelvin - helmholtz turbulence can not be reached this way . in this paper we present computer simulations where we have let the dust grains , represented by particles each with an individual velocity vector and position , move relative to the gas . this allows us to obtain a state of self - sustained kelvin - helmholtz turbulence from which we can measure quantities such as the diffusion coefficient and the maximum dust density . a better knowledge of these important characteristics of kelvin - helmholtz turbulence is vital for our understanding of planet formation .

a two - dimensional corotating slice in the azimuthal vertical plane of the disk is considered where we include the coriolis force and the radial advection of the keplerian rotation flow . the azimuthally averaged state of the self - sustained kelvin - helmholtz turbulence is found to have a constant richardson number in the region around the mid - plane where the dust - to - gas ratio is significant .

we perform numerical simulations of the kelvin - helmholtz instability in the mid - plane of a protoplanetary disk . a two - dimensional corotating slice in the azimuthal vertical plane of the disk is considered where we include the coriolis force and the radial advection of the keplerian rotation flow . dust grains , treated as individual particles , move under the influence of friction with the gas , while the gas is treated as a compressible fluid . the friction force from the dust grains on the gas leads to a vertical shear in the gas rotation velocity . as the particles settle around the mid - plane due to gravity , the shear increases , and eventually the flow becomes unstable to the kelvin - helmholtz instability . the kelvin - helmholtz turbulence saturates when the vertical settling of the dust is balanced by the turbulent diffusion away from the mid - plane . the azimuthally averaged state of the self - sustained kelvin - helmholtz turbulence is found to have a constant richardson number in the region around the mid - plane where the dust - to - gas ratio is significant . nevertheless the dust density has a strong non - axisymmetric component . we identify a powerful clumping mechanism , caused by the dependence of the rotation velocity of the dust grains on the dust - to - gas ratio , as the source of the non - axisymmetry . our simulations confirm recent findings that the critical richardson number for kelvin - helmholtz instability is around unity or larger , rather than the classical value of 1/4 .

0707.2636

i

we have observed ngc 6946 as part of a program of deep mapping of extended , cold co in nearby spiral galaxies @xcite . large , fully - sampled , and deep images of co in nearby galaxies are relatively uncommon . interferometric images either do not contain zero spacing data ( e.g. , * ? ? ? * ) or if they do , the images typically do not go deep enough to detect cold , molecular disk gas . the detection of cold , extended molecular gas at large galactocentric radii in large spirals is a goal of this mapping project with the nrao 12 meter telescope . a gas - rich sc galaxy at a distance of 6 mpc @xcite , ngc 6946 is known for its bright and asymmetric optical spiral arms @xcite . the pronounced asymmetry may be caused by interactions with neighboring dwarf galaxies , ugc 11583 and l149 @xcite . ngc 6946 has a 9 optical diameter on the sky and an atomic hydrogen ( hi ) gas disk extending to 25 @xcite . like many spirals , ngc 6946 has a central @xmath5 diameter hole " in its hi disk . radio continuum , far infrared ( fir ) , optical line , and x - ray observations indicate vigorous star formation in the disk and an interstellar medium ( ism ) stirred by supernovae and stellar winds @xcite . the high level of star formation in the disk of ngc 6946 has been attributed both to its strong spiral density wave @xcite , and to stochastic , self - propagating , star formation @xcite . until recently , the co morphology beyond the inner 3 radius ( @xmath65 kpc ) was not well known . like many spiral galaxies , ngc 6946 has a bright nuclear co peak which falls off nearly exponentially with galactocentric radius @xcite . a 3 diameter map of @xcite revealed substantial co emission in the inner disk with several emission peaks . @xcite mapped two 2 circular fields in the disk of ngc 6946 in a study of arm and interarm regions . the inner 2 region has been mapped with higher resolution single dish and aperture synthesis telescopes in transitions of co and co isotopomers @xcite . a larger interferometer+single dish map of ngc 6946 has been made by the bima song team @xcite , although with lower sensitivity to cold extended emission than the maps we present here . @xcite presented 21 resolution fully - sampled single dish maps of co(1 - 0 ) . we compare our results to theirs . the on the fly " ( otf ) observing mode at the nrao was ideally suited to the imaging of extended gas in galaxies . here , we present 16 by 10 maps of co(1 - 0 ) and co(2 - 1 ) covering the optical disk of ngc 6946 made with the 12 meter telescope . the maps are deep enough to detect cold , extended interarm co to levels of @xmath7 , or @xmath8 . we use these primary tracers of the molecular gas phase and combine them with archival and published data in order to study the molecular gas disk , the total neutral gas surface density , and their relationship to star formation in the disk of ngc 6946 .

we present on the fly " maps of the co(1 - 0 ) and co(2 - 1 ) emission covering a 10 10 region of the ngc 6946 . using our co maps and archival vla hi observations we create a total gas surface density map , , for ngc 6946 . the predominantly molecular inner gas disk transitions smoothly into an atomic outer gas disk , with equivalent atomic and molecular gas surface densities at r = 3.5 ( 6 kpc ) . star formation tracers are better correlated with the total neutral gas disk than with the molecular gas by itself implying . using the 100 and 21 cm continuum from ngc 6946 as star formation tracers , we arrive at a gas consumption timescale of 2.8 gyr , which is relatively uniform across the disk .

we present on the fly " maps of the co(1 - 0 ) and co(2 - 1 ) emission covering a 10 10 region of the ngc 6946 . using our co maps and archival vla hi observations we create a total gas surface density map , , for ngc 6946 . the predominantly molecular inner gas disk transitions smoothly into an atomic outer gas disk , with equivalent atomic and molecular gas surface densities at r = 3.5 ( 6 kpc ) . we estimate that the total mass is , roughly 1/3 of the interstellar hydrogen gas mass , and about 2% of the dynamical mass of the galaxy at our assumed distance of 6 mpc . the value of the co(2 - 1)/co(1 - 0 ) line ratio ranges from 0.35 to 2 ; 50% of the map is covered by very high ratio , , gas . the very high ratios are predominantly from interarm regions and appear to indicate the presence of wide - spread optically thin gas . star formation tracers are better correlated with the total neutral gas disk than with the molecular gas by itself implying . using the 100 and 21 cm continuum from ngc 6946 as star formation tracers , we arrive at a gas consumption timescale of 2.8 gyr , which is relatively uniform across the disk . the high star formation rate at the nucleus appears to be due to a large accumulation of molecular gas rather than a large increase in the star formation efficiency . the mid - plane gas pressure in the outer ( r 10 kpc ) hi arms of ngc 6946 is close to the value at the radial limit ( 10 kpc ) of our observed co disk . if the mid - plane gas pressure is a factor for the formation of molecular clouds , these outer hi gas arms should contain molecular gas which we do not see because they are beyond our detection limit .

0707.2636

c

we have obtained deep co(1 - 0 ) and co(2 - 1 ) images of a 10 @xmath0 10 region centered on the sc galaxy , ngc 6946 , with 55 and 27 resolution using the nrao 12 meter telescope . we have combined these deep co images with vla hi images to obtain images of the neutral gas in this galaxy . to summarize our findings : + 1 ) the co(1 - 0 ) and co(2 - 1 ) disks shows many of the same features seen in the optical disk : a strong nuclear peak of emission , molecular gas arms spatially coincident with the asymmetric optical arm pattern , and a 10 diameter inner disk filled with co emission . the co disk is roughly exponential with a scale length of @xmath77 or @xmath134 kpc . + 2 ) we obtain @xmath135 for the molecular hydrogen gas mass of ngc 6946 . this can be compared to the total atomic hydrogen gas mass of @xmath136 ( which includes a vla missing flux estimate ) . molecular hydrogen constitutes 1/3 of the interstellar hydrogen gas mass and 2% of the dynamical mass of ngc 6946 . at the nucleus , = 170 , the molecular gas surface density is @xmath47 along the strong optical arm pattern , and @xmath137 regions beyond . + 3 ) the mean value for @xmath138 in the the nucleus and optical arm pattern is @xmath139 . this is consistent with optically thick gas , with an excitation temperature of 10 - 15 k. interarm co shows the largest variation in , with a range of 0.35 to 2 . only 4% of the co ratio map has @xmath68 which is likely to be cold or subthermally excited co ( @xmath140 k ) . fully half of the co disk of ngc 6946 has @xmath65 . the high appears to indicate the presence of widespread , optically - thin gas in between the spiral arms . + 4 ) molecular gas dominates the total gas surface density within a galactocentric radius of 3.5 ( 6 kpc ) . co emission fills the central depression in the hi disk . beyond the nuclear region , there is a good correlation between the co and hi gas , which both peak on the spiral arms . the inner disk arm pattern transitions smoothly from predominantly co arms to predominantly hi arms at the outskirts of the co disk . at the nucleus , = 240 , the total gas surface density is @xmath141 along the optical arm pattern , and @xmath142 along the outer gas arms ( @xmath143 ) . co(2 - 1 ) emission peaks are coincident with clumps of hi , , and fir emission outside the nuclear region , presumably an indication of warmer molecular gas and dissociated . + 5 ) we find strong correlations between star formation tracers ( except for ) and co emission in ngc 6946 . the correlation coefficients are @xmath144 . when is used instead of just the molecular gas tracer in a radial comparison , the dependence on galactic radius is reduced . star formation is more closely related to the total gas surface density than the molecular gas surface density alone , a schmidt law relationship . this also implies star formation is likely well beyond the radial limits of our observed co disk in regions of predominantly hi gas . surprisingly , is not as well correlated with co or as are the other star formation tracers . the correlation coefficient is significantly less , @xmath88 , and the ratio formed from and co or declines more rapidly with galactic radius . + 6 ) the star formation efficiency in the disk of ngc 6946 is relatively uniform ; 2.8 gyr when expressed in terms of the gas consumption time - scale . this uniformity implies @xmath3 . the star formation rate at the nucleus is due to the high molecular surface density rather than an elevated star formation efficiency . + 7 ) the edge of the observed co disk represents a detection limit rather than a threshold for molecular cloud formation . we suspect the exponential fall of the co disk continues smoothly past our detection limit . the mid - plane gas pressure in the outer hi arm structure , @xmath133 is close to the value at the limits of the observed co disk . these gas arms should support the formation of molecular clouds which could be detectable in higher resolution observations . the authors would like to thank the nrao 12 m telescope operators , duane clark , paul hart , victor gasho and harry stahl , for their help and company during the many observing sessions in which the data for ngc 6946 and other galaxies were acquired . we thank l. tacconi giving for us permission to retrieve and use the hi 21 cm line observations from the vla archive . and we appreciate the careful reading and critical suggestions made by the anonymous referee . this work was supported in part by nsf grants ast00 - 71276 and ast03 - 07950 . we made use of the nasa / ipac / iras hires data reduction facilities as well as stsci digital sky survey facilities . the digitized sky surveys were produced at the space telescope science institute under u.s . government grant nag w-2166 . adler , d. s. , allen , r. j. , & lo , k. y. 1991 , , 283 , 475 arp , h. 1966 , , 14 , 1 arsenault , r. 1989 , , 217 , 66 ball , r. , et al . 1985 , , 298 , l21 bell , e.f . 2003 , , 586 , 794 belley , j. & roy , j. 1992 , , 78 , 61 bicay , m.d . , helou , g. & condon , j.j . 1989 , , 338 , l53 bicay , m.d . & helou , g. 1990 , , 362 , 59 boulanger , f. & viallefond , f. 1992 , , 266 , 37 carignan , c. , et al . 1990 , , 234 , 43 casoli , f. , et al . 1990 , , 233 , 357 castets , a. , et al . 1990 , , 234 , 469 catalogued galaxies and quasars in the iras survey 1985 , prepared by lonsdale , c.j . ( jpl ) combes , f. 1991 , , 29 , 195 condon , j.j . 1992 , , 30 , 575 crosthwaite , l.p . , et al . 2001 , , 122 , 797 crosthwaite , l.p . , 2002 , , 123 , 1892 dahmen , g. , et al . 1998 , , 331 , 959 degioia - eastwood , k. , et al . 1984 , , 278 , 564 de vaucouleurs , g. , de vaucouleurs , a. & corwin , h. 1976 , _ second reference catalogue of bright galaxies _ , univ . of texas press dickman , r.l . , snell , r.l . & schloerb , f.p . 1986 , , 309 , 326 eastman , r.g . , et al . 1996 , , 466 , 911 elmegreen , b.g . 1989 , , 338 , 178 elmegreen , b.g . 1993 , , 411 , 170 elmegreen , b.g . 2002 , , 577 , 206 elmegreen , b.g . & parravano , a. 1994 , , 435 , l121 engargiola , g. 1991 , , 76 , 875 ferguson , a.m.n . , 1998 , , 506 , l19 gordon , k.j . , remage , n.h . & roberts , m.s . 1968 , , 154 , 845 helou , g. & bicay , m.d . 1993 , , 415 , 93 hodge , p.w . & kennicutt , r.c . 1983 , , 88 , 296 honma , m. , sofue , y. & arimoto , n. 1995 , , 304 , 1 hunter , s.d . , et al . 1997 , , 481 , 205 ishizuki , s. et al . 1990 , , 355 , 436 kamphuis , j. & sancisi , r. 1993 , , 273 , l31 karachentsev , i.d . , sharina , m.e . & hutchmeier , w.k . 2000 , , 362 , 544 kennicutt , r.c . 1998a , , 498 , 541 kennicutt , r.c . 1998b , , 36 , 189 kulkarni , s.r . , & heiles , c. 1987 , in _ interstellar processes _ , ed . hollenbach & thronson , p. 87 kutner , m.l . & ulich , b.l . 1981 , , 250 , 341 lacey , c. , duric , n. & goss , w.m . 1997 , , 109 , 417 lequeux , j. , allen , r.j . & guilloteau , s. 1993 , , 280 , l23 leroy , a . 2005 , , 625 , 763 madden , s.c . 1993 , , 407 , 579 maloney , p. & black , j.h . 1988 , , 325 , 389 mangum , j.g . user s manual for the nrao 12 m millimeter - wave telescope _ , nrao 12 m publication mangum , j.g . 1996 , _ on the fly observing at the 12 m _ , nrao 12 m publication mangum , j.g . 1997 , _ equipment and calibration status for the nrao 12 m telescope _ , nrao 12 m publication mason , a. m. , & wilson , c. d. 2004 , , 612 , 860 mead , k.n . & kutner , m.l . 1988 , , 330 , 399 meier , d.s . & turner , j.l . 2004 , , 127 , 2069 meier , d.s . & turner , j.l . 2001 , , 551 , 687 meier , d.s . , turner , j.l . , & beck , s. c. 2002 , , 124 , 877 morris , m. & lo , k.y . 1978 , , 223 , 803 murgia , m. , et al . 2002 , , 385 , 412 nikiforov , i.i . , petrovskaya , i.v . & ninkovic , s. 2000 , _ small galaxy groups , asp conference serier _ , 209 , 399 oka , t. , et al . 1996 , , 460 , 334 oka , t. , et al . 1998 , , 493 , 730 pisano , d.j . & wilcots , e.m . 2000 , , 319 , 821 regan , m.w . & vogel , s.n . 1995 , , 452 , l21 regan , m.w . 2001 , , 561 , 218 rogstad , d.h . , shostak , g.s . & rots , a.h . 1973 , , 22 , 111 rownd , b.k . & young , j.s . 1999 , , 118 , 670 saha , a , claver , j. & hoessel , j.g . 2002 , , 124 , 839 sakamoto , s. 1996 , , 462 , 215 sakamoto , s. , et al . 1995 , , 100 , 125 sakamoto , s. , et al . 1997 , , 486 , 290 sakamoto , s. , et al . 1999 , , 124 , 403 sauvage , m. & thuan , t.x . 1992 , , 396 , l69 schlegel , e. m. 1994 , , 434 , 523 schmidt , m. 1959 , , 129 , 243 scoville , n. z. & sanders , d.b . 1987 , in _ interstellar processes _ , hollenbach & thronson , p 21 scoville , n. z. , et al . 1987 , , 63 , 821 sharina , m.e . , karachentsev , i.d . & tikhonov , n.a . 1997 , astl , 23 , 373 sofue , y. , et al . 1988 , , 40 , 511 solomon , p.m. , et al . 1987 , , 319 , 730 spaans , m. , et al . 1994 , , 437 , 270 strong , a.w . 1988 , , 207 , 1 suchkov , a. , allen , r.j . & heckman , t.m . 1993 , , 412 , 542 tacconi , l.j . & young , j.s . 1986 , , 308 , 600 tacconi , l.j . & young , j.s . 1989 , , 71 , 455 tacconi , l.j . & young , j.s . 1990 , , 352 , 595 tilanus , r.p.j . & allen , r.j . 1989 , , 339 , l57 tilanus , r.p.j . & allen , r.j . 1993 , , 274 , 707 tuffs , r.j . 1996 , , 315 , l149 ulich , b.l . & hass , r.w . 1976 , , 30 , 247 van der kruit , p.c . , allen , r.j . & rots , a.h . 1977 , , 55 , 421 wall , w.f . 1993 , , 414 , 98 walsh , w. , et al . 2002 , , 388 , 7 weliachew , l. , casoli , f. & combes , f. 1988 , , 199 , 29 wiklind , t. , et al . 1990 , , 232 , l11 wilson , c.d . 1995 , , 448 , l97 wilson , c.d . & scoville , n. 1991 , , 370 , 184 wong , t. & blitz , l. 2002 , , 569 , 157 wouterloot , j.g.a . , brand , j. & henkel , c. 1988 , , 191 , 323 young , j.s . & scoville , n.z . 1982 , , 258 , 467 young , j.s . & scoville , n.z . 1991 , , 29 , 581 young , j.s . , 1995 , , 98 , 219 hubble type & scd + r.a . ( j2000 ) & @xmath145 + dec . ( j2000 ) & @xmath146 + @xmath147 & @xmath148 + distance & 6 mpc + @xmath149 & 50.5 @xmath150 + inclination & 40@xmath151 + position angle & 242@xmath152 + @xmath153 & 170 @xmath154 + @xmath155 & 5.1@xmath156 ( 9.0 kpc ) + @xmath157 & 1.07 @xmath158 + @xmath159 & @xmath160 + @xmath161 & @xmath162 + @xmath163 & @xmath164 + @xmath165 & @xmath166 + @xmath167 & @xmath168 + & hubble type & scd & scd & sbc & sbbc + distance ( mpc ) & 3.3 & 6 & 4 & ... + d@xmath84 ( arcminutes ) & 20 & 11 & 11 & ... + d@xmath84/d@xmath169 & 0.23 & 0.38 & 0.15 & ... + d@xmath84/d@xmath170 & 1.1 & 1.0 & 1.0 & ... + m@xmath171 ( @xmath172 ) & 1.8 & 1.9 & 0.7 & @xmath57 + @xmath173 ( @xmath174 ) & 16 & 33 & 23 & 25 + @xmath173 ( @xmath174 ) & 57 & 70 & 62 & 48 + / & 0.27 & 0.47 & 0.37 & 0.52 + m@xmath82/m@xmath83 & 0.055 & 0.074 & 0.16 & 0.03 + @xmath175 ( ) & 140 & 170 & 220 & 400 + @xmath176 ( ) & 11 & 13 & 23 & 3.4 + & 100 & -0.196 ( 0.015 ) + & 60 & -0.189 ( 0.015 ) + & 21 cm continuum & -0.170 ( 0.014 ) + & & -0.224 ( 0.015 ) + & & + & 100 & -0.102 ( 0.008 ) + & 60 & -0.095 ( 0.009 ) + & 21 cm continuum & -0.076 ( 0.009 ) + & & -0.130 ( 0.014 ) + & & + & 100 & -0.002 ( 0.006 ) + & 60 & 0.005 ( 0.007 ) + & 21 cm continuum & 0.024 ( 0.007 ) + & & -0.030 ( 0.013 ) +

we estimate that the total mass is , roughly 1/3 of the interstellar hydrogen gas mass , and about 2% of the dynamical mass of the galaxy at our assumed distance of 6 mpc . the high star formation rate at the nucleus appears to be due to a large accumulation of molecular gas rather than a large increase in the star formation efficiency . the mid - plane gas pressure in the outer ( r 10 kpc ) hi arms of ngc 6946 is close to the value at the radial limit ( 10 kpc ) of our observed co disk . if the mid - plane gas pressure is a factor for the formation of molecular clouds , these outer hi gas arms should contain molecular gas which we do not see because they are beyond our detection limit .

we present on the fly " maps of the co(1 - 0 ) and co(2 - 1 ) emission covering a 10 10 region of the ngc 6946 . using our co maps and archival vla hi observations we create a total gas surface density map , , for ngc 6946 . the predominantly molecular inner gas disk transitions smoothly into an atomic outer gas disk , with equivalent atomic and molecular gas surface densities at r = 3.5 ( 6 kpc ) . we estimate that the total mass is , roughly 1/3 of the interstellar hydrogen gas mass , and about 2% of the dynamical mass of the galaxy at our assumed distance of 6 mpc . the value of the co(2 - 1)/co(1 - 0 ) line ratio ranges from 0.35 to 2 ; 50% of the map is covered by very high ratio , , gas . the very high ratios are predominantly from interarm regions and appear to indicate the presence of wide - spread optically thin gas . star formation tracers are better correlated with the total neutral gas disk than with the molecular gas by itself implying . using the 100 and 21 cm continuum from ngc 6946 as star formation tracers , we arrive at a gas consumption timescale of 2.8 gyr , which is relatively uniform across the disk . the high star formation rate at the nucleus appears to be due to a large accumulation of molecular gas rather than a large increase in the star formation efficiency . the mid - plane gas pressure in the outer ( r 10 kpc ) hi arms of ngc 6946 is close to the value at the radial limit ( 10 kpc ) of our observed co disk . if the mid - plane gas pressure is a factor for the formation of molecular clouds , these outer hi gas arms should contain molecular gas which we do not see because they are beyond our detection limit .

1511.00218

i

the understanding of low - energy electronic structures and excitations in real materials is an important subject of condensed - matter physics and material science . interesting phenomena such as a non - fermi - liquid behavior and unconventional superconductivity are caused by the instability of electronic structures near the fermi level . a common feature is often found in the band structures showing such phenomena ; isolated bands appear near the fermi level . the width of these isolated bands is typically the order of 1 ev , which is comparable to local electronic interactions . thus , in these isolated bands , the kinetic and potential energies compete with each other , and the competition is often discussed within a local - interaction approximation , as in the hubbard model . in real materials , however , there exist various elementally excitations not described by the local electronic interaction . the plasmon in metallic systems or the exciton in insulating systems are well known examples of such nonlocal excitations , which result from the long - range coulomb interaction . in the above - mentioned isolated - band systems , the plasmon excitation can occur in this band , and its energy scale can be very small ( of the order of 1 ev ) , which is comparable to the bandwidth and the size of the local coulomb interaction . in this study , we investigate the effect of the low - energy - plasmon fluctuation on the electronic structure of real isolated - band systems from first principles . for this purpose , we choose two materials , a quasi - one dimensional organic conductor ( tmtsf)@xmath1pf@xmath2 ( ref . ) , where tmtsf stands for tetramethyltetraselenafulvalene , and a three - dimensional perovskite transition - metal oxide srvo@xmath3 ( ref . ) . these materials are typical isolated - band systems and are studied as benchmark materials of the correlated metal , where the local coulomb - interaction effect on the electronic properties is investigated with much interest . @xcite in the present work , we focus on the low - energy plasmon effect on the electronic structure . @xcite through the comparison between theoretical and experimental results on the plasmon - related properties and spectral functions , we verify low - energy plasmon effects on the electronic structure of the real system . the organic conductor ( tmtsf)@xmath1pf@xmath2 is a representative quasi - one - dimensional material , @xcite and basically behaves as a good metallic conductor . @xcite at low temperature ( around 12 k ) , it undergoes a transition to a spin - density - wave phase . @xcite in the high - temperature metallic region , photoemission spectroscopy has observed small spectral weight near the fermi level . @xcite from this observation and the quasi - one - dimensional nature , the origin of the renormalization has been discussed in view of the tomonaga - luttinger liquid . @xcite on the other hand , this material exhibits clear low - energy plasma edges around 0.1 - 1 ev in the reflectance spectra . @xcite therefore , this plasmon excitation would also be a prominent renormalization factor of the electronic structure . the transition - metal oxide srvo@xmath3 is another well known correlated metal . @xcite many high - resolution photoemission measurements @xcite including the bulk - sensitive version @xcite were performed and clarified a strong renormalization of the low - energy isolated band and the satellite peak just below this band . the origin has actively been discussed in terms of the local electronic correlation . @xcite on the other hand , the reflectance and electron - energy - loss spectra have clarified low - energy plasmon excitations around 1.4 ev . @xcite in density - functional band structure , srvo@xmath3 has isolated bands of the @xmath4 orbitals around the fermi level , whose bandwidth is about 2.7 ev . @xcite also , the constrained random phase approximation gives an estimate of the local electronic interaction of @xmath52 - 3 ev . @xcite thus , the energy scale of the experimentally observed plasmon excitation is comparable to the bandwidth and the local coulomb interaction . hence , the low - energy plasmon fluctuation would certainly be relevant to the low - energy properties . to study how the plasmon excitation affects the electronic structure , we perform _ ab initio _ calculations based on the @xmath0 approximation . @xcite the @xmath0 calculation considers the self - energy effect due to the plasmon fluctuation and describes properly quasiparticle energies in the valence region . on the other hand , the description for the plasmon satellite in the spectral function is known to be less accurate . @xcite in order to improve this deficiency , _ ab initio _ @xmath0 plus cumulant ( @xmath0+@xmath6 ) calculations have recently been performed . @xcite the accuracy of the @xmath0+@xmath6 method has been verified in bulk silicon @xcite and simple metals , @xcite where the satellite property is satisfactorily improved . in the present study , we apply the @xmath0+@xmath6 method to the study of the above mentioned isolated - band systems , and show that the low - energy plasmon fluctuations modify substantially the low - energy electronic structure . the present paper is organized as follows . in sec . ii , we describe the @xmath0 and @xmath0+@xmath6 methods to calculate dielectric and spectral properties . computational details and results for ( tmtsf)@xmath1pf@xmath2 and srvo@xmath3 are given in sec . we also discuss the comparison between theory and experiment , focusing on the renormalization of the electronic structure due to the plasmon excitation . the summary is given in sec .

we present _ ab initio _ these materials exhibit characteristic low - energy band structures around the fermi level , which bring about interesting low - energy properties ; the low - energy bands near the fermi level are isolated from the other bands and , in the isolated bands , unusually low - energy plasmon excitations occur . to study the effect of this low - energy - plasmon fluctuation on the electronic structure , we calculate spectral functions and photoemission spectra using the _ ab initio _ cumulant expansion of the green s function based on the self - energy .

we present _ ab initio _ plus cumulant - expansion calculations for an organic compound ( tmtsf)pf and a transition - metal oxide srvo . these materials exhibit characteristic low - energy band structures around the fermi level , which bring about interesting low - energy properties ; the low - energy bands near the fermi level are isolated from the other bands and , in the isolated bands , unusually low - energy plasmon excitations occur . to study the effect of this low - energy - plasmon fluctuation on the electronic structure , we calculate spectral functions and photoemission spectra using the _ ab initio _ cumulant expansion of the green s function based on the self - energy . we found that the low - energy plasmon fluctuation leads to an appreciable renormalization of the low - energy bands and a transfer of the spectral weight into the incoherent part , thus resulting in an agreement with experimental photoemission data .

1511.00218

c

we have performed _ ab initio _ @xmath0 plus cumulant - expansion calculations for an organic conductor ( tmtsf)@xmath1pf@xmath2 and a transition - metal oxide srvo@xmath3 to study the low - energy plasmon - fluctuation effect on the electronic structure . the bands around the fermi level of these materials are isolated from the other bands , and the low - energy plasmon excitations derived from these isolated bands exist . our calculated reflectance spectra well identify the experimental low - energy plasmon peaks . by calculating the cumulant - expanded green s function based on the @xmath0 approximation to the self - energy , we have simulated spectral functions and compare them with photoemission data . we have found agreements between them , indicating that the low - energy plasmon excitation certainly affects the low - energy electronic structure ; it reduces the quasiparticle spectral weight around the fermi level and leads to the weight transfer to the satellite parts . this effect was found to be more or less appreciable in ( tmtsf)@xmath1pf@xmath2 than in srvo@xmath3 . in particular , in ( tmtsf)@xmath1pf@xmath2 , the spectrum at the standard @xmath0 level exhibits a clear plasmaron state , but considering the plasmon - fluctuation effects not treated in the standard @xmath0 calculation leads to the disappearance of the state in the @xmath7 . since the low - energy isolated - band structure is commonly found in various materials , the low - energy plasmon effect pursued in the present work can provide a basis for understanding the electronic structure of real systems . the recent progress in the photoemission experiment for the correlated materials ( refs . ) requires a concomitant progress on the theoretical side and the present _ ab initio _ many - body calculations would provide firm basis for them . in the present study , we have focused on the long - range correlation and treated it effectively with the cumulant - expansion method , while the short - ranged correlation , which is appropriately described by the @xmath185-matrix framework , is neglected . this is a future challenge which remains to be explored . in addition , recent photoemission spectroscopy reveals appreciable differences in the electronic structure between the bulk and thin film systems . @xcite the electronic structure of the surface is very sensitive to the atomic configurations at the surface @xcite and therefore careful analyses for the atomic structure and its effect on electronic structure are required . the _ ab initio _ calculations for the surface effect are clearly important for the deep understanding of the spectroscopy of the real materials . this is also a future challenge . we would like to thank norikazu tomita and teppei yoshida for useful discussions . calculations were done at supercomputer center at institute for solid state physics , university of tokyo . this work was supported by grants - in - aid for scientific research ( no . 22740215 , 22104010 , 23110708 , 23340095 , 23510120 , 25800200 ) from mext , japan , and consolidator grant correlmat of the european research council ( project number 617196 ) .

plus cumulant - expansion calculations for an organic compound ( tmtsf)pf and a transition - metal oxide srvo . we found that the low - energy plasmon fluctuation leads to an appreciable renormalization of the low - energy bands and a transfer of the spectral weight into the incoherent part , thus resulting in an agreement with experimental photoemission data .

we present _ ab initio _ plus cumulant - expansion calculations for an organic compound ( tmtsf)pf and a transition - metal oxide srvo . these materials exhibit characteristic low - energy band structures around the fermi level , which bring about interesting low - energy properties ; the low - energy bands near the fermi level are isolated from the other bands and , in the isolated bands , unusually low - energy plasmon excitations occur . to study the effect of this low - energy - plasmon fluctuation on the electronic structure , we calculate spectral functions and photoemission spectra using the _ ab initio _ cumulant expansion of the green s function based on the self - energy . we found that the low - energy plasmon fluctuation leads to an appreciable renormalization of the low - energy bands and a transfer of the spectral weight into the incoherent part , thus resulting in an agreement with experimental photoemission data .

astro-ph0703769

i

we present a catalog of variable stars within @xmath6 pc of sgr a * in the near - infrared @xmath1 and @xmath2 bands over a three - year baseline ; @xmath153% of these variables are previously unidentfied . this is the first photometric variability study of the galactic center to use the technique of image subtraction . we find several new periodic sources . of these , psd j174535.60 - 290035.4 and psd j174538.34 - 290036.7 have the shortest periods , with @xmath154 days . we believe psd j174535.60 - 290035.4 is a pulsating variable , though it is unlikely to be a cepheid . psd j174538.34 - 290036.7 appears to be a nearly edge - on contact eclipsing binary system , in which case the period is actually @xmath155 days and the individual components are likely to be quite large . another periodic source , psd j174540.16 - 290055.7 ( @xmath156 days ) , is coincident with the x - ray variable cxogc j174540.1 - 290055 ; we suspect that this source is a cataclysmic variable . among the previously identified sources in the catalog , we associate irs 10 * with a bright x - ray point source and oh , h@xmath4o , and sio masers , suggesting that it is an agb star with an accreting companion . and @xmath2 , sorted by ra as in table [ tbl : vary ] . for each source , the extent of the @xmath1- and @xmath2-band ranges plotted is the same . points lying below the magnitude limit ( see [ sec : analysis ] ) are not included.[fig : lcs ] ] lrrrrl psd j174535.45 - 290007.6 & @xmath157 & @xmath158 & 12.6 & & + psd j174535.89 - 290113.5 & @xmath159 & @xmath160 & 14.2 & 10.9 & + psd j174535.98 - 290011.5 & @xmath161 & @xmath162 & 12.1 & 9.7 & + psd j174536.18 - 290047.6 & @xmath163 & @xmath164 & 12.4 & 9.6 & + psd j174536.53 - 290125.9 & @xmath165 & @xmath166 & 11.9 & 9.5 & + psd j174536.63 - 290015.4 & @xmath167 & @xmath168 & 10.7 & 8.8 & + psd j174536.85 - 290114.5 & @xmath169 & @xmath170 & 13.1 & 10.1 & + psd j174536.97 - 290138.0 & @xmath171 & @xmath172 & 15.0 & 10.2 & + psd j174536.97 - 290019.4 & @xmath173 & @xmath174 & 12.5 & 10.0 & + psd j174537.10 - 290103.9 & @xmath175 & @xmath176 & 13.9 & 10.8 & + psd j174537.24 - 290045.7 & @xmath177 & @xmath178 & 12.6 & 9.9 & v4910 sgr + psd j174537.42 - 290042.2 & @xmath179 & @xmath180 & 13.8 & 10.6 & + psd j174537.52 - 290046.2 & @xmath181 & @xmath182 & 14.5 & 10.6 & + psd j174537.52 - 290041.7 & @xmath183 & @xmath184 & 12.8 & 10.0 & + psd j174537.66 - 285955.2 & @xmath185 & @xmath186 & 12.7 & & + psd j174537.74 - 290009.2 & @xmath187 & @xmath188 & 13.1 & 10.2 & + psd j174537.91 - 290045.9 & @xmath189 & @xmath190 & 12.7 & 10.0 & + psd j174537.98 - 290134.4 & @xmath191 & @xmath192 & 9.5 & 8.6 & cxogc j174537.9 - 290134 + psd j174537.98 - 290004.4 & @xmath191 & @xmath193 & 12.0 & 9.2 & + psd j174538.02 - 290102.6 & @xmath194 & @xmath195 & 12.5 & 9.3 & v4911 sgr ; blended + psd j174538.05 - 290058.1 & @xmath196 & @xmath197 & 12.7 & 10.2 & + psd j174538.10 - 290033.5 & @xmath198 & @xmath199 & 13.3 & 10.3 & + psd j174538.19 - 290005.0 & @xmath200 & @xmath201 & 12.5 & 9.7 & + psd j174538.29 - 290011.1 & @xmath202 & @xmath203 & 13.4 & 9.9 & + psd j174538.61 - 290054.3 & @xmath204 & @xmath205 & 13.0 & 9.5 & + psd j174538.74 - 290012.7 & @xmath206 & @xmath207 & 11.1 & 9.2 & + psd j174538.79 - 290004.6 & @xmath208 & @xmath209 & 11.8 & 9.3 & + psd j174538.95 - 290139.6 & @xmath210 & @xmath211 & & 9.4 & + psd j174538.99 - 290102.2 & @xmath212 & @xmath213 & 14.3 & 12.4 & + psd j174539.09 - 290039.3 & @xmath214 & @xmath215 & 12.7 & 9.8 & blended + psd j174539.23 - 290036.3 & @xmath216 & @xmath217 & 14.1 & 10.8 & + psd j174539.31 - 290016.3 & @xmath218 & @xmath219 & 11.6 & 9.2 & cxogc j174539.3 - 290016 + psd j174539.31 - 290054.1 & @xmath218 & @xmath220 & 12.6 & 10.0 & + psd j174539.36 - 290055.6 & @xmath221 & @xmath222 & 13.0 & 10.5 & + psd j174539.39 - 290014.6 & @xmath223 & @xmath224 & 10.9 & 8.4 & blended + psd j174539.45 - 290056.6 & @xmath225 & @xmath226 & 12.4 & 9.1 & blended + psd j174539.56 - 290057.3 & @xmath227 & @xmath228 & & 12.5 & + psd j174539.58 - 290049.5 & @xmath229 & @xmath230 & 12.4 & 9.6 & + psd j174539.60 - 290107.6 & @xmath231 & @xmath232 & 15.3 & 11.8 & @xmath1 past threshold + psd j174539.73 - 290050.2 & @xmath233 & @xmath234 & 12.5 & 9.9 & + psd j174539.75 - 290055.5 & @xmath235 & @xmath236 & 12.5 & 9.6 & irs 22 + psd j174539.76 - 290026.4 & @xmath237 & @xmath238 & 14.1 & 10.6 & irs 34w + psd j174539.77 - 290043.9 & @xmath239 & @xmath240 & 12.1 & 9.6 & + psd j174539.79 - 290008.5 & @xmath241 & @xmath242 & 11.4 & 8.5 & + psd j174539.79 - 290029.9 & @xmath241 & @xmath243 & 13.2 & 10.6 & irs 13e + psd j174539.83 - 290053.9 & @xmath244 & @xmath245 & 9.6 & 7.6 & blended + psd j174539.86 - 290024.7 & @xmath246 & @xmath247 & & 10.5 & + psd j174539.86 - 290019.7 & @xmath248 & @xmath249 & 12.8 & 10.0 & + psd j174539.90 - 290034.5 & @xmath250 & @xmath251 & 13.3 & 10.6 & blended + psd j174539.91 - 290011.0 & @xmath247 & @xmath252 & 12.5 & 9.9 & cxogc j174539.9 - 290012 + psd j174539.93 - 290024.9 & @xmath253 & @xmath254 & 13.9 & 10.5 & + psd j174539.95 - 285951.1 & @xmath255 & @xmath256 & 12.2 & & + psd j174539.95 - 290110.8 & @xmath257 & @xmath258 & 14.2 & 10.2 & + psd j174539.99 - 290016.5 & @xmath259 & @xmath260 & 12.1 & 9.8 & irs 15sw + psd j174540.02 - 290037.2 & @xmath261 & @xmath262 & 11.8 & 9.4 & irs 14sw + psd j174540.04 - 290018.0 & @xmath263 & @xmath264 & 12.3 & 10.3 & irs 15ne + psd j174540.04 - 290022.7 & @xmath263 & @xmath263 & 9.0 & & irs 7 , saturated in @xmath2 + psd j174540.04 - 290027.0 & @xmath263 & @xmath235 & 11.6 & 9.3 & irs 16nw + psd j174540.04 - 290019.5 & @xmath265 & @xmath266 & 12.6 & & + psd j174540.06 - 290112.1 & @xmath267 & @xmath268 & 12.0 & 9.9 & + psd j174540.11 - 290036.4 & @xmath269 & @xmath270 & 11.8 & 9.1 & irs 14ne + psd j174540.21 - 290043.2 & @xmath271 & @xmath272 & 14.1 & 10.7 & + psd j174540.21 - 290056.3 & @xmath271 & @xmath273 & 15.6 & & + psd j174540.31 - 285953.8 & @xmath274 & @xmath275 & 13.0 & & + psd j174540.31 - 290039.5 & @xmath274 & @xmath276 & 12.3 & 9.4 & + psd j174540.37 - 285954.0 & @xmath277 & @xmath278 & 11.7 & & blended + psd j174540.38 - 290033.6 & @xmath279 & @xmath199 & & 11.4 & + psd j174540.45 - 290036.3 & @xmath260 & @xmath217 & 13.8 & 11.1 & cxogc j174540.4 - 290036 + psd j174540.47 - 290034.6 & @xmath219 & @xmath280 & 11.3 & 8.6 & irs 9 , blended + psd j174540.58 - 290026.7 & @xmath281 & @xmath239 & 11.9 & 9.2 & irs 1ne + psd j174540.64 - 290023.6 & @xmath282 & @xmath255 & & 10.2 & irs 10 * , blended + psd j174540.83 - 290034.0 & @xmath283 & @xmath284 & 12.2 & 9.1 & irs 28 + psd j174540.84 - 290027.0 & @xmath285 & @xmath286 & 12.0 & 9.5 & irs 32 + psd j174541.04 - 290022.7 & @xmath287 & @xmath263 & 11.1 & 8.5 & + psd j174541.09 - 285952.6 & @xmath288 & @xmath289 & 13.2 & & + psd j174541.17 - 290046.9 & @xmath290 & @xmath291 & 10.7 & 7.9 & irs 19 , blended + psd j174541.20 - 290039.0 & @xmath292 & @xmath210 & 11.6 & 9.1 & + psd j174541.27 - 290049.9 & @xmath293 & @xmath294 & 12.8 & 9.6 & + psd j174541.35 - 290033.0 & @xmath295 & @xmath296 & 13.2 & 10.2 & + psd j174541.36 - 290022.7 & @xmath297 & @xmath298 & 11.9 & 9.4 & + psd j174541.39 - 290126.6 & @xmath299 & @xmath300 & 13.9 & 9.6 & + psd j174541.56 - 285949.6 & @xmath301 & @xmath302 & 12.4 & & + psd j174541.59 - 290023.0 & @xmath303 & @xmath304 & 13.7 & 10.8 & + psd j174541.75 - 290004.6 & @xmath305 & @xmath209 & 12.1 & 9.0 & + psd j174541.75 - 290013.0 & @xmath305 & @xmath306 & 12.3 & 9.3 & + psd j174542.01 - 290024.4 & @xmath307 & @xmath308 & 13.2 & 10.0 & + psd j174542.13 - 290045.8 & @xmath309 & @xmath310 & 13.5 & 10.3 & + psd j174542.39 - 285950.7 & @xmath311 & @xmath312 & 10.8 & & + psd j174542.60 - 290103.9 & @xmath313 & @xmath176 & 14.3 & 11.0 & + psd j174542.72 - 285957.4 & @xmath314 & @xmath315 & 10.5 & 7.7 & v4928 sgr , blended + psd j174542.76 - 290125.5 & @xmath316 & @xmath317 & 10.7 & 8.7 & + psd j174542.88 - 285951.3 & @xmath318 & @xmath319 & 12.7 & & + psd j174542.92 - 285958.4 & @xmath320 & @xmath321 & 14.2 & 11.8 & + psd j174543.01 - 290011.9 & @xmath322 & @xmath323 & 10.8 & 8.5 & + psd j174543.19 - 290013.0 & @xmath324 & @xmath306 & 10.9 & 8.2 & v4930 sgr , blended + psd j174543.29 - 290118.8 & @xmath325 & @xmath326 & & 11.9 & + psd j174543.31 - 290014.3 & @xmath327 & @xmath328 & 14.8 & 12.6 & + psd j174543.50 - 290005.3 & @xmath329 & @xmath292 & & 9.0 & + lrrrrrl psd j174535.60 - 290035.4 & @xmath330 & @xmath331 & 14.5 & 12.3 & @xmath332 & cxogc j174535.6 - 290034 + psd j174537.11 - 290033.1 & @xmath333 & @xmath296 & 10.0 & 12.8 & @xmath334 & + psd j174538.28 - 290006.8 & @xmath335 & @xmath336 & 11.9 & 9.3 & @xmath337 & + psd j174538.35 - 290036.7 & @xmath338 & @xmath339 & 14.9 & 12.1 & @xmath340 & + psd j174538.97 - 290123.7 & @xmath341 & @xmath161 & 14.6 & 10.8 & @xmath342 & + psd j174538.98 - 290007.7 & @xmath343 & @xmath287 & 13.0 & 9.9 & @xmath344 & + psd j174539.39 - 290039.5 & @xmath345 & @xmath346 & 11.9 & 9.8 & @xmath347 & + psd j174539.79 - 290035.2 & @xmath348 & @xmath349 & 11.8 & 8.8 & @xmath350 & irs 12n , very blended + psd j174540.12 - 290029.6 & @xmath351 & @xmath352 & 10.8 & 8.3 & @xmath353 & irs 16sw + psd j174540.13 - 290016.8 & @xmath354 & @xmath355 & 12.1 & 9.4 & @xmath356 & irs 15n , blended + psd j174540.16 - 290055.7 & @xmath357 & @xmath358 & 12.1 & 9.3 & @xmath359 & cxogc j174540.1 - 290055 + psd j174542.36 - 290011.2 & @xmath360 & @xmath203 & 13.4 & 10.6 & @xmath361 & + psd j174540.25 - 290027.2 & @xmath266 & @xmath233 & 10.5 & 8.3 & @xmath362 & irs 16ne + psd j174543.14 - 290050.2 & @xmath363 & @xmath234 & 12.0 & 9.1 & @xmath364 & blended + lllcl @xcite & @xmath365 & 1991 , 1992 , 1993 ; & @xmath2 & irs 9 , 10 * , 12n , 14sw , 28 + & & 3 epochs & & + @xcite & central @xmath366 & may 1993april 1995 ; & @xmath367 , @xmath1 , @xmath2 , @xmath368 & irs 7 , 9 , 12n + & & 58 epochs & & + @xcite & @xmath369 & august 1992may 1998 ; & @xmath2 & irs 7 , 9 , 10 * , 12n , 16sw and + & & @xmath370 epochs & & potentially irs 1ne , 6 , 14sw + @xcite & @xmath371 & 19941997 ; @xmath372 epochs & @xmath2 & only large - amplitude , + & & & & periodic variability + @xcite & @xmath373 & 19952005 ; 50 epochs & @xmath2 & irs 16sw , 16nw , 16cc , 29n , + & & & & and 12 s stars lllll irs 1ne & [ fig : lcs](d ) & & & psd j174540.58 - 290026.7 + irs 7 & [ fig : lch ] & & lpv & psd j174540.04 - 290022.7 + irs 9 & [ fig : lcs](d ) & & lpv & psd j174540.47 - 290034.6 + irs 10 * & [ fig : lck ] & & oh / ir , x - ray source & psd j174540.64 - 290023.6 + irs 12n & [ fig : periodic ] & @xmath350 & lpv & psd j174539.79 - 290035.2 + irs 14sw & [ fig : lcs](d ) & & & psd j174540.02 - 290037.2 + irs 14ne & [ fig : lcs](d ) & & agb & psd j174540.11 - 290036.4 + irs 15sw & [ fig : lcs](c ) & & wr & psd j174539.99 - 290016.5 + irs 16ne & [ fig : periodic ] & @xmath362 & lbv ? & psd j174540.25 - 290027.2 + irs 16sw & [ fig : periodic ] & @xmath374 & binary & psd j174540.12 - 290029.6 + irs 28 & [ fig : lcs](d ) & & & psd j174540.83 - 290034.0 + irs 34w & [ fig : lcs](c ) & & lbv ? & psd j174539.76 - 290026.4 + v4910 sgr & [ fig : lcs](a ) & & & psd j174537.24 - 290045.7 + v4911 sgr & [ fig : lcs](b ) & & sio maser ? & psd j174538.02 - 290102.6 + v4928 sgr & [ fig : lcs](e ) & & oh / ir & psd j174542.72 - 285957.4 + v4930 sgr & [ fig : lcs](e ) & & mira & psd j174543.19 - 290013.0 + psd j174535.60 - 290035.4 & [ fig : periodic ] & @xmath332 & & + psd j174538.34 - 290036.7 & [ fig : periodic ] & @xmath340 & binary ? ? & + psd j174538.98 - 290007.7 & [ fig : periodic ] & @xmath344 & lpv ? & + psd j174537.98 - 290134.4 & [ fig : lcs](a ) & & x - ray source & + psd j174539.31 - 290016.3 & [ fig : lcs](b ) & & x - ray source & + psd j174540.16 - 290055.7 & [ fig : periodic ] & @xmath359 & variable x - ray & + psd j174541.39 - 290126.6 & [ fig : lcs](e ) & & oh / ir & + psd j174542.36 - 290011.2 & [ fig : periodic ] & @xmath361 & & + psd j174543.14 - 290050.2 & [ fig : periodic ] & @xmath364 & & + , d. l. , atwood , b. , belville , s. r. , brewer , d. f. , byard , p. l. , gould , a. , mason , j. a. , obrien , t. p. , pappalardo , d. p. , pogge , r. w. , steinbrecher , d. p. , & teiga , e. j. 2003 , in instrument design and performance for optical / infrared ground - based telescopes . edited by iye , masanori ; moorwood , alan f. m. proceedings of the spie , volume 4841 , pp . 827 - 838 ( 2003 ) . , m. iye & a. f. m. moorwood , 827838 , f. , genzel , r. , alexander , t. , abuter , r. , paumard , t. , ott , t. , gilbert , a. , gillessen , s. , horrobin , m. , trippe , s. , bonnet , h. , dumas , c. , hubin , n. , kaufer , a. , kissler - patig , m. , monnet , g. , strbele , s. , szeifert , t. , eckart , a. , schdel , r. , & zucker , s. 2005 , , 628 , 246 , t. , genzel , r. , martins , f. , nayakshin , s. , beloborodov , a. m. , levin , y. , trippe , s. , eisenhauer , f. , ott , t. , gillessen , s. , abuter , r. , cuadra , j. , alexander , t. , & sternberg , a. 2006 , , 643 , 1011 , m. f. , cutri , r. m. , stiening , r. , weinberg , m. d. , schneider , s. , carpenter , j. m. , beichman , c. , capps , r. , chester , t. , elias , j. , huchra , j. , liebert , j. , lonsdale , c. , monet , d. g. , price , s. , seitzer , p. , jarrett , t. , kirkpatrick , j. d. , gizis , j. e. , howard , e. , evans , t. , fowler , j. , fullmer , l. , hurt , r. , light , r. , kopan , e. l. , marsh , k. a. , mccallon , h. l. , tam , r. , van dyk , s. , & wheelock , s. 2006 , , 131 , 1163

we present a catalog of 110 variable stars within of sgr a * based on image subtraction of near - infrared ( and ) photometry . we associate irs 10 * with oh , sio , and ho masers and a bright x - ray point source ; this analysis suggests irs 10 * is an agb star with an accreting companion . among the newly discovered sources are a probable cataclysmic variable , a potential edge - on contact 84 day period eclipsing binary , and a possible 41 day period pulsating variable

we present a catalog of 110 variable stars within of sgr a * based on image subtraction of near - infrared ( and ) photometry . our images were obtained over 133 nights from 2000 to 2002 in-band and over 134 nights from 2001 to 2002 in-band ; the typical fwhm is . we match the catalog to other near - infrared , x - ray , and radio ( i.e. , maser ) data , and we discuss some of the more interesting objects . the catalog includes 14 sources with measurable periods , several known long - period variables and three new lpv candidates . we associate irs 10 * with oh , sio , and ho masers and a bright x - ray point source ; this analysis suggests irs 10 * is an agb star with an accreting companion . among the newly discovered sources are a probable cataclysmic variable , a potential edge - on contact 84 day period eclipsing binary , and a possible 41 day period pulsating variable

astro-ph9808206

i

with the recent development of large format infrared array detectors , high quality photometric surveys are routinely conducted at wavelengths between 12.5 @xmath3 m . soon the completion of the 2 micron all sky survey ( 2mass ; skrutskie 1997 ) and denis ( epchtein et al . 1997 ) will provide comprehensive catalogues of near infrared sources with detection limits sensitive to a wide variety of stellar and non stellar objects . infrared spectra will be required for appropriate identification of many of these sources , and for further study of their astrophysical properties . the pioneering study of johnson and mendez ( 1970 ) was the first to explore the spectra of a large sample of normal stars in the near infrared . however many years passed before improvements in instrumentation made possible similar observations of large numbers of targets of astrophysical interest . the majority of the work done in near infrared spectroscopy to date has been focused on the k band , in large part because intrinsically cool or heavily obscured objects are typically brighter at k band than in the j or h bands . in 1986 , kleinmann and hall ( 1986 ; kh86 ) provided the first comprehensive medium resolution atlas ( @xmath4 ) of stellar spectra in the k - band , covering all luminosity classes , but restricted to spectral types between f8-m7 . more recently , wallace and hinkle ( 1997 ; wh97 ) have extended the kh96 k band atlas , using the same fts spectrograph on the kpno 4 m with @xmath4 , but including stellar spectra spanning spectral types o - m and luminosity classes i - v . they also summarize the considerable body of work directed toward k band spectroscopy in the last decade . while in many situations , the k band will be the wavelength selection of choice for spectroscopic studies of highly obscured or very cool objects , the presence of circumstellar dust ( @xmath5 ; pollack et al . 1994 ) often results in significant excess continuum emission longward of @xmath6 m . this continuum excess is commonly found in two important classes of objects : young stars with circumstellar disks ( e.g. meyer , calvet , & hillenbrand , 1997 ) and evolved stars with extensive envelopes from mass - loss ( e.g. le bertre 1997 ) . near infrared excess due to warm dust can also complicate spectroscopic studies of composite stellar systems aimed at discerning the stellar populations of other galaxies ( e.g. schinnerer et al . continuum excess longward of @xmath6 m will weaken or even render invisible the photospheric features in the k band , while the photosphere will dominate at shorter wavelengths . in such a situation , near infrared spectra shortward of 2 @xmath3 m will be required to see the stellar photosphere too obscured to be detected optically . to date , there has been relatively little work in the h band ( 1.551.75 @xmath3 m ) . recent publications include : i ) observations of 40 g , k , and m stars of luminosity class i and iii at @xmath7 by origlia , moorwood , and oliva ( 1993 ) ; ii ) the library of 56 spectra o m of luminosity class i , ii , and v at @xmath8 ( lancon & rocca volmerange , 1992 ) ; iii ) a library of 37 stars of luminosity classes i , iii , v at @xmath9 ( dallier , boisson , & joly 1996 ) over a limited portion of the h band ; and iv ) a study of 9 ob stars at @xmath10 ( blum et al . 1997 ) . here we present an h band spectral atlas for 85 stars of nearly solar abundance with spectral types on the mk classification system ranging from 07-m5 and luminosity classes i - v . these @xmath4 spectra were collected with the same fts at the kpno 4 m as the k band atlases of kh86 and wh97 . in section 2 , we describe the sample selection and in section 3 we describe the observations and calibration of the data . in section 4 we discuss the dependence of the spectral features on temperature and luminosity and suggest a two dimensional classification appropriate for late - type stars . in section 5 we discuss near ir spectral classification with regard to wavelength range / spectral resolution , and conclude with a summary of our results .

we present a catalogue of h band spectra for 85 stars of approximately solar abundance observed at a resolving power of 3000 with the kpno mayall 4 m fts . the atlas covers spectral types o7m5 and luminosity classes i - v as defined on the mk system . the line ratios permit spectral classification in the presence of continuum excess emission , which is commonly found in pre main sequence and evolved stars . we demonstrate that with spectra of obtained at it is possible to derive spectral types within subclasses for late type stars .

we present a catalogue of h band spectra for 85 stars of approximately solar abundance observed at a resolving power of 3000 with the kpno mayall 4 m fts . the atlas covers spectral types o7m5 and luminosity classes i - v as defined on the mk system . we identify both atomic and molecular indices and line ratios which are temperature and luminosity sensitive allowing spectral classification to be carried out in the h band . the line ratios permit spectral classification in the presence of continuum excess emission , which is commonly found in pre main sequence and evolved stars . we demonstrate that with spectra of obtained at it is possible to derive spectral types within subclasses for late type stars . these data are available electronically through the astronomical data center in addition to being served on the world wide web . # 1#2#3#4#5#6#7 to = -0.2 truein

1601.00243

i

fluid flows in two spatial dimensions have been the subject of substantial research efforts in recent decades . for the greater part of the twentieth century , it was generally considered that two - dimensional ( 2d ) flows were merely a theoretical idealization with limited practical relevance . this conception has changed drastically since the 1980s , when experiments in thin electrolyte layers @xcite , soap films @xcite , and liquid metals @xcite demonstrated that nearly 2d flows can indeed be realized in the laboratory . today , experimental approximations of 2d flows are widely employed as models of atmospheric and oceanic flows @xcite . being theoretically and experimentally more amenable than their three - dimensional ( 3d ) counterparts , 2d flows have also served as platforms for studying new phenomena such as turbulent cascades @xcite , coherent structures @xcite , and mixing @xcite . hence , experimental studies of 2d flows , with quantitative comparisons to theoretical predictions , are important for improving our understanding of fluid flows , in both 2d and 3d . perhaps one of the best known 2d fluid flow models is a flow introduced by andrey kolmogorov in 1959 as a mathematical problem for studying hydrodynamic stability @xcite . the kolmogorov flow represents the motion of a viscous fluid in two dimensions ( we will refer to these as @xmath0 and @xmath1 ) driven by a forcing that points along the @xmath0-direction and varies sinusoidally along the @xmath1-direction . the fluid flow is considered incompressible , @xmath2 , and is governed by the 2d navier - stokes equation , @xmath3 here , @xmath4 is the velocity field , @xmath5 is the 2d pressure field , and @xmath6 represents the driving force with amplitude @xmath7 and wavenumber @xmath8 . the parameters @xmath9 and @xmath10 are the density and the kinematic viscosity of the fluid , respectively . kolmogorov flow has served as a convenient model for understanding a wide variety of hydrodynamic phenomena in 2d , such as fluid instabilities @xcite , 2d turbulence @xcite , and coherent structures @xcite . practically realizable flows , however , are never strictly 2d . experimental approximations of kolmogorov flow have often been carried out in either shallow layers of electrolytes @xcite or in soap films @xcite , wherein geometric confinement suppresses the component of velocity along one of the spatial directions ( @xmath11 ) . the remaining two velocity components , however , generally depend on both extended and confined coordinates , making the flow `` quasi - two - dimensional '' ( q2d ) . to account for the dependence on the confined coordinate , q2d flows in shallow layers have often been modelled by adding a linear friction term to the 2d navier - stokes equation ( [ eq:2dns ] ) : @xmath12 where @xmath13 is a constant . the addition of this term was first suggested by @xcite to model a q2d flow generated in a homogeneous shallow electrolyte layer . in such a flow , the bottom of the fluid layer is constrained to be at rest because it is in contact with the solid surface of the container holding the fluid . this no - slip constraint at the bottom of the fluid layer causes a gradient in the magnitude of the horizontal velocity * u * along the confined direction @xmath11 . @xcite rationalized that the dissipation due to this shear , for sufficiently shallow fluid layers , is captured by the linear friction term . in the context of q2d flows in electrolyte layers , this term has come to be known as `` rayleigh friction . '' experimental flows in thin layers as well as models that employ equation ( [ eq:2dns_wf ] ) are now commonly referred to as `` kolmogorov - like , '' when the forcing profile is nearly sinusoidal . experimental realizations of q2d flows in recent years have employed two - fluid - layer setups : either a setup with _ miscible _ layers comprised of a heavy electrolyte fluid ( salt water ) beneath a lighter nonconducting fluid ( pure water ) @xcite , or a setup with _ immiscible _ layers comprised of a heavy dielectric fluid beneath a lighter electrolyte @xcite . the rationale behind these modifications was that in addition to confinement , density stratification and immiscibility should enhance two - dimensionality in the top layer . theoretical models of q2d experimental flows realized in stratified layers of fluids , however , have not accurately modelled the effect of the gradient in the magnitude of horizontal velocity @xmath14 along the confined direction @xmath11 . consequently , experiments were compared with simulations based on the 2d model ( [ eq:2dns_wf ] ) with empirically estimated parameters @xcite . to address this deficiency , @xcite have investigated the variation in the horizontal velocity @xmath14 along the confined direction @xmath11 for a stratified two - immiscible - layer setup . the following modified version of equation ( [ eq:2dns_wf ] ) which accounts for the inherent three - dimensionality of a q2d flow and describes its evolution quantitatively was proposed : @xmath15 where @xmath16 is a numerical prefactor to the advection term that was not accounted for previously . this prefactor reflects the decrease of the mean inertia of the fluid layer due to the variation of the horizontal velocity in the vertical direction . equation ( [ eq:2dns_mod ] ) has been derived from first principles by depth - averaging the full 3d navier - stokes equation under the q2d approximation ( cf . appendix [ sec : sensitivity ] ) . the coefficients @xmath13 , @xmath10 , and @xmath17 are estimated by including effects of both the inhomogeneity in fluid properties as well the vertical profile of @xmath14 along @xmath11 . equations ( [ eq:2dns ] ) and ( [ eq:2dns_wf ] ) can be treated as special cases of ( [ eq:2dns_mod ] ) with suitable choices of the parameters @xmath13 and @xmath17 . furthermore , equation ( [ eq:2dns_wf ] ) can also be obtained from ( [ eq:2dns_mod ] ) by rescaling the variables ( cf . appendix [ sec : nondim ] ) . in this article we study instabilities of a q2d kolmogorov - like flow realized in a setup with two immiscible fluid layers . we compare experiments with results from direct numerical simulations ( dns ) of equation ( [ eq:2dns_mod ] ) . most previous studies of kolmogorov - like flow compared experiments to theoretical analysis of a perfectly sinusoidal shear flow on unbounded or periodic domain @xcite . while some of these studies reported quantitative agreement between theory and experiment in regards to the primary instability , none where able to match simultaneously both the critical reynolds number and the critical wavenumber . even matching one of these required treating the rayleigh friction coefficient @xmath13 as an adjustable parameter . to address these difficulties , we systematically investigate the effects of lateral confinement , which was not properly addressed by the previous studies , using numerical simulations with three different sets of boundary conditions . furthermore , we investigate how the observed flow patterns and their stability are affected by the deviations in the forcing profile from perfect periodicity in the extended directions and by the variation of the forcing profile in the confined direction . finally , we compare the results of numerical simulations with the experimental observations for the secondary instability , which introduces time - dependence into the flow . the prospect of developing a 2d model that quantitatively describes the evolution of a q2d flow is particularly appealing in view of the recent theoretical advances in understanding transitional flows and weak turbulence as a dynamical process guided by nonchaotic , unstable solutions of the navier - stokes equation , often referred to as exact coherent structures ( ecs ) @xcite . the bulk of numerical studies have explored the role of ecs in 3d flows simulated on periodic domains with simple geometries , such as pipe flow , plane couette flow , and plane poiseuille flow . however , experimental evidence for the role of ecs is scarce @xcite , in part due to technical limitations in obtaining spatially and temporally resolved 3d velocity fields . q2d flows , on the other hand , can be quantified using 2d planar velocity fields which are relatively easy to measure . hence , experimental q2d flows that can be described by a 2d model serve as an ideal platform for testing and developing an ecs - based approach to understanding turbulence . recently , @xcite and @xcite have identified dozens of ecs in numerical simulations of a weakly turbulent 2d kolmogorov flow , governed by equation ( [ eq:2dns ] ) with periodic boundary conditions . we hope that this paper helps provide a rigorous foundation for future studies in the same spirit , but with a focus on experimental validation of theoretical predictions . this article is organized as follows . in [ sec : exp_setup ] , we describe the experimental setup employed to generate a q2d kolmogorov - like flow . in [ sec : model ] , we introduce a realistic model of the forcing in the experiment and discuss different types of lateral boundary conditions which are used to study the effects of confinement theoretically . in [ sec : comparison ] , we compare the flow fields obtained from experimental measurements with those from the numerical simulations for different flow regimes and characterize the bifurcations associated with increasing the forcing strength . conclusions are presented in [ sec : conclusion ] .

we also introduce a realistic numerical model approximating the electromagnetic body force in the experiment , which is validated against experimental measurements . we then explore the effect of lateral boundaries by performing direct numerical simulations ( dns ) subject to different types of boundary conditions . we compare the flow regimes observed in the experiment with those obtained from the dns and find that only the simulation subject to physical no - slip boundary conditions provides close , quantitative agreement . our analysis offers additional validation of the 2d model proposed by , as well as a demonstration of the importance of properly modelling the forcing and boundary conditions .

we present a combined experimental and numerical study of the primary and secondary instabilities in a kolmogorov - like flow . the experiment uses electromagnetic forcing with an approximately sinusoidal spatial profile to drive a quasi - two - dimensional ( q2d ) shear flow in a thin layer of electrolyte suspended on a thin lubricating layer of a dielectric fluid . the experimental flow is modelled using a recently proposed 2d equation , derived from first principles by depth - averaging the full three - dimensional navier - stokes equations . we also introduce a realistic numerical model approximating the electromagnetic body force in the experiment , which is validated against experimental measurements . we then explore the effect of lateral boundaries by performing direct numerical simulations ( dns ) subject to different types of boundary conditions . as the strength of the forcing is increased , this q2d flow undergoes a series of bifurcations . we compare the flow regimes observed in the experiment with those obtained from the dns and find that only the simulation subject to physical no - slip boundary conditions provides close , quantitative agreement . our analysis offers additional validation of the 2d model proposed by , as well as a demonstration of the importance of properly modelling the forcing and boundary conditions .

1601.00243

c

in this article , we have presented a combined experimental and numerical study of bifurcations in a q2d kolmogorov - like flow . this flow is realized in the laboratory by electromagnetically driving a stratified layer of electrolyte above an immiscible layer of dielectric . this q2d flow is described using a 2d model derived from first principles by depth - averaging the 3d navier - stokes equation . in contrast , virtually all previous studies have modelled q2d flows using equation ( [ eq:2dns_wf ] ) , a semi - empirical variation of the 2d navier - stokes equation with the addition of a linear friction . also unlike previous studies of kolmogorov - like flows which have assumed a perfectly sinusoidal forcing profile , we have introduced a realistic model of the forcing which has been validated against 3d experimental measurements . to test the importance of lateral confinement , we have compared experimental measurements with numerical simulations using different boundary conditions . we have found that by incorporating realistic , no - slip boundary conditions at all lateral boundaries and a realistic forcing profile , quantitative agreement between the experiment and simulation can be achieved with no adjustable parameters . in particular , the reynolds number @xmath153 for the primary instability can be predicted to within about 5% and the critical wavenumber @xmath178 can be predicted to an accuracy higher than the measurement accuracy . these are significant improvements compared with previous studies , none of which were able to predict both @xmath153 and @xmath178 with this level of accuracy , _ despite using adjustable parameters_. furthermore , we have demonstrated that the model reasonably accurately predicts the modulated flow pattern ( both the wavenumber and the amplitude ) beyond the onset of the primary instability . moreover , this is the first study , experimental or theoretical , to provide a quantitative analysis of the secondary instability of a kolmogorov - like flow which generates a time - dependent pattern of vortices . even for the secondary instability the numerical predictions of the critical reynolds number @xmath175 and the critical period @xmath176 are in general agreement with the experiment , although the accuracy of the numerical predictions decreases with increasing @xmath139 . the discrepancy between the numerical predictions and experiments has been traced back to two factors : ( 1 ) the dependence of the parameters of the 2d model on the average wavenumber of the flow pattern , which in turn depends on the reynolds number , and ( 2 ) the variation of the forcing profile with height . both of these issues point to the limitations of the 2d model of what in reality is a 3d flow , albeit with a strongly suppressed vertical component of the velocity . nonetheless , for a select range of @xmath139 the experimental flow can be reproduced with extremely good quantitative accuracy by making fairly small adjustments to the model parameters , compared with their depth - average values computed for a simple straight flow . the ability of the model to closely reproduce the experimental flow is crucial for the utility of q2d flows for testing the geometrical description of weakly turbulent flows and studying the dynamical role of exact coherent structures . such tests will be the main focus of a follow up study .

we present a combined experimental and numerical study of the primary and secondary instabilities in a kolmogorov - like flow . the experiment uses electromagnetic forcing with an approximately sinusoidal spatial profile to drive a quasi - two - dimensional ( q2d ) shear flow in a thin layer of electrolyte suspended on a thin lubricating layer of a dielectric fluid . the experimental flow is modelled using a recently proposed 2d equation , derived from first principles by depth - averaging the full three - dimensional navier - stokes equations .

we present a combined experimental and numerical study of the primary and secondary instabilities in a kolmogorov - like flow . the experiment uses electromagnetic forcing with an approximately sinusoidal spatial profile to drive a quasi - two - dimensional ( q2d ) shear flow in a thin layer of electrolyte suspended on a thin lubricating layer of a dielectric fluid . the experimental flow is modelled using a recently proposed 2d equation , derived from first principles by depth - averaging the full three - dimensional navier - stokes equations . we also introduce a realistic numerical model approximating the electromagnetic body force in the experiment , which is validated against experimental measurements . we then explore the effect of lateral boundaries by performing direct numerical simulations ( dns ) subject to different types of boundary conditions . as the strength of the forcing is increased , this q2d flow undergoes a series of bifurcations . we compare the flow regimes observed in the experiment with those obtained from the dns and find that only the simulation subject to physical no - slip boundary conditions provides close , quantitative agreement . our analysis offers additional validation of the 2d model proposed by , as well as a demonstration of the importance of properly modelling the forcing and boundary conditions .

1104.0656

i

the rapid development of quantum information science has brought together several areas of theoretical and experimental physics @xcite . much effort has been concentrated in the search for solutions to sensitive problems that prevent the efficient realization of quantum information processing @xcite . we first mention the system - environment coupling which induces the decoherence of quantum states @xcite , apart from other barriers such as scalability @xcite and optimal control of individual systems @xcite . these challenges motivate both fundamental physical phenomena and outstanding technological issues such as individually addressing quantum systems , separated by only few @xmath3 m , with small errors @xcite . potential platforms for the implementation of quantum logic operations appeared in many fields such as condensed matter , quantum optics , and atomic physics @xcite . however , the problems mentioned above are faced by all the different communities when employing their particular techniques . in the particular case of the dissipation and decoherence phenomena in which we focus in the present work the redfield formalism redfield , slichter and the master equation @xcite have been the most applied approaches to address the environment effects on the proposed protocols for quantum information processing . whereas the semiclassical redfield formalism relies on a classical noise source , a quantum environment is assumed in the master equation approach . in this article , considering the general case of a time - dependent system , we discuss general similarities and differences between both approaches and show that they are equivalent , in the sense that they lead to the same phenomenological bloch equations bloch . consequently , both of them result in the same characteristic relaxation times @xmath1 and @xmath2 associated with the longitudinal and transverse relaxations , respectively @xcite . from this identification we show how these characteristic times are related to the operator - sum representation @xcite and the phenomenological - operator approach @xcite . the redfield formalism was intended to offer a microscopic description of the relaxation phenomenon , thus providing a deeper understanding of the parameters @xmath1 and @xmath2 . whereas the classical noise source employed in the redfield theory suffices to derive both relaxation times , two distinct quantum environments must be adopted to derive these time scales from the master equation formalism . on this regard , an amplitude and a phase damping environment are assumed to define the longitudinal and the transverse relaxation times , respectively . these quantum environments represent an energy - draining and a phase - shuffle channel by which the system loses excitations and phase relations . after presenting a detailed derivation of the redfield theory and comparing the derived characteristics times with those obtained from the master equation , we finally apply these equivalent formulations to the problem of state protection . we note that several distinct techniques have been proposed to control the effects of decoherence on quantum states , aiming to enlarge the fidelity of quantum information protocols . among others , we mention the quantum - error correction codes @xcite , environments engineering @xcite , decoherence - free subspaces @xcite , and dynamical decoupling @xcite . we finally mention that in a previous work @xcite , addressing the energy draining and decoherence of a harmonic oscillator , it was demonstrated that the inevitable action of the environment can be substantially weakened when considering appropriate non - stationary quantum systems . reasoning by analogy with the technique presented in ref . lucas , we show how to enlarge the longitudinal relaxation time associated with the amplitude - damping channel focusing a spin-@xmath0 system . the ideas presented here for decoherence control can be easily implemented in the nuclear magnetic resonance ( nmr ) context . this article is organized as follows : in sec . ii we present the derivation of the redfield equation for the general scenario of a time - dependent system . in sec . iii we apply the master equation approach to the same case . in sec . iv we show the equivalence between the redfield and the master equation formalisms by deriving the bloch equations from both approaches . focusing on a time - dependent spin-@xmath0 system , in secs . v and vi we present the operator - sum representation and the phenomenological operator approaches and their relation with the previous techniques . as an application of the theory , in sec . vii we address the state protection of a non - stationary spin-@xmath0 system . finally , sec . viii is dedicated to our final remarks where we discuss the generalization of the methods presented in this article for larger systems . in general , we will adopt the language of the nmr quantum information processing @xcite , although the theory presented here is valid for several other platforms , as quantum dots @xcite , superconducting artificial atoms @xcite , etc . throughout the article we will use natural units such that @xmath4 .

we present a derivation of the redfield formalism for treating the dissipative dynamics of a time - dependent quantum system coupled to a classical environment . we compare such a formalism with the master equation approach where the environments are treated quantum mechanically . focusing on a time - dependent spin- system these characteristic times are shown to be related to the operator - sum representation and the equivalent phenomenological - operator approach .

we present a derivation of the redfield formalism for treating the dissipative dynamics of a time - dependent quantum system coupled to a classical environment . we compare such a formalism with the master equation approach where the environments are treated quantum mechanically . focusing on a time - dependent spin- system we demonstrate the equivalence between both approaches by showing that they lead to the same bloch equations and , as a consequence , to the same characteristic times and ( associated with the longitudinal and transverse relaxations , respectively ) . these characteristic times are shown to be related to the operator - sum representation and the equivalent phenomenological - operator approach . finally , we present a protocol to circumvent the decoherence processes due to the loss of energy ( and thus , associated with ) . to this end , we simply associate the time - dependence of the quantum system to an easily achieved modulated frequency . a possible implementation of the protocol is also proposed in the context of nuclear magnetic resonance .

cond-mat9411034

c

in this paper we presented a detailed , quantitative study of the physical effects caused by the simultaneous presence of interactions and randomness in a system of lattice electrons . to this end we investigated the anderson - hubbard model with diagonal disorder at half filling in the limit of infinite spatial dimensions , i.e. within a dynamical mean - field theory , for three different disorder distributions . numerical results were obtained by employing quantum monte carlo techniques that provide an explicit finite - temperature solution of the model in @xmath9 . no further approximation were used . to construct the thermodynamic phase diagram we derived and evaluated the appropriate averaged two - particle correlation function , i.e. a dynamical response function , whose poles determine the magnetic instabilities of the disordered , interacting system . only this function and not the averaged staggered susceptibility @xmath144 itself , which is only a weighted sum of the response function over the ( matsubara ) frequencies fulfills a closed equation that determines the two - particle spectrum . the value of the response function at the lowest frequency is the most sensitive indicator for an instability of the system . in the temperature dependence of @xmath144 two distinct disorder regimes are observed : ( i ) for weak disorder the curie - weiss law holds , while ( ii ) at strong disorder @xmath144 acquires a maximum at a temperature below which a crossover from the temperature- to the disorder- dominated regime takes place . we demonstrated that at low temperatures and sufficiently strong interaction there always exists a phase with antiferromagnetic long - range order ( aflro ) . furthermore we discovered a new strong - coupling anomaly , namely that the nel - temperature @xmath194 is not always a monotonously decreasing function of disorder . indeed , at strong coupling and not too large disorder @xmath194 is always found to be an _ increasing _ function of disorder , i.e. disorder favors the formation of aflro in this regime . this implies the existence of an unexpected _ disorder - induced _ transition to a phase with aflro . under the assumption that @xmath194 is proportional to the effective exchange coupling between spins even in the disordered system we proved that for diagonal disorder the anomalous behavior is generic , i.e. is independent of the type of disorder distribution . it is a consequence of the fact that for diagonal disorder the difference between the local energies of neighboring sites becomes _ smaller _ on average , thus leading to stronger effective exchange coupling . we then studied the existence of metal - insulator transitions in the anderson - hubbard model . although in @xmath9 anderson localization does not take place the presence of disorder may well have other strong effects . in particular , binary - alloy disorder is able to cause band - splitting ( thereby resembling the effect of genuine interactions ) and hence may induce a metal - insulator transition all by itself . special attention was given to the question whether or not the disorder allows for the stabilization of an antiferromagnetic ( af ) _ metal_. to this end the average compressibility @xmath96 was evaluated both in the paramagnetic ( p ) and af phase . contrary to our expectation the presence of disorder was found to enhance the metallicity of the af - phase close to the p - af transition . this enhancement strongly suggests the existence of an af metal in the low - temperature phase of the anderson - hubbard mode . to investigate whether this af phase persists to be metallic down to @xmath320 , at least in @xmath9 , we studied the frequency components of the one - particle green function and of @xmath96 , respectively , down to @xmath290 , our lowest temperature . however , a reliable answer to this question can only be found at still lower temperatures , which at present are beyond the reach of the finite - temperature monte carlo techniques used here . in any case , the transition scenario involving metal - insulator and p - af transitions obtained for the anderson - hubbard model is remarkably rich . for alloy - type disorder with @xmath265 ( split - band limit for @xmath324 ) at @xmath320 an increase of the interaction @xmath57 from zero will probably first lead to a transition from a paramagnetic insulator to a paramagnetic metal at @xmath266 , then , at @xmath277 , to an antiferromagnetic metal and finally , at @xmath270 , to an antiferromagnetic insulator . no compelling evidence was found that @xmath277 and @xmath270 coincide . in the case of a continuous disorder distribution one has @xmath267 since band splitting never occurs . the above findings prove that the interplay between electronic interactions and scattering from disorder leads to interesting , and even novel , physical effects . in particular , the strong coupling anomaly discovered here call for an experimental verification . we are grateful to a. altland , p. van dongen , m. jarrell , h. mller - krumbhaar and especially e. mller - hartmann for very useful discussions and comments . this work was supported in part by the sonderforschungsbereich 341 of the deutsche forschungsgemeinschaft . in fact in their exact solution of the falicov - kimball model in @xmath9 brandt and mielsch [ z. phys . * b75 * , 365 , ( 1989 ) ; ibid . * b79 * , 295 ( 1990 ) ; ibid * b82 * , 37 ( 1991 ) ] were the first to show that the local self - energy may be used to map the lattice problem in @xmath9 onto its atomic counterpart with an additional `` kadanoff - baym field '' which has to be determined self - consistently ; for an application to other hubbard - like models see g. hlsenbeck and f. stephan ( preprint ) . x. y. zhang , m. j. rozenberg and g. kotliar , phys . lett . * 70 * , 1666 ( 1993 ) ; m. j. rozenberg , g. kotliar and x.y . zhang , phys . rev . * b49 * , 10181 ( 1994 ) ; q. si , m.j . rozenberg , g. kotliar , and a. e. ruckenstein , preprint . in principle the grand potential ( 6 ) may equally be viewed as a functional of the set of complex functions @xmath48 _ only _ in which case , however , one has to _ supplement _ ( 6 ) by the definition of @xmath47 as the local ( i.e. site - diagonal ) element of the full temperature green function ( see below ) . in the initial formulation of the theory it is advisable to treat _ both _ set of legendre - conjugate variables @xmath48 and @xmath47 as variational parameters , since only then does the expression ( 6 ) contain the full information about the thermodynamic equilibrium of the system @xcite . an equilibrium corresponds to a saddle point where the variations of @xmath51 w.r.t . @xmath48 and @xmath47 vanish , i.e. @xmath325 and @xmath326 . these variational conditions determine the physical values of @xmath47 and @xmath48 as the local green function and self - energy of the electrons , respectively . a theory generated in this way is guaranteed to be thermodynamically consistent and conserving in the spirit of the baym - kadanoff formalism [ g. baym , phys . rev . * 127 * , 836 ( 1962 ) ] . fig.1 : : local green function @xmath160 for @xmath324 and @xmath328 ( a ) , @xmath329 ( b ) obtained by usual fourier transformation ( dotted line ) and by redefinition acc . to eq.(13 ) ( solid line ) fig.2 : : dynamical antiferromagnetic response function @xmath185 vs. temperature @xmath330 . binary alloy with disorder strength @xmath331 and @xmath332 . here and in the following figures lines are usually guides for the eye , and error - bars are roughly of the size of the symbols unless shown explicitly . fig.3 : : inverse averaged antiferromagnetic susceptibility @xmath333 vs. @xmath330 for the binary alloy with @xmath334 and several values of @xmath57 . fig.4a : : @xmath202-phase diagram for the binary alloy with @xmath335 obtained from the zeroes of @xmath333 ( see fig . 4 ) . the af - phase is stable below the curves . the dotted lines at @xmath320 depict the regimes where the curie law would give negative transition temperatures . below the crosses @xmath333 has no zeroes but a minimum and an af - phase can no longer be expected . fig.4b : : @xmath202-phase diagram for the semi - elliptic distribution of the random energies with width @xmath336 . fig.5 : : @xmath330 vs. @xmath227 for the binary - alloy distribution at @xmath337 and 5.5 . dashed lines : quadratic increase of @xmath194 according to eq . ( 30b ) with @xmath338 ; dotted lines are guides to the eye only . fig.6 : : scaling plot of @xmath339 for the binary - alloy distribution according to eq . ( 30 ) ; symbols as in fig . 4a . fig.7 : : @xmath340-phase diagram for the percolation - type disorder for several values of @xmath57 . the dashed line depicts the hartree - fock ( hf ) result for the smallest value of @xmath341 . fig.8a : : averaged quadratic local moment @xmath239 vs. u at inverse temperature @xmath342 . without disorder ( full circle ) ; binary alloy with @xmath334 ( full square ) ; semi - elliptic distribution with @xmath334 ( open circle ) and @xmath343 ( open square ) . fig.8b : : averaged quadratic local moment @xmath344 ( normalized to the concentration of sites with random energy zero ) vs. @xmath99 for the percolation - type disorder at @xmath342 for several values of @xmath57 . fig.9 : : staggered magnetization @xmath249 and inverse averaged susceptibility @xmath333 vs. @xmath330 for the binary alloy with @xmath331 , @xmath332 , @xmath345 . dashed line : square - root fit of the last two points of @xmath249 below the transition ; dotted line : linear fit of @xmath333 . the arrow indicates the extrapolated nel temperature . fig.10 : : averaged compressibility @xmath96 vs. @xmath57 at @xmath342 . a ) binary alloy , @xmath335 ; b ) semi - elliptic distribution with @xmath346 and 4 . arrows indicate the transition to the antiferromagnetic state . fig.11 : : averaged compressibility @xmath96 vs. @xmath57 at @xmath342 in the paramagnetic phase ( dashed lines ) and in the antiferromagnetic phase ( dotted lines ) . a ) binary alloy , b ) semi - elliptic distribution . fig.12 : : averaged compressibility @xmath96 vs. @xmath57 for the binary alloy at @xmath342 and 40 . arrows indicate the transition to the antiferromagnetic state . fig.13 : : real part of @xmath347 vs. @xmath45 , a ) for @xmath348 , @xmath349 and @xmath350 in the paramagnetic phase ( p ) and the antiferromagnetic phase ( af ) ; b ) for the binary alloy with @xmath334 , @xmath351 in the af - phase at @xmath352 and 64 . fig.14 : : imaginary part of the one - particle green function @xmath161 at the lowest matsubara frequency @xmath353 vs. @xmath57 for a system without disorder ( @xmath279 ) and with binary - alloy disorder ( @xmath280 at @xmath342 ; p - phase ( dashed lines ) , af - phase ( dotted lines ) . fig.15 : : imaginary part of @xmath161 vs. @xmath45 in the p- and af - phase ; a ) no disorder , u=1.75 , @xmath350 ; b ) binary alloy with @xmath334 , @xmath351 , @xmath354 ( the two curves at @xmath355 in the middle are undistinguishable ) . fig.16 : : imaginary part of @xmath356 vs. @xmath330 for the binary alloy with @xmath334 for several values of @xmath57 .

we present a detailed , quantitative study of the competition between interaction- and disorder - induced effects in electronic systems . for this the anderson - hubbard model with diagonal disorder is investigated analytically and numerically in the limit of infinite spatial dimensions , i.e. within a dynamical mean - field theory , at half filling . numerical results are obtained for three different disorder distributions by employing quantum monte carlo techniques which provide an explicit finite - temperature solution of the model in this limit . the magnetic phase diagram is constructed from the zeros of the inverse averaged staggered susceptibility . we find that at low enough temperatures and sufficiently strong interaction there always exists a phase with antiferromagnetic long - range order . the existence of metal - insulator transitions is studied by evaluating the averaged compressibility both in the paramagnetic and antiferromagnetic phase . a rich transition scenario , involving metal - insulator and magnetic transitions , is found and its dependence on the choice of the disorder distribution is discussed .

we present a detailed , quantitative study of the competition between interaction- and disorder - induced effects in electronic systems . for this the anderson - hubbard model with diagonal disorder is investigated analytically and numerically in the limit of infinite spatial dimensions , i.e. within a dynamical mean - field theory , at half filling . numerical results are obtained for three different disorder distributions by employing quantum monte carlo techniques which provide an explicit finite - temperature solution of the model in this limit . the magnetic phase diagram is constructed from the zeros of the inverse averaged staggered susceptibility . we find that at low enough temperatures and sufficiently strong interaction there always exists a phase with antiferromagnetic long - range order . a novel strong coupling anomaly , i.e. an _ increase _ of the nel - temperature for increasing disorder , is discovered . an explicit explanation is given which shows that in the case of diagonal disorder this is a generic effect . the existence of metal - insulator transitions is studied by evaluating the averaged compressibility both in the paramagnetic and antiferromagnetic phase . a rich transition scenario , involving metal - insulator and magnetic transitions , is found and its dependence on the choice of the disorder distribution is discussed .

astro-ph0506121

r

as the dynamical evolution calculations begins , the disks experience an initial transient , which is essentially numerical in origin , and is due to the fact that the configuration is not in strict hydrostatic equilibrium ( recall that it is obtained from azimuthally averaging the results of 3d calculations ) . this transient , and its effects , are negligible , and it is essentially over in a hydrodynamical timescale ( @xmath106 ms ) . after that , the disk proceeds to evolve on a much longer timescale , determined by accretion onto the central black hole . we first present the details of the spatial structure of the disk due to fundamental physical effects , and then proceed to show the temporal evolution and associated transitions . the fundamental variable affecting the instantaneous spatial structure of the disk is the optical depth to neutrinos , @xmath107 , since it determines whether the fluid cools efficiently or not . all quantities show a radical change in behavior as the threshold @xmath108 is passed . figure [ compa0.01t100 ] shows color coded contours in a meridional slice of the thermodynamical variables , while figure [ compa0.01t100r ] displays the run of density and entropy per baryon along the equator , @xmath109 , where the density and temperature are highest ( we will refer here mainly to run a2 m , unless noted otherwise ) . the density increases as one approaches the black hole , varying as @xmath110 in the outer regions , where @xmath111 . the critical density for opaqueness is reached at @xmath112 cm , and for @xmath113 , @xmath114 . since the energy that would otherwise be lost via neutrino emission remains in the disk if the cooling is suppressed by scattering , the density does not rise as fast . in fact , one can note this change also by inspecting the scale height @xmath115 , which scales as @xmath116 for @xmath117 and @xmath118 const for @xmath113 ( and thus in terms of surface density the change is from @xmath119 to @xmath120 const at @xmath121 ) . the opaque region of the disk remains inflated to a certain extent , trapping the internal energy it holds and releasing it only on a diffusion timescale . at densities @xmath122 g @xmath10 the optical depth can reach @xmath123 , and thus the suppression of cooling is dramatic , as can be seen in figure [ compa0.01t100 ] , where the contours of @xmath107 and @xmath124 are shown . the flattening of the entropy profile is related to the change in composition in the optically thick region and the occurrence of convection ( see [ nuconvection ] below ) . as the density rises in the inner regions of the disk , the electron fraction @xmath28 initially drops as neutronization becomes more important , with the equilibrium composition being determined by the equality of electron and positron captures onto free neutrons and protons ( see equation ( [ yethin ] ) and the discussion preceding it ) . the lowest value is reached at @xmath125 , where @xmath126 . thereafter it rises again , reaching @xmath127 close to the horizon . thus flows that are optically thin everywhere will reach a higher degree of neutronization close to the black hole than those which experience a transition to the opaque regime . the numerical values for the electron fraction at the transition radius and at the inner boundary are largely insensitive to @xmath5 , as long as the transition does occur . the baryons in the disk are essentially in the form of free neutrons and protons , except at very large radii and low densities ( @xmath128 cm and @xmath129 g @xmath10 ) where @xmath5 particles form . figure [ rtye ] shows the region in the density temperature plane where the fluid lies , for runs a2 m and a1 m at two different times ( and color coded according to the electron fraction ) . most of the gas lies close to the line determining the formation of helium nuclei ( see equation [ nuc ] ) , but does not cross over to lower temperatures . the reason for this is that the energy that would be released by the creation of one helium nucleus ( 28.3 mev ) would not leave the disk ( recall that we consider neutrino emission as the only source of cooling ) and thus immediately lead to the photodisintegration of another @xmath5 particle . an equilibrium is thus maintained in which the gas is close to helium synthesis , but this does not occur . in the opaque regime , the gas moves substantially farther from the transition line to helium , since it can not cool efficiently and , as mentioned before , the density does not rise as quickly . figure [ rtye ] also shows clearly why one must make use of an equation of state which takes into account properly the effects of arbitrary degeneracy of the electrons and positrons . the solid straight line marks the degeneracy temperature as a function of density , given by @xmath130 mev . the disk straddles this line , with a degeneracy parameter @xmath131 in the inner regions , and @xmath132 at lower densities . thus making the approximation that the electrons in the flow are fully degenerate is not accurate . modeling the accretion flow in two dimensions @xmath101 , without the assumption of equatorial symmetry allows one to solve clearly for the vertical motions in the disk , something which is not possible when considering vertically integrated flows . it was pointed out by @xcite that a standard @xmath5 viscosity could lead to meridional flows in which @xmath133 changed sign as a function of height above the mid plane , leading to inflows as well as outflows . @xcite , @xcite and @xcite later considered a similar situation , and found that for a range of values in @xmath5 , the gas flowed inward along the surface of the disk , and outward in the equatorial regions . in previous work @xcite , we found this solution in disk flow , with large scale circulations exhibiting inflows and outflows for high viscosities ( @xmath134 ) , and small scale eddies at lower values ( @xmath135 ) . how vertical motions affect the stability of accretion disks and may lead to the transport of angular momentum is a question that has become of relevance in this context @xcite . here we show in figure [ velcont ] the velocity field and magnitude of meridional velocity for run a2 m . the small scale eddies are clearly visible , with their strength usually diminishing as the disk is drained of matter and the density drops . the effect of a different value of @xmath5 can be seen in the top row of figure [ compa0.01a00.1t100t200 ] , where the comparison between @xmath136 and @xmath137 is made . the large , coherent lines of flow aimed directly at the origin for @xmath136 correspond to inflow , while the lighter shades along the equator at larger radii show outflowing gas . the instantaneous structure of the disk concerning the density , temperature and composition profiles is largely independent of the viscosity , as long as there is enough mass to produce the optically thin / optically thick transition . even in the optically thin regime , the structure can not be determined analytically following the standard arguments applied to thin , cool disks of the shakura sunyaev type . the reason for this is the following : a central assumption in the standard solution is that the disk is cool , i.e. , @xmath138 . this comes from the requirement that all the energy dissipated by viscosity , @xmath139 be radiated away efficiently , and produce the observed flux , @xmath140 . in our case , the disk is most certainly not cool , as it originates from the tidal disruption of a neutron star by a black hole ( or possibly the coalescence of two neutron stars , and subsequent collapse of the central mass to a black hole ) . the gas that constitutes the disk is dynamically hot because it was in hydrostatic equilibrium in a self gravitating configuration , where @xmath141 and has not been able to release this internal energy ( the merger process itself may lead to additional heating ) . a further deviation from the standard solution is that the pressure support in the disk leads to a rotation curve that is slightly sub keplerian . the dissipation by viscosity is in fact smaller than the cooling rate over much of the disk , and the released luminosity comes from a combination of viscous dissipation and the store of internal energy given to the disk at its conception . the classical requirement for convective instability in the presence of entropy and composition gradients , as well as rotation , is the solberg hoiland criterion @xcite : @xmath142 where @xmath143 is the radial epicyclic frequency ( @xmath144 being the angular velocity ) , and @xmath145 \label{bv}\ ] ] is the brunt visla frequency ( see * ? ? ? * ; * ? ? ? the adiabatic index @xmath146 in our case , since @xmath147 is the most important contribution to the total pressure . thus a region may be convectively unstable because of a composition gradient , or an entropy gradient , or both . strictly speaking , here one should consider the total lepton fraction @xmath148 . we do not consider the transport of neutrinos in detail , as already mentioned ( see [ nuopacities ] ) , and do not calculate explicitly the contribution of neutrinos to this @xmath148 . for what follows , we will thus simply take @xmath28 to represent the behavior of the full @xmath148 . the origin of convection in the lepton inversion zone can be understood as follows @xcite . consider a fluid element in the lepton inversion zone that is displaced in the outward direction and then comes to pressure equilibrium with its surroundings . the displaced element , which is lepton rich relative to its new surroundings , attains the ambient pressure at a lower density than the surrounding fluid , because the pressure depends directly on the lepton number , and thus tends to drift outwards . by the same token , an inwardly displaced fluid element in the lepton inversion zone is depleted in leptons relative to its new surroundings and thus tends to sink . when temperature gradients are allowed for , the displaced fluids , in general , have temperatures which differ from those of their surroundings . these variations tend to promote stability or instability depending on whether the existing temperature gradient is less than or greater than the adiabatic gradient , respectively . in the present scenario then , the transition to the optically thick regime leads to instability , because initially @xmath149 , and also @xmath150 ( see also figure [ compa0.01t100r ] ) . the entropy profile is then flattened by efficient convective mixing , and the lepton gradient inversion is due to the different way in which the composition is determined through weak interactions once the neutrinos become trapped . in a dynamical situation such as the one treated here , convection tends to erase the gradients which give rise to it . accordingly , the sum of the corresponding terms in equation ( [ sh ] ) tends to zero once the simulation has progressed and convection has become established in the inner disk ( note that in our case the rotation curve there is sub keplerian , with @xmath151 , so that @xmath144 and @xmath152 are not equal ) . this is analogous to what occurs in proto neutron stars following core collapse @xcite and has actually been confirmed in numerical simulations of such systems @xcite , where the convection leads to strong mixing . for a neutrino driven convective luminosity @xmath153 at radius @xmath154 and density @xmath155 , the convective velocity may be estimated by standard mixing length theory as ( see e.g. , * ? ? ? * ) @xmath156 where we have used the fact that gas pressure dominates in the fluid . the assumption of spherical symmetry implicit in this expression is not strictly met in our case , but it may provide us nevertheless with a useful guide . the overturn time of a convective cell is then given by @xmath157 ms , where @xmath158 km is the mixing length , usually set to a pressure scale height . in our calculations , we see that the magnitude of the meridional velocity @xmath159 decreases as @xmath4 decreases , then rises again as the opaque region is reached , reaching @xmath160 cm s@xmath8 , at @xmath161 km . the associated turnover times are thus @xmath162 ms , in good agreement with the estimate made above . to isolate this effect from that due to viscosity ( which generates the meridional circulations mentioned previously ) we have performed a simulation in which @xmath163 ( run aim in table [ table : ics ] ) . this calculation shows essentially no accretion onto the black hole except for a small amount of gas transferred at early times ( because the initial condition is not in strict equilibrium ) . the result concerning the profile of @xmath28 as a function of radius is as we have described above , i.e. , increasing neutronization as @xmath4 decreases , until the opaque region is reached , followed by an increase in @xmath28 as the black hole boundary is approached . the revealing difference lies in the time evolution of this profile . as mentioned above , convection generates motions which tend to eliminate the composition gradient which drives it . in the presence of viscosity , matter is continuously transported radially , and the gradient is not erased entirely . height integrated profiles of @xmath164 at 50 ms and 200 ms show a similar behavior at small radii . when viscosity is removed , the initial composition gradient is gradually softened until it disappears in the innermost regions ( see figure [ convection ] ) . the disk then has a nearly constant electron fraction along the equator for @xmath113 , with a sharp transition region leading to an increase with radius in the transparent regime . this may resemble the convection dominated accretion solution found analytically by @xcite and numerically by several groups @xcite , where convection transports angular momentum inward , energy outward , and gives a radial profile in density @xmath165 . we find a different power law , with @xmath114 . there are several factors which may account for this difference : the disk vs. spherical geometry ; the initial condition with a large amount of internal energy ; and the limited , but present cooling rate on the boundaries of the flow . since the neutrinos are diffusing out of the fluid ( the mean free path is small compared to the size of the system in the optically thick regime ) , one would expect a corresponding viscosity through @xmath166 ( see , e.g. * ? ? ? this needs to be compared to the corresponding viscosity generated by the assumed @xmath5 prescription for consistency . we may assume @xmath167 , where @xmath168 is the energy density associated with neutrinos , and @xmath169 is their mean free path ( see [ nuopacities ] ) . thus @xmath170 for conditions near the equatorial plane , @xmath109 . this value will increase in the outer regions , since the mean free path becomes larger as the density drops . at the neutrino surface , where @xmath171 , the mean free path is @xmath172 cm , and so an averaged value of the viscosity over the entire neutrino opaque region will result in an increase of about one order of magnitude over the estimate given in equation ( [ nuvisc ] ) ( note however , that strictly speaking the diffusion approximation is no longer valid in the outer regions , so this must be interpreted with care ) . for the @xmath5 prescription , @xmath173 . the rotation curve is not too far from keplerian , and scaling this expression to typical values we find @xmath174 the effects on angular momentum transport are thus a full two orders of magnitude below those arising from our viscosity prescription ( for @xmath175 ) . to put it another way , a lower limit for the viscosity under these conditions ( and in the optically thick portion of the disk ) would be @xmath176 . an alternative analysis would be to consider the corresponding timescales induced by this viscosity . since @xmath177 , smaller viscosities imply longer timescales , as expected . aside from convection , several general criteria for stability can be analyzed in our case . evidently , since we are performing dynamical calculations , any instability that arises will quickly lead to a change in structure . it is nevertheless instructive to consider the corresponding conditions within the disk . we first consider the toomre criterion , comparing gravitational and internal energies , with @xmath178 where @xmath152 is the local epicyclic frequency ( essentially equal in this case to the angular frequency @xmath144 ) , @xmath179 is the local sound speed and @xmath180 is the surface density . we find @xmath181 in all cases throughout the calculations , and thus that the disks are stable in this respect ( figure [ qthermalstability]a shows a typical profile of @xmath182 $ ] ) . previous studies @xcite have shown that neutrino cooled disks are thermally unstable if radiation pressure dominates in the flow , and stable otherwise , although the effects of neutrino opacities were not considered . the criterion for stability can be written in this case as @xmath183 where @xmath184 and @xmath185 are the volume heating and cooling rates . essentially , in order to be stable the disk must be able to get rid of any excess internal energy generated by an increase in the heating rate . we show in figure [ qthermalstability]b typical values for @xmath184 and @xmath185 as functions of the central ( equatorial ) temperature , for run a2 m at @xmath186 ms . in the optically thin region the disks are thermally stable , because the pressure is dominated by the contribution from free nucleons and cooling is efficient . in the optically thick region the cooling is greatly suppressed and the criterion would indicate that the disk becomes thermally unstable . this is reasonable , since with optical depths @xmath187 , essentially no direct cooling takes place , and dissipation is not suppressed . this is why the disk is geometrically thick in the inner regions , as already described above ( [ structure ] ) . the balance that gives a quasi steady state structure is achieved mainly through the diffusion of neutrinos , since the diffusion timescale is shorter than the accretion timescale . the evolution of the disk on long timescales is determined by the balance between two competing effects : on one hand , viscosity transports angular momentum outwards , matter accretes and the disk drains into the black hole on an accretion timescale @xmath188 . the trend in this respect is towards lower densities and temperatures . on the other hand , cooling reduces pressure support and leads to vertical compression , increasing the density . the internal energy of the fluid is released on a cooling timescale , @xmath189 . the presence of an optically thick region in the center of the disk limits the luminosity ( dominated by electron and positron capture onto free nucleons ) to @xmath190erg s@xmath8 , and the initial internal energy is @xmath191 erg , so @xmath192 s. if @xmath193 , the disk will cool before its mass or internal energy reservoir is significantly affected by accretion onto the black hole . the maximum density ( shown in figure [ rhomax ] for runs a1 m , a2 m and a3 m ) actually increases slightly due to vertical compression , and subsequently drops once mass loss through accretion dominates ( the @xmath194 ms delay at the start , during which it is approximately constant , is simply the sound crossing time across the optically thick region of the disk ) . the accretion rate onto the black hole , @xmath195 , the accretion timescale @xmath196 and the total neutrino luminosity @xmath197 are shown in figures [ mdot ] , [ tacc ] and [ lum ] . they all show the same qualitative behavior , remaining fairly constant ( or changing slowly ) for an accretion timescale , and abruptly switching thereafter . the accretion timescales are approximately 0.5 s and 5 s for @xmath198 respectively . for high viscosity , @xmath136 , the transport of angular momentum is vigorous , and the black hole quickly accretes a substantial amount of mass ( @xmath199 within the first 100 ms ) . the accretion timescale is @xmath200 ms . the circulation pattern consists of large scale eddies , with @xmath6 . in fact there is essentially one large eddy on each side of the equatorial plane , with mass inflow along the surface of the disk , and an equatorial outflow . part of the outflowing gas moves away from the equator and reverses direction close to the surface of the disk , contributing to the inflow . for an intermediate viscosity , @xmath137 , the intensity of the circulations is smaller , but also , the eddies are smaller , with several of them clearly occurring in the disk at once . the transport of angular momentum being less vigorous , the accretion rate is substantially smaller than in the previous case . finally , for a yet lower viscosity , @xmath201 , the trend continues , and the eddies become smaller still . in these last two cases , angular momentum transport is so low that very high densities are maintained in the central regions of the disk for a large number of dynamical times , contrary to what is seen in the high viscosity case . a simple way to quantify how much of the mass flow is actually making it to the central black hole is to measure the fraction of mass at any given radius that is moving inwards , @xmath202 , where we have divided the mass flow rate into two height integrated parts , one with @xmath203 and another with @xmath204 . for example , for run a2 m at @xmath205 ms and @xmath206 it is approximately 1/3 . the ratio tends to unity only in the innermost regions of the disk , and shows the true black hole mass accretion rate , plotted in figure [ mdot ] for the same cases . we now turn our attention to the neutrinos , which are the main source of cooling . the results in this case are markedly different from what we initially found @xcite , simply because of the new equation of state , the more realistic cooling rates and the approximate computation of opacities . they are in general agreement with the preliminary results we presented before @xcite , which used a less detailed equation of state than the one shown here . the optically thick region is present in every disk at the start of the calculation . as already mentioned , this has two important effects : enhancing the pressure and suppressing the neutrino emission . this reflects upon the total luminosity , since the suppression occurs precisely in the hottest regions , where most of the energy would otherwise be released . the most important qualitative difference between the runs presented here is that for a high viscosity ( @xmath136 , runs a1 m and a1 m ) , the disk is drained of mass so fast that it has no chance to cool ( @xmath207 ) and release most of its internal energy . it is in fact advected into the black hole . moreover , the optically thick region disappears entirely from the disk ( see figure [ compa0.01a00.1t100t200 ] ) by @xmath208 ms . the emission is then no longer suppressed ( see the contours of @xmath124 in the same figure ) , and the disk radiates at the maximum possible cooling rate . the thermal energy content of the disk is so large , that as it thins , the luminosity actually increases briefly around @xmath209 ms before dropping again . the drop at late times follows an approximate power law , with @xmath210 . in this interval the energy release comes from a combination of residual internal energy and viscous dissipation within the disk . for lower values of the viscosity , @xmath211 , the central regions of the disk remain optically thick throughout the calculations . note that reducing the disk mass by a factor of five ( as was done for runs a1 m , a2 m , a3 m ) does not affect these overall conclusions . the fluid is simply compressed into a smaller volume , and thus the densities and temperatures that are reached are similar than for the high mass runs . for run a3 m the neutrino luminosity is practically constant at @xmath212 erg s@xmath8 for @xmath213 ms . table [ table : evol ] summarizes the typical disk mass , energy density , accretion rate , luminosity , and duration and energetics of neutrino emission for all runs .

this transition produces an inversion of the lepton gradient in the innermost regions of the flow which drives convective motions , and affects the density and disk scale height radial profiles . the electron fraction remains low in the region close to the black hole , and if preserved in an outflow , could give rise to heavy element nucleosynthesis . our specific initial conditions arise from the binary merger context , and so we explore the implications of our results for the production of gamma ray bursts .

we present a detailed , two dimensional numerical study of the microphysical conditions and dynamical evolution of accretion disks around black holes when neutrino emission is the main source of cooling . such structures are likely to form after the gravitational collapse of massive rotating stellar cores , or the coalescence of two compact objects in a binary ( e.g. , the hulse taylor system ) . the physical composition is determined self consistently by considering two regimes : neutrino opaque and neutrino transparent , with a detailed equation of state which takes into account neutronization , nuclear statistical equilibrium of a gas of free nucleons and alpha particles , blackbody radiation and a relativistic fermi gas of arbitrary degeneracy . various neutrino emission processes are considered , with capture onto free nucleons providing the dominant contribution to the cooling rate . we find that important temporal and spatial scales , related to the optically thin / optically thick transition are present in the disk , and manifest themselves clearly in the energy output in neutrinos . this transition produces an inversion of the lepton gradient in the innermost regions of the flow which drives convective motions , and affects the density and disk scale height radial profiles . the electron fraction remains low in the region close to the black hole , and if preserved in an outflow , could give rise to heavy element nucleosynthesis . our specific initial conditions arise from the binary merger context , and so we explore the implications of our results for the production of gamma ray bursts .

astro-ph0506121

i

we have performed two - dimensional hydrodynamical simulations of accretion disks in the regime of hypercritical accretion , where neutrino emission is the main cooling agent . the disks are assumed to be present around a stellar mass black hole , and are evolved for a few hundred dynamical timescales . we have paid particular attention to the relevant microphysical processes under the conditions at hand , and used a detailed equation of state which includes an ideal gas of @xmath5 particles and free baryons in nuclear statistical equilibrium and a relativistic fermi gas of arbitrary degeneracy . the composition of the fluid is determined by weak interactions . the density and temperature are such that the inner regions of the disk become opaque to neutrinos , and this is taken into account in a simple approximation . * once the fluid becomes photodisintegrated into free nucleons , neutronization becomes important and lowers the electron fraction substantially below 1/2 , with the electron fraction @xmath28 reaching @xmath214 at its minimum . this value , however , does not occur in the immediate vicinity of the black hole , but rather at the transition radius where the fluid becomes optically thick to its own neutrino emission , @xmath215 cm . at smaller radii , the electron fraction rises again , reaching @xmath216 close to the horizon . * neutrino trapping produces a change in composition and an inversion in the electron fraction . the associated negative gradient in @xmath164 induces convective motions in the optically thick region of the disk . this is analogous to what presumably occurs following core collapse , in a proto neutron star and its surrounding envelope . due to the radial flows induced by viscosity , convection is unable to suppress this composition gradient . it would appear , however , that the entropy per baryon in the optically thick region is very close to being constant , with @xmath217 . * the spatial structure of the disk is characterized by the transition radius @xmath121 where the material becomes optically thick . for @xmath117 the disk cools efficiently , whereas for @xmath113 the emission is suppressed and the fluid is unable to cool directly ( although it does so on a neutrino diffusion timescale ) . this leads to larger pressures and a more moderate rise in density . * the temporal evolution of the disk is determined by the balance between accretion and neutrino emission . for low viscosities ( @xmath218 ) , the disk is able to cool in a quasi steady state and radiate its internal energy reservoir . this lasts for approximately 0.1 - 0.4 s , with @xmath219 erg s@xmath8 . thereafter the typical luminosity and density quickly decay . for large viscosities ( @xmath220 ) the disk is drained of mass on an accretion timescale which is shorter than the cooling timescale , and the internal energy of the disk is essentially advected into the black hole . an interesting result in this case is that as the disk becomes transparent before being engulfed by the hole , it undergoes a re - brightening , as some of the stored energy escapes . * the total energy output in neutrinos is @xmath221 erg , over a timescale of @xmath222 ms . the typical accretion rates are @xmath223 s@xmath8 , and neutrino energies are @xmath224 mev at @xmath121 . energy densities in the inner regions of the disk are @xmath225erg @xmath10 . there are two main ingredients in the results presented here that contrast with those available previously in the literature concerning the steady state structure of neutrino cooled accretion disks . the first is the ability to dynamically model the evolution of the system for hundreds of dynamical timescales , taking the previous history of the fluid into consideration through the choice of initial conditions . this allows us to consider the energetics on more relevant timescales ( cooling , viscous ) than the dynamical one accessible in three dimensional studies . the second is the realization that the structure of the disks is affected qualitatively by the presence of an optically thick region at high densities . in this sense the situation is similar to that encountered following massive core collapse , where convection occurs ( recent work assessing the relative importance of the mri , neutrino and convection driven viscosity in rotating collapsing cores has been reported by * ? ? ? * ; * ? ? ? * we comment further on this issue in [ magnetic ] ) . clearly there is room for improvement in the results presented here . to begin with , our expressions for the effect of a finite optical depth on cooling and pressure are too simple , and attempt only to capture the essential physical behavior of the system . there is an important region of the disk where the optical depth is not large enough to consider the diffusion approximation , and where more detailed transport effects ought to be considered . we have not separated the neutrino variables into three species , which would be more rigorous ( e.g. , * ? ? ? * have noted that in such flows , the electron neutrinos @xmath226 might be preferentially absorbed with respect to electron anti neutrinos @xmath227 , thus affecting the neutrino annihilation luminosity above the surface of the disk in an important way ) . we have only considered the effects of coherent scattering off nucleons and @xmath5 particles for opacity purposes . this specifically ignores absorptive scattering , which would : ( i ) lead to a modification of the neutrino emergent spectrum ; ( ii ) produce heating of the fluid . the latter may be particularly important concerning the driving of powerful winds off the surface of the disk , and needs to be addressed more carefully in the context of grbs . as in previous work , we have chosen to maintain the newtonian expression for the potential well of the central black hole for ease of comparison to previous results ( this will be addressed in future work ) . the sudden release of gravitational binding energy of a neutron star is easily sufficient to power a grb . the minimum energy requirement is @xmath228 erg if the burst is beamed into a solid angle @xmath229 . as discussed in [ intro ] , a familiar possibility , especially for bursts belonging to the short duration category , is the merger of a neutron star - black hole , or double neutron star binary by emission of gravity waves , which , as illustrated here , is likely to generate a black hole surrounded by a lower mass accretion disk . how is the available rotational and gravitational energy converted into an outflowing relativistic plasma ? a straightforward way is that some of the energy released as thermal neutrinos is reconverted , via collisions outside the disk , into electron - positron pairs or photons . the neutrino luminosity emitted when disk material accretes via viscous ( or magnetic ) torques on a timescale @xmath230 s is roughly @xmath231 for a canonical radiation efficiency of 0.1 . during this time , the rate of mass supply to the central black hole is of course much greater than the eddington rate . although the gas photon opacities are large , the disk becomes sufficiently dense and hot to cool via neutrino emission . there is in principle no difficulty in dissipating the disk internal energy , but the problem is in allowing these neutrinos to escape from the inflowing gas . at sufficiently low accretion rates , @xmath232 , we find that the energy released by viscous dissipation is almost completely radiated away on a timescale given by @xmath233 s. in contrast , for a higher mass supply , @xmath234 , energy advection remains important until the entire disk becomes optically thin . the restriction on the cooling rate imposed by high optical depths is key because it allows the energy loss to be spread over an extended period of time during which the neutrino luminosity stays roughly constant . this gives a characteristic timescale for energy extraction and may be essential for determining the duration of neutrino - driven short grbs @xcite . neutrinos could give rise to a relativistic pair - dominated wind if they converted into pairs in a region of low baryon density ( e.g. along the rotation axis , away from the equatorial plane of the disk ) . the @xmath235 process can tap the thermal energy of the torus produced by viscous dissipation . for this mechanism to be efficient , the neutrinos must escape before being advected into the hole ; on the other hand , the efficiency of conversion into pairs ( which scales with the square of the neutrino density ) is low if the neutrino production is too gradual . typical estimates suggest a lower bound of @xmath236 ( e.g * ? ? ? if the pair - dominated plasma were collimated into a solid angle @xmath229 then of course the apparent isotropized energy would be larger by a factor @xmath237 , but unless @xmath229 is @xmath238 this may fail to satisfy the apparent isotropized energy of @xmath13 ergs implied by a redshift @xmath239 for short grbs . one attractive mechanism for extracting energy that could circumvent the above efficiency problem is a relativistic magneto hydrodynamic ( mhd ) wind @xcite . such a wind carries both bulk kinetic energy and ordered poynting flux , and it possible that gamma - ray production occurs mainly at large distances from the source @xcite . a rapidly rotating neutron star ( or disk ) releases energy via magnetic torques at the rate @xmath240 , where @xmath241 s is the spin period , and @xmath242 g is the strength of the poloidal field at a radius @xmath243 cm . the last stable orbit for a schwarzschild hole lies at a coordinate distance @xmath244 km , to be compared with @xmath245 km for an extremal kerr hole . thus the massive neutron disk surrounding a schwarzschild black hole of approximately @xmath246 should emit a spin - down luminosity comparable to that emitted by a millisecond neutron star . a similar mhd outflow would result if angular momentum were extracted from a central kerr hole via electromagnetic torques @xcite . the field required to produce @xmath247 is colossal , and may be provided by a helical dynamo operating in hot , convective nuclear matter with a millisecond period @xcite . a dipole field of the order of @xmath248 g appears weak compared to the strongest field that can in principle be generated by differential rotation ( @xmath249^{-1}\;{\rm g}$ ] ) , or by convection ( @xmath250 ) , although how this may come about in detail is not resolved . we examine in more detail the possible generation of strong magnetic fields below in [ magnetic ] . computer simulations of compact object mergers and black hole formation can address the fate of the bulk of the matter , but there are some key questions that they can not yet tackle . in particular , high resolution of the outer layers is required because even a tiny mass fraction of baryons loading down the outflow severely limits the attainable lorentz factor - for instance a poynting flux of @xmath251 erg could not accelerate an outflow to @xmath252 if it had to drag more than @xmath253 solar masses of baryons with it . one further effect renders the computational task of simulating jet formation even more challenging . this stems from the likelihood that the high neutrino fluxes ablate baryonic material from the surface of the disk at a rate @xcite @xmath254 a rest mass flux @xmath255 limits the bulk lorentz factor of the wind to @xmath256 if one assumes that the external poloidal field strength is limited by the vigor of the convective motions , then the spin - down luminosity scales with neutrino flux as @xmath257 , where @xmath258 is the convective velocity . the ablation rate given in equation ( [ ablation ] ) then indicates that the limiting bulk lorentz factor @xmath259 of the wind decreases as @xmath260 . thus the burst luminosity emitted by a magnetized neutrino cooled disk may be self - limiting . the mass loss would , however , be suppressed if the relativistic wind were collimated into a jet . this suggests that centrifugally driven mass loss will be heaviest in the outer parts of the disk , and that a detectable burst may be emitted only within a certain solid angle centered on the rotation axis ( see e.g. , * ? ? ? there is also the question of magnetic fields , which we have not included , but should obviously be considered . the field in a standard disk is probably responsible for viscous stresses and dissipation , through the mri . in this respect the current scenario should exhibit these characteristics . the mri operates on an orbital timescale , and so the field would grow in a few tens of milliseconds . it may also be amplified by the convective motions described above , [ nuconvection ] . the saturation value for the field can be naively estimated as that at which its energy density is in equipartition with the gas , @xmath261 , or when the alfven speed , @xmath262 is comparable with the azimuthal velocity , @xmath263 . this gives @xmath264 g. it is not clear at all , however , that the field amplitude will reach such high levels , because the magnetic reynolds number is far beyond its critical value ( where diffusion balances dynamo driven growth ) , and amplification can lead to field expulsion from the convective region , thereby destroying the dynamo @xcite . precious little is known about the growth of magnetic fields at such overcritical levels @xcite , and a definitive answer will require the self consistent inclusion of full mhd into the problem at hand ( but with a level of resolution which may be well above present computational capabilities ) . it is not clear either that the magnetic shearing instability can generate a mean poloidal field as strong as @xmath265 g , since to first order it does not amplify the total magnetic flux . the non - linear evolution of the instability depends sensitively on details of magnetic reconnection , and it has indeed been suggested that this can smooth reversals in the field on very small scales , pushing the dominant growing mode to much larger scales @xcite . it is certainly possible , as shown here , that compact binary mergers do form a neutron disk that is hot enough to be optically thick to neutrinos , and convective instability is a direct consequence of the hot nuclear equation of state . a neutron disk is likely to be convective if the accretion luminosity exceeds @xmath266 . note that even if the accretion luminosity is lower , a hot , massive disk ( such as those forming in collapsars , * ? ? ? * ) would undergo a brief period of convection as a result of secular cooling ( notice that convection is driven by secular neutrino cooling , whereas the mri is powered by a release of shear kinetic energy ) . if the dense matter rotates roughly at the local keplerian angular velocity , @xmath267 , then @xmath268 is approximately independent of radius , and the required poloidal field for a given luminosity is @xmath269 . if a period of convection is a necessary step in the formation of a strong , large - scale poloidal field , an acceptable model thus requires that the surrounding torus should not completely drain into the hole on too short of a timescale . whether a torus of given mass survives clearly depends on its thickness and stratification , which in turn depends on internal viscous dissipation and neutrino cooling . a large amount of differential rotation ( as may occur in newborn neutron stars or those in x ray binaries , and is definitely the case in toroidal structures supported mainly by centrifugal forces ) , combined with short periods , may produce substantial magnetic field amplification @xcite . the energy transferred to the magnetic field is released in episodic outbursts when the buoyancy force allows the field to rise to the surface of the star or disk . the amplification of a magnetic field to such strong values would clearly have important consequences on the evolution and time variability of the disk and its energy output . it would probably lead to strong flaring and reconnection events accompanied by the release of large amounts of energy , if the growth time for field amplification , @xmath270 , is shorter than the accretion timescale , @xmath188 ( otherwise the disk would drain of matter before the field had the chance to reach large values ; in this case , the survival of a massive , rapidly rotating neutron star as the end - point of binary ns merger might be preferred over the prompt formation of a bh ) . an effective helical dynamo of the @xmath271 type should be favored by a low effective viscosity , @xmath5 , because , as stated in [ nuconvection ] , the overturn time ms . ] is @xmath272ms ( this applies only if the disk is not fed with matter externally for a time longer than @xmath188 , otherwise convection would also be able to amplify the magnetic field ) . core collapse sne and compact object mergers are natural astrophysical sites for the production of heavy elements @xcite . in particular , nucleosynthesis in neutrino driven winds is an issue that may be relevant for iron group elements , as well as for heavier nuclei through the r process @xcite . initial investigations into this matter @xcite determined that the entropy in the outflow arising from a newborn neutron star was probably too low to give rise to the r process efficiently . more recently , this problem has been addressed again in the specific case of collapsar or post merger accretion disks , based on the results of analytical calculations of neutrino cooled disks in one dimension @xcite . if the wind consists of a uniform outflow driven from the surface of the disk , the entropy is too low , and essentially only iron group elements are synthesized , in agreement with the earlier results . however , their results also indicate that in a bubble - type outflow ( against a background steady wind ) , where episodic expulsion of material from the inner regions takes place , material may be ejected from the disk and preserve its low electron fraction , thus allowing the r - process to occur . the convective motions reported here , occurring in the optically thick portion of the disk , would represent one way such cells could be transported to the disk surface if they can move fast enough to preserve the neutron excess . since the existence of a convection region is dependent upon the densities reached in the inner disk ( so that it becomes opaque ) , the synthesized nuclei ( iron group vs. heavier , r process elements ) could be a reflection of its absence / presence . we have benefited from many useful discussions and correspondence with u. geppert , n. itoh , k. kohri , p. kumar , a. macfadyen , p. mszros , m. prakash , m. rees , t. thompson , s. woosley and a. socrates . financial support for this work was provided in part by conacyt 36632e ( whl , dp ) and by nasa through a chandra postdoctoral fellowship award pf3 - 40028 ( err ) . part of this work was done during visits to the institute for advanced study ( whl ) and instituto de astronoma , unam ( err ) , whose hospitality is gratefully acknowledged . we thank the anonymous referee for helpful comments and suggestions on the initial manuscript . abramowicz , m.a . , in growing black holes : accretion in a cosmological context `` , in press : eds . a. merloni , s. nayakshin and r. sunyaev , ' ' eso astrophysics symposia series " , springer - verlag , berlin , 2004 ( astro - ph/0411185 ) clccc aim & 0.0 & 3.85 & 0.308 & 19,772 + a1 m & 0.1 & 3.85 & 0.308 & 19,772 + a2 m & 0.01 & 3.85 & 0.308 & 19,772 + a3 m & 0.001 & 3.85 & 0.308 & 19,772 + a1 m & 0.1 & 3.85 & 0.062 & 19,772 + a2 m & 0.01 & 3.85 & 0.062 & 19,772 + a3 m & 0.001 & 3.85 & 0.062 & 19,772 + cccccccc aim & 0 & 6@xmath273 & 2@xmath274 & 1.5@xmath274 & 8 & @xmath275140 & 0.308 + a1 m & 7 & 2@xmath276 & 2@xmath277 & 2@xmath274 & 8 & 60 & 0.098 + a2 m & 0.7 & 2@xmath276 & 6@xmath274 & 2@xmath274 & 8 & 150 & 0.2 + a3 m & 0.05 & 1.6@xmath276 & 2.@xmath274 & 1.55@xmath274 & 8 & @xmath275140 & 0.287 + a1 m & 0.7 & 4@xmath273 & 2@xmath274 & 5.4@xmath278 & 8 & 25 & 0.022 + a2 m & 0.05 & 3@xmath273 & 1@xmath274 & 5.5@xmath278 & 8 & 75 & 0.045 + a3 m & 0.005 & 3@xmath273 & 1.5@xmath274 & 7@xmath278 & 8 & @xmath275160 & 0.059 +

we present a detailed , two dimensional numerical study of the microphysical conditions and dynamical evolution of accretion disks around black holes when neutrino emission is the main source of cooling . the physical composition is determined self consistently by considering two regimes : neutrino opaque and neutrino transparent , with a detailed equation of state which takes into account neutronization , nuclear statistical equilibrium of a gas of free nucleons and alpha particles , blackbody radiation and a relativistic fermi gas of arbitrary degeneracy .

we present a detailed , two dimensional numerical study of the microphysical conditions and dynamical evolution of accretion disks around black holes when neutrino emission is the main source of cooling . such structures are likely to form after the gravitational collapse of massive rotating stellar cores , or the coalescence of two compact objects in a binary ( e.g. , the hulse taylor system ) . the physical composition is determined self consistently by considering two regimes : neutrino opaque and neutrino transparent , with a detailed equation of state which takes into account neutronization , nuclear statistical equilibrium of a gas of free nucleons and alpha particles , blackbody radiation and a relativistic fermi gas of arbitrary degeneracy . various neutrino emission processes are considered , with capture onto free nucleons providing the dominant contribution to the cooling rate . we find that important temporal and spatial scales , related to the optically thin / optically thick transition are present in the disk , and manifest themselves clearly in the energy output in neutrinos . this transition produces an inversion of the lepton gradient in the innermost regions of the flow which drives convective motions , and affects the density and disk scale height radial profiles . the electron fraction remains low in the region close to the black hole , and if preserved in an outflow , could give rise to heavy element nucleosynthesis . our specific initial conditions arise from the binary merger context , and so we explore the implications of our results for the production of gamma ray bursts .

1011.0798

c

our photometric and spectroscopic analyses of the components of 2mass j0850 + 1057 and 2mass j1728 + 3948 have revealed several unusual traits , particularly in their primaries : under- and overluminous fluxes and unusually early and late near - infrared spectral types . these peculiar traits can be related to their unique atmospheric properties . in the case of 2mass j1728 + 3948a , we hypothesize that condensate cloud effects are responsible for shifting this component toward both an earlier near - infrared spectral classification and toward slightly fainter @xmath19-band fluxes . grey extinction from condensate cloud grains in l dwarf photospheres dominate the opacity at the @xmath19- and @xmath36-band flux peaks , as these windows in molecular gas opacity probe deeper into the atmosphere and sample a larger column depth of cloud material @xcite . the @xmath62-band peak , on the other hand , is modulated by both cloud opacity and collision - induced @xmath67 absorption @xcite . as a result , greater condensate opacity tends to produce redder @xmath20 colors and fainter @xmath19-band fluxes @xcite . in addition , contrast in molecular absorption bands is reduced , particularly for the near - infrared h@xmath6o and feh bands to which near - infrared schemes are commonly tied @xcite . veiling of these features can skew near - infrared classifications toward earlier types @xcite . hence , thicker condensate clouds leads to redder near - infrared colors , reduced flux at @xmath19-band , and systematically earlier near - infrared classifications . these conditions accurately reflect the properties of 2mass j1728 + 3948a , and to a lesser degree 2mass j0850 + 1057b . the best - fit primary components for 2mass j1728 + 3948 are consistently red , mid - type l dwarfs whose near - infrared spectral types are consistently earlier than their optical types ( table [ tab_fits_all ] ) . these include the templates 2mass j01443536 - 0716142 ( hereafter 2mass j0144 - 0716 ; l5 optical classification , l4 spex classification ; @xcite ) and 2massw j2224438 - 015852 ( hereafter 2mass j2224 - 0158 ; l4.5 optical classification , l3.5 near - infrared classification ; @xcite ) . both of these sources exhibit indications of thick clouds , based on the detection of linear polarization in the case of 2mass j0144 - 0716 ( [email protected]% at @xmath68-band ; @xcite ) , and spectral model fits and pronounced silicate grain absorption at 911 @xmath9 in the case of 2mass j2224 - 0158 @xcite . unusually red dwarfs such as these have been found to have cooler t@xmath2s for their spectral types @xcite , which can contribute to fainter magnitudes . the older age of the 2mass j1728 + 3948 system , based on the absence of li i absorption , argues that thick condensate clouds , rather than low surface gravity , gives rise to its unusually red color @xcite . thick clouds in 2mass j1728 + 3948a may also explain its comparable brightness at 1 @xmath9 compared to its later - type companion . as discussed above , 2mass j1728 + 3948b appears to be at the threshold of the l dwarf / t dwarf transition , a phase in which condensate clouds are inexplicably dispersed and @xmath19-band fluxes increase . in a sample of unresolved l / t transition binary candidates , @xcite found that sources with comparable component types but redder primaries showed a more pronounced flux reversal than sources with normal or blue primaries . they argued that @xmath19-band fluxes in the primaries of these systems were more suppressed . we may be seeing a similar effect in the 2mass j1728 + 3948ab pair . alternatively , we may be observing the tops of thick clouds that are constrained to the same temperature layer in both sources ( i.e. , t@xmath69 = constant ; @xcite ) . these possibilities should be explored with detailed modeling of resolved component spectra , rather than spectral templates . while thick clouds can explain the unusual faintness of 2mass j1728 + 3948a , thin clouds do not readily explain the unusual brightness of 2mass j0850 + 1057a . thin - cloud l dwarfs , also known as unusually blue l dwarfs ( ubls ) , do exhibit discrepancies between optical and near - infrared spectral types @xcite . there is also evidence that ubls have systematically larger t@xmath2s and/or absolute fluxes for their spectral classifications @xcite . however , 2mass j0850 + 1057a has a normal near - infrared color for its spectral type , and best - fit templates exhibit none of the spectroscopic hallmarks of a ubl ( e.g. , deep h@xmath6o and feh bands , blue near - infrared sed ; @xcite ) . it is therefore unlikely that this source is a ubl with unusually thin clouds . youth may play a role in the unusual brightness of 2mass j0850 + 1057a , arising from the enlarged radius of a still - contracting brown dwarf . the radius of a 1500 k brown dwarf is 25% larger at an age of 250 myr the minimum bound cited by @xcite than at 2 gyr , corresponding to an increase in brightness of roughly 0.5 mag . given the current uncertainties in the distance of this source , such a shift may be sufficient to move 2mass j0850 + 1057a back onto absolute magnitude / spectral type tracks . however , assuming that the two components of this system are coeval , 2mass j0850 + 1057b would also have to be overluminous by roughly the same factor , which does not appear to be the case . moreover , such a correction fails to explain the significant brightness difference between these two comparably - classified sources . we propose an alternative explanation : 2mass j0850 + 1057a is itself an unresolved , near - equal mass binary . such a scenario would explain how this component could appear both brighter and warmer than its equivalently - typed companion , but have an otherwise normal spectral energy distribution . the components of an equal - mass 2mass j0850 + 1057a pair would have absolute magnitudes , luminosities and t@xmath2s fully consistent with empirical trends ( table [ tab_components ] ) . the fact that existing high angular resolution images have not resolved this source requires an angular separation @xmath1150100 mas , or a projected separation @xmath1124 pc at this distance of this system . in fact , long - term dynamic stability requires an even tighter binary . a hierarchical triple is generally stable if the ratio of outer periastron and inner apastron distances : @xmath70 satisfies @xmath71 @xcite . here , @xmath72 @xmath73 1 and @xmath74 are the inner and outer mass ratios . assuming that all three components have nearly equal masses , and that the inner and outer orbits have eccentricities @xmath75 = 0 ( 0.5 ) , equation [ eqn_orbit ] requires a limit on the inner orbit semimajor axis of @xmath76 ( @xmath770.07@xmath78 ) and hence @xmath79 25 mas ( 8 mas ) , or 1 au ( 0.3 au ) , based on the semimajor axis determination of @xcite and the @xcite parallax . such a tight separation is not unusual for brown dwarf multiples ; brown dwarf spectroscopic binaries with comparable separations have already been identified ( e.g. , @xcite ) . moreover , three other tight , hierarchical , brown dwarf triple candidates identified in the literature gliese 569bcd , @xcite , denis j0205 - 1159abc @xcite , and kelu 1abc @xcite exhibit evidence that one component is an unresolved pair , based on radial velocity variations , psf - fitting residuals and spectroscopic features , respectively . this scenario is also consistent with orbital mass constraints from @xcite , as the estimated total mass of a triple 2mass j0850 + 1057 system , [email protected] m@xmath12 , is closer to the mean ( but weakly constrained ) value of 0.2 m@xmath12 found in that study . indeed , tighter constraints on the total mass of this system from ongoing astrometric monitoring may affirm or refute the presence of a third body . the triple hypothesis can also be tested though radial velocity monitoring ; a @xmath115 km s@xmath22 ( @xmath118 km s@xmath22 ) line shift arising from a pair of 0.05 m@xmath12 brown dwarfs separated by 1.0 au ( 0.3 au ) can be readily detected with current near - infrared instrumentation @xcite . if 2mass j0850 + 1057a is confirmed as a binary , it would complete a remarkable , young , low - mass , hierarchical quintuple system with the double m dwarf nltt 20346ab , encompassing 4 orders of magnitude in separation and composed entirely of objects less massive than 0.15 m@xmath12 @xcite .

we deduce that thick condensate clouds are likely responsible for the unusual properties of 2mass j1728 + 3948a , while 2mass j0850 + 1057a is either an inflated young brown dwarf or a tight unresolved binary , making it potentially part of a wide , low - mass , hierarchical quintuple system .

we present a detailed examination of the brown dwarf multiples 2mass j08503593 + 1057156 and 2mass j17281150 + 3948593 , both suspected of harboring components that straddle the l dwarf / t dwarf transition . resolved photometry from _ hubble space telescope_/nicmos show opposite trends in the relative colors of the components , with the secondary of 2mass j0850 + 1057 being redder than its primary , while that of 2mass j1728 + 3948 is bluer . we determine near - infrared component types by matching combined - light , near - infrared spectral data to binary templates , with component spectra scaled to resolved nicmos and photometry . combinations of l7 + l6 for 2mass j0850 + 1057 and l5 + l6.5 for 2mass j1728 + 3948 are inferred . remarkably , the primary of 2mass j0850 + 1057 appears to have a later - type classification compared to its secondary , despite being 0.81.2 mag brighter in the near - infrared , while the primary of 2mass j1728 + 3948 is unusually early for its combined - light optical classification . comparison to absolute magnitude / spectral type trends also distinguishes these components , with 2mass j0850 + 1057a being mag brighter and 2mass j1728 + 3948a.5 mag fainter than equivalently - classified field counterparts . we deduce that thick condensate clouds are likely responsible for the unusual properties of 2mass j1728 + 3948a , while 2mass j0850 + 1057a is either an inflated young brown dwarf or a tight unresolved binary , making it potentially part of a wide , low - mass , hierarchical quintuple system .

1011.0798

i

we have presented photometric and spectroscopic analyses of the late - type l dwarf binaries 2mass j0850 + 1057 and 2mass j1728 + 3948 , aimed at assessing component spectral types , absolute magnitudes and near - infrared colors . multi - band hst / nicmos photometry have revealed distinct trends in the relative colors of these two systems , with 2mass j0850 + 1057b being redder than its primary and 2mass j1728 + 3948b being bluer . neither secondary exhibits narrow - band colors consistent with being a t dwarf . these results are borne out in spectral template fits , using nicmos and @xmath0 resolved photometry , which also determine component near - infrared spectral types of l7 + l6 for 2mass j0850 + 1057 and l5 + l6.5 for 2mass j1728 + 3948 . the early classification of 2mass j1728 + 3948a , its relative faintness at @xmath19 , and its unusually red color can be explained by the presence of thick condensate clouds in its photosphere . the secondary of this system , in contrast , may be losing its photospheric cloud deck as it transitions onto the t dwarf sequence . for 2mass j0850 + 1057 , the surprisingly later spectral type of its bright primary may stem from youth ( inflated radius ) and/or unresolved multiplicity . the latter hypothesis , which would make 2mass j0850 + 1057 part of a low - mass hierarchical quintuple , can be tested through ongoing astrometric monitoring and/or resolved spectroscopic monitoring to search for rv variations . as two resolved ( or partly - resolved ) coeval systems spanning the end of the l dwarf sequence and exhibiting a broad range of cloud properties , 2mass j0850 + 1057 and 2mass j1728 + 3948 remain important laboratories for studying cloud formation and evolution in low - temperature atmospheres . improved parallactic distance measurements including resolution of current distance discrepancies for 2mass j0850 + 1057resolved component spectroscopy , and ongoing photometric and astrometric monitoring will aid in characterizing the clouds , spectral properties and multiplicity of these benchmark brown dwarf systems . the authors would like to thank telescope operators dave griep and bill grolisch and instrument specialist john rayner for their assistance during the irtf observations . we acknowledge helpful comments from trent dupuy and jacqueline faherty on our original manuscript , and thank our referee sandy leggett for her prompt and helpful review . ajb acknowledges support from the chris and warren hellman fellowship ; dcbg acknowledges funding from the john reed fund . this publication makes use of data from the two micron all sky survey , which is a joint project of the university of massachusetts and the infrared processing and analysis center , and funded by the national aeronautics and space administration and the national science foundation . 2mass data were obtained from the nasa / ipac infrared science archive , which is operated by the jet propulsion laboratory , california institute of technology , under contract with the national aeronautics and space administration . this research has also made use of the simbad database , operated at cds , strasbourg , france ; the m , l , and t dwarf compendium housed at dwarfarchives.org and maintained by chris gelino , davy kirkpatrick , and adam burgasser ; the spex prism spectral libraries , maintained by adam burgasser at http://www.browndwarfs.org/spexprism ; and the vlm binaries archive maintained by nick siegler at http://www.vlmbinaries.org . the authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of mauna kea has always had within the indigenous hawaiian community . we are most fortunate to have the opportunity to conduct observations from this mountain . , h. , martn , e. l. , brandner , w. , forveille , t. , delfosse , x. , hulamo , n. , basri , g. , girard , j. , zapatero osorio , m .- r . , stumpf , m. , ghez , a. , valdivielso , l. , marchis , f. , burgasser , a. j. , & cruz , k. 2008 , , 481 , 757 , c. c. , harris , h. c. , vrba , f. j. , guetter , h. h. , canzian , b. , henden , a. a. , levine , s. e. , luginbuhl , c. b. , monet , a. k. b. , monet , d. g. , pier , j. r. , stone , r. c. , walker , r. l. , burgasser , a. j. , gizis , j. e. , kirkpatrick , j. d. , liebert , j. , & reid , i. n. 2002 , , 124 , 1170 , t. r. , knapp , g. r. , leggett , s. k. , fan , x. , golimowski , d. a. , anderson , s. , brinkmann , j. , csabai , i. , gunn , j. e. , hawley , s. l. , hennessy , g. , henry , t. j. , hill , g. j. , hindsley , r. b. , ivezi , . , lupton , r. h. , mcdaniel , a. , munn , j. a. , narayanan , v. k. , peng , e. , pier , j. r. , rockosi , c. m. , schneider , d. p. , smith , j. a. , strauss , m. a. , tsvetanov , z. i. , uomoto , a. , york , d. g. , & zheng , w. 2002 , , 564 , 466 , g. r. , leggett , s. k. , fan , x. , marley , m. s. , geballe , t. r. , golimowski , d. a. , finkbeiner , d. , gunn , j. e. , hennawi , j. , ivezi , z. , lupton , r. h. , schlegel , d. j. , strauss , m. a. , tsvetanov , z. i. , chiu , k. , hoversten , e. a. , glazebrook , k. , zheng , w. , hendrickson , m. , williams , c. c. , uomoto , a. , vrba , f. j. , henden , a. a. , luginbuhl , c. b. , guetter , h. h. , munn , j. a. , canzian , b. , schneider , d. p. , & brinkmann , j. 2004 , , 127 , 3553 , s. k. , geballe , t. r. , fan , x. , schneider , d. p. , gunn , j. e. , lupton , r. h. , knapp , g. r. , strauss , m. a. , mcdaniel , a. , golimowski , d. a. , henry , t. j. , peng , e. , tsvetanov , z. i. , uomoto , a. , zheng , w. , hill , g. j. , ramsey , l. w. , anderson , s. f. , annis , j. a. , bahcall , n. a. , brinkmann , j. , chen , b. , csabai , i. , fukugita , m. , hennessy , g. s. , hindsley , r. b. , ivezi , . , lamb , d. q. , munn , j. a. , pier , j. r. , schlegel , d. j. , smith , j. a. , stoughton , c. , thakar , a. r. , & york , d. g. 2000 , , 536 , l35 , d. 1986 , in society of photo - optical instrumentation engineers ( spie ) conference series , vol . 627 , society of photo - optical instrumentation engineers ( spie ) conference series , ed . d. l. crawford , 733+ , f. j. , henden , a. a. , luginbuhl , c. b. , guetter , h. h. , munn , j. a. , canzian , b. , burgasser , a. j. , kirkpatrick , j. d. , fan , x. , geballe , t. r. , golimowski , d. a. , knapp , g. r. , leggett , s. k. , schneider , d. p. , & brinkmann , j. 2004 , , 127 , 2948

hubble space telescope_/nicmos show opposite trends in the relative colors of the components , with the secondary of 2mass j0850 + 1057 being redder than its primary , while that of 2mass j1728 + 3948 is bluer . combinations of l7 + l6 for 2mass j0850 + 1057 and l5 + l6.5 for 2mass j1728 + 3948 are inferred .

we present a detailed examination of the brown dwarf multiples 2mass j08503593 + 1057156 and 2mass j17281150 + 3948593 , both suspected of harboring components that straddle the l dwarf / t dwarf transition . resolved photometry from _ hubble space telescope_/nicmos show opposite trends in the relative colors of the components , with the secondary of 2mass j0850 + 1057 being redder than its primary , while that of 2mass j1728 + 3948 is bluer . we determine near - infrared component types by matching combined - light , near - infrared spectral data to binary templates , with component spectra scaled to resolved nicmos and photometry . combinations of l7 + l6 for 2mass j0850 + 1057 and l5 + l6.5 for 2mass j1728 + 3948 are inferred . remarkably , the primary of 2mass j0850 + 1057 appears to have a later - type classification compared to its secondary , despite being 0.81.2 mag brighter in the near - infrared , while the primary of 2mass j1728 + 3948 is unusually early for its combined - light optical classification . comparison to absolute magnitude / spectral type trends also distinguishes these components , with 2mass j0850 + 1057a being mag brighter and 2mass j1728 + 3948a.5 mag fainter than equivalently - classified field counterparts . we deduce that thick condensate clouds are likely responsible for the unusual properties of 2mass j1728 + 3948a , while 2mass j0850 + 1057a is either an inflated young brown dwarf or a tight unresolved binary , making it potentially part of a wide , low - mass , hierarchical quintuple system .

1104.3090

c

we have introduced a framework of removable pairings to find eulerian multigraphs . this framework proved to be useful to obtain an approximation algorithm for graph - tspwith an approximation ratio smaller than @xmath0 and to obtain a tight upper bound on the integrality gap of the held - karp relaxation for a restricted class of graphs that contains degree three bounded and claw - free graphs . in particular , we showed that in subcubic @xmath11-vertex - connected graphs we can always find a solution to graph - tspof at most @xmath12 edges , which settles a conjecture from @xcite affirmatively . our framework is not restricted to graph - tsp . with the same techniques and a more detailed analysis , our result translates to the traveling salesman path problem on graphic metrics with prespecified start and end vertex . in this way , one is guaranteed to obtain an approximation ratio smaller than @xmath2 and , for the degree three bounded case , the approximation ratio gets arbitrarily close to @xmath4 . we note that the framework of removable pairings is straightforward to generalize to general metrics , but the problem of finding a large enough removable pairing in such graphs in order to improve on christofides algorithm remains open . 10 sanjeev arora . polynomial time approximation schemes for euclidean traveling salesman and other geometric problems . 45:753782 , 1998 . francisco barahona . fractional packing of t - joins . , 17:661669 , 2004 . sylvia boyd , rene sitters , suzanne van der ster , and leen stougie . on cubic and subcubic graphs . in _ proc . of the 15th conference on integer programming and combinatorial optimization ( ipco 2011 ) _ , 2011 . to appear . nicos christofides . worst - case analysis of a new heuristic for the travelling salesman problem . technical report 388 , graduate school of industrial administration , carnegie - mellon university , 1976 . grard cornujols , jean fonlupt , and denis naddef . the traveling salesman problem on a graph and some related integer polyhedra . , 33:127 , 1985 . jack edmonds . maximum matching and a polyhedron with @xmath322 vertices . , 69:125130 , 1965 . greg n. frederickson and joseph jaja. on the relationship between the biconnectivity augmentation and travelling salesman problems . , 19(2):189 201 , 1982 . david gamarnik , moshe lewenstein , and maxim sviridenko . an improved upper bound for the tsp in cubic 3-edge - connected graphs . , 33(5):467474 , 2005 . shayan oveis gharan , amin saberi , and mohit singh . a randomized rounding approach to the traveling salesman problem . preprint , 2011 . michel x. goemans and dimitris j. bertsimas . on the parsimonious property of connectivity problems . in _ proceedings of the 1st annual acm - siam symposium on discrete algorithms ( soda 1990 ) _ , pages 388396 , 1990 . michel x. goemans . worst - case comparison of valid inequalities for the tsp . , 69:335349 , 1995 . michelangelo grigni , elias koutsoupias , and christos h. papadimitriou . an approximation scheme for planar graph tsp . in _ proc . of the 36th annual symposium on foundations of computer science ( focs 1995 ) _ , pages 640645 , 1995 . martin grtschel , lszlo lovsz , and alexander schrijver . , volume 2 of _ algorithms and combinatorics_. springer , 1988 . michael held and richard m. karp . the traveling - salesman problem and minimum spanning trees . , 18:11381162 , 1970 . j. a. hoogeveen . analysis of christofides heuristic : some paths are more difficult than cycles . , 10(5):291295 , 1991 . atsushi kaneko , alexander kelmans , and tsuyoshi nishimura . on packing 3-vertex paths in a graph . , 36(4):175197 , 2001 . joseph s. b. mitchell . guillotine subdivisions approximate polygonal subdivisions : a simple polynomial - time approximation scheme for geometric tsp , k - mst , and related problems . , 28:12981309 , march 1999 . c. l. monma , b. s. munson , and w. r. pulleyblank . minimum - weight two - connected spanning networks . , 46:153171 , 1990 . denis naddef and wiliam r. pulleyblank . matchings in regular graphs . , 34(3):283291 , 1981 . christos h. papadimitriou and santosh vempala . on the approximability of the traveling salesman problem . , 26(1):101120 , 2006 . alexander schrijver . . springer , 2003 . david b. shmoys and david p. williamson . analyzing the held - karp tsp bound : a monotonicity property with application . , 35(6):281285 , 1990 . laurence a. wolsey . heuristic analysis , linear programming and branch and bound . in _ combinatorial optimization ii _ , volume 13 of _ mathematical programming studies _ , pages 121134 . springer , 1980 .

the approach yields a-approximation algorithm with respect to the held - karp lower bound . for graph - tsprestricted to a class of graphs that contains degree three bounded and claw - free graphs , we show that the integrality gap of the held - karp relaxation matches the conjectured ratio .

we present a framework for approximating the metric tsp based on a novel use of matchings . traditionally , matchings have been used to add edges in order to make a given graph eulerian , whereas our approach also allows for the removal of certain edges leading to a decreased cost . for the tsp on graphic metrics ( graph - tsp ) , the approach yields a-approximation algorithm with respect to the held - karp lower bound . for graph - tsprestricted to a class of graphs that contains degree three bounded and claw - free graphs , we show that the integrality gap of the held - karp relaxation matches the conjectured ratio . the framework allows for generalizations in a natural way and also leads to a-approximation algorithm for the traveling salesman path problem on graphic metrics where the start and end vertices are prespecified .

0907.2016

i

in this study , we consider nonlinear evolution equations of the form @xmath9 , \qquad { \bf x}\in\mathbb{r}^d , \quad d>1.\ ] ] examples for such equations are the nonlinear schrdinger equation , the biharmonic nonlinear schrdinger equation , the nonlinear heat equation , and the biharmonic nonlinear heat equation . it is well known that these equations admit solutions that become singular at a point . recently , it was discovered that the nonlinear schrdinger equation with a quintic nonlinearity admits solutions that become singular on a @xmath0-dimensional sphere @xcite , see figure [ fig : standingringilustration ] . following @xcite , we refer to these solutions as _ singular standing - ring _ solutions . the main goal of this study is to present a general framework for constructing singular standing - ring solutions of nonlinear evolution equations of the form ( [ eq : nl_pde_form ] ) . in order to understand the basic idea , let us assume that equation admits a singular standing - ring solution . then , near the singularity , equation reduces to the one - dimensional equation @xmath10 , \qquad r = |{\bf x}|.\ ] ] hence , equation ( [ eq : nl_pde_form_one_dimensional ] ) `` should '' admit a solution that becomes singular at a point . conversely , if the one - dimensional equation ( [ eq : nl_pde_form_one_dimensional ] ) admits a solution that becomes singular at a point , then equation `` should '' admit a standing - ring singular solution . moreover , the asymptotic profile and blowup rate of the standing - ring solutions of `` should '' be the same as those of the corresponding solution of the one - dimensional equation ( [ eq : nl_pde_form_one_dimensional ] ) . the above argument is obviously very informal . nevertheless , in what follows we will provide numerical evidence in support of the relation between standing - ring singular solutions of and singular solutions of the one - dimensional equation ( [ eq : nl_pde_form_one_dimensional ] ) .

we present a general framework for constructing singular solutions of nonlinear evolution equations that become singular on a -dimensional sphere , where . we provide a detailed numerical investigation of these new singular solutions for the following equations : the nonlinear schrdinger equation with , the biharmonic nonlinear schrdinger equation with , the nonlinear heat equation with , and the nonlinear biharmonic heat equation with .

we present a general framework for constructing singular solutions of nonlinear evolution equations that become singular on a -dimensional sphere , where . the asymptotic profile and blowup rate of these solutions are the same as those of solutions of the corresponding one - dimensional equation that become singular at a point . we provide a detailed numerical investigation of these new singular solutions for the following equations : the nonlinear schrdinger equation with , the biharmonic nonlinear schrdinger equation with , the nonlinear heat equation with , and the nonlinear biharmonic heat equation with .

1511.05145

c

using 38 sne ii in the hubble flow we develop a technique based solely on photometric data ( pcm ) to build a hubble diagram based on sne ii . in summary : 1 . using pcm we find a dispersion of 0.44 mag using the @xmath2 band and 0.43 mag with the @xmath25 band , thus using nir filters the improvement is not so significant for the pcm . 2 . the @xmath11 plays a useful role , allowing us to reduce the dispersion from 0.58 mag to 0.50 mag for @xmath2 band . 3 . the colour term does not have so much influence on the nir filters because it is related to the host - galaxy extinction . we find very low ( @xmath61 ) values ( the colour - magnitude coefficient ) . if @xmath61 is purely extrinsic , it implies very low @xmath56 values . the hubble diagram derived from the csp sample using the scm yields to a dispersion of 0.29 mag , some what better than those found in the literature and emphasising the potential of scm in cosmology . it is interesting also to obtain more data and sne for which the initial decline rate and the plateau are clearly visible to try to reduce this dispersion . the pcm is very promising , and more efforts must be done in this direction , i.e. , trying to use only photometric parameters . in the coming era of large photometric wide field surveys like lsst , having spectroscopy for every sne will be impossible hence the pcm which is the first purely photometric method could be very useful . the referee is thanked for their through reading of the manuscript , which helped clarify and improve the paper . support for t. d. , s. g. , l. g. , m. h. , c. g. , f. o. , h. k. , is provided by the ministry of economy , development , and tourism s millennium science initiative through grant ic120009 , awarded to the millennium institute of astrophysics , mas . s. g. , l. g. , h. k. and f.o . also acknowledge support by conicyt through fondecyt grants 3130680 , 3140566 , 3140563 and 3140326 , respectively . the work of the csp has been supported by the national science foundation under grants ast0306969 , ast0607438 , and ast1008343 . m. d. s. , c. c. and e. h. gratefully acknowledge generous support provided by the danish agency for science and technology and innovation realized through a sapere aude level 2 grant . the authors thank f. salgado for his work done with the csp . this research has made use of the nasa / ipac extragalactic database ( ned ) which is operated by the jet propulsion laboratory , california institute of technology , under contract with the national aeronautics and space administration and of data provided by the central bureau for astronomical telegrams . figure [ bootstrapping_scm ] ( left ) presents the evolution of the rms versus the number of sne for both methods ( pcm and scm ) using the @xmath25 band . for both methods , after a constant median value the rms decreases when the number of sne is lower than 1012 sne because the model starts diverging . indeed if we look at the figure [ bootstrapping_scm ] on the right where the evolution of the fit parameters versus the number of sne for one single epoch ( optd*0.55 ) and the @xmath25 band are presented , we see that for the pcm , @xmath53 , @xmath61 and @xmath55 change significantly when the number of sne is around 12 . the values start diverging for a number of sne smaller than 12 , so this implies that the rms is driven by the reduced number of objects and therefore it will be difficult to conclude between the fact that @xmath61 and @xmath55 are better because we have a better rms or because it is due to a statistical effect . note that the figure does not present directly the value of the fit parameters but a fraction of the value , i.e. , the value divided by the first value plus an offset corresponding to the first value . + band and the @xmath1 colour . the black squares represent the evolution for the pcm whereas the black circles are used for the scm . _ right figure : _ we present the evolution of our fit parameters ( @xmath53 , @xmath61 , and @xmath55 ) versus the number of sne . the black colour represents the @xmath61 , the red is for @xmath55 , and the blue for @xmath53 . the circles are used for the scm and the squares for the pcm.,title="fig:",width=340 ] band and the @xmath1 colour . the black squares represent the evolution for the pcm whereas the black circles are used for the scm . _ right figure : _ we present the evolution of our fit parameters ( @xmath53 , @xmath61 , and @xmath55 ) versus the number of sne . the black colour represents the @xmath61 , the red is for @xmath55 , and the blue for @xmath53 . the circles are used for the scm and the squares for the pcm.,title="fig:",width=340 ]

we obtain a dispersion of 0.44 mag using a combination of the colour and the band and we are able to reduce the dispersion to 0.39 mag using our golden sample . a comparison of our photometric colour method ( pcm ) with the standardised candle method ( scm ) is also performed . the dispersion obtained for the scm ( which uses both photometric and spectroscopic information ) the construction of a photometric hubble diagram is of high importance in the coming era of large photometric wide - field surveys , which will increase the detection rate of supernovae by orders of magnitude .

we present a hubble diagram of type ii supernovae using corrected magnitudes derived only from photometry , with no input of spectral information . we use a data set from the carnegie supernovae project i ( csp ) for which optical and near - infrared light - curves were obtained . the apparent magnitude is corrected by two observables , one corresponding to the slope of the plateau in the band and the second a colour term . we obtain a dispersion of 0.44 mag using a combination of the colour and the band and we are able to reduce the dispersion to 0.39 mag using our golden sample . a comparison of our photometric colour method ( pcm ) with the standardised candle method ( scm ) is also performed . the dispersion obtained for the scm ( which uses both photometric and spectroscopic information ) is 0.29 mag which compares with 0.43 mag from the pcm , for the same sn sample . the construction of a photometric hubble diagram is of high importance in the coming era of large photometric wide - field surveys , which will increase the detection rate of supernovae by orders of magnitude . such numbers will prohibit spectroscopic follow - up in the vast majority of cases , and hence methods must be deployed which can proceed using solely photometric data .

astro-ph0508457

i

the magellanic clouds ( mcs ) are the closest , easily observable galaxies from our vantage point in the milky way . they are of fundamental importance for studies of stellar populations , the interstellar medium , and the cosmological distance scale . moreover , due to their locations at 50 kpc ( lmc ) and 61 kpc ( smc ) from the sun , and @xmath4 kpc from the galactic plane , they provide one of our best probes of the composition and properties of the galactic dark halo . precise measurements for the proper motions ( pms ) of the magellanic clouds combined with observational knowledge of the distribution and velocities of hi gas in the magellanic stream can constrain models of the shape and radial density distribution of the galactic dark halo as well as theoretical models for the formation of the magellanic stream . few developments in modern astronomy are as important as the discovery of cosmic dark matter . most of the matter in the universe is invisible to astronomers . in the milky way , the dark component is about twenty times more massive than the visible disk of stars and gas . the dark matter forms a vast , diffuse halo that occupies more than a thousand times the volume of the stellar disk . the composition of this dark halo is unknown , but it may comprise a mixture of ancient degenerate dwarf stars and exotic , hypothetical elementary particles ( for a review see alcock 2000 ) . the investigation of the galactic halo is complicated by our immersion deep inside it . interior to the solar circle the rotation curve may be measured with confidence ( see e.g. fich & tremaine 1991 ) . while the estimates out to 20 kpc are not as robust , these observations indicate that the rotation curve is flat ( merrifield 1992 ; yoshiaki & rubin 2001 ) . much has been learned about the halo in the direction out of the plane from the study of the motions of gas clouds and stars . these imply that the halo is more spherical than disk - like , but no precise measurement for its axial ratio exists ( van der marel 2001 ) . analyses of the kinematics of some globular clusters and small satellite galaxies suggest that the halo extends out to @xmath5 kpc ( e.g. , wilkinson & evans 1999 ) . for the large majority of these objects there is only radial velocity data , and where available , proper motions are much less precisely determined than the radial velocities . so while these studies constrain the total halo mass , they say little about its shape or distribution . spherical symmetry and a particular density distribution are generally _ assumed_. the magellanic stream is an hi emission feature that spans more than @xmath6 across the sky ( br " uns et al . it consists of gas that trails the magellanic clouds as they orbit the milky way and provides an opportunity for detailed study of tidal disruption of galaxies as well as the milky way dark halo . many detailed theoretical models have been constructed for the magellanic stream ( gardiner & noguchi 1996 ; lin , jones & klemola 1995 ; heller & rohlfs 1994 ) . these models describe how the magellanic clouds orbit the milky way , and how they lose material through tidal effects and other physical mechanisms . the model parameters are adjusted to best reproduce the observed density , morphology and velocities of the hi gas seen along the magellanic stream . the most sophisticated calculations in the class of models that invoke tidal stripping are those by gardiner & noguchi ( 1996 ) . in their models the lmc and smc form a gravitationally bound system that orbits the milky way . the magellanic stream represents material that was stripped from the smc @xmath7 gyr ago . this was the time of the previous perigalactic passage , which coincided with a close encounter between the clouds . in magellanic stream models , the orbit of the clouds is one of the prime variables that is adjusted to fit the data . as a result of newton s law , the orbit is fully determined by : ( a ) the position of the clouds on the sky ; ( b ) the present distance of the clouds from the sun ; ( c ) the radial velocities of the clouds ; ( d ) the proper motions of the clouds ; and ( e ) the gravitational potential of the galactic dark halo . the first three of these are known relatively well ; the latter two are not . even small differences in the proper motions can give vastly different orbits for the clouds , especially when considered over the long period of time over which the magellanic stream developed . the usual approach has therefore been to estimate the proper motions of the clouds from the properties of the magellanic stream , assuming a fixed galactic halo gravitational potential ( usually a simple spherical isothermal halo ) . the results imply that the magellanic clouds are just past their pericenter ( itself at @xmath8 kpc from the galactic center ) , that the apocenter to pericenter ratio is @xmath9 , and the orbital period is @xmath7 gyr . the inferred tangential velocity in the galactocentric rest - frame has differed substantially from model to model , ranging from , e.g. , @xmath10 km / s ( heller & rohlfs 1994 ) to @xmath11 km / s ( gardiner & noguchi 1996 ) . it was realized by heller & rohlfs ( 1994 ) and lin , jones & klemola ( 1995 ) that the arguments used to model the magellanic stream can be turned around . if the proper motions of the magellanic clouds are known , then the gravitational potential of the galactic dark halo can be determined from models of the magellanic stream . they explored this using the available proper motion data , but found that its accuracy was insufficient to obtain strong constraints on the dark halo properties . kroupa et al . ( 1994 ) found that a proper motion better than @xmath12 is required for strong results . so far the proper motion of the lmc has not been known accurately enough to strongly constrain properties of the milky way dark halo . measurements , however , are steadily improving and determinations are available from the following sources : kroupa et al . ( 1994 ) , using stars from the ppm catalogue ; jones et al . ( 1994 ) , using photographic plates with a 14 year epoch span ; kroupa & bastian ( 1997 ) , using hipparcos data ; drake et al . ( 2002 ) , using data from the macho project ; anguita , loyola & pedreros ( 2000 ) and pedreros et al . ( 2002 ) , using ccd frames with an 11 year epoch span ; and momany & zaggia ( 2005 ) , using the usno ccd astrograph all - sky catalog ( ucac2 ) . the measurements are all consistent with each other to within the error bars with the exception of two outliers : anguita et al . ( 2000 ) and momany & zaggia ( 2005 ) present values of @xmath13 . when these two results are ignored , the weighted average of the remaining measurements yields proper motions towards the west and north of @xmath14 ( van der marel et al . 2002 ) i.e. to @xmath15% accuracy . this implies @xmath16 km / s , which is consistent with most of the published magellanic stream models , but is not accurate enough to discriminate between them . the proper motion of the smc is much less well known than that of the lmc , and only one reasonably accurate measurement exists . kroupa & bastian ( 1997 ) obtain a value of @xmath17 for the smc with an error of ( 0.84 , 0.75)@xmath18 . a sound measurement of the proper motion of the lmc requires all of the following key factors : ( 1 ) an instrument which can perform the astrometry with adequate precision ; ( 2 ) a reference frame consisting of point sources , such as quasars , distributed widely behind the clouds ; ( 3 ) secure determination of the membership of stars in the lmc ; and ( 4 ) a reliable kinematic model of the internal rotation of the clouds . the previous estimates of the motions ( referred to above ) do not satisfy all of these requirements . however , a much improved measurement is now possible . some teams have identified qsos behind our neighboring galaxies in order to provide good inertial reference frames against which the motions of these galaxies can be measured ( geha et al . 2003 ; dobrzycki et al . 2002 ; 2003 ) ; the @xmath0 has been shown to be very stable for astrometry ( anderson & king 2003 ) and in particular the advanced camera for surveys ( acs ) has higher resolution , and is better calibrated and more stable than even wfpc2 ( anderson & king 2004 ; hereafter ak04 ) . we were awarded two epochs on the high resolution camera ( hrc ) on the acs in cycle 11 and cycle 13 for a study of the proper motions of the mcs . we report our proper motion result for the lmc in this paper . smc results as well as implications for the mc system as a whole will be presented in subsequent papers . the paper is outlined as follows : 2 deals with the sample and the observations ; 3 describes our analysis ; 4 presents the center of mass proper motion of the lmc ; 5 is a discussion and 6 is a summary .

we present a measurement of the systemic proper motion of the large magellanic cloud ( lmc ) from astrometry with the high resolution camera ( hrc ) of the advanced camera for surveys ( acs ) on the hubble space telescope ( ) . the qsos are distributed homogeneously behind the central few degrees of the lmc . with 2 epochs of hrc data and a year baseline

we present a measurement of the systemic proper motion of the large magellanic cloud ( lmc ) from astrometry with the high resolution camera ( hrc ) of the advanced camera for surveys ( acs ) on the hubble space telescope ( ) . we observed lmc fields centered on 21 background qsos that were discovered from their optical variability in the macho database . the qsos are distributed homogeneously behind the central few degrees of the lmc . with 2 epochs of hrc data and a year baseline we determine the proper motion of the lmc to better than% accuracy : . this is the most accurate proper motion measurement for any milky way satellite thus far . when combined with hi data from the magellanic stream this should provide new constraints on both the mass distribution of the galactic halo and models of the stream .

astro-ph0508457

i

we undertook a project using two epochs of @xmath0/acs data of magellanic cloud fields centered on background qsos to determine the systemic proper motion of the clouds . the lmc results are presented in this paper . we have determined the proper motion of the lmc to be @xmath158 . this is accurate to better than 5% . when combined with hi data for the magellanic stream , this should allow improved constraints on both the mass distribution in the galactic halo , and theoretical models for the origin of the magellanic stream . our data provides the most accurate proper motion measurement for any milky way satellite . in future papers , we will present results for the smc as well as implications for the lmc - smc system and the magellanic stream . improvements to our work may be possible by using @xmath0 with a longer baseline . a long baseline would facilitate measurements of the internal motions of the clouds as well as an even more accurate measurement of their systemic motions . in addition we may be able to make a distance determination using the method of rotational parallax ( olling & peterson , 2000 ) . rotational parallax is the method of determining distances to local group galaxies by measuring a rotation curve using both proper motions and radial velocities ( see brunthaler et al . 2005 for an application of this method to m33 ) . since the former are distance - dependent and the latter are not , using both methods provides a distance measurement . more generally , considerable improvement may be expected when the next generation of astrometric satellites such as @xmath159 and @xmath160 come on - line . the authors would like to thank jay anderson and ivan king for their geometric distortion calibration of the hrc , and in particular , jay anderson for making his analysis software available to the community . support for this work was provided by nasa through grant numbers go-09462 and go-10130 from the space telescope science institute ( stsci ) , which is operated by the association of universities for research in astronomy , inc . , under nasa contract nas5 - 26555 . m. g. is supported by nasa through hubble fellowship grant hf-01159.01-a awarded by stsci . khc s work was performed under the auspices of the u.s . department of energy , national nuclear security administration by the university of california , lawrence livermore national laboratory under contract no . w-7405-eng-48 .

we determine the proper motion of the lmc to better than% accuracy : . this is the most accurate proper motion measurement for any milky way satellite thus far . when combined with hi data from the magellanic stream this should provide new constraints on both the mass distribution of the galactic halo and models of the stream .

we present a measurement of the systemic proper motion of the large magellanic cloud ( lmc ) from astrometry with the high resolution camera ( hrc ) of the advanced camera for surveys ( acs ) on the hubble space telescope ( ) . we observed lmc fields centered on 21 background qsos that were discovered from their optical variability in the macho database . the qsos are distributed homogeneously behind the central few degrees of the lmc . with 2 epochs of hrc data and a year baseline we determine the proper motion of the lmc to better than% accuracy : . this is the most accurate proper motion measurement for any milky way satellite thus far . when combined with hi data from the magellanic stream this should provide new constraints on both the mass distribution of the galactic halo and models of the stream .

1106.5739

i

feynman integrals in perturbative quantum field theory are generically expressed in terms of the classical polylogarithm functions @xmath3 and the nielsen polylogarithms @xmath4 @xcite . in the late nineties , it was realized that these classes of functions are too restricted when going beyond one - loop level in the perturbative expansion , where new functions appear that can no longer be expressed in terms of the classical polylogarithm functions . while a completely generic generalization of polylogarithms has been studied in the mathematical literature ( going under the name of multiple polylogarithms @xcite ) , it is mostly only a specific subset of multiple polylogarithms , the so - called harmonic polylogarithms @xcite and their two - dimensional and cyclotomic generalizations @xcite , that make their appearance in the theoretical predictions of physical quantities beyond leading - order . in this paper we concentrate exclusively on harmonic polylogarithms ( hpl s ) up to weight four , which since their introduction have found many applications in computations up to two - loop order in the perturbative expansion , _ e.g. _ , @xcite . in order to confront the theoretical next - to - next - to - leading order ( nnlo ) predictions to experiment , it is mandatory to be able to evaluate hpl s numerically in a fast and accurate way . the requirements to such a numerical code are twofold : first , the evaluation should be fast , because the use of nnlo matrix elements in monte carlo integration codes may require thousands , if not millions , of function calls . second , it is desirable to be able to compute hpl s for arbitrary complex arguments , which appear for example when the complex mass scheme is employed or for certain kinematic configurations in loop calculations involving massive particles @xcite . in the last decade , various codes have been developed to evaluate hpl s numerically . while the code hplog @xcite , written in fortran , is restricted to the evaluation of hpl s up to weight four and for real values of the arguments , the code hpl ( mathematica)@xcite and the implementation of the harmonic polylogarithms into the ginac framework ( c++ ) @xcite are generic and allow to evaluate in principle any harmonic polylogarithm with arbitrary precision for any complex argument . the focus of this paper is the chaplin ( complex harmonic polylogarithms in fortran ) library , a new fortran code that allows to evaluate numerically all harmonic polylogarithms up to weight four for arbitrary complex arguments . while chaplin is similar in spirit to the aforementioned codes , its main advantages lie in speed , through the use of fortran as a programming language , and in its capability to compute hpl s numerically for any point in the complex plane . chaplin reduces each of the 120 hpl s up to weight four to a set of 32 basis functions @xcite , which are entirely expressed through only two new functions of weight four besides the classical polylogarithms . the basis functions are then mapped to the interior of the unit circle , where they are computed numerically using suitably chosen series expansions that allow to obtain a fast numerical convergence . the paper is organized as follows : in section [ sec : hpl_review ] we give a short review of harmonic polylogarithms and of their main algebraic and analytic properties . in section [ reductionchapter ] we review the reduction of all hpl s up to weight four to the set of basis functions introduced in ref . the series expansions used by chaplin to compute the basis functions numerically are derived in section [ sec : series ] , while the chaplin library itself , together with comparisons to hplog , hpl and ginac , is presented in section [ sec : chaplin ] .

pacs : 12.38.bx , perturbative calculations harmonic polylogarithms , fortran , loop computations . ippp/11/36 , dcpt/11/72 , , * program summary * _ manuscript title : _ chaplin - complex harmonic polylogarithms in fortran + _ authors : _ stephan buehler , claude duhr + _ program title : _ chaplin + _ journal reference : _ + _ catalogue identifier : _ + _ licensing provisions : _ + _ programming language : _ fortran 77 + _ computer : _ computing systems on which fortran 77 compilers are available . + _ keywords : _ harmonic polylogarithms , fortran , loop computations . + _ pacs : _ 12.38.bx , perturbative calculations . + _ solution method : _ inside the unit circle : series expansion . + _ restrictions : _ only harmonic polylogarithms up to weight four are supported . + _ unusual features : _ allows to evaluate hpl s numerically for any point in the complex plane . + _ running time : _ depending on the weight vector and argument of the hpl , between and . + +

we present a new fortran library to evaluate all harmonic polylogarithms up to weight four numerically for any complex argument . the algorithm is based on a reduction of harmonic polylogarithms up to weight four to a minimal set of basis functions that are computed numerically using series expansions allowing for fast and reliable numerical results . pacs : 12.38.bx , perturbative calculations harmonic polylogarithms , fortran , loop computations . ippp/11/36 , dcpt/11/72 , , * program summary * _ manuscript title : _ chaplin - complex harmonic polylogarithms in fortran + _ authors : _ stephan buehler , claude duhr + _ program title : _ chaplin + _ journal reference : _ + _ catalogue identifier : _ + _ licensing provisions : _ + _ programming language : _ fortran 77 + _ computer : _ computing systems on which fortran 77 compilers are available . + _ operating system : _ operating systems on which fortran 77 compilers are available . + _ keywords : _ harmonic polylogarithms , fortran , loop computations . + _ pacs : _ 12.38.bx , perturbative calculations . + _ classification : _ 11.1 general , high energy physics and computing + _ nature of problem : _ numerical evaluation of harmonic polylogarithms . + _ solution method : _ inside the unit circle : series expansion . outside the unit circle : inversion relations . + _ restrictions : _ only harmonic polylogarithms up to weight four are supported . + _ unusual features : _ allows to evaluate hpl s numerically for any point in the complex plane . + _ running time : _ depending on the weight vector and argument of the hpl , between and . + +

1106.5739

i

in this paper we have presented chaplin , a new fortran library to compute harmonic polylogarithms up to weight four for arbitrary complex argument . the algorithm is based on a reduction of hpl s to a set of basis functions which are then evaluated numerically using series expansions allowing for a very fast numerical convergence , hence rendering the computational cost of a function call quite modest . we have checked our numerical results against well - established codes @xcite and found agreement to at least 14 digits for any argument in the complex plane .

we present a new fortran library to evaluate all harmonic polylogarithms up to weight four numerically for any complex argument . the algorithm is based on a reduction of harmonic polylogarithms up to weight four to a minimal set of basis functions that are computed numerically using series expansions allowing for fast and reliable numerical results .

we present a new fortran library to evaluate all harmonic polylogarithms up to weight four numerically for any complex argument . the algorithm is based on a reduction of harmonic polylogarithms up to weight four to a minimal set of basis functions that are computed numerically using series expansions allowing for fast and reliable numerical results . pacs : 12.38.bx , perturbative calculations harmonic polylogarithms , fortran , loop computations . ippp/11/36 , dcpt/11/72 , , * program summary * _ manuscript title : _ chaplin - complex harmonic polylogarithms in fortran + _ authors : _ stephan buehler , claude duhr + _ program title : _ chaplin + _ journal reference : _ + _ catalogue identifier : _ + _ licensing provisions : _ + _ programming language : _ fortran 77 + _ computer : _ computing systems on which fortran 77 compilers are available . + _ operating system : _ operating systems on which fortran 77 compilers are available . + _ keywords : _ harmonic polylogarithms , fortran , loop computations . + _ pacs : _ 12.38.bx , perturbative calculations . + _ classification : _ 11.1 general , high energy physics and computing + _ nature of problem : _ numerical evaluation of harmonic polylogarithms . + _ solution method : _ inside the unit circle : series expansion . outside the unit circle : inversion relations . + _ restrictions : _ only harmonic polylogarithms up to weight four are supported . + _ unusual features : _ allows to evaluate hpl s numerically for any point in the complex plane . + _ running time : _ depending on the weight vector and argument of the hpl , between and . + +

0802.0377

i

in recent years , the dynamics of heavy ion collisions at intermediate energy has been extensively investigated within the framework of transport theories , such as the nordheim approach , in which the vlasov equation for the one - body phase space density , @xmath1 , is supplemented with a pauli - blocked boltzmann collision term @xcite . the basic ingredients that enter the resulting transport equation , often called boltzmann - uehling - uhlenbeck ( buu ) equation , are the self - consistent mean - field potential and the two - body scattering cross sections . these transport models hence describe the time evolution of the reduced one - body density in phase - space and , consequently , they are suited for the description of one - body observables , such as inclusive particle spectra in nuclear collisions , average collective flows and excitations . however , they can not provide a reliable description of fluctuation phenomena , such as multi - fragmentation processes , i.e. the break - up of excited nuclear systems into many pieces . in fact , neither fluctuations of one - body observables nor many - body correlations can be addressed with this class of mean - field models . hence suitable extensions , including fluctuations of the one - body density , have to be considered . an intense theoretical work on fluctuations in nuclear dynamics has started in the past years , also stimulated by the availability of large amounts of experimental data on fragment formation in intermediate energy heavy ion collisions and the possibility to observe a liquid - gas phase transition @xcite . in order to introduce fluctuations in transport theories , a number of different avenues have been taken , that can be essentially reconducted to two different classes of models . one is the class of molecular dynamics models @xcite while the other kind is represented by stochastic mean - field approaches @xcite . in molecular dynamics models the many - body state is represented by a simple product wave function , with or without antisymmetrization . the single particle wave functions are assumed to have a fixed gaussian shape . in this way , though nucleon wave functions are supposed to be independent ( mean - field approximation ) , the use of localized wave packets induces many - body correlations both in mean - field propagation and hard two body scattering ( collision integral ) , which is treated stochastically . hence this way to introduce many - body correlations and produce a trajectory branching is essentially based on the use of empirical gaussian wave packets . if wave functions were allowed to assume any shape , the method would become identical to standard mean - field descriptions . while the wave function localization appears appropriate to describe final fragmentation channels , where each single particle wave function should be localized within a fragment , the use of fixed shape localized wave packets in the full dynamics could affect the correct description of one - body effects , such as spinodal instabilities and zero sound propagation @xcite . on the other side , in the so - called stochastic mean - field approaches , the stochastic extension of the transport treatment for the one - particle density is obtained by introducing a stochastic term representing the fluctuating part of the collision integral @xcite , in close analogy with the langevin equation for a brownian motion . this can be derived as the next - order correction , in the equation describing the time evolution of @xmath2 , with respect to the standard average collision integral , leading to the boltzmann - langevin ( bl ) equation . thus , the system is still described solely in terms of the reduced one - body density @xmath2 , but this function experiences a stochastic time evolution in response to the random effect of the fluctuating collision term . in this way density fluctuations are introduced , that are amplified when instabilities or bifurcations occur in the dynamics . this procedure is suitable also for addressing multifragmentation phenomena , since fragments can be associated with the regions where the spacial density becomes larger , which finally can be reconstructed by sampling the one - body distribution function . a specific method for solving the boltzmann - langevin equation by direct numerical simulation was introduced in refs.@xcite . in this numerical implementation , the one - particle density @xmath3 is represented on a lattice of grid points in phase space and the collision integral is treated by considering all possible transitions between phase space cells , adding a noise term whose features are related to the average rate of transitions between two specified initial cells and two final cells . the numerical implementation of this method has only been possible in two dimensions ( 2d ) because it requires too large computer resources in 3d . hence several approximate solutions of the bl equations have been formulated , mostly based on the projection of the bl noise only on a given dynamical variable ( such as the local quadrupole tensor of the momentum distribution ) or on @xmath4 space only @xcite . more tractable fluctuating terms , such as a stochastic force added to the mean - field potential , have also been proposed and extensively applied to multifragmentation studies @xcite . however , the implementation of the full structure in phase space of the original bl term can still be considered as an important goal to reach . in fact , this would allow to treat a more general class of phenomena , where the correct description of fluctuations and correlations in @xmath5 space is essential ( such as particle production and fragment velocity correlations for instance ) . moreover , also in the multi - fragmentation mechanism , that is dominated by spacial density fluctuations , a more accurate representation of the full phase space dynamics , including fluctuations , would allow to improve the description of the fragment kinematical properties and correlations . we also stress the general interest of this effort . indeed transport phenomena occur in many physical systems , for which a more precise description of the time evolution of the one - body distribution function , including the effect of many - body correlations , would be important . a first attempt to introduce a fluctuating collision term in a 3d transport approach was made by bauer et al . this method can be implemented relatively easily into standard transport codes that adopt the scattering of pseudo - particles ( or test particles ) as a method of solution of the collision integral and it consists essentially in forcing similar two - body collisions to occur for neighboring test particles , defined according to a given distance in phase space , so that effectively two nucleons are involved in each particular collision event . the distance should reproduce the phase - space shape of the nucleon wave function . in this way the random nature of the two - body scattering , that in the standard codes applies only to test particles and is washed - out when using a huge number of them , is transferred to entire nucleons . however , in the procedure proposed in ref.@xcite the pauli blocking is checked only for the collision of the two original test particles and not for the entire swarm affected , leading to some unpleasant features in the description of fermionic systems . indeed the pauli - blocking violation introduces important inaccuracies in the fluctuations of the one - body density . in the present manuscript we present a new method to reconstruct the phase space nucleon wave function in mean - field approaches , in such a way that the pauli - blocking is checked for the entire cloud of moved particles this will improve the description of the fluctuation variance , approaching the one expected for fermionic systems . we will also pay special attention to the definition of the phase - space metric that would optimize the value of this variance . the paper is organized as follows : we will first recall the main ingredients of the bl theory , in order to connect the formalism with the numerical implementation adopted ( section 2 ) . then we will discuss in more detail the methods that have been proposed so far to solve the boltzmann - langevin equation ( section 3 ) . the new procedure that we follow to build fluctuations is presented in section 4 . several results demonstrating and analyzing in detail the method are discussed from section 5 to 8 . conclusions and perspectives are drawn in section 9 .

we present a new method to introduce phase - space fluctuations in transport theories , corresponding to a full implementation of the boltzmann - langevin equation for fermionic systems . it is based on the procedure originally developed by bauer et al . for transport codes employing the test particle method . in the new procedure , accurate tests are carried out in one and two dimensional idealized systems , and finally results for a full 3d application are shown . we stress the reliability of this method , which can be easily plugged into existing tranport codes using test particles , and its general applicability to systems characterized by instabilities , like for instance multifragmentation processes . fluctuations ; stochastic collision integral ; fermionic systems ; transport theories . + pacs numbers : 24.10.cn ; 24.60.ky , 25.70.pq

we present a new method to introduce phase - space fluctuations in transport theories , corresponding to a full implementation of the boltzmann - langevin equation for fermionic systems . it is based on the procedure originally developed by bauer et al . for transport codes employing the test particle method . in the new procedure , the pauli principle is carefully checked , leading to a good reproduction of the correct fluctuations in the `` continuum limit '' ( ) . accurate tests are carried out in one and two dimensional idealized systems , and finally results for a full 3d application are shown . we stress the reliability of this method , which can be easily plugged into existing tranport codes using test particles , and its general applicability to systems characterized by instabilities , like for instance multifragmentation processes . fluctuations ; stochastic collision integral ; fermionic systems ; transport theories . + pacs numbers : 24.10.cn ; 24.60.ky , 25.70.pq

1312.2025

i

atmospheric pressure is a fundamental parameter for characterizing the environment and habitability of an extrasolar planet . water s stability on a planetary surface as a liquid depends on both the surface temperature and pressure . while the freezing point of water is not strongly dependent on pressure , pressure does affect water s boiling point and sublimation . thus , a reliable estimate of the surface pressure is an important part of the measurement suite required to determine the habitability of an exoplanet . despite the importance of atmospheric pressure , current proposed methods for measuring pressure using remote - sensing techniques that could be applicable to exoplanet atmospheres are challenging . the existing techniques include the use of rayleigh scattering @xcite or the widths of individual absorption lines @xcite or absorption bands @xcite . the presence and location of a blue rayleigh scattering tail in a spectrum can provide information about the existence and pressure of an atmosphere . however , strong blue absorbers in the atmosphere ( e.g. o@xmath3 , so@xmath0 , no@xmath0 and many others ) or surface features can mask this tail @xcite . furthermore , the rayleigh scattering tail is most prominent shortward of 0.6 @xmath1 m , below the short wavelength cutoff of the james webb space telescope ( jwst ) @xcite . lastly , planets around m dwarfs are likely to be the first to be characterized @xcite , and m dwarfs have relatively less visible - flux to rayleigh scatter than solar - type stars . the rayleigh tail would be more difficult to detect and characterize for planets orbiting stars of this stellar class . it is also possible to use the widths of absorption features to estimate pressure . pressure increases the widths of vibration rotation lines of gases . this method has been successful for the earth using high - resolution spectra of the o@xmath0 a band @xcite , for mars using co@xmath0 features near 2 @xmath1 m @xcite and the cloud tops of venus using the 1.6 @xmath1 m co@xmath0 band @xcite . this method provides unambiguous results when the spectral resolution is sufficiently high to resolve the profiles of individual spectral lines . it can also be used at lower spectral resolution , but requires prior knowledge of the mixing ratio of the absorbing gas . here we explore the feasibility of a new method to directly measure the pressure of an earth - like atmosphere that combines the absorption features of dimers with those of monomer vibration - rotation bands to yield estimates of the atmospheric pressures even when the mixing ratio of the monomer is uncertain . previous pressure estimates using dimer absorption have been made for the cloud tops of earth , but these techniques have required prior knowledge of the gas mixing ratio profile @xcite . dimers are bound or quasi - bound states between two molecules driven together by molecular interactions . for example , the o@xmath0-o@xmath0 or o@xmath4 dimer consists of two o@xmath0 molecules temporarily bound to each other by van der waals forces . this dimer has its own rotational and vibrational modes , and produces spectral features distinct from its constituent o@xmath0 monomers . additionally , absorption from dimer molecules is more sensitive to pressure than that of monomers . the optical depth ( how much absorption occurs ) for dimers and monomers can be expressed by the following equations : @xmath5 where @xmath6 and @xmath7 are the monomer and dimer differential optical depths , @xmath8 is the monomer cross section , @xmath9 is the number density of the gas , @xmath10 is the dimer cross section , @xmath11 is the pressure , @xmath12 is the temperature , and @xmath13 is the path length . while the monomer ( e.g. o@xmath0 ) optical depth is directly proportional to pressure , the optical depth of the dimer ( e.g. o@xmath0-o@xmath0 ) is dependent on the square of the density ( and hence square of the pressure ) . this difference in pressure dependence allows us to estimate atmospheric pressure by comparing the dimer and monomer absorption features for an oxygen - rich , earth - like atmosphere , the best combination of bands to use for pressure determination at near infrared wavelengths ( @xmath20.6 @xmath1 m ) would be the ( o@xmath0 ) a band at 0.76 @xmath1 m , and the 1.06 @xmath1 m o@xmath4 dimer band . the 0.76 @xmath1 m o@xmath0 a band is the strongest o@xmath0 feature in the visible - near - infrared spectral region , and is found in a relatively clean region of the spectrum between two water vapor bands . we have chosen the dimer feature at 1.06 @xmath1 m as the likely best option , due to its combination of band strength and its location in a relatively uncluttered region of the planetary spectrum . other o@xmath4 features overlap with water features ( dimer feature between 5.5 and 7 @xmath1 m ) or o@xmath0 vibration - rotation bands ( 0.63 @xmath1 m , 0.76 @xmath1 m and 1.27 @xmath1 m dimer features ) or are weaker than the 1.06 @xmath1 m dimer feature ( 0.477 @xmath1 m and 0.57 @xmath1 m dimer features ) . nevertheless , some other features , in particular the strong 1.27 @xmath1 m feature , could be used if the 1.06 @xmath1 m dimer feature is not detectable . in the proposed technique , the o@xmath0 0.76@xmath1 m ( monomer ) band is used to provide an estimate of the atmospheric concentration of o@xmath0 , and combined with the o@xmath4 1.06 @xmath1 m ( dimer ) band to constrain the atmospheric pressure . this method can be used with either transmission spectroscopy , or directly - detected reflection or emission spectra , and serves as a complement to pressure determination techniques such as rayleigh scattering , which only work in the visible . although the proof of concept is shown here with oxygen , this technique is not limited to the oxygen dimer in an earth - like atmosphere in the visible to near - infrared . the same technique is applicable to pairs of monomer and dimer absorption features across a wider range of planetary atmospheric composition and spectral wavelength range .

we present a new method to probe atmospheric pressure on earthlike planets using ( o-o ) dimers in the near - infrared . the absorption by dimers changes more rapidly with pressure and density than that of monomers , and can therefore provide additional information about atmospheric pressures . by comparing the absorption strengths of rotational and vibrational features to the absorption strengths of dimer features ,

we present a new method to probe atmospheric pressure on earthlike planets using ( o-o ) dimers in the near - infrared . we also show that dimer features could be the most readily detectable biosignatures for earthlike atmospheres , and may even be detectable in transit transmission with the james webb space telescope ( jwst ) . the absorption by dimers changes more rapidly with pressure and density than that of monomers , and can therefore provide additional information about atmospheric pressures . by comparing the absorption strengths of rotational and vibrational features to the absorption strengths of dimer features , we show that in some cases it may be possible to estimate the pressure at the reflecting surface of a planet . this method is demonstrated by using the o a band and the 1.06 m dimer feature , either in transmission or reflected spectra . it works best for planets around m dwarfs with atmospheric pressures between 0.1 and 10 bars , and for o volume mixing ratios above 50% of earth s present day level . furthermore , unlike observations of rayleigh scattering , this method can be used at wavelengths longer than 0.6 m , and is therefore potentially applicable , although challenging , to near - term planet characterization missions such as jwst . we have also performed detectability studies for jwst transit transmission spectroscopy and find that the 1.06 m and 1.27 m dimer features could be detectable ( snr ) for an earth - analog orbiting an m5v star at a distance of 5 pc . the detection of these features could provide a constraint on the atmospheric pressure of an exoplanet , and serve as biosignatures for oxygenic photosynthesis . we have calculated the required signal - to - noise ratios to detect and characterize o monomer and dimer features in direct imaging reflected spectra and find that signal - to - noise ratios greater than 10 at a spectral resolving power of r=100 would be required .

1312.2025

m

in this paper we generated transit transmission and direct imaging reflected spectra for cloud and aerosol - free earthlike exoplanets . the models we have used to do this are described below . when an extrasolar planet transits or occults its host star , the planetary atmosphere is backlit and some of the star s light traverses the planet s atmosphere on limb trajectories . this transmitted light can be used to characterize the planet s atmosphere @xcite . this has been done for a number of jupiter- and neptune - sized planets ( e.g. @xcite ) and the super - earth / mini - neptune gj1214b ( e.g. @xcite ) the transmission spectroscopy model used here is based on smart ( spectral mapping atmospheric radiative transfer ) @xcite , which is a spectrum resolving ( line - by - line ) , multi - stream , multiple scattering radiative transfer model . we have modified smart to generate transit transmission spectra by combining the monochromatic absorption and scattering opacities for each atmospheric layer calculated by smart with the limb path lengths inherent in a transit transmission event . the transmission model included gas absorption , rayleigh scattering , interaction - induced absorption , extinction from clouds and aerosols , refraction and limb darkening . because multiple scattering was not included in the transit transmission component of the model , we have considered only cloud and aerosol - free atmospheres for this work . an important characteristic of the transit transmission spectroscopy model is the inclusion of refraction . as described in @xcite , refraction sets a fundamental limit on the range of pressures that can be probed during a transit , independent of absorption and scattering . light refracts as it passes through an atmosphere , with a larger refraction angle as higher pressures and densities are probed . for every planet - star system there will be a maximum tangent pressure in the planet s atmosphere that can be probed , because at greater pressures the light will be refracted by too large of an angle to be able to reach a distant observer during the transit . for each tangent height in each atmosphere , we calculated the total angle of refraction for a beam of light emitted from the host star using a modified version of the method described in @xcite . their method was developed for calculating refraction for astronomical observations on the earth given a tangent altitude , apparent zenith angle and an atmospheric density profile . we calculated the angle of refraction over a range of zenith angles to determine if a path exists to connect the host star to a distant observer via the planetary atmosphere . transit transmission spectroscopy can not probe the tangent altitudes at which no such path exists . the radius of the star , planet - star distance and composition of the planet s atmosphere determine the maximum tangent pressure . the radius of the star and planet - star distance control the apparent angular size of the star from the planet s perspective . the larger the angular size , the greater the range of pressures that can be probed . for an earth - analog orbiting a sun - like star the angular size of the star is @xmath140.5@xmath15 while for an earth - analog orbiting an m5 dwarf and receiving the same total flux , the angular size of the star is @xmath142@xmath15 . therefore , transit transmission spectroscopy can probe higher pressures , i.e. see deeper into the atmosphere , for the planet orbiting an m dwarf . figure [ fig : refraction ] shows this effect by comparing all possible paths at one tangent height for a planet around an m dwarf and a planet around a sun - like star , with each planet receiving the same total flux . the composition of the atmosphere determines the refractivity ( index of refraction - 1 ) of the atmosphere . atmospheres with greater refractivities will have lower maximum tangent pressures . the refractivity at stp ( standard temperature and pressure ) can vary from @xmath141.5 times the refractivity of air for co@xmath0 to slightly less than half the refractivity of air for h@xmath0 , when considering only the common bulk atmospheric gases in the solar system . therefore , in general it will be possible to probe higher pressures for an h@xmath0 atmosphere than a co@xmath0 or air atmosphere . the reflected spectra were generated using the standard version of smart , which can include multiple scattering from clouds and aerosols . however , to maintain consistency with the transmission spectrum model , clouds were not included in this study . the reflected spectra were generated assuming a surface with a constant albedo of 0.16 , which is the average albedo of the cloud - free earth @xcite . we also assumed the surface is a lambertian scatterer . other surface types could have introduced an error into any quantitative estimates in this paper , unless explicitly included in a retrieval attempt . figure [ fig : albedo ] shows the wavelength - dependence of a variety of surface types from the aster ( advanced spaceborne thermal emission and reflection radiometer ) spectral library @xcite and the usgs ( united states geological survey ) digital spectral library @xcite . surface albedos are nearly constant in the bands considered here , though , for example , snow fluctuates by @xmath1420% within the 1.06 @xmath1 m dimer band . we modeled a test case with smart with the snow surface to determine the error different surfaces can introduce . we found a difference of 15% between the seawater and snow cases when measuring the equivalent width ( defined in section [ sec : measurements ] ) of the 1.06 @xmath1 m dimer band . this difference was significantly less than the difference in measured equivalent widths for the cases considered in this paper . therefore , we have considered any discrepancies due to variations in surface albedo to be minimal for the present work . the cloud - free model atmospheres were generated by a one - dimensional ( altitude ) photochemical code with an extensive history in early earth @xcite , modern earth @xcite and exoplanet @xcite research . the planetary radius and surface gravity used were the radius of the earth ( 6371 km ) and surface gravity of the earth ( 9.87 m / s@xmath16 ) . the vertical grid consists of 200 plane - parallel layers that are each 0.5 km thick in altitude , in which radiative transfer , atmospheric transport , and photochemical production and loss are solved simultaneously , subject to upper ( stellar flux , atmospheric escape ) and lower ( volcanos and biology ) boundary conditions . the model calculates the mixing ratios of each species in each layer by solving the coupled mass - continuity / flux equations with the reverse euler method ( appropriate for stiff systems ) and a variable time - stepping algorithm . dimer concentrations were computed using estimations from quantum mechanical calculations . first , we fit the temperature dependence of the equilibrium constant for o@xmath0-dimer formation using the following equations : @xmath17 where p(o@xmath0)@xmath0 and po@xmath0 are the dimer and monomer partial pressures and @xmath12 is the temperature @xcite . the dimer mixing ratio was then computed as : @xmath18 where @xmath11 is the pressure in atmospheres . the choice of quantum parameters in our fit of equation [ eqn : kpt ] ensures that our calculation is a lower limit to the dimer concentrations , an assumption which matches modern atmospheric data well @xcite . the model atmospheres used in this study started with boundary conditions that reproduce earth s modern atmospheric chemistry . we then replaced the solar spectrum with the m dwarf spectrum described in [ sec : stellar - prop ] and decreased the surface albedo to 0.16 to account for cloud - free conditions . this `` modified earth around an m dwarf '' model was then perturbed to examine changes to both total pressure and oxygen concentrations . total atmospheric pressures of 0.1 , 0.5 , 1.0 , 3.0 , 5.0 and 10.0 bar were examined . at each of these total pressures , lower boundary conditions on o@xmath0 mixing ratios were set at 0.1 , 0.5 , 1.0 and 2.0 times earth s present level . this corresponds to oxygen mixing ratios from 2% to 42% , which is roughly the range of oxygen concentrations experienced throughout the past 2.5 gyr of earth s history @xcite . boundary conditions for all other species were held fixed at modern values . figure [ fig : profiles ] shows the pressure - temperature profile and volume mixing ratio profiles for the 1.0 times the present atmospheric level ( pal ) o@xmath0 cases for spectrally active gases in the wavelength region examined here . stable steady - state solutions were analyzed in all but the three cases with the highest o@xmath0 surface partial pressures . for those three cases , we extrapolated from our converged results at lower o@xmath0 mixing ratios by increasing o@xmath0 concentrations , scaling the dimer concentrations with the square of the o@xmath0 mixing ratio , and keeping all other gas mixing ratios constant . for all tests presented here , we assumed earth - like temperature and water vapor profiles , precluding the need for costly climate simulations . more specifically , we took the modern earth temperature profile as a function of altitude , and computed the corresponding pressure levels assuming hydrostatic equilibrium . this simplification provides a useful baseline , but limits the range of validity of these results somewhat because the gas absorption cross sections and number densities depend on these atmospheric properties . the impact of this assumption is assessed in section [ sec : sensitivity ] . for pressure / temperature regions corresponding to earth s mesosphere , we adopted a temperature of 180 k which likely overestimates the temperature in these regions . dimer absorption preferentially occurs in higher pressure regions and so is insensitive to assumed conditions in the tenuous upper atmosphere . our model grid was capped at 100 km altitude ensuring optically thin conditions for nearly all species , and made allowances for co@xmath0 and n@xmath0 photolysis above the upper boundary . we used the modern measured eddy - diffusion profile to simulate convective motion for all atmospheres regardless of pressure . while this simplification would affect the prediction of trace gas concentrations , all species analyzed here are either well - mixed ( co@xmath0 , o@xmath0 , n@xmath0 ) or short - lived ( dimers ) so are not sensitive to changes in turbulent mixing . we used the hitran 2008 database line lists @xcite to generate opacities in smart . the dimer absorption cross sections for o@xmath0 were taken from @xcite . in this paper we have assumed that the planet orbits an m5v star , or an m dwarf with stellar radius of 0.20 r@xmath19 and a luminosity of 0.0022 l@xmath19 @xcite . the planet was placed at a distance of 0.047 au , so that the total integrated incoming stellar flux was equal to the total flux the earth receives from the sun today . an m5 dwarf was chosen for these tests because there is a high probability that the terrestrial planets which will be the most easily characterized in the near future will be orbiting this class of star @xcite . additionally , transit transmission spectroscopy can probe pressures up to the @xmath141 bar level in an atmosphere for an earth analog around an m dwarf while for the earth around the sun , it can only probe pressures as great as @xmath140.2 bars . at these pressures and lower , the 1.06 @xmath1 m dimer feature is very weak , even at 200% pal o@xmath0 . thus , using earth - like atmospheres around an m dwarf instead of around the sun provides a better demonstration of the pressure - dependence of dimer features . to simulate the spectrum of the star , we used a phoenix nextgen synthetic spectrum @xcite for all wavelengths greater than @xmath14300 nm . for the shorter wavelengths we used the uv spectrum of ad leo @xcite . the phoenix spectrum was normalized so that the total integrated flux was equal to 1373 w / m@xmath16 , which is the integrated flux the earth receives from the sun . the ad leo spectrum from @xcite was left unchanged , as it was already normalized to equal the amount of flux a planet near the inner edge of the habitable zone would receive . we have used a spectral resolving power of 100 to provide relevance to the james webb space telescope ( jwst ) , which will provide new opportunities for characterizing the atmospheres of transiting exoplanets @xcite . once the jwst is launched , the near - infrared spectrograph ( nirspec ) will provide spectra with a spectral resolving power ( r=@xmath20 ) of @xmath21100 between 0.6 and 5.0 @xmath1 m in single prism mode @xcite . we examine the effect of varying the spectral resolution in sections [ sec : resolution1 ] and [ sec : resolution2 ] . to make a quantitative estimate of absorption strengths we measured equivalent widths for the reflected spectra and measured parts per million ( ppm ) differences in flux for the transmission spectra . equivalent widths were calculated using the following equation : @xmath22 where @xmath23 is the equivalent width , f@xmath24 is the flux at each wavelength @xmath25 , and f@xmath26 is the continuum flux at each wavelength . to obtain the equivalent widths , we first measured the area of the spectral band below the continuum . for each absorption band , we define the continuum by hand . for the o@xmath0 a band , the continuum was assumed to be linear with wavelength , and for the 1.06 @xmath1 m feature the continuum was assumed to be constant with wavelength because in several of the cases the continuum at the longer wavelengths was difficult to define due to h@xmath0o absorption . the equivalent width is the width , in units of wavelength , of a rectangle measured from the continuum to the level of zero flux with the same total area as the spectral band . we could not use this type of measurement for the transmission spectra because of the difficulty in defining the zero flux level . therefore , we quantified the absorption strengths for the transmission spectra by measuring the change in flux from the continuum to the point of greatest absorption within a band . we performed detectability studies @xcite for the model spectra , assuming the exoplanet - star system is at a distance of 5 pc . we calculated the expected signal - to - noise ratio ( snr@xmath27 ) using the jwst exposure time calculator ( etc ) . the jwst etc includes background noise from sky , dark , thermal and zodiacal sources along with read - out noise and photon noise . the estimates provided are expected to be within 20% of the mission requirements . we note that for all the cases considered , the noise is dominated by photon noise . we do not include noise from detector intrapixel variations @xcite , but in principle calibration time could be devoted to mapping the pixels , as has been done with spitzer space telescope infrared array camera @xcite . we also note that noise levels within @xmath1420% of the photon noise limit have been obtained for transit transmission spectra with hst in spatial scan mode , wherein the target star is trailed during each exposure by telescope motion perpendicular to the direction of dispersion @xcite . spatial scan mode is being considered for jwst , and so assuming photon - limited noise in our calculations , while optimistic , provides a reasonable estimate of the detectability of absorption features ( drake deming , private communication , october 13 , 2013 ) . we assumed that every possible transit is observed in jwst s 5 year mission lifetime , ignoring decreases in integration time due to non - zero impact parameters ( where the planet does not traverse the center of the stellar disk and therefore has a lower transit duration ) and limits on visibility based on the ecliptic latitude of the exoplanet ( see @xcite ) . for the earth - sun analog , this corresponds to a total integration time of @xmath142.3 * 10@xmath28 seconds , and for the earth - m5v analog , an integration time of @xmath1410@xmath29 seconds . we normalize the solar spectrum ( using the phoenix g2v model available on the site ) to a johnson v magnitude of 3.32 and the m5v spectrum to 2.1 * 10@xmath30 erg @xmath31 s@xmath32 @xmath32 at 1 @xmath1 m . at each wavelength within an absorption band we measured the signal as the magnitude of the difference from the continuum flux . the noise is expressed in parts per million ( ppm , 10@xmath29/snr@xmath27 ) at each wavelength . the final snr of the transit transmission spectrum is the square root of the sum of the squares of the snr at each wavelength in the absorption band divided by @xmath33 , which is included because the transit transmission spectrum must be calibrated against the out of transit spectrum of the star . we also performed detectability studies for the direct imaging reflected spectra that could be relevant to proposed direct imaging planet detection and characterization missions . for these calculations , we did not use an instrument simulator because the exact specifications for these missions are not currently defined . for each absorption band , we calculated two snrs , one for detecting the spectral feature ( snr@xmath34 ) and one for measuring the flux at the center of the band to a precision of 3@xmath8 ( snr@xmath35 ) . we calculated two snrs because obtaining information about pressure from a spectral feature requires more than detection ; it also requires a quantitative estimate of the strength of that spectral feature . to calculate snr@xmath34 , we divided the reflected flux by the stellar flux , defined a continuum and calculated the signal as the difference between the continuum and the normalized reflected flux . we assumed the noise is constant over the entire absorption band , and then calculated the noise level required to detect the spectral feature with a snr@xmath36 of 3 in the absorption band . the final snr@xmath34 is the mean of the continuum reflected flux level divided by the calculated noise . to calculate snr@xmath35 , we selected the wavelength within the band with the lowest radiance . we set the value of a second noise level as the lowest radiance divided by 3 . snr@xmath35 is the continuum flux level divided by the noise required to obtain a snr of 3 at the lowest radiance in the absorption band .

it works best for planets around m dwarfs with atmospheric pressures between 0.1 and 10 bars , and for o volume mixing ratios above 50% of earth s present day level . furthermore , unlike observations of rayleigh scattering , this method can be used at wavelengths longer than 0.6 m , and is therefore potentially applicable , although challenging , to near - term planet characterization missions such as jwst .

we present a new method to probe atmospheric pressure on earthlike planets using ( o-o ) dimers in the near - infrared . we also show that dimer features could be the most readily detectable biosignatures for earthlike atmospheres , and may even be detectable in transit transmission with the james webb space telescope ( jwst ) . the absorption by dimers changes more rapidly with pressure and density than that of monomers , and can therefore provide additional information about atmospheric pressures . by comparing the absorption strengths of rotational and vibrational features to the absorption strengths of dimer features , we show that in some cases it may be possible to estimate the pressure at the reflecting surface of a planet . this method is demonstrated by using the o a band and the 1.06 m dimer feature , either in transmission or reflected spectra . it works best for planets around m dwarfs with atmospheric pressures between 0.1 and 10 bars , and for o volume mixing ratios above 50% of earth s present day level . furthermore , unlike observations of rayleigh scattering , this method can be used at wavelengths longer than 0.6 m , and is therefore potentially applicable , although challenging , to near - term planet characterization missions such as jwst . we have also performed detectability studies for jwst transit transmission spectroscopy and find that the 1.06 m and 1.27 m dimer features could be detectable ( snr ) for an earth - analog orbiting an m5v star at a distance of 5 pc . the detection of these features could provide a constraint on the atmospheric pressure of an exoplanet , and serve as biosignatures for oxygenic photosynthesis . we have calculated the required signal - to - noise ratios to detect and characterize o monomer and dimer features in direct imaging reflected spectra and find that signal - to - noise ratios greater than 10 at a spectral resolving power of r=100 would be required .

1312.2025

r

we have generated transit transmission spectra and direct imaging spectra for atmospheres with o@xmath0 concentrations of 10% , 50% , 100% and 200% of the pal and pressures of 0.1 , 0.5 , 1 , 3 , 5 and 10 bars . figure [ fig : plottran ] shows the resulting transit transmission spectra for 100% pal o@xmath0 . transmission spectra of atmospheres with pressures @xmath371 bar are nearly identical for a given o@xmath0 concentration because refraction limits to depth of penetration to @xmath381 bar . the dimer feature is weak for pressures of 0.1 and 0.5 bars . figures [ fig : plottran-0.1 ] , [ fig : plottran-0.5 ] , and [ fig : plottran-2.0 ] show the resulting transit transmission spectra for o@xmath0 concentrations of 10% , 50% and 200% pal . the 1.06 @xmath1 m dimer feature is very weak at all pressures for o@xmath0 concentrations at 10% pal and weak for 50% pal o@xmath0 , while it is very strong for pressures above 0.1 bars at 200% pal o@xmath0 . figure [ fig : plotrad ] shows the reflected spectra for modern o@xmath0 mixing ratios in atmospheres with a range of total pressure from 0.1 to 10 bars . the 0.5 and 5.0 bar cases are omitted from the plots to increase clarity , but are still included in the equivalent width and snr calculations . the 1.06 @xmath1 m dimer feature is extremely weak for the present - day atmosphere , but in contrast to its behavior in the transmission spectra , it is a very prominent feature for the 3 , 5 and 10 bar atmospheres . figures [ fig : plotrad01 ] , [ fig : plotrad05 ] and [ fig : plotrad2 ] show the reflected spectra at pressures between 0.1 and 10 bars for 10% , 50% , and 200% pal of o@xmath0 , respectively . in the 10% pal case , the o@xmath0 dimer feature at 1.06 @xmath1 m is very weak because the total amount of o@xmath0 is very low , even for a 10 bar atmosphere . the 1.06 @xmath1 m dimer feature is stronger in the 3 , 5 and 10 bar cases for atmospheres with 50% pal o@xmath0 . finally for the 200% o@xmath0 atmospheres , the 1.06 @xmath1 m dimer feature is one of the strongest spectral features , even in the 1 bar atmosphere . figure [ fig : equivwidths ] shows the flux change for the transmission spectra and equivalent widths for reflected spectra at different pressures and o@xmath0 concentrations for the o@xmath0 a band and the 1.06 @xmath1 m dimer feature . for a given o@xmath0 concentration the transit transmission flux differences for the o@xmath0 a band are roughly constant for pressures @xmath371 bar due to refraction . the 1.06 @xmath1 m dimer feature flux differences increase slightly with pressure but are also constant for pressures @xmath371 bar for a given o@xmath0 concentration . the dimer feature does not appear in transmission for cases with 10% pal o@xmath0 . for the direct imaging ( reflected ) spectra , both the o@xmath0 a band and 1.06 @xmath1 m dimer feature equivalent widths increase with pressure and increased o@xmath0 concentrations . however , the dimer feature equivalent widths are much more sensitive to pressure . at higher pressures the dimer feature is strong except for cases with 10% pal o@xmath0 , in which the 1.06 @xmath1 m dimer feature is too weak to quantify . figures [ fig : pressure - tran - only]-[fig : refl - only ] show the relationships between the quantitative absorption measurements described above and atmospheric quantities including the o@xmath0 mixing ratio and the o@xmath0 partial pressure at the surface . the ppm flux difference measured in transit transmission for the o@xmath0 a band could be used to constrain the o@xmath0 mixing ratio , as shown in figure [ fig : pressure - tran - only]b . the o@xmath0 partial pressure at the surface can be estimated using the ratio between the 1.06 @xmath1 m dimer and o@xmath0 a band absorption measurements . these ratios are shown in figure [ fig : pressure - tran - only]a using ppm flux differences , and in figure [ fig : refl - only - toto2 ] using equivalent widths . the o@xmath0 mixing ratio and o@xmath0 partial pressure at the surface can be combined to provide a unique estimate of the surface pressure of the planet . a more detailed description of the pressure measurement technique is given in section [ sec : measure ] . to quantify the errors introduced by assuming the modern day temperature profile , we generated spectra to test the sensitivity of our models to changes in the temperature profile and in changes to the water vapor profiles . we compared our 1.0 bar , 1.0x pal o@xmath0 spectra to a spectrum generated using the same volume mixing ratio profiles but with an isothermal atmosphere at 250 k. we also compared our 1.0 bar , 1.0x pal o@xmath0 transit transmission spectrum to spectra generated with atmospheres with 0.1 and 10.0 times the h@xmath0o levels . we perform a similar comparison for the 1.0 bar , 2.0x pal o@xmath0 reflected spectrum . figure [ fig : iso - comp ] shows the sensitivity of the spectra to the temperature profile , with our earth - like profile and an isothermal approximation profile compared . both the transit transmission and reflected spectra show little sensitivity to the temperature profile . the 1.06 @xmath1 m dimer band also shows little sensitivity to the temperature profile , despite the dependence of the dimer optical depth on the square of the temperature because the isothermal temperature profile approximates the average temeprature of the troposphere , where the majority of dimer absorption occurs . figures [ fig : h2o - test-1 ] , [ fig : h2o - test-2 ] and [ fig : h2o - test-3 ] show the results for the h@xmath0o sensitivity tests . the o@xmath0 a band equivalent widths and ppm flux differences are not strongly affected by changes in the h@xmath0o profiles . however , changes in the h@xmath0o mixing ratios affect the continuum flux in the wings of the 1.06 micron dimer feature , complicating measurements of the equivalent width of this feature . for the transit transmission spectra , the change in the total ppm flux difference ( across the entire band ) is less than 20% between the 0.1 and 10.0x h@xmath0o cases . for the reflected spectra , the change in the equivalent width of the 1.06 @xmath1 m dimer feature is less than 20% . for the reflected spectra , the equivalent widths can be much greater than that for the 1.0 bar , 2.0x o@xmath0 case . for these greater equivalent widths , the effect of increasing or decreasing h@xmath0o levels will diminish as the difference from the continuum flux ( affected by h@xmath0o ) and the flux within the absorption band will increase . table [ tab : detect ] shows the snrs for observations by jwst for the o@xmath0 a band , 1.06 @xmath1 m feature and the 1.27 @xmath1 m feature for the range of pressures and o@xmath0 concentrations considered here for an earth analog at a distance of 5 pc . the snr are calculated assuming that every transit of an earth analog orbiting an m5v star is observed over jwst s 5 year mission lifetime . in transit transmission , the o@xmath0 a band snrs are no greater than 1.1 . the 1.06@xmath1 m dimer feature is detectable at a snr of @xmath23 for many of the 2.0x pal o@xmath0 cases . the 1.27 @xmath1 m feature is the most detectable o@xmath0 feature in this wavelength range , with snrs greater than 5 for many of the 1.0x pal o@xmath0 cases and greater than 7 for many of the 2.0x pal o@xmath0 cases . jwst will not be able to detect o@xmath0 species for earth - like exoplanets in secondary eclipse in the visible and near - infrared , as shown by the secondary eclipse snr levels . even for the highest pressure cases , the snrs are no greater than @xmath140.2 . table [ tab : tpf ] shows the snrs necessary to detect and characterize the o@xmath0 a band , 1.06 @xmath1 m band and the 1.27 @xmath1 m feature for the range of pressures and o@xmath0 concentrations considered here for the direct imaging reflected spectra at r=100 . the snrs in table [ tab : tpf ] would be relevant to a direct imaging characterization mission . while these snrs were calculated for an earth analog orbiting an m5v star , the results should be independent of stellar spectral type because we divided out the stellar flux in our calculations . the o@xmath0 a band , 1.06 @xmath1 m dimer feature and 1.27 @xmath1 m feature are detectable at an average snr@xmath34 of 14 , 9 and 14 , respectively , for the cases when the features are strong enough to identify in the model spectra . the average required snr@xmath35 to use the features for pressure estimation are 11 , 31 and 34 . we examined the effect of spectral resolving power on the detectability of spectral features by generating transit transmission spectra for the 1.0 bar , 1.0x pal o@xmath0 case with spectral resolving powers of 500 , 200 , 100 , 80 , 60 , 40 , 30 and 20 . figure [ fig : specres1 ] shows the spectra for each of these cases . we also measured the snr of each feature at each resolving power . the snrs were calculated assuming a noise profile equivalent to the jwst nirspec noise profile , but with the noise level at each wavelength divided by ( 100/r)@xmath16 , so that the noise at each wavelength decreased as the resolving power decreased . figure [ fig : specres2 ] shows how the total snr in each absorption band changes with resolving power . in general , the snrs drop off rapidly as resolving power decreases for r@xmath3860 . we also generated direct imaging reflected spectra for the 1.0 bar , 2.0x pal o@xmath0 case at resolving powers of 500 , 200 , 100 , 80 , 60 , 40 , 30 and 20 , as shown in figure [ fig : specres - imaging ] . figures [ fig : specres - imaging - snrd ] and [ fig : specres - imaging - snrp ] show how the snr for detection and precision vary with resolving power for the o@xmath0 a band , 1.06 @xmath1 m dimer feature and the 1.27 @xmath1 m feature . the snr required to detect each feature increases as resolving power decreases . at r=20 , the snrs are not shown because no spectral features could be identified , and at r=30 only the o@xmath0 a band was identified . the snr required to quantify the flux at the center of each spectral feature decreases as resolving power decreases , because the lowest radiance level increases as the spectral resolving power decreases .

we have also performed detectability studies for jwst transit transmission spectroscopy and find that the 1.06 m and 1.27 m dimer features could be detectable ( snr ) for an earth - analog orbiting an m5v star at a distance of 5 pc . we have calculated the required signal - to - noise ratios to detect and characterize o monomer and dimer features in direct imaging reflected spectra and find that signal - to - noise ratios greater than 10 at a spectral resolving power of r=100 would be required .

we present a new method to probe atmospheric pressure on earthlike planets using ( o-o ) dimers in the near - infrared . we also show that dimer features could be the most readily detectable biosignatures for earthlike atmospheres , and may even be detectable in transit transmission with the james webb space telescope ( jwst ) . the absorption by dimers changes more rapidly with pressure and density than that of monomers , and can therefore provide additional information about atmospheric pressures . by comparing the absorption strengths of rotational and vibrational features to the absorption strengths of dimer features , we show that in some cases it may be possible to estimate the pressure at the reflecting surface of a planet . this method is demonstrated by using the o a band and the 1.06 m dimer feature , either in transmission or reflected spectra . it works best for planets around m dwarfs with atmospheric pressures between 0.1 and 10 bars , and for o volume mixing ratios above 50% of earth s present day level . furthermore , unlike observations of rayleigh scattering , this method can be used at wavelengths longer than 0.6 m , and is therefore potentially applicable , although challenging , to near - term planet characterization missions such as jwst . we have also performed detectability studies for jwst transit transmission spectroscopy and find that the 1.06 m and 1.27 m dimer features could be detectable ( snr ) for an earth - analog orbiting an m5v star at a distance of 5 pc . the detection of these features could provide a constraint on the atmospheric pressure of an exoplanet , and serve as biosignatures for oxygenic photosynthesis . we have calculated the required signal - to - noise ratios to detect and characterize o monomer and dimer features in direct imaging reflected spectra and find that signal - to - noise ratios greater than 10 at a spectral resolving power of r=100 would be required .

1312.2025

c

the 1.06 @xmath1 m dimer feature is prominent in transit transmission for atmospheres with @xmath3750% pal o@xmath0 and surface pressures @xmath370.5 bars . it is a prominent feature in the reflected spectra for atmospheres with @xmath3750% pal o@xmath0 and surface pressure @xmath373 bars . for the reflected spectra , the dimer feature equivalent width is highly dependent on surface pressure when compared to the o@xmath0 a band and therefore the dimer feature can be used to constrain pressure . here we discuss how to do this for the cases investigated here . figure [ fig : equivwidths ] confirms that dimer absorption features are more strongly dependent on pressure than monomer features . therefore dimer features can be combined with monomer features to determine pressure , even if the mixing ratio of the absorbing gas is not known . with only transit transmission spectroscopy , it is impossible to probe pressures over @xmath141 bar for the cases examined here , but it may be possible to constrain the o@xmath0 mixing ratio and set a lower bound for pressure . in reflected spectra , it is possible to determine the surface partial pressure of o@xmath0 and set a lower bound for pressure using the 1.06 @xmath1 m dimer feature as an on / off pressure gauge . with both a transit transmission spectrum and a reflected spectrum , it should be possible to determine total atmospheric surface pressure for an earth - like exoplanet . transmission spectroscopy provides only a lower bound on the atmospheric pressure because refraction provides a fundamental limit to which pressures can be probed using this technique . for the spectra presented here , a lower limit of @xmath141 bar can be set for the high pressure atmospheres . for a given o@xmath0 concentration a unique estimate of the pressure can be retrieved from the ratio between the ppm flux differences of the 1.06 @xmath1 m dimer feature and the o@xmath0 a band . figure [ fig : pressure - tran - only]a shows the relationship between this ratio and the total amount of o@xmath0 above 0.9 bars , which is the highest pressure that can be probed in this particular case . there is a clear trend between this ratio and the amount of o@xmath0 in the atmosphere . when combined with an o@xmath0 mixing ratio this relationship can provide a quantitative estimate of a lower level of the surface or cloud - top pressure . the o@xmath0 mixing ratio can be estimated from the flux difference of the o@xmath0 a band in transmission . figure [ fig : pressure - tran - only]b shows the relationship between the o@xmath0 mixing ratio ( which is constant throughout the atmosphere ) and the o@xmath0 a band flux difference . for pressure @xmath390.1 bars the o@xmath0 flux difference is roughly constant for a given o@xmath0 concentration , meaning that a measurement of the o@xmath0 a band flux difference should correlate with the o@xmath0 mixing ratio . reflected spectra alone can provide an estimate of the surface partial pressure of o@xmath0 by examination of the ratio of the 1.06 @xmath1 m dimer equivalent width and the o@xmath0 a band equivalent width . figure [ fig : refl - only - toto2 ] shows the relationship between this ratio and the surface partial pressure of o@xmath0 . the strength of the 1.06 @xmath1 m dimer feature could also be used as an on / off gauge to set a lower bound for pressure . determining total atmospheric pressure with only a reflected spectrum is difficult due to degeneracies between the o@xmath0 concentration and total atmospheric pressure , as shown in figure [ fig : refl - only ] . for large equivalent widths of the dimer feature , the pressure will certainly be above 1 bar . however , it is difficult to differentiate between atmospheres with the same o@xmath0 surface partial pressure . nevertheless , it appears that it is possible to set a lower limit on pressure by measuring the 1.06 @xmath1 m dimer feature equivalent width . for example , a 1.06 @xmath1 m dimer feature equivalent width greater than @xmath1410 nm would imply a surface pressure @xmath21 bar . if both a transit transmission spectrum and a reflected spectrum are available , it should be possible to directly measure the total atmospheric surface pressure in the absence of clouds . transit transmission spectroscopy can provide an estimate of the o@xmath0 mixing ratio as described previously . a reflected spectrum can theoretically probe to the reflecting surface and therefore can be used to constrain the o@xmath0 partial pressure at the surface . by combining the o@xmath0 mixing ratio and o@xmath0 partial pressure we can determine the total pressure at the reflecting surface , which could either be a reflective cloud layer or the physical surface of the planet . the methods described here could be used in the near future by the jwst nirspec instrument , which will potentially be able to characterize transiting planets between 0.6 and 5.0 @xmath1 m . the o@xmath0 a band will likely not be detectable for a nearby earth - analog with jwst . although this feature is strong in the spectrum , the sensitivity of nirspec is poor at shorter wavelengths . the 1.06 @xmath1 m dimer feature is detectable at the 3@xmath8 level in transit transmission for cases with 2.0x pal o@xmath0 and high surface pressures . thus , the detection of the 1.06 @xmath1 m dimer feature would imply a surface or cloud - top pressure greater than or equal to 1.0 bars . for cases in which the 1.06 @xmath1 m dimer feature is not detectable , the 1.27 @xmath1 m feature could be used to constrain the pressure . this feature is not as strongly dependent on pressure as the 1.06 @xmath1 m dimer feature , but it is more detectable in all cases explored here . tpf , darwin or a similar direct imaging mission will be required to characterize the reflected spectra of nearby earth analogs in the visible and near - infrared . the snr values for secondary eclipse using jwst are all less than 1 , so jwst will not be able to characterize the reflected spectra of earth - analogs in secondary eclipse . table [ tab : tpf ] shows the necessary snrs to detect and characterize spectral features for a direct imaging planet characterization mission . while the snrs were calculated for an earth analog orbiting an m5v star , the results should be largely independent of spectral type because we have divided the reflected flux by the stellar flux in our calculations . the required snr values suggest that a snr of @xmath210 would be necessary to detect and quantify the o@xmath0 a band , 1.06 @xmath1 m dimer feature and 1.27 @xmath1 m feature for a true earth analog . however , because continuum brightness changes with pressure , a different snr criteria would be necessary for higher pressure atmospheres . for example , the continuum brightness near the o@xmath0 a band is three times lower for the 10.0 bar cases than it is for the 0.1 bar cases . for most cases , a snr of @xmath27 would likely be sufficient to set a lower limit on the surface pressure using the 1.06 @xmath1 m dimer feature . clouds and aerosols will also affect the detectability of absorption features . in transit transmission , clouds can effectively mask the highest pressures of the atmospheres at which dimer absorption is most prominent . however , in partially cloudy atmospheres some of the paths will probe pressures as high as the maximum tangent pressure , and thus the planetary transmission spectrum could show evidence of dimer absorption . furthermore , absorption in the 1.27 @xmath1 m dimer band can be detected with snr@xmath21 in even some 0.1 bar cases , meaning that there could be a detectable dimer absorption signal for even a completely cloud - covered planet if the cloud deck pressure was @xmath370.1 bars . for reflected spectra , clouds will truncate paths before they reach the surface and limit the dimer absorption for those paths . however , for partially cloudy atmospheres , dimer absorption could be detectable in the paths that do reach the surface . additionally , because cloud albedos are typically greater than surface albedos , the presence of clouds will increase the continuum brightness levels and the brightness in the center of absorption bands . the increase in brightness has been shown to decrease the required snr to detect and characterize o@xmath0 monomer absorption in cloudy atmospheres when compared to cloud - free cases @xcite , though the effect of cloud albedo on the detectability of dimer absorption features has not been heretofore examined . therefore , while clouds will impact the detectability of dimer features , using dimers to determine pressure and as biosignatures may still be feasible for cloudy atmospheres . figures [ fig : specres2 ] shows the snrs for spectral features at varying resolving powers for transit transmission spectra of a 1.0 bar , 1.0x pal o@xmath0 atmosphere . the snrs for each band are greatest at the highest resolving powers , and then gradually decrease until r@xmath1460 or 80 , at which the snrs decrease strongly . this dramatic decrease with resolving power occurs because at the lowest resolving powers , the absorption bands are indistinguishable from the continuum . additionally , the highest flux levels in the continuum can not be resolved at lower spectral resolving powers , decreasing the total signal . this effect can be seen most easily for the spectra with r=20 , in which no absorption features can be identified . figures [ fig : specres - imaging - snrd ] and [ fig : specres - imaging - snrp ] show the snrs for the direct imaging reflected spectra . in contrast to figure [ fig : specres2 ] , these two figures show the required snr to detect and characterize an absorption band , not the snr that could be obtained with jwst . the required snr to detect spectral features increases as resolving power decreases . however , this effect would be mitigated because the expected noise at each wavelength should decrease as resolving power decreases . the snr required to quantify each absorption band decreases as resolving power decreases because the lowest radiance level increases . at r@xmath3840 , however , spectral features are very difficult to identify , making these resolving powers unsuitable for detecting and characterizing o@xmath0-related absorption features . in addition to their utility as pressure probes , the 1.06 @xmath1 m dimer feature and 1.27 @xmath1 m feature could potentially be detectable biosignatures for nearby earth - like planets . the o@xmath0 a band has long been considered the most viable o@xmath0 biosignature , but it is unlikely to be the most detectable biosignature for an earth - like planet in transit transmission . as initially described in @xcite , lunar eclipse observations show that the 1.06 @xmath1 m and 1.27 @xmath1 m dimer features are more detectable than o@xmath0 monomer features like the a band , which is corroborated by our model spectra and detectability calculations . the 1.27 @xmath1 m o@xmath0 feature has been examined as a potential biosignature for ground - based telescopes by @xcite , but to our knowledge detectability studies of neither the 1.06 @xmath1 m dimer feature nor the 1.27 @xmath1 m feature have been undertaken for jwst . our results show that the 1.27 @xmath1 m feature would be detectable with a snr of 5 for a cloud - free earth - analog at 5 pc . therefore , we conclude that o@xmath0 features , especially the 1.06 @xmath1 m dimer feature and the 1.27 @xmath1 m feature , could be detectable biosignatures for oxygenic photosynthesis with jwst . clouds and aerosols will make estimating pressure using dimer features more difficult . a direct imaging observation of a partially cloud - covered planet will be able to probe to the surface for a fraction of the paths , such that the dimer feature will be weaker than for a cloud - free planet . cloud and aerosol extinction can also be wavelength dependent , which may complicate using equivalent widths or ppm flux differences to determine pressure . nevertheless dimer features can still provide a lower bound for pressure if clouds and aerosols can not be explicitly included in the retrieval method . higher h@xmath0o or co@xmath0 abundances in an atmosphere could also make this method more challenging . higher h@xmath0o abundances will make it more difficult to define a continuum for the 1.06 @xmath1 m dimer feature , as shown in figures [ fig : h2o - test-2 ] and [ fig : h2o - test-3 ] . however , as discussed in section [ sec : sensitivity ] , the magnitude of this error should typically be less than 20% for most cases in which the 1.06 @xmath1 m dimer feature could be detectable . co@xmath0 has an absorption feature near 1.06 @xmath1 m @xcite , which could make using the 1.06 @xmath1 m dimer feature difficult . this could be overcome by modeling out absorption from h@xmath0o and co@xmath0 or by using other dimer features to supplement information from the 1.06 @xmath1 m dimer feature , such as the 1.27 @xmath1 m dimer feature . lastly , not knowing the mixing ratio of o@xmath0 will make estimating pressure difficult when only a reflected spectrum is available . this is similar to the problem in using the absorption widths of rotation - vibration features to constrain pressure . however , the 1.06 @xmath1 m dimer feature is more sensitive to pressure than a monomer feature , and therefore can provide a better estimate of pressure than a monomer feature alone . furthermore , monomer features can not be used as an on / off pressure gauge , while dimer features can .

we also show that dimer features could be the most readily detectable biosignatures for earthlike atmospheres , and may even be detectable in transit transmission with the james webb space telescope ( jwst ) . we show that in some cases it may be possible to estimate the pressure at the reflecting surface of a planet . the detection of these features could provide a constraint on the atmospheric pressure of an exoplanet , and serve as biosignatures for oxygenic photosynthesis .

we present a new method to probe atmospheric pressure on earthlike planets using ( o-o ) dimers in the near - infrared . we also show that dimer features could be the most readily detectable biosignatures for earthlike atmospheres , and may even be detectable in transit transmission with the james webb space telescope ( jwst ) . the absorption by dimers changes more rapidly with pressure and density than that of monomers , and can therefore provide additional information about atmospheric pressures . by comparing the absorption strengths of rotational and vibrational features to the absorption strengths of dimer features , we show that in some cases it may be possible to estimate the pressure at the reflecting surface of a planet . this method is demonstrated by using the o a band and the 1.06 m dimer feature , either in transmission or reflected spectra . it works best for planets around m dwarfs with atmospheric pressures between 0.1 and 10 bars , and for o volume mixing ratios above 50% of earth s present day level . furthermore , unlike observations of rayleigh scattering , this method can be used at wavelengths longer than 0.6 m , and is therefore potentially applicable , although challenging , to near - term planet characterization missions such as jwst . we have also performed detectability studies for jwst transit transmission spectroscopy and find that the 1.06 m and 1.27 m dimer features could be detectable ( snr ) for an earth - analog orbiting an m5v star at a distance of 5 pc . the detection of these features could provide a constraint on the atmospheric pressure of an exoplanet , and serve as biosignatures for oxygenic photosynthesis . we have calculated the required signal - to - noise ratios to detect and characterize o monomer and dimer features in direct imaging reflected spectra and find that signal - to - noise ratios greater than 10 at a spectral resolving power of r=100 would be required .

astro-ph0003128

c

we have developed a simple model to compute the effects of dust on the integrated spectral properties of galaxies , based on an idealized prescription of the main features of the ism . our model includes the ionization of hii regions in the interiors of the dense clouds in which stars form . emission lines from hii regions and the non - ionizing continuum from young stars are attenuated in the same way by dust in the outer hi envelopes of the birth clouds and the ambient ism . however , since the model also includes the finite lifetimes of the birth clouds , the non - ionizing continuum radiation from stars that live longer than the birth clouds is attenuated only by the ambient ism . we show that this can fully resolve the apparent discrepancy between the attenuation of line and continuum photons in starburst galaxies . this enables us , in turn , to interpret in a consistent way all the observations of a homogeneous sample of nearby ultraviolet - selected starburst galaxies , including the ratio of far - infrared to ultraviolet luminosities ( @xmath151 ) , the ratio of h@xmath0 to h@xmath1 luminosities ( @xmath150 ) , the h@xmath0 equivalent width ( @xmath139 ) , and the ultraviolet spectral slope ( @xmath1 ) . the different parameters in our model , including the effective age of the starburst , the lifetime and effective optical depth of the stellar birth clouds , the effective optical depth in the ambient ism , and the fraction of dust in the ionized gas , each have a specific influence on the integrated spectral properties @xmath137 , @xmath270 , @xmath136 , and @xmath1 . this provides new insights into the origin of the mean relations defined by the data and the scatter about these relations . in particular , the relation between the ratio of far - infrared to ultraviolet luminosities and the ultraviolet spectral slope in starburst galaxies reflects the wavelength dependence of the effective absorption in the ambient ism . we find that a power law of the form @xmath271 accounts remarkably well for all the observations . the relation between @xmath137 and @xmath1 can then be interpreted as a sequence in the overall dust content of the galaxies . interestingly , this relation is accompanied by much weaker trends of the oxygen abundance and the optical luminosity with the ultraviolet spectral slope . the fact that our model reproduces the observed spectral properties of nearby starburst galaxies relatively easily leads us to suspect that , with suitable adjustment of the parameters , it could also reproduce those of more quiescent ( but still star - forming ) galaxies . the effective absorption curve required by the observations is much greyer than would be produced by a foreground screen of dust like that in the milky way , the lmc , or the smc . we have explored whether this could be accounted for by a much steeper wavelength dependence of the optical depth , i.e. @xmath75 with @xmath272 , combined with a more realistic spatial distribution of the dust . we find that a mixed slab model for the ambient ism can produce the required effective absorption curve for low dust content but can not explain the observations of starburst galaxies with very reddened ultraviolet spectra . in contrast , we show that a random distribution of discrete clouds provides a consistent interpretation of all the observed integrated spectral properties of starburst galaxies . while these results were anticipated in some previous studies , we have shown here for the first time how to reconcile them with the large h@xmath148h@xmath1 ratios and other observations . we also find that the optical depths of the clouds favored by our analysis are similar to those inferred from the statistics of stellar reddening in the milky way . the model we have developed for computing the absorption of starlight by dust in galaxies can be combined easily , by design , with any population synthesis model . the observed mean relations for starburst galaxies can also be reproduced by the following simple recipe : use an effective absorption curve proportional to @xmath3 to attenuate the line and continuum radiation from each stellar generation , and lower the normalization of the curve typically by a factor of 3 after @xmath4yr to account for the dispersal of the birth clouds . this recipe accounts at least as well as the one by calzetti et al . ( 1994 , and as modified by calzetti 1997 , 1999 ) for the effects of dust on the non - ionizing continuum radiation . in addition , it fully resolves the apparent discrepancy between the attenuation of line and continuum photons in starburst galaxies . we believe , therefore , that our model and the recipe derived from it provide simple yet versatile tools to interpret the integrated spectral properties of starburst and possibly other types of galaxies . in future work , we plan to apply them to the growing body of observations of high - redshift galaxies . we thank p. boiss , a. ferrara , and n. panagia for valuable discussions . s.c . appreciates the hospitality of the stsci , and s.m.f . that of the iap , during the course of several visits . this research was supported in part by the national science foundation through grant no . phy94 - 07194 to the institute for theoretical physics . bianchi , s. , ferrara , a. , & giovanardi , c. 1996 , apj , 465 , 127 blitz , l. , & shu , f. h. 1980 , apj , 238 , 148 bottorff , m. , lamothe , j. , momjian , e. , verner , e. , vinkovi , & ferland , g. 1998 , pasp , 110 , 1040 bruzual , a. g. , & charlot , s. 1993 , apj , 405 , 538 bruzual , a. g. , magris , g. , & calvet , n. 1988 , apj , 333 , 673 calzetti , d. 1997 , aj , 113 , 162 . 1999 , private communication calzetti , d. , bohlin , r. c. , gordon , k. d. , witt , a. n. , & bianchi , l. 1995 , apj , 446 , l97 calzetti , d. , kinney , a. l. , & storchi - bergmann , t. 1994 , apj , 429 , 582 . 1996 , apj , 458 , 132 caplan , j. , & deharveng , l. 1986 , a&a , 155 , 297 di bartolomeo , a. , barbaro , g. , & perinotto , m. 1995 , mnras , 277 , 1279 disney , m. , davies , j. , & phillipps , s. 1989 , mnras , 239 , 939 draine , b. t. , & lee , h. m. 1984 , apj , 285 , 89 fanelli , m. n. , oconnell , r. w. , & thuan , t. x. 1988 , apj , 334 , 665 ferland , g. j. 1996 , in hazy , a brief introduction to cloudy , univ . of kentucky dept . of physics & astronomy int . gordon , k. d. , calzetti , d. , & witt , a. n. 1997 , apj , 487 , 625 helou , g. , khan , i. r. , malek , l. , & boehmer , l. 1988 , apjs , 68 , 151 hummer , d. , & storey , p. 1992 , mnras , 254 , 277 humphreys , r. m. 1978 , apjs , 38 , 309 kim , s .- h . , martin , p. g. , & hendry , p. d. 1994 , apj , 422 , 164 kinney , a. l. , bohlin , r. c. , calzetti , d. , panagia , n. , & wyse , r. f. g. 1993 , apjs , 86 , 5 mcquade , k. , calzetti , d. , & kinney , a. l. 1995 , apjs , 97 , 331 meurer , g. r. , heckman , t. m. , & calzetti , d. 1999 , apj , 521 , 64 meurer , g. r. , heckman , t. m. , leitherer , c. , kinney , a. , robert , c. , & garnett , d. r. 1995 , aj , 110 , 2665 mezger , p. g. , & smith , l. f. 1977 , in iau symp . 75 , star formation , ed . t. de jong & a. maeder ( dordrecht : reidel ) , 133 natta , a. , & panagia , n. 1984 , apj , 287 , 228 oey , m. s. , & kennicutt , r. c. 1997 , mnras , 291 , 827 osterbrock , d. e. 1989 , astrophysics of gaseous nebulae and active galactic nuclei ( mill valley : university science books ) pei , y. c. 1992 , apj , 395 , 130 petrosian , v. , silk , j. , & field , g. b. 1972 , apj , 117 , l69 puxley , p. j. , & brand , p. w. j. l. 1994 , mnras , 266 , 431 rybicki , g. b. , & lightman , a. p. 1979 , radiative processes in astrophysics ( new york : wiley ) , 36 silva , l. , granato , g. l. , bressan , a. , & danese , l. 1998 , apj , 509 , 103 spitzer , l. 1978 , physical processes in the interstellar medium ( new york : wiley ) , 154 storchi - bergmann , t. , kinney , a. l. , & challis , p. 1995 , apjs , 98 , 103 witt , a. n. , & gordon , k. d. 1996 , apj , 463 , 681 witt , a. n. , thronson , h. a. , & capuano , j. m. 1992 , apj , 393 , 611

we present a new model to compute the effects of dust on the integrated spectral properties of galaxies , based on an idealized prescription of the main features of the interstellar medium ( ism ) . the model includes the ionization of hii regions in the interiors of the dense clouds in which stars form and the influence of the finite lifetime of these clouds on the absorption of radiation . this enables us to interpret simultaneously all the observations of a homogeneous sample of nearby ultraviolet - selected starburst galaxies , including the ratio of far - infrared to ultraviolet luminosities , the ratio of h to h luminosities , the h equivalent width , and the ultraviolet spectral slope . . a noteworthy outcome of our detailed analysis is that the observed mean relations for starburst galaxies can be closely approximated by the following simple recipe : use an effective absorption curve proportional to to attenuate the line and continuum radiation from each stellar generation , and lower the normalization of the curve typically by a factor of 3 afteryr to account for the dispersal of the birth clouds .

we present a new model to compute the effects of dust on the integrated spectral properties of galaxies , based on an idealized prescription of the main features of the interstellar medium ( ism ) . the model includes the ionization of hii regions in the interiors of the dense clouds in which stars form and the influence of the finite lifetime of these clouds on the absorption of radiation . we compute the production of emission lines and the absorption of continuum radiation in the hii regions and the subsequent transfer of line and continuum radiation in the surrounding hi regions and the ambient ism . this enables us to interpret simultaneously all the observations of a homogeneous sample of nearby ultraviolet - selected starburst galaxies , including the ratio of far - infrared to ultraviolet luminosities , the ratio of h to h luminosities , the h equivalent width , and the ultraviolet spectral slope . we show that the finite lifetime of stellar birth clouds is a key ingredient to resolve an apparent discrepancy between the attenuation of line and continuum photons in starburst galaxies . in addition , we find that an effective absorption curve proportional to reproduces the observed relation between the ratio of far - infrared to ultraviolet luminosities and the ultraviolet spectral slope . we interpret this relation most simply as a sequence in the overall dust content of the galaxies . the shallow wavelength dependence of the effective absorption curve is compatible with the steepness of known extinction curves if the dust has a patchy distribution . in particular , we find that a random distribution of discrete clouds with optical depths similar to those in the milky way provides a consistent interpretation of all the observations . a noteworthy outcome of our detailed analysis is that the observed mean relations for starburst galaxies can be closely approximated by the following simple recipe : use an effective absorption curve proportional to to attenuate the line and continuum radiation from each stellar generation , and lower the normalization of the curve typically by a factor of 3 afteryr to account for the dispersal of the birth clouds . this recipe or our full model for absorption can be incorporated easily into any population synthesis model .

gr-qc0301006

i

in this paper we introduce a numerical code designed to solve the einstein field equations for axisymmetric spacetimes . even though the predominant focus in numerical relativity in recent years has been to study situations of relevance to gravitational wave detection , and hence lacking symmetries , there are still numerous interesting problems , both physical and computational , that can be tackled with an axisymmetric code . the advantages of restricting the full 3d problem to axisymmetry ( 2d ) are that the complexity and number of equations are reduced , as are the computational requirements compared to solving a similar problem in 3d . prior numerical studies of axisymmetric spacetimes include head - on black hole collisions @xcite , collapse of rotating fluid stars @xcite , the evolution of collisionless particles applied to study the stability of star clusters @xcite and the validity of cosmic censorship @xcite , evolution of gravitational waves @xcite , black hole - matter - gravitational wave interactions @xcite , and the formation of black holes through gravitational wave collapse @xcite and corresponding critical behavior at the threshold of formation @xcite . our goals for creating a new axisymmetric code are not only to explore a wider range of phenomena than those studied before , but also to provide a framework to add adaptive mesh refinement ( amr ) and black hole excision to allow more thorough and detailed investigations than prior works . the outline of the rest of the paper is as follows . in section [ sec : formalism ] we describe the @xmath0 decomposition @xcite of spacetime that we adopt to arrive at our system of equations . the @xmath0 formalism is the familiar adm space + time decomposition ( in this case @xmath1 ) applied to a dimensionally reduced spacetime obtained by dividing out the axial killing vector , following a method devised by geroch @xcite . in section [ sec : coords_and_vars ] we specialize the equations to our chosen coordinate system , namely cylindrical coordinates with a conformally flat @xmath2metric . at this stage we do not model spacetimes with angular momentum , and we include a massless scalar field for the matter source . in section [ sec : apph ] we discuss how we search for apparent horizons during evolution . in section [ sec : implementation ] we describe our numerical implementation of the set of equations derived in section [ sec : coords_and_vars ] . a variety of tests of our code are presented in section [ sec : test ] , which is followed by conclusions in section [ sec : conclusion ] . some details concerning our finite difference approximations , solution of elliptic equations via the multi - grid technique , and a spherically symmetric code used for testing purposes are given in appendices a and b. unless otherwise specfied , we use the units and conventions adopted by misner , thorne and wheeler @xcite .

we present a new numerical code designed to solve the einstein field equations for axisymmetric spacetimes . the long term goal of this project is to construct a code that will be capable of studying many problems of interest in axisymmetry , including gravitational collapse , critical phenomena , investigations of cosmic censorship , and head - on black hole collisions .

we present a new numerical code designed to solve the einstein field equations for axisymmetric spacetimes . the long term goal of this project is to construct a code that will be capable of studying many problems of interest in axisymmetry , including gravitational collapse , critical phenomena , investigations of cosmic censorship , and head - on black hole collisions . our objective here is to detail the ( 2 + 1)+1 formalism we use to arrive at the corresponding system of equations and the numerical methods we use to solve them . we are able to obtain stable evolution , despite the singular nature of the coordinate system on the axis , by enforcing appropriate regularity conditions on all variables and by adding numerical dissipation to hyperbolic equations .

cond-mat9810403

c

we have presented a dynamic density functional approach for the relaxation of a classical system in terms of its equilibrium free energy density functional @xmath0 $ ] . the approach is valid only when the velocity correlations ( not included in @xmath0 $ ] ) are irrelevant , which excludes hydrodynamic modes and temperature gradients but may include problems like the brownian motion of colloidal particles and the molecular rearrangements of highly packed systems , when the collision time is much shorter than the relaxation time . in the derivation of the approach we start with the stochastic equations for the langevin dynamics of brownian particles and get a deterministic ddf equation for the time dependence of the density distribution , @xmath43 , which has to be interpreted as the ensemble average over the realizations of the random noise in the langevin dynamics . the equilibrium df formalism is used to go from a formal bbgky hierarchy , for the coupled dynamics of the n - particle distribution functions , to a closed ddf equation for @xmath43 . the assumption leading to this result is that the correlation structure in the system out of equilibrium is replaced by that in the equilibrium system with the same density distribution . moreover , we do not need to calculate the correlation structure explicitly because its effects on the dynamics are given directly in terms of the functional derivative of @xmath0 $ ] . the main advantage of the approach , for practical purposes , is being able to use the good approximations developed for the free energy functional of hard - core molecules to include the packing constraints in the dynamics of dense systems . we have presented several examples of one - dimensional hard rods to compare the ddf with the average of langevin simulations . the results , with the exact @xmath0 $ ] , are always in qualitatively good agreement , even in those cases in which the dynamics becomes extremely slow , when the relaxation requires very unlikely correlations between the particles . the main source of discrepancy is probably due to the fact that the free energy density functional developed for equilibrium always refers to the grand - canonical ensemble , so that the ddf includes the relaxation through changes in the number of particles , while the langevin dynamics keeps @xmath2 fixed . with the use of approximate density functionals , of the same type as developed for realistic models in three dimensions , we get an approximate description of the relaxation dynamics of a quality comparable to that for the equilibrium properties arising from the same @xmath0 $ ] . in some cases , when the role of the correlations is not too important and the relaxation times are not too long , the approximate @xmath0 $ ] is fairly accurate for the dynamic properties . in the third case , in which strong effects of the hard core packing lead to very long relaxation times , the difference between the results of the exact and the approximate @xmath0 $ ] are qualitative both for the equilibrium and for the dynamic properties . the presence of different local minima in the equilibrium free energy density functional produces , in the approximate ddf equation , the permanent freezing of system in any of these states , in contrast with the results of the exact free energy functional . in an interesting series of papers kawasaki @xcite derived a dynamic equation , in terms of a density functional hamiltonian @xmath155 $ ] , by a method different from ours , but his resulting fokker - planck equation for the probability distribution of the density : @xmath156}{\partial t}= -\int d{\bf r}\frac{\delta}{\delta \rho({\bf r})}{\bf \nabla}\cdot\rho({\bf r } ) { \bf \nabla } \bigl[t\frac{\delta}{\delta\rho({\bf r})}+\frac{\delta h[\rho]}{\delta \rho({\bf r})}\bigr ] p[\rho({\bf r}),t ] \label{eq : fpe}\ ] ] is equivalent to our eq . ( [ eq : df ] ) , apart from the presence of a term due to the stochastic noise . if @xmath57 in kawasaki equation ( [ eq : fpe ] ) is interpreted as a density operator , @xmath157 , his approach is equivalent to dean s equation ( [ eq : otto ] ) ; leading to a stochastic equation for this density operator . the main qualitative difference of this approach with our deterministic equation ( [ eq : df ] ) is that , when the approximate density functional for the equilibrium free energy has different local minima , the random noise in the ddf equation would always give a chance for changing from one minimum to another . the long time average of the density would always be a superposition of the density in the different local minima and the long relaxation times would appear as the result of high barriers between the local minima . all these features may appear to be physically correct and to represent a qualitative improvement over the deterministic ddf developed here . however , the use of the equilibrium free energy density functional requires always a density defined as a thermal ensemble average , while ( [ eq : otto ] ) refers to the instantaneous density operator . the hamiltonian @xmath158 in ( [ eq : fpe ] ) should also be a functional of the density operator , @xmath28 , and it is a completely different mathematical object that the excess free energy @xmath32 $ ] , as a functional of the equilibrium density @xmath159 . thus , equations ( [ eq : otto ] ) and ( [ eq : fpe ] ) are correct , for the density operator @xmath28 but impossible to translate in terms of the equilibrium free energy density functional . on the other hand , if these equations are interpreted as equations for the ensemble averaged density , @xmath43 , the random noise term would lead to a double counting of thermal fluctuations in the equation for @xmath43 . in particular , it would lead to a wrong equilibrium distribution , as can be verified in the simple case of non interacting particles for which the fpe ( [ eq : fpe ] ) converges to a probability distribution : @xmath160\sim \exp(-f[\rho]/k_b t)= \exp \left [ -\int d r \ \rho({\bf r } ) \ ( \ln \rho({\bf r})-1)+ \beta v_{ext}({\bf r } ) \right ] , \label{eq : false}\ ] ] where the first term in the exponential is clearly due to the overcounting of thermal fluctuations . moreover , we have shown that if we use the exact functional @xmath0 $ ] the ddf equation does not require random noise to give the correct results : the system always flows towards the true and unique equilibrium state . the ddf relaxation time may become very long when the system has to go through highly correlated states , but this effect corresponds to systems in which the true relaxation time ( in the langevin description ) is also very long . when we use an approximate @xmath0 $ ] these long relaxation times may become infinite , with the system trapped at a local minima of @xmath0 $ ] . an attempt to avoid this effect would require the intensity of random noise not to be proportional to the temperature , but to the error made by the approximation to @xmath0 $ ] , an error which is obviously unknown until we make a better approximation . otherwise , the relaxation time given by the stochastic ddf in these cases would just be a direct result of the uncontrolled level of noise kept in the functional equation . within our deterministic ddf , the existence of frozen states in the local minima of the approximate @xmath0 $ ] should be interpreted as the signature of very long relaxation times , but the only way to calculate how long these times are is to improve the approximation for the equilibrium free energy . nevertheless , knowing the existence and the approximate structure of these states , as given by workable approximations for @xmath0 $ ] , is already an interesting use of the ddf , together with its use to study the relaxation process in those cases in which there are no problems with different local minima . finally , as a plan for future work , we can consider a systematic way to improve the use of the equilibrium @xmath0 $ ] to estimate the correlation structure in the following terms : we tag particle number 1 and follow its position @xmath161 separately , while all the other particles ( @xmath162 ) are included in a density description , with @xmath163 as the noise - averaged conditional probability of finding a particle at position @xmath131 , and time @xmath29 , if the tagged particle is at position @xmath164 . now , we may consider that the @xmath165 particles are moving in an effective external potential , @xmath166 , which also contains the interaction with the tagged particle . the equivalent to the ddf equation ( [ eq : df ] ) may be applied to the conditional density @xmath167 , which is now coupled to the stochastic equation for @xmath161 . the advantage is that the correlation structure becomes partially described at the level of an effective one - particle density in an effective external potential , which is in principle much easier to describe with approximate free energy density functionals . this type of description would be similar to the reaction path description of chemical reactions , in which one , or a few , variables are used to describe the relevant functional directions for the changes in the molecular conformations . however , the formal and the practical use of density functional approximations along this line is still an open problem .

we present a new time - dependent density functional approach to study the relaxational dynamics of an assembly of interacting particles subject to thermal noise . starting from the langevin stochastic equations of motion for the velocities of the particles we are able by means of an approximated closure to derive a self - consistent deterministic equation for the temporal evolution of the average particle density . the closure is equivalent to assuming that the equal - time two - point correlation function out of equilibrium has the same properties as its equilibrium version . in particular the static solutions of the equation for the density correspond to the exact equilibrium profiles provided one is able to determine the exact form of $ ] . in order to assess the validity of our approach we performed a comparison between the langevin dynamics and the dynamic density functional method for a one - dimensional hard - rod system in three relevant cases and found remarkable agreement , with some interesting exceptions , which are discussed and explained .

we present a new time - dependent density functional approach to study the relaxational dynamics of an assembly of interacting particles subject to thermal noise . starting from the langevin stochastic equations of motion for the velocities of the particles we are able by means of an approximated closure to derive a self - consistent deterministic equation for the temporal evolution of the average particle density . the closure is equivalent to assuming that the equal - time two - point correlation function out of equilibrium has the same properties as its equilibrium version . the changes in time of the density depend on the functional derivatives of the grand canonical free energy functional $ ] of the system . in particular the static solutions of the equation for the density correspond to the exact equilibrium profiles provided one is able to determine the exact form of $ ] . in order to assess the validity of our approach we performed a comparison between the langevin dynamics and the dynamic density functional method for a one - dimensional hard - rod system in three relevant cases and found remarkable agreement , with some interesting exceptions , which are discussed and explained . in addition , we consider the case where one is forced to use an approximate form of $ ] . finally we compare the present method with the stochastic equation for the density proposed by other authors [ kawasaki , kirkpatrick etc . ] and discuss the role of the thermal fluctuations .

1309.0061

c

in this paper , we have presented the scattering - state description of the inelastic shot noise in a regime of a weak el - vib coupling and equilibrated vibrons . as discussed in the inelastic current , the inelastic shot noise is determined by the interplay of elastic and inelastic scattering processes . the elastic and inelastic contributions to the current noise are further decomposed into current correlations of electrons at the same energy and those of electrons at two energies that differ by the vibrational energy @xmath118 . applied to single - channel systems , our description enables to find out two ranges of transmission at which the crossover in the inelastic noise signal can take place . in particular , for mirror - symmetric systems , we have shown that even parity modes lead to the crossover at @xmath4 , while the crossover occurs at @xmath3 for odd parity ones . considering the ratio @xmath277 , we have confirmed that the ratio @xmath277 of the even parity mode is indeed an upper bound to ratios of general cases in a high transmission regime as speculated in ref . , and further we have predicted that the ratio @xmath277 of the odd parity mode is a lower bound . our scattering - state description is formulated for general situations involving many electronic states , many vibrational modes , and multiple conducting channels , so that it can be used to analyze first - principle calculation results , especially when specification of inter - channel and intra - channel scattering processes is crucial to understand the results .

we present a scattering theory description for the inelastic current noise in the presence of electron - vibration interactions . in this description , we specify elastic and inelastic scattering contributions to the shot noise by examining charge transfers between scattering states and energy exchange between electrons and vibrations . the elastic and inelastic scattering processes are further decomposed into current correlations of electrons at the same energy and those of electrons at different energies . focusing on the inelastic noise signals defined as steps in the voltage derivative of the shot noise , we show that single - channel systems have two ranges of transmission at which the inelastic noise signals exhibit the crossover between positive and negative signs . in a high transmission regime ,

we present a scattering theory description for the inelastic current noise in the presence of electron - vibration interactions . in this description , we specify elastic and inelastic scattering contributions to the shot noise by examining charge transfers between scattering states and energy exchange between electrons and vibrations . the elastic and inelastic scattering processes are further decomposed into current correlations of electrons at the same energy and those of electrons at different energies . focusing on the inelastic noise signals defined as steps in the voltage derivative of the shot noise , we show that single - channel systems have two ranges of transmission at which the inelastic noise signals exhibit the crossover between positive and negative signs . in a high transmission regime , even and odd vibrational modes of mirror - symmetric systems provide upper and lower bounds to the ratio of the inelastic noise signal to the conductance step . this can be a theoretical justification for models used to understand the recent noise experiment [ phys . rev . lett . * 108 * , 146602 ( 2012 ) ] and numerical calculations on gold atomic chains [ phys . rev . b * 86 * , 155411 ( 2012 ) ] .

cond-mat0601004

i

the matter wave beam splitters are nowadays the cornerstone of a wide range of experiments , from atomic clocks and gravito - inertial sensors to laser cooling and ultracold atoms characterization , quantum computing and cavity qed experiments , atom lithography and chemical reaction dynamics , detection of tiny effects of general relativity and test of fundamental theories , measurement of atom surface interactions ... in view of the recent progress in non - dissipative atom optics ( coherent beam splitters , mirrors , lenses ... ) as well as in dissipative atom optics ( slowing , trapping and cooling of atoms and molecules ) , it is needed to deepen our comprehension of light - matter interactions in the presence of other external potentials , like gravito - inertial or trapping potentials . in particular , the precision and stability of atom interferometers are now so outstanding @xcite that it is necessary to go beyond the former modeling of their main component , namely the beam splitters . in fact , the concept of atomic beam splitter is not confined to light - matter interactions , and can be extended to any interaction process between matter waves . it is thus possible to write the action of such atom optical elements as a s matrix between the incident and diffused matter waves , where the s matrix depends mainly on the splitting potential , which can be material ( slits , periodic microstructures ... ) or electromagnetic ( magnetic or electric static fields , laser fields ... ) . this s matrix description is particularly useful for the modeling of atom interferometers , and more generally for any set up having a succession of such beam splitters @xcite . however , for a long time , the precision of atom optics experiments has remained low enough not to require an accurate study of matter wave beam splitters . thus , in the most common simplified modeling of these elements , only the following effects were considered : 1 ) the splitting of an incident atomic wave packet into several wave packets , 2 ) among them one was equal to the incident wave packet , up to a change of amplitude , 3 ) and where the others could differ from the incident wave packet in their central momentum , internal state , amplitude and phase . however , this practical modeling - sometimes called infinitely thin because it amounts to neglecting the duration of the interaction - does not take into account several important effects , like the dispersive structuring of the incident wave packet ( velocity selection and sidebands , borrmann effect , anomalous dispersion ... ) , the time and space dependency of the splitting potential , the effect of relaxation processes , or the effect of other external fields during the splitting ( time - dependent gravito - inertial effects , trapping potentials ... ) . during the last past two decades , several authors studied some of these problems , namely : the effect of a non - trivial time dependency of the splitting potential ( for running and standing laser beam splitters ) @xcite ; the atomic borrmann effect and anomalous dispersion effect without any other external potential @xcite or with a constant and uniform acceleration ( wkb solution @xcite ) ; the atomic splitting in a constant and uniform acceleration ( exact solution in the temporal case and wkb solution in the spatial case ) @xcite ; a common modeling for both spatial and temporal beam splitters to the first order in the splitting potential ( weak field theory ) @xcite ... in the light of what happened in neutron optics , where the beam splitters modeling proved to be crucial to understand properly the origin of the interferometer phase shifts @xcite , it appears to be necessary to go beyond these studies , so as to provide a comprehensive modeling of the true action of a matter wave beam splitter ( strong fields theory for all the involved external fields ) . this paper is organized as follows . first , we give some details on our framework and explain how to put in equation the problem of the triple interaction matter - splitting potential - other external fields . then , we detail how to transform the obtained equation in a simpler one thanks to unitary transformations ( interaction picture ) and passage into the rotating frames . we expound then how to solve this equation ( analytically or numerically , with different developments or relevant approximations ) , and we go back to the initial representation to explain how to write the effect of such beam splitters as an effective instantaneous interaction ( generalized @xmath0 scheme ) . finally , we study the atomic borrmann effect and other anomalous dispersive properties and model them in the general framework detailed in the second part .

the dispersive structuring of an incident atomic wave packet - due to such generalized beam splitters - is studied and modeled , and several important dynamical features of the solutions are detailed ( generalized rabi oscillations , velocity selection , anomalous dispersion , generalized borrmann effect and anomalous gravitational bending ) .

we present a strong field theory of matter wave splitting in the presence of various gravitational , inertial and trapping potentials . the effect of these potentials on the resonance condition ( between the splitting potential and the considered effective two - level system ) and on the atomic borrmann effect is investigated in detail . the dispersive structuring of an incident atomic wave packet - due to such generalized beam splitters - is studied and modeled , and several important dynamical features of the solutions are detailed ( generalized rabi oscillations , velocity selection , anomalous dispersion , generalized borrmann effect and anomalous gravitational bending ) . finally , we show how to express this triple interaction matter - splitting potential - gravito - inertial and trapping potentials as an equivalent instantaneous interaction which turns out to be a very efficient tool for the modeling of atom interferometers . pacs number(s ) : 03.75.-b , 32.80.-t , 33.80.-b , 39.25.+k , 42.50.-p

cond-mat0601004

c

in conclusion , we have shown in this paper how to solve the problem of matter wave splitting in the presence of various gravitational , inertial and trapping potentials . in particular , we have seen how the resonance condition between the splitting potential and the effective two - level atoms has to be changed . then , we have shown how to express this triple interaction matter - splitting potential - other external potentials as an equivalent instantaneous interaction ( generalized @xmath0 scheme ) . finally , we have investigated in detail what is the dispersive structuring of an incident atomic wave packet inside such beam splitters , both in the free case ( for which @xmath12 is reduced to @xmath132 ) and in the general case of @xmath12 . several significant features of the solutions have been studied : group velocities , generalized rabi oscillations , velocity selection , anomalous dispersion effects ... in the light of this study , the generalized @xmath0 scheme leads to a very practical and efficient ( gaussian ) modeling of atomic beam splitters which is particularly relevant for atom interferometric signal calculations @xcite . however , several points still have to be cleared up : for instance , the problem of the effective mass change which occurs when the atomic internal state is changed , and which leads to non - trivial ( small ) relativistic corrections . furthermore , it is necessary to extend our formalism to additional external potentials which are more than quadratic ( in position and momentum ) if we want to investigate the effect of van - der - waals , casimir or yukawa - type potentials on the matter wave splitting . more generally speaking , it would be interesting to go beyond the various approximations listed in part [ part2 ] , and in particular beyond the two - beam approximation . t.l . gustavson , a. landragin and m.a . kasevich . , class . quantum grav . 17 , 2385 ( 2000 ) ; a. peters , k.y . chung and s. chu , metrologia 38 , 25 ( 2001 ) ; j.m . mcguirk , g.t . foster , j.b . fixler , m.j . snadden , and m.a . kasevich , phys . rev . a 65 , 033608 ( 2002 ) ; g. wilpers , t. binnewies , c. degenhardt , u. sterr , j. helmcke , and f. riehle , phys . 89 , 230801 ( 2002 ) . b. eiermann , p. treutlein , th . anker , m. albiez , m. taglieber , k .- marzlin and m. k. oberthaler , phys . 91 , 060402 ( 2003 ) ; b. eiermann , th . anker , m. albiez , m. taglieber , p. treutlein , k .- marzlin and m. k. oberthaler , phys . 92 , 230401 ( 2004 ) . j. robert , ch . miniatura , s. le boiteux , j. reinhardt , v. bocvarski , j. baudon , j. europhys . , 16 , 29 ( 1991 ) ; ch . miniatura , j. robert , s. le boiteux , j. reinhardt and j. baudon , appl . b 54 , 347 ( 1992 ) . g.i . opat , s.j . wark , and a. cimmino , appl . b 54 , 396 ( 1992 ) ; t.m . roach , h. abele , m.g . boshier , h.l . grossman , k.p . zetie and e.a . hinds , phys . 75 , 629 ( 1995 ) ; a.i . sidorov , r.j . mclean , g.i . opat , w.j . rowlands , d.c . lau , j.e . murphy , m. walkiewicz and p. hannaford , quantum semiclass . opt . 8 , 713 ( 1996 ) . e. arimondo , h. lew and t. oka , phys . 43 , 753 ( 1979 ) ; p.e . moscowitz , p.l . gould , s.r . atlas and d.e . pritchard , phys . 51 , 370 ( 1983 ) ; e. m. rasel , m.k . oberthaler , h. batelaan , j. schmiedmayer and a. zeilinger , phys . rev . lett . 75 , 2633 ( 1995 ) ; d. m. giltner , r.w . mcgowan and s.a . lee , phys . 75 , 2638 ( 1995 ) ; m. kozuma , l. deng , e.w . hagley , j. wen , r. lutwak , k. helmerson , s.l . rolston and w.d . phillips , phys . 82 , 871 ( 1999 ) ; s. wu , y .- j . wang , q. diot and m. prentiss , phys . a 71 , 043602 ( 2005 ) . baklanov , b.y . dubetsky and v.p . chebotayev , appl . phys . a 9 , 171 ( 1976 ) ; ch.j . bord , ch . salomon , s. avrillier , a. van lerberghe , ch . brant , d. bassi and g. scoles , phys . a 30 , 1836 ( 1984 ) . cook and r.k . hill , opt . 43 , 258 ( 1982 ) ; j.v . hajnal and g.i . opat , opt . 71 , 119 ( 1989 ) ; m. christ , a. scholz , m. sciffer , r. deutschmann and w. ertmer , opt . 107 , 211 ( 1994 ) ; a. steane , p. szriftgiser , p. desbiolles and j. dalibard , phys . 74 , 4972 ( 1995 ) ; r. brouri , r. asimov , m. gorlicki , s. feron , j. reinhardt , v. lorent and h. haberland , opt . 124 , 448 ( 1996 ) ; a. landragin , g. labeyrie , c. henkel , r. kaiser , n. vansteenkiste , c. i. westbrook and a. aspect , opt . 21 , 1591 ( 1996 ) ; c. henkel , k. molmer , r. kaiser , n. vansteenkiste , c.i . westbrook and a. aspect , phys . a 55 , 1160 ( 1997 ) . j. oreg , f.t . hioe and j. h. eberly , phys . a 29 , 690 ( 1984 ) ; p. marte , p. zoller and j. l. hall , phys . a 44 , r4418 ( 1991 ) ; p. pillet , c. valentin , r .- yuan and j. yu , phys . rev . a 48 , 845 ( 1993 ) ; j. lawall and m. prentiss , phys . 72 , 993 ( 1994 ) ; p.d . featonby , g.s . summy , j.l . martin , h. wu , k.p . zetie , c.j . foot and k. burnett , phys . a 53 , 373 ( 1996 ) . letokhov and v.g . minogin , zh . 74 , 1318 ( 1978 ) ; y. castin and j. dalibard , europhys . 14 , 761 ( 1991 ) ; m. bchner , r. delhuille , a. miffre , c. robilliard , j. vigu and c. champenois , phys . rev . a 68 , 013607 ( 2003 )

we present a strong field theory of matter wave splitting in the presence of various gravitational , inertial and trapping potentials . the effect of these potentials on the resonance condition ( between the splitting potential and the considered effective two - level system ) and on the atomic borrmann effect is investigated in detail . finally , we show how to express this triple interaction matter - splitting potential - gravito - inertial and trapping potentials as an equivalent instantaneous interaction which turns out to be a very efficient tool for the modeling of atom interferometers .

we present a strong field theory of matter wave splitting in the presence of various gravitational , inertial and trapping potentials . the effect of these potentials on the resonance condition ( between the splitting potential and the considered effective two - level system ) and on the atomic borrmann effect is investigated in detail . the dispersive structuring of an incident atomic wave packet - due to such generalized beam splitters - is studied and modeled , and several important dynamical features of the solutions are detailed ( generalized rabi oscillations , velocity selection , anomalous dispersion , generalized borrmann effect and anomalous gravitational bending ) . finally , we show how to express this triple interaction matter - splitting potential - gravito - inertial and trapping potentials as an equivalent instantaneous interaction which turns out to be a very efficient tool for the modeling of atom interferometers . pacs number(s ) : 03.75.-b , 32.80.-t , 33.80.-b , 39.25.+k , 42.50.-p

0902.0004

i

recent studies of galaxy evolution have invoked a significant component of energetic feedback from an active galactic nucleus ( agn ) to attempt to explain many of the observed properties of galaxies . specifically , agn feedback has been implicated as the primary cause behind the differences between observed and theoretical galaxy luminosity functions , the observed color bimodality of galaxies @xcite in which galaxies are divided into distinct populations of blue star - forming galaxies and red dead galaxies ( e.g. * ? ? ? * ) with a pronounced deficit of galaxies with intermediate color , and the very tight correlations observed between the masses of supermassive black holes and their host galaxy spheroid s velocity dispersion @xcite and luminosity @xcite . the relative lack of galaxies with intermediate star formation rates suggests that star formation ceases abruptly rather than gradually @xcite , while the tight correlations of black hole masses with the properties of their host bulges are taken as evidence that the two are strongly co - eval . in the various proposed feedback scenarios , the supermassive black holes grow primarily by accretion until they become sufficiently massive and energetic that thermal and/or radiative feedback from their activity either heats the surrounding interstellar gas or mechanically pushes it out of the host galaxy @xcite . this simultaneously deprives the host galaxy of raw material from which to form new stars @xcite and starves the black hole , truncating its growth @xcite . despite the considerable utility of agn feedback , there is scant observational data to inform us about how or even if agn feedback operates in actual galaxies . arguably some of the most promising cases are instances of mechanical feedback , such as the outflows that have been detected as blue - shifted uv and x - ray intrinsic absorbers in @xmath050% of seyfert 1 galaxies and quasars @xcite . while it is clear that nuclear outflows are common , it is still difficult to determine the mechanical luminosity ( critical for determining if the outflow can unbind circumnuclear gas ) and mass outflow rate to determine their impact . a good example of this difficulty is the work of @xcite on the time evolution of the ionization state of the x - ray absorbers relative to the x - ray continuum in the warm absorber of ngc4051 . from the absorption line variability they obtained good constraints on the density and location of the absorbers from the black hole . assuming a biconical geometry , they measured a low outflow velocity with respect to the escape speed from the black hole , and a corresponding low mass outflow rate with respect to the accretion rate of the black hole . from these values , they conclude that while the outflows might disrupt the hot ism , they are not capable of ejecting large amounts of interstellar gas from the host . @xcite reached a similar conclusion for the nearby seyfert galaxy markarian78 using visible - wavelength and radio data to measure the kinematics and ionization state of extended emission - line regions surrounding the active nucleus . they concluded that the biconical outflow in mrk78 is weak , slow and heavy , and hence insufficiently energetic to significantly heat or unbind surrounding host galaxy gas . the results for these two galaxies present an apparent problem for agn feedback : in both cases the measured agn outflows are too weak to thermally or mechanically disrupt star formation in the host galaxy . nevertheless , these are just two objects , and analyses of other agns are critical for further understanding the role of agn feedback in galaxy evolution . seyfert galaxies are of special interest for detailed studies of agn feedback as many are luminous enough to be accreting near their maximal eddington rate , yet close enough to study at relatively high angular resolution ( scales of 10100pc for nearby examples ) using the _ hubble space telescope _ and ground - based adaptive optics . a very promising candidate for detailed study is the nearby seyfert galaxy markarian573 , which is well - known for its extended , richly structured circumnuclear emission - line regions . mrk573 features a prominent ionization bicone and bright arcs and knots of emission - line gas @xcite that are strongly aligned and interacting with a kiloparsec - scale low - power radio outflow @xcite . the association between the radio lobes and the emission regions was studied in detail by @xcite , who proposed that the arcs result from gas cooling after passing through a radio - induced radiative bow shock . in contrast , @xcite argued that the arcs are instead a morphological artifact of dust lanes in the galaxy being illuminated by the nuclear ionizing continuum . this is supported by visible and near - infrared _ hst _ images that show the emission - line arcs as extensions of larger circumnuclear dust lanes that are illuminated when they pass into the ionization cone . circumnuclear dust lanes like these are commonly found in active and inactive galaxies alike @xcite , and in the specific case of mrk573 they appear to be shaped by the presence of a nuclear bar @xcite , not unlike what is seen in galaxies with inner bars regardless of the presence of nuclear activity . @xcite examined several possibilities for the origin of the arcs and extended emission - line regions using a combination of _ hst _ images and ground - based , integral - field spectroscopy . they modeled the excitation in the arcs as either shock features , linked to radio jets , or pre - existing structures photoionized by the nucleus . using standard emission - line diagnostics , they showed that the inner arcs are excited by the central continuum source rather than by fast , photoionizing shocks , citing as evidence that they detected no signs of a strong kinematic interaction through a radio - induced shock . the excitation of the outer arcs was not completely explained by nuclear photoionization ; these arcs require an external source of photons in addition to the nucleus to account for the excitation levels . however , their analysis suggested that both the inner and outer arcs show little evidence that they are shock excited . in this paper we present new long - slit spectrophotometry of mrk573 obtained with the hubble space telescope imaging spectrograph ( stis ) . the high angular resolution of stis , combined with good velocity resolution in the bright h@xmath1 + [ ] emission lines allows us to examine the excitation and kinematics in these regions in greater detail than possible in previous studies . our analysis of mrk573 follows a threefold approach . first , we examine the kinematics of the circumnuclear emission - line regions to separate kinematically - disturbed gas from quiescent gas in the rotating disk of the galaxy ( see sec.[sec : los_vel ] ) . the existence of these two kinematical components was clear from the earlier ground - based work , but the greater angular resolution provided by _ hubble _ lets us pinpoint the regions of disturbed and undisturbed gas on a scale of 10s of parsecs . next we measure the densities and temperatures in the emission - line regions using standard nebular diagnostic emission lines . we examine the excitation state of these regions , comparing them with the predictions of published photoionization and shock excitation models to clarify the nature of the ionization mechanism in the bright emission line areas ( see sec.[sec : spectrophot ] ) . finally , we undertake a quantitative analysis of the outflow energetics of mrk573 using techniques developed by @xcite to examine the degree of feedback on the host galaxy surroundings ( see sec.[sect : whittle_quant ] ) . throughout this paper we adopt a distance of 74mpc for mrk573 ( @xmath2 km s@xmath3 , @xmath4 km s@xmath3mpc@xmath3 ) , giving us a projected linear scale of @xmath0360pcarcsec@xmath3 . for reference , this is approximately half the distance of mrk78 ( @xmath5 km s@xmath3 ) .

we present a study of outflow and feedback in the well - known seyfert 2 galaxy markarian 573 using high angular resolution long - slit spectrophotometry obtained with the _ hubble space telescope _ imaging spectrograph ( stis ) . these results are at odds with the picture of strong agn feedback that has been invoked to explain certain aspects of galaxy evolution .

we present a study of outflow and feedback in the well - known seyfert 2 galaxy markarian 573 using high angular resolution long - slit spectrophotometry obtained with the _ hubble space telescope _ imaging spectrograph ( stis ) . through analysis of the kinematics and ionization state of a biconical outflow region emanating from the nucleus , we find that the outflow does not significantly accelerate the surrounding host - galaxy interstellar gas and is too weak to be a strong ionization mechanism in the extended emission regions . instead , the excitation of the extended regions is consistent with photoionization by the active nucleus . from energetics arguments we show that the nuclear outflow is slow and heavy and has a mechanical luminosity that is only% of the estimated bolometric luminosity of the system . the energy in the outflow is able to mildly shape the gas in the extended regions but appears to be insufficient to unbind it , or even to plausibly disrupt star formation . these results are at odds with the picture of strong agn feedback that has been invoked to explain certain aspects of galaxy evolution .

1207.0767

i

the phase space of the galaxy s disk , as observed near the sun , contains substructure within the larger structures known as the thin and the thick disk . substructures include open clusters and stellar streams , also often referred to as moving groups . moving groups , as promoted by eggen ( see e.g. , @xcite and references therein ) , are considered to be stars having a common space motion and a common chemical composition , and originating from a disssolving open cluster . while some of the well known moving groups do share a common composition ( e.g. , the hr 1614 group - @xcite ) suggesting they come from a tidally - disrupted open cluster , other groups contain stars having very different compositions , a result demanding an origin more complex than disruption of an open cluster . the term ` stream ' is now applied to some entities in galactic phase space . examples in the thin disk include the hercules stream @xcite and the hyades stream or supercluster ( @xcite , @xcite ) . for streams in the thin disk , a likely explanation involves dynamical interactions of disk stars with the central bar @xcite or spiral density waves @xcite . other possibilities arise for streams belonging to the thick disk including accretion from external galaxies . in this paper , we present chemical compositions for subdwarfs belonging to two streams in the thick disk . @xcite undertook a search for fine structure in the phase space populated by subdwarfs from the large sample of f and g subdwarfs considered by @xcite for which @xcite refined data on stellar distances and kinematics . two clumps in phase space were noted by @xcite . one with @xmath1 @xmath2 km s@xmath3 and @xmath4 km s@xmath3 is referred to as the arcturus stream . an arcturus moving group had been previously identified by @xcite . the second stream af06 with a stronger presence in phase space than the arcturus stream is at @xmath5 km s@xmath3 and @xmath6 km s@xmath3 . here , we report chemical compositions of stars from these two streams . we show that stars in both streams span a range in metallicity but with relative abundances which match closely the ratios reported for field thick disk stars . this result serves to constrain greatly explanations for the origins of the two streams .

results show both streams are metal - poor and very old ( 10 gyrs ) with kinematics and abundances overlapping with the properties of local field thick disk stars . both streams exhibit a range in metallicity but with relative elemental abundances that are identical to those of thick disk stars of the same metallicity . these results show that neither stream can result from dissolution of an open cluster .

we present abundances for 20 elements for stars in two stellar streams identified by : 18 stars from the arcturus stream and 26 from a new stream , which we call af06 stream , both from the galactic thick disk . results show both streams are metal - poor and very old ( 10 gyrs ) with kinematics and abundances overlapping with the properties of local field thick disk stars . both streams exhibit a range in metallicity but with relative elemental abundances that are identical to those of thick disk stars of the same metallicity . these results show that neither stream can result from dissolution of an open cluster . it is highly unlikely that either stream represents tidal debris from an accreted satellite galaxy . both streams most probably owe their origin to dynamical perturbations within the galaxy . [ firstpage ] stars : abundances stars : moving groups galaxy : kinematics and dynamics galaxy : disk

1207.0767

r

since their discovery ( cf . @xcite ) , various theories have been put forward to explain the substructures in phase space known variously as moving groups or stellar streams . the most prominent theories identify a stellar stream as either a ) a dissolved open cluster , b ) debris from an accreted satellite galaxy , or c ) the result of dynamical perturbations within the galaxy . in establishing the correct explanation for a particular stream , the importance of the chemical signatures or chemical tagging of stream members has been recognized ( @xcite , @xcite , @xcite , @xcite ) not only to pinpoint their origin but also to understand the formation and evolution of the disk . quantitative abundances of twenty elements has been extracted here for members of the two streams . there are three main groups of elements : iron peak ( v , cr , mn , fe , ni , co ) , @xmath49-elements ( o , mg , si , ca , ti ) , and heavy elements such as @xmath50-process ( ba , nd ) and @xmath51-process ( eu ) elements . for the present ( thin ) disk , iron peak elements mostly come from type ia supernovae explosions ( snia ) and the @xmath49-elements are primarily produced in type ii supernovae ( snii ) . the @xmath50-process elements are known to be produced mainly in evolved stars of low and intermediate mass ( 1 @xmath0 8m@xmath33 ) with the lighter @xmath50-process elements also contributed by snii . eu , an @xmath51-process element , is most proably produced in snii . therefore , an abundance ratio of products of snia and snii helps to track the chemical history of a stellar system . abundances of elements from o to eu in the form of [ x / fe ] ratios , where x is any element , are given in table 2 and 3 and shown in figures 5 , 6 , and 7 as the run of [ x / fe ] against [ fe / h ] . the metallicity distribution of the streams is shown in figure 8 for the arcturus and af06 streams . next , we discuss these abundances for the arcturus and af06 streams in the light of the proposed scenarios for the origin of stellar streams . dispersal of an open cluster as the origin of a stellar stream is one explanation from the suite of potential explanations that is directly testable from photometric and spectroscopic determinations of stellar compositions . given that open clusters exhibit chemical homogeneity , if a stream represents a dissolved cluster , one requires chemical homogeneity also among stream members . chemical homogeneity is not a characteristic of these two streams . results shown in figure 8 and table 2 show that stellar metallicities span a wide range for both streams . for the arcturus stream , [ fe / h ] runs from @xmath01.40 to @xmath00.37 . for af06 stream , the range is from @xmath52 to @xmath53 but the range is @xmath01.69 to @xmath00.17 for those 15 members with a high probability of belonging to the thick disk . for a strictly homogenous populations such as open and globular clusters the degree of chemical homogeneity is quite high and which is about 0.05 dex ( @xcite , @xcite ) for a number of elements . in the case of [ fe / h ] dispersions are found to be in the range of 0.02 @xmath0 0.1 and in some extreme cases dispersions are of the order of 0.2 dex @xcite . the wide range in metallicity clearly shows that the systems from which the stream members originated had a relatively long history with a multiple episodes of star formation . we have also inspected the sample for an evidence of a sub - group with chemical homegenity . results shown in figure 5 , 6 and 7 indicate existence of no such a group among arcturus stream sample , however , for the stream af06 , we find a hint of clustering of stars at [ fe / h ] = @xmath00.4 . about 8 stars ( out of total 26 ) show metallicity [ fe / h ] = @xmath00.4 within the dispersion of 0.04 dex . does it mean that part of the sample stars originated from the disrupted cluster ? probably not , this may be a manifestation of thick disk metallicity distribution which peaks at about @xmath00.6 dex , and to some extent due to a smaller sample size . thus , neither the arcturus nor the af06 stream , as shown by metallicity distribution of member stars of the two streams ( see figure 8) , is a dissolved open cluster . this conclusion about the arcturus stream was reached also by @xcite who selected 134 stream members by selection criteria different from those used by @xcite and applied to stellar catalogues other than that compiled by @xcite . their analyses of high - resolution spectra led to a [ fe / h ] range similar to that quoted above ( see also @xcite ) . among thick disk field stars in the solar neighborhood , there is a strikingly very small dispersion in elemental abundance ratios , i.e. , [ x / fe ] at a given [ fe / h ] . indeed , @xcite found the [ x / fe ] to be gaussian - like with a dispersion @xmath54 of less than 0.10 dex for the most of the elements except for v , y , and zr for which @xmath54 is slightly more than 0.1 dex . furthermore , such dispersions were uncorrected for measurement uncertainties so that the intrinsic or ` cosmic ' dispersion must be very small . in sharp contrast to the very similar [ x / fe ] ratios at a given [ fe / h ] for local field stars , element ratios reflecting different contributions from major processes of stellar nucleosynthesis do vary from stellar system to stellar system . for example , ratios in the galactic bulge are not uniformly identical at a given [ fe / h ] to those among local thick or thin disk field stars . similarly , ratios of @xmath49-elements ( [ @xmath49/fe ] ) at a given [ fe / h ] among stars of dwarf spheroidal galaxies differ appreciably from those of local stars and from galaxy to galaxy ( @xcite , @xcite , @xcite ) . this has been illustrtated in figure 9 where ratio of [ @xmath49-process / fe ] ( mean of @xmath49-process elements mg , si , ca and ti ) is compared with that of thick disk @xcite and a number of dwarf spheroidal galaxies for which data was taken from ( @xcite , @xcite ) . therefore , the chemical signatures or tags in the form of [ x / fe ] for those key elements from the major processes of nucleosynthesis may test some proposed origins for the streams . to address the question ` are these two streams chemically identical to thick disk field stars ? ' , we show plots of [ x / fe ] versus [ fe / h ] in figures 5 to 7 with field stars from @xcite shown as by grey symbols , and the arcturus and af06 stream members shown as black filled circles . to quantify possible systematic offsets between the field thick disk and stream members , we compute mean values and dispersions for the arcturus and af06 streams over the [ fe / h ] interval @xmath00.30 to @xmath01.0 . dispersions about the trends is a combination of both the cosmic scatter as well as errors associated with the model parameters . the values ( @xmath55 , @xmath56 ) are computed as the standard deviation of the residuals from straight line fits to abundance trends of [ x / fe ] against [ fe / h ] . similarly , @xmath57 for the thick disk abundance trends of @xcite are computed over the same interval . all results are provided in table 7 . in figure 10 , we made comparison of @xmath57 with those of arcturus and af06 streams . dispersion values about the abundance trends of arcturus stream show very good agreement with that of the thick disk within about 0.02 . however , for af06 stream , in most cases , dispersions are lower by 0.01 @xmath0 0.03 compared to thick disk dispersions . with the current limited sample it would be far fetching to attribute this to the different chemical evolution for the af06 stream and hence external origin to it . mean abundances of arcturus and af06 streams are quite similar and the differences are within 0.05 dex except three elements ( o , cu , eu ) for which difference is 0.06 - 0.07 dex . dispersions about the trends for all the elements are comparable to the estimated scatter ( @xmath58 ) due to uncertainties in model parameters . diferences between dispersions and @xmath58 are within 0.05 dex . thus , we conclude that the streams are chemically identical to within high precision with the field thick disk stars . given that external galaxies of the local group have different element ratios across the [ fe / h ] range sampled by these two streams , especially for @xmath49-elements ( see , for example , @xcite , figure 11 ) , it seems most unlikely that either stream represents debris from an accreted satellite galaxy .

we present abundances for 20 elements for stars in two stellar streams identified by : 18 stars from the arcturus stream and 26 from a new stream , which we call af06 stream , both from the galactic thick disk . it is highly unlikely that either stream represents tidal debris from an accreted satellite galaxy . both streams most probably owe their origin to dynamical perturbations within the galaxy .

we present abundances for 20 elements for stars in two stellar streams identified by : 18 stars from the arcturus stream and 26 from a new stream , which we call af06 stream , both from the galactic thick disk . results show both streams are metal - poor and very old ( 10 gyrs ) with kinematics and abundances overlapping with the properties of local field thick disk stars . both streams exhibit a range in metallicity but with relative elemental abundances that are identical to those of thick disk stars of the same metallicity . these results show that neither stream can result from dissolution of an open cluster . it is highly unlikely that either stream represents tidal debris from an accreted satellite galaxy . both streams most probably owe their origin to dynamical perturbations within the galaxy . [ firstpage ] stars : abundances stars : moving groups galaxy : kinematics and dynamics galaxy : disk

1307.1830

r

first we present the statistical result concerning the distribution of leading and trailing modes of galaxies in the total sample and subsamples . then , we study the distribution of the leading and trailing arm galaxies in the unit area of the sky and the groups . the equatorial pa - distribution of galaxies in the total sample and subsamples is presented . at the end , a general discussion and a comparison with the previous results will be presented . a statistical comparison between the total sample and subsamples of the leading and trailing arm field galaxies is given in table 1 . 2 shows this comparison in the histogram . the @xmath14(% ) in table 1 and fig . 2 represent the percentage difference between the number of trailing and leading arm galaxies . we studied the standard deviation ( @xmath39 ) of the major diameters ( @xmath17 ) of galaxies in the total sample and subsamples for both the leading and trailing modes . in table 1 , @xmath40 represents the difference between the standard deviation of the major diameters of leading and the trailing arm galaxies . an insignificant difference ( 0.4% @xmath1 0.2% ) between the total number of trailing and the leading arm galaxies are found ( table 1 ) . the difference between the standard deviation of the major diameters ( @xmath40 ) of the trailing and leading arm galaxies is found less than 0.019 ( seventh column , table 1 ) . interestingly , the sum of the major diameters of total trailing and leading arm galaxies coincide . this result strongly suggests the existence of chirality of field galaxies having rv 3000 km s@xmath0 to 4000 km s@xmath0 . in fig . 2 , the slanting - line ( grey - shaded ) region corrosponds to the region showing @xmath18 10% ( 5% ) @xmath14 value . almost all subsamples lie within this region , suggesting the existence of chirality within 10% error limit . we present the distribution of the subsamples of leading and trailing arm galaxies classified according as their morphology , rvs , area and the groups below . in the spirals , leading structural modes are found 3.7% @xmath1 1.8% more than that of trailing modes . the chirality is found stronger for the late - type spirals ( sc , scd , sd and sm ) than that of early - type ( sa , sab , sb and sbc ) : @xmath14 value turned out to be 9.5% ( @xmath1 4.8% ) and 1.8% ( @xmath1 1.0% ) for early- and late - types ( table 1 ) . thus , the late - type spirals are the best candidate of the chiral object in our database . ( % ) = @xmath41-@xmath42 , where @xmath41 and @xmath42 represent the number of trailing and leading arm galaxies , respectively . the statistical error bars @xmath5(% ) shown in the figure are calculated as : @xmath5(% ) = @xmath5/(@xmath43+@xmath44)@xmath45100 , where @xmath5 = ( @xmath43-@xmath44 ) . the grey - shaded and the slanting - line region represent the @xmath18@xmath15% and @xmath18@xmath110% @xmath14 value , respectively . , title="fig:",height=166 ] ( % ) = @xmath41-@xmath42 , where @xmath41 and @xmath42 represent the number of trailing and leading arm galaxies , respectively . the statistical error bars @xmath5(% ) shown in the figure are calculated as : @xmath5(% ) = @xmath5/(@xmath43+@xmath44)@xmath45100 , where @xmath5 = ( @xmath43-@xmath44 ) . the grey - shaded and the slanting - line region represent the @xmath18@xmath15% and @xmath18@xmath110% @xmath14 value , respectively . , title="fig:",height=166 ] @xmath46 p0.23rccrrrrrr sample / subsample & @xmath47 & @xmath42 & @xmath41 & @xmath48 & @xmath46@xmath49 & @xmath46@xmath49 & @xmath46@xmath50a@xmath51 & @xmath46@xmath50a@xmath52 + total & @xmath41 & 814 & 807 & 7 & 0.4 & 0.2 & 0.019 & 0.0 + spiral & @xmath53 & 395 & 367 & -28 & -3.7 & -1.8 & 0.031 & 3.0 + spiral ( early - type ) & @xmath54 & 150 & 124 & -26 & -9.5 & -4.8 & 0.058 & 8.1 + spiral ( late - type ) & @xmath55 & 131 & 126 & -5 & -1.9 & -1.0 & 0.031 & 0.3 + barred spiral & @xmath56 & 191 & 269 & 78 & 17.0 & 8.5 & 0.062 & 15.2 + barred spiral ( early - type ) & @xmath57 & 97 & 140 & 43 & 18.1 & 9.1 & 0.091 & 14.6 + barred spiral ( late - type ) & @xmath58 & 68 & 83 & 15 & 9.9 & 5.0 & 0.066 & 9.2 + unknownmorphology & @xmath59 & 31 & 37 & 6 & 8.8 & 4.4 & 0.095 & 3.9 + irregular & @xmath60 & 20 & 14 & -6 & -17.6 & -8.9 & 0.105 & 5.2 + 3000@xmath15rv ( km s@xmath0)@xmath183500 & @xmath61 & 194 & 208 & -14 & -3.5 & -1.7 & 0.046 & 2.6 + 3500@xmath15rv ( km s@xmath0)@xmath184000 & @xmath62 & 201 & 201 & 0 & 0.0 & 0.0 & 0.031 & 3.1 + 4000@xmath15rv ( km s@xmath0)@xmath184500 & @xmath63 & 182 & 172 & 10 & 2.8 & 1.4 & 0.034 & 0.6 + 4500@xmath15rv ( km s@xmath0)@xmath185000 & @xmath64 & 237 & 226 & 11 & 2.4 & 1.2 & 0.041 & 1.0 + grid 1 & @xmath65 & 21 & 20 & 1 & 2.4 & 1.2 & 0.068 & 6.8 + grid 2 & @xmath66 & 116 & 121 & -5 & -2.1 & -1.1 & 0.014 & 1.6 + grid 3 & @xmath67 & 112 & 88 & 24 & 12.0 & 6.0 & 0.076 & 9.3 + grid 4 & @xmath68 & 14 & 12 & 2 & 7.7 & 3.9 & 0.647 & 9.7 + grid 5 & @xmath69 & 8 & 11 & -3 & -15.8 & -7.9 & 0.042 & 22.3 + grid 6 & @xmath70 & 80 & 75 & 5 & 3.2 & 1.6 & 0.081 & 4.5 + grid 7 & @xmath71 & 56 & 62 & -6 & -5.1 & -2.5 & 0.095 & 8.9 + grid 8 & @xmath72 & 33 & 22 & 11 & 20.0 & 10.1 & 0.073 & 12.9 + grid 9 & @xmath73 & 31 & 20 & 11 & 21.6 & 10.9 & 0.028 & 18.1 + grid 10 & @xmath74 & 108 & 124 & -16 & -6.9 & -3.5 & 0.004 & 5.4 + grid 11 & @xmath75 & 52 & 61 & -9 & -8.0 & -4.0 & 0.025 & 6.2 + grid 12 & @xmath76 & 20 & 20 & 0 & 0.0 & 0.0 & 0.409 & 7.4 + grid 13 & @xmath77 & 66 & 78 & -12 & -8.3 & -4.2 & 0.039 & 2.9 + grid 14 & @xmath78 & 44 & 37 & 7 & 8.6 & 4.3 & 0.356 & 10.3 + grid 15 & @xmath79 & 44 & 47 & -3 & -3.3 & -1.6 & 0.050 & 5.2 + grid 16 & @xmath80 & 9 & 9 & 0 & 0.0 & 0.0 & 0.191 & 8.6 + group 1 & @xmath81 & 30 & 37 & -7 & -10.4 & -5.2 & 0.032 & 7.1 + group 2 & @xmath82 & 70 & 48 & 22 & 18.6 & 9.4 & 0.097 & 12.6 + group 3 & @xmath83 & 37 & 31 & 6 & 8.8 & 4.4 & 0.027 & 4.3 + group 4 & @xmath84 & 40 & 34 & 6 & 8.1 & 4.1 & 0.031 & 3.9 + group 5 & @xmath85 & 85 & 107 & -22 & -11.5 & -5.7 & 0.089 & 11.2 + group 6 & @xmath86 & 45 & 42 & 3 & 3.4 & 1.7 & 0.024 & 1.6 + @xmath46 the dominance of trailing structural modes are significant ( 17% @xmath1 8.5% ) in spiral barred galaxies . the @xmath14 value is found @xmath37 9% for both early- ( sba , sbab , sbb and sbbc ) and late - type ( sbc , sbcd , sbd , sbm ) barred spirals . we suspect that the field sb galaxies are not the best candidates of chiral objects . similar result ( i.e. , @xmath14 @xmath37 8% ) is found for the irregulars and the morphologically unidentified galaxies . one interesting similarity is noticed between the late - type spirals and barred spirals . the @xmath14 value for both the late - types are found less than that of early - types ( see table 1 ) . thus , the chirality is favourable for the late - types rather than the early - types . the difference between the standard deviation of the major diameters ( @xmath40 ) for trailing and leading arm galaxies is found less than 0.050 arc minute for the total sample , spirals and the late - type spirals ( seventh column , table 1 ) . these samples showed @xmath14 value @xmath15 5% ( grey - shaded region , fig . thus , we found a good correlation between the @xmath14(% ) and @xmath40 value . probably , this result hints the fact that the size of the non - superimposable mirror images should lie within a limit . in our database , this limit should not exceed 0.050 arc minute for @xmath40 . the difference between the sum of the major diameters ( in percentage ) are found greater than 10% for the barred spirals and early - type barred spirals . interestingly , these two subsamples showed @xmath14 value greater than 15% ( fig . we suspect that the sb galaxies possess chiral symmetry breaking . thus , we found that the chirality between the total leading and trailing arm galaxies exist in our database . this behavior is found prominent for the spirals , mainly for early - type spirals . a very good correlation between the number of leading and trailing arm galaxies can be seen in the rv classifications ( fig . all 4 subsamples show the @xmath14 and @xmath40 value less than 5% and 0.050 , respectively ( table 1 ) . in addition , @xmath14(@xmath17 ) is found to be @xmath15 5% . this result is important in the sense that the statistics in these subsamples is rich ( number of galaxies @xmath37 170 ) enough . thus , we could not observe the violation of chirality in the low and high rv galaxies in our database . a difference is noticed : dominance of leading and trailing modes in low ( rv1 ) and high ( rv3 , rv4 ) rv subsamples , respectively . however , this dominance is not significant ( i.e. , @xmath14 @xmath15 5% ) . an equal number of leading and trailing arm galaxies are found in the subsample rv2 ( 3500 @xmath15 rv ( km s@xmath0 ) @xmath18 4000 ) ( table 1 ) . this might be a coincidence . in order to check the binning effect , we further classify the total galaxies in 6 ( @xmath14rv = 333 km s@xmath0 ) and 8 bins ( @xmath14rv = 250 km s@xmath0 ) and study the statistics . no significant dominance of either trailing or leading structural modes are noticed . thus , it is found that the chirality of field galaxies remain invariant with the global expansion ( i.e , expansion of the universe ) . this is an important result . we further discuss this result below . we study the distribution of leading and trailing arm galaxies by dividing the sky into 16 equal parts ( fig . the area of the grid ( g ) is 90@xmath2 @xmath45 45@xmath2 ( ra @xmath45 dec ) . the area distribution of leading and trailing arm galaxies are plotted , that can be seen in fig . 3a. the statistical parameters are given in table 1 . 1@xmath5 is shown . the positions of the clusters abell 0426 and abell 3627 are shown by the symbol @xmath45 " ( a , b).,title="fig:",height=173 ] 1@xmath5 is shown . the positions of the clusters abell 0426 and abell 3627 are shown by the symbol @xmath45 " ( a , b).,title="fig:",height=113 ] 1@xmath5 is shown . the positions of the clusters abell 0426 and abell 3627 are shown by the symbol @xmath45 " ( a , b).,title="fig:",height=173 ] 1@xmath5 is shown . the positions of the clusters abell 0426 and abell 3627 are shown by the symbol @xmath45 " ( a , b).,title="fig:",height=113 ] a significant dominance ( @xmath372@xmath5 ) of trailing structures is noticed in grid 3 ( ra : 0@xmath2 to 90@xmath2 , dec : 0@xmath2 to 45@xmath2 ( j2000 ) ) ( fig . 3a , a ) . an elongated subcluster like structure can be seen in this grid . in this grid , @xmath14 , @xmath40 and @xmath87 are found to be 12% @xmath1 6% , 0.076 and 9.3% , respectively . these figures suggest that the galaxies in g3 lost their chiral property . probably , this might be due to the apparent subgroupings or subclusterings of the galaxies . the trailing arm galaxies dominate in the grids g8 and g9 ( fig . 3a ) . however , the statistics is poor ( @xmath15 40 ) in these grids ( table 1 ) . in addition , no groupings or subclustering are noticed . thus , we conclude nothing for these area grids . a dominance ( @xmath881.5@xmath5 ) of leading structures is noticed in g10 ( ra : 180@xmath2 to 270@xmath2 , dec : 45@xmath2 to 0@xmath2 ( j2000 ) ) and g13 ( ra : 270@xmath2 to 360@xmath2 , dec : 90@xmath2 to 45@xmath2 ( j2000 ) ) ( fig . 3a , a ) . in both the grids , a large aggregation of the galaxies can be seen . a subcluster - like aggregation can be seen in g10 . an elongated structure can be seen in g13 . in both the grids , @xmath14 value is found to be greater than 5% ( table 1 ) . thus , we suspect that the galaxies in these grids ( g10 , g13 ) are loosing chiral symmetry . this result might reveal the effects of the cluster evolution on chiral symmetry of galaxies . no dominance of either leading or trailing structures is noticed in the groups g1 , g2 , g4 , g5 , g6 , g7 , g11 , g12 , g14 , g15 and g16 . thus , the chirality is found intact in @xmath88 80% area of the sky . we suspect that the groupings or subclusterings of the galaxies lead the violation of chirality in the grids g3 , g10 and g13 . we study the existence of chirality in the groups below . in all - sky map , several groups of galaxies can be seen ( fig . it is interesting to study the existence of chirality in these groups . for this , we systematically searched for the groups fulfilling following selection criteria : ( a ) major diameter @xmath37 30@xmath2 , ( b ) cutoff diameter @xmath15 2 times the background galaxies , ( c ) number of galaxies @xmath37 50 . we found 6 groups fulfilling these criteria ( fig . all 6 groups ( gr ) are inspected carefully . in 3 groups ( gr2 , gr5 and gr6 ) , subgroups can be seen . the number of galaxies in the groups gr2 and gr5 are found more than 100 . the clusters abell 0426 and abell 3627 are located close to the groups gr2 and gr6 . the symbol @xmath45 " represents the cluster center in fig . the mean radial velocities of these clusters are 5366 km s@xmath0 and 4881 km s@xmath0 , respectively . however , we have removed the member galaxies of these clusters from our database . a significant dominance ( @xmath372@xmath5 ) of trailing structures is noticed in the group gr2 ( fig . 3b , b ) . the @xmath14 , @xmath40 and @xmath87 values are found to be 18.6% @xmath1 9.4% , 0.097 and 12.6% , respectively ( table 1 ) . these values indicate that the galaxies in this group might lost their chiral symmetry . we suspect that the galaxies in this group is under the influence of the cluster abell 0426 , due to which apparent subclustering of the galaxies can be seen in this group . this subclustering lead the violation of chiral symmetry . the galaxies in gr5 shows an opposite preference : a significant dominance of the leading arm galaxies ( @xmath372@xmath5 ) ( fig . 3b , b ) . in this group , @xmath14 , @xmath40 and @xmath87 are found to be 11.5% @xmath1 5.7% , 0.089 and 11.2% , suggesting the violation of chirality ( table 1 ) . no humps or dips can be seen in the groups gr1 , gr3 , gr4 and gr6 ( fig . 3.2b , b ) . thus , the galaxies in these groups show chiral property . it is interesting that the number of galaxies in these groups are less than 100 . the groups gr2 and gr5 , which showed a significant difference between the number of leading and trailing arm galaxies , have a very good statistics ( i.e. , @xmath37 100 ) . thus , we conclude that the large aggregation of the galaxies lead the violation of chirality . in the group 6 , we could not notice the influence of the cluster abell 3627 . this might be due to the off location of the cluster center from the group center . 1@xmath5 error bars are shown . pa = 90@xmath2@xmath145@xmath2 ( grey - shaded region ) corresponds to the galactic rotation axes tend to be oriented perpendicular with respect to the equatorial plane.,title="fig:",height=151 ] 1@xmath5 error bars are shown . pa = 90@xmath2@xmath145@xmath2 ( grey - shaded region ) corresponds to the galactic rotation axes tend to be oriented perpendicular with respect to the equatorial plane.,title="fig:",height=151 ] we study the equatorial position angle ( pa ) distribution of trailing and leading arm galaxies in the total sample and the subsamples . a spatially isotropic distribution is assumed in order to examine non - random effects in the pa - distribution . in order to discriminate the deviation from the randomness , we use three statistical tests : chi - square , auto correlation and the fourier . the bin size was chosen to be 20@xmath2 ( 9 bins ) in all these tests . the statistically poor bins ( number of solution @xmath15 5 ) are omitted in the analysis . the conditions for anisotropy are the following : the chi - square probability p(@xmath19 ) @xmath15 0.050 , correlation coefficient @xmath89/@xmath90 @xmath37 1 , first order fourier coefficient @xmath27/@xmath91 ) @xmath37 1 and the first order fourier probability p(@xmath92)@xmath150.150 as used by godlowski ( 1993 , 1994 ) . table 2 lists the statistical parameters for the total samples and subsamples . in the fourier test , @xmath27 @xmath15 0 ( i.e. , negative ) indicates an excess of galaxies with the galactic plane parallel to the equatorial plane . in other words , a negative @xmath27 suggests that the rotation axis of galaxies tend to be oriented perpendicular with respect to the equatorial plane . because the galactic plane is perpendicular to the rotation axis of the galaxy . similarly , @xmath27 @xmath37 0 ( i.e. , positive ) indicates that the rotation axis of galaxies tend to lie in the equatorial plane . in the histograms ( see figs . 4 - 7 ) , a hump at 90@xmath2@xmath145@xmath2 ( grey - shaded region ) suggests that the galactic planes of galaxies tend to lie in the equatorial plane . in other words , the rotation axes of galaxies tend to be oriented perpendicular with respect to the equatorial plane when there is excess number of solutions in the grey - shaded region in the histograms . all three statistical tests show isotropy in the total trailing arm galaxies . thus , no preferred alignment is noticed for the total trailing arm field galaxies ( solid circles in fig . interestingly , all three statistical tests show anisotropy in the total leading arm galaxies . the chi - square and fourier probabilities ( p@xmath93 , p(@xmath92 ) ) are found 1.5% ( @xmath15 5% limit ) and 8.5% ( @xmath15 15% limit ) , respectively ( table 2 ) . the auto correlation coefficient ( c / c(@xmath5 ) ) turned 3.2 ( @xmath941 ) . the @xmath27/@xmath91 ) value is found to be negative at @xmath88 2@xmath5 level , suggesting that the rotation axes of leading arm galaxies tend to be oriented the equatorial plane . three humps at 50@xmath2 ( @xmath371.5@xmath5 ) , 90@xmath2 ( @xmath372@xmath5 ) and 130@xmath2 ( 1.5@xmath5 ) can be seen ( fig . 4a ) . all these humps lie in the grey - shaded region . we checked the binning biasness in the statistics by increasing the number of bins to 12 and 16 . a similar statistical result is found for both structural modes . 4b shows the pa - distribution histogram for the total sample in 18 bins . the leading arm galaxies show three significant humps in the grey - shaded region , supporting the above mentioned result . thus , we conclude isotropy for trailing whereas anisotropy for leading arm galaxies in the total sample . @xmath46 p0.1ccccccccc sample & & @xmath95 & & & & @xmath96 & & + & @xmath36(>^2)@xmath46 & @xmath97@xmath98 & @xmath46_11@xmath99@xmath50_11@xmath98 & @xmath100>_1@xmath98 & @xmath36(>^2)@xmath46 & @xmath97@xmath98 & @xmath46_11@xmath99@xmath50_11@xmath98 & @xmath100>_1@xmath98 + total & 0.666 & + 0.0 & -0.9 & 0.381 & 0.015 & -3.2 & -1.9 & 0.085 + s & 0.511 & -0.7 & -1.2 & 0.434 & 0.225 & + 0.4 & + 0.8 & 0.383 + se & 0.973 & + 0.1 & -0.9 & 0.569 & 0.031 & + 2.0 & + 2.8 & 0.015 + sl & 0.234 & + 0.5 & + 0.8 & 0.209 & 0.460 & -0.1 & -0.5 & 0.345 + sb & 0.729 & + 0.3 & + 1.0 & 0.454 & 0.285 & -1.0 & -0.2 & 0.497 + sbe & 0.739 & + 0.1 & -0.5 & 0.566 & 0.230 & -0.7 & + 0.1 & 0.521 + sbl & 0.043 & + 1.8 & + 1.7 & 0.046 & 0.620 & -0.9 & -0.2 & 0.872 + rv1 & 0.910 & + 0.3 & + 0.8 & 0.362 & 0.369 & -0.9 & -0.6 & 0.285 + rv2 & 0.790 & + 0.3 & -1.0 & 0.496 & 0.925 & -0.4 & -0.2 & 0.887 + rv3 & 0.050 & + 1.6 & -1.5 & 0.083 & 0.033 & -1.8 & -2.3 & 0.046 + rv4 & 0.043 & -2.3 & -1.5 & 0.116 & 0.636 & + 0.2 & -0.7 & 0.692 + gr2 & 0.455 & + 0.6 & + 0.8 & 0.861 & 0.033 & -1.8 & + 1.7 & 0.116 + gr5 & 0.033 & -1.4 & -2.0 & 0.085 & 0.516 & + 0.4 & -0.4 & 0.548 + @xmath46 in the spirals , the chi - square and auto correlation tests show isotropy for both the trailing and leading modes . the first order fourier probability is found greater than 35% , suggesting no preferred alignment . however , the @xmath27 value exceeds 1@xmath5 limit ( 1.2@xmath5 ) in the trailing spirals . a hump at 90@xmath2 is not enough to turn the @xmath27/@xmath91 ) @xmath37 1.5 ( fig . 5a ) . similarly , a hump at 150@xmath2 is not enough to make the @xmath27/@xmath91 ) @xmath37 1.5 in the leading spirals . hence , the preferred alignment is not profounded in both the leading and trailing spirals . thus , we conclude a random orientation of trailing and leading arm spirals . early- and late - type trailing arm spirals show isotropy in all three statistical tests ( table 2 ) . no humps and the dips are seen in the histograms ( solid circles in fig . thus , the trailing arm spirals show a random alignment in the pa - distribution . in the subsample se , all three statistical tests show anisotropy ( table 2 ) . two significant humps at @xmath37 150@xmath2 cause the first order fourier coefficient ( @xmath27 ) @xmath37 + 2.5@xmath5 ( hollow circle in fig . thus , a preferred alignment is noticed in the early - type leading arm spirals : the galactic rotation axes tend to lie in the equatorial plane . the late - type leading spirals show a random alignment . the spiral barred galaxies show a random alignment in both the trailing and leading modes . in fig . 5d , no deviation from the expected distribution can be seen . all three statistical tests support this result ( table 2 ) . a similar result is found for the early - type sb galaxies in both structural modes ( table 2 , fig . the p@xmath93 and p(@xmath92 ) are found less than 5% , suggesting a preferred alignment for the late - type sb galaxies having trailing arm ( table 2 ) . the auto correlation coefficient ( c / c(@xmath5 ) ) and the hump at @xmath37 150@xmath2 support this result ( fig . 5f ) . the @xmath27/@xmath91 ) is found to be positive at 1.7@xmath5 level , suggesting that the trailing arm sbl galaxies tend to lie in the equatorial plane . thus , the late - type trailing and the leading arm sb galaxies show preferred and random alignments , respectively . the subsamples rv1 and rv2 show isotropy in all three statistical tests ( table 2 ) . no humps or dips can be seen in figs . thus , the galaxies having radial velocity in the range 3000 km s@xmath0 to 4000 km s@xmath0 show a random alignment for both the leading and the trailing structural modes . the humps at 90@xmath2 ( @xmath372@xmath5 ) and 110@xmath2 ( @xmath372@xmath5 ) are found in the leading and trailing arm rv3 galaxies , respectively ( fig . these two significant humps lead the subsample show anisotropy in the statistical tests ( table 2 ) . the @xmath27 values are found negative at @xmath1011.5 level , suggesting a similar preferred alignment for both modes : the galaxy rotation axes tend to be directed perpendicular to the equatorial plane . a hump at 70@xmath2 ( @xmath371.5@xmath5 ) and a dip at 150@xmath2 ( @xmath882@xmath5 ) cause the trailing arm rv4 galaxies to show anisotropy in all three statistical tests ( fig . thus , the trailing arm galaxies having radial velocity in the range 4500 km s@xmath0 to 5000 km s@xmath0 show a similar alignment as shown by the subsample rv3 : galactic planes of galaxies tend to lie in the equatorial plane . the leading arm galaxies in the subsample rv4 show a random alignment ( table 2 , fig . 6d ) . we do not study pa - distribution of leading and trailing arm galaxies in the groups gr1 , gr3 , gr4 and gr6 because of poor statistics ( number @xmath15 50 ) . fortunately , a very good correlation between the leading and trailing arm galaxies are noticed in these groups . in other words , we noticed that the chirality is not violated in these groups . we study the pa - distribution of leading and the trailing arm galaxies in the groups gr2 and gr5 , where the chiral symmetry seem to be violated . in addition , the statistics is relatively better in these two groups . in the group gr2 , leading arm galaxies dominate the trailing arm galaxies . in this group , the leading arm galaxies show a preferred alignment whereas trailing arm galaxies show a random alignment in the pa - distribution . all three statistical tests suggest anisotropy in the leading arm galaxies ( table 2 ) . the humps at @xmath37 150@xmath2 cause the @xmath27 value to be positive at @xmath37 1.5@xmath5 level ( fig . 7a ) , suggesting that the rotation axes of leading arm galaxies in gr2 tend to be oriented parallel the equatorial plane . the trailing arm galaxies dominate in the group gr5 . interestingly , a preferred alignment of trailing arm galaxies is noticed in the pa - distribution . in fig . 7b , two significant humps at 90@xmath2 ( @xmath882@xmath5 ) and 110@xmath2 ( @xmath372@xmath5 ) can be seen . these humps lead the subsample ( trailing gr5 ) to show anisotropy in the statistical tests ( table 2 ) . no preferred alignment is noticed in the leading arm galaxies in this group . thus , the structural modes ( leading or trailing ) whose population dominates in the groups show a preferred alignment in the pa - distribution . this is an interesting result . 8a shows a comparison between the number ( @xmath14 ) and position angle ( @xmath27/@xmath91 ) ) distribution of leading and trailing arm galaxies in the total sample and subsamples . this plot deals the possible correlation between the chirality ( non - chirality ) and the random ( preferred ) alignment in the subsamples . the grey - shaded region represents the region of isotropy and chirality for the @xmath27/@xmath91 ) and @xmath14(% ) , respectively . twenty five ( out of 39 , 64% ) subsamples lie in the grey - shaded region ( fig . 8a ) , suggesting a good agreement between the chiral property and the random alignment of the rotation axes of galaxies . in four subsamples ( se , sbl , gr2 and gr5 ) , a good correlation between the preferred alignment and the achiral ( i.e. , non - chiral ) property is noticed ( fig . thus , it is found that the random alignment of the pas of galaxies lead the existence of chiral property of galaxies . % ) and the position angle ( @xmath27/@xmath91 ) ) distribution of leading and trailing arm galaxies in the total sample and subsamples . ( b ) a comparison with the previous work ( aryal , acharya & saurer 2007 ) . error bars as in fig . 2.,title="fig:",height=181 ] % ) and the position angle ( @xmath27/@xmath91 ) ) distribution of leading and trailing arm galaxies in the total sample and subsamples . ( b ) a comparison with the previous work ( aryal , acharya & saurer 2007 ) . error bars as in fig . 2.,title="fig:",height=188 ] here we discuss our results with the results obtained by aryal et al . aryal et al . ( 2007 , hereafter paper 1 ) compiled a database of 667 leading and trailing arm galaxies in the lsc ( rv @xmath15 3000 km s@xmath0 ) and studied the chiral property in the total sample and 17 subsamples . in the present study , the database is compiled from the field galaxies having rv in the range 3000 to 5000 km s@xmath0 . thus , we are moving deep inside the sky and studying the existence of chiral property in the present work . the distribution of the total trailing and leading structures in the lsc ( paper 1 ) as well as in the field galaxies ( present work ) are found homogeneous ( fig . the difference between the leading and trailing arm galaxies are turned well within 5% error limit . this result indicate the fact that the chiral property is a global phenomenon rather than a local phenomenon . for the spirals , we noticed no deviation from the chirality . however , paper 1 found a slight deviation because of the presence of the virgo cluster . interestingly , we noticed the violation of chirality in the barred spirals ( sb ) whereas paper 1 concluded the sb galaxies as a chiral object . for other morphological types ( irregulars , morphologically unidentified galaxies ) , our result is similar to that of the previous result . in the pa - distribution , we noticed a random alignment for the trailing arm galaxies whereas paper 1 concludes no preferred alignments for the leading arm galaxies . in this work , a preferred alignment is noticed for the leading arm galaxies . in contrast to this , a preferred alignment for the trailing arm galaxies is concluded in paper 1 . this contradiction is interesting in the sense that the nature of the databases are different in the past ( lsc galaxies having rv @xmath15 3000 km s@xmath0 ) and the present ( field galaxies having 3000 @xmath15 rv ( km s@xmath0 ) @xmath18 5000 ) study . the preferred alignment in both the studies are obviously different . for leading arm lsc galaxies , they found that the galactic rotation axes tend to lie in the equatorial plane . in our case , we noticed that the rotation axes of the trailing arm galaxies tend to be oriented perpendicular the equatorial plane . paper 1 concludes a similar preferred alignment in the leading and trailing structures in the virgo cluster galaxies : the galactic rotation axes tend to lie in the equatorial plane . we noticed a similar result for the early - type spirals ( leading mode ) , late - type barred spirals ( trailing mode ) and the group gr2 ( leading mode ) . an important similarity is found in the past and present study : late - type galaxies show chiral property , whereas this property is found to be violated in the early - type galaxies . in this way , we noticed few inconsistencies in the present and the previous work . it should be remembered that these inconsistencies are profounded either due to the poor statistics or because of the bias in the sample classification . we combined the database of the paper 1 and the present work , and studied the basis statistics . a strong chirality ( @xmath14 @xmath15 1% , @xmath40 @xmath15 0.010 and @xmath87 @xmath15 5% ) is found in the total and the spiral galaxies . however , the @xmath14 value is found to be greater than 10% for the barred spirals . thus , the total and the spiral galaxies having rvs less than 5000 km s@xmath0 strongly exhibit chiral symmetry . it seems that the chirality loss sequence ( spiral @xmath102 barred spiral @xmath102 elliptical ) as predicted by capozziello and lattanzi ( 2006 ) might be true . aryal & saurer ( 2004 , 2005b , 2006 ) and aryal , paudel & saurer ( 2007 ) studied the spatial orientation of galaxies in 32 abell clusters of bm type i ( 2004 ) , bm type iii ( 2005b ) , bm type ii - iii ( 2006 ) and bm type ii ( 2007 ) and found a significant preferred alignment in the late - type cluster ( bm type ii - iii , bm type iii ) . they concluded that the randomness decreases systematically in galaxy alignments from early - type ( bm type i , ii ) to late - type ( bm type ii - iii , iii ) clusters . thus , the existance of chirality in bm type i cluster , as predicted by capozziello and lattanzi ( 2006 ) might be true . because we noticed a very good correlation between the randomness and the chirality . probably , this result reveals the fact that the progressive loss of chirality might have some connection with the rotationally supported ( spirals , barred spirals ) to the randomized ( lenticulars , ellipticals ) system . thus , we suspect that the dynamical processes in the cluster evolution ( such as late - type clusters ) give rise to a dynamical loss of chirality . in other word , a good correlation between the achirality and anisotropy can be suspected for the late - type clusters . it would be interesting to test this prediction by analysing the chiral property of spirals in the late - type clusters in the future .

is studied using chi - square , auto - correlation and the fourier tests . we noticed a good agreement between the random alignment of the position angle ( pa ) distribution and the existence of chirality in both the leading and trailing arm galaxies . chirality is found stronger for the late - type spirals ( sc , scd , sd and sm ) than that of the early - types ( sa , sab , sb and sbc ) . a significant dominance ( 17% 8.5% ) of trailing modes in addition , chirality of field galaxies is found to remain invariant under the global expansion . it is found that the rotation axes of leading arm galaxies tend to be oriented perpendicular the equatorial plane . example.eps gsave newpath 20 20 moveto 20 220 lineto 220 220 lineto 220 20 lineto closepath 2 setlinewidth gsave .4 setgray fill grestore stroke grestore

we present an analysis of the chiral property of 1621 field galaxies having radial velocity 3000 km s to 5000 km s . a correlation between the chiral symmetry breaking and the preferred alignment of galaxies in the leading and trailing structural modes is studied using chi - square , auto - correlation and the fourier tests . we noticed a good agreement between the random alignment of the position angle ( pa ) distribution and the existence of chirality in both the leading and trailing arm galaxies . chirality is found stronger for the late - type spirals ( sc , scd , sd and sm ) than that of the early - types ( sa , sab , sb and sbc ) . a significant dominance ( 17% 8.5% ) of trailing modes is noticed in the barred spirals . in addition , chirality of field galaxies is found to remain invariant under the global expansion . the pa - distribution of the total trailing arm galaxies is found to be random , whereas preferred alignment is noticed for the total leading arm galaxies . it is found that the rotation axes of leading arm galaxies tend to be oriented perpendicular the equatorial plane . a random alignment is noticed in the pa - distribution of leading and trailing spirals . example.eps gsave newpath 20 20 moveto 20 220 lineto 220 220 lineto 220 20 lineto closepath 2 setlinewidth gsave .4 setgray fill grestore stroke grestore

0803.4032

r

emission lines can be used to separate star - burst and active galactic nuclei ( agn ) galaxies using the [ oiii]@xmath305007/h@xmath29[nii]/h@xmath28 diagnostic . values of the [ oiii]@xmath305007 , [ nii ] and h@xmath28 emission strengths were estimated from the 6dfgs spectra in a traditional manner ( section [ meas_em_lines ] ) . the strength of h@xmath29 emission , however , was taken to be the deviation between the measured value of the h@xmath29 index and the best - fit ssp model value . due to the relatively small observed dynamical range of h@xmath29 in absorption compared to that observed in emission , this is a fairly robust estimate of h@xmath29 emission . this is evidenced by the classic ` y'-shape evident in the diagnostic plot of [ oiii]@xmath305007/h@xmath29 against [ nii]/h@xmath28 shown in in fig . [ emission ] . the position of the line dividing star - burst galaxies from agn varies between studies in the literature . in fig . [ emission ] we show the lines given by kewley et al . ( 2001 ) ( solid line ) and kauffmann ( 2003a ) ( dashed line ) . in order to ensure reliable identifications of agn we use the definition of kewley et al . in the following . about 2200 of the 2700 emission galaxies in our sample have reliable estimates of all four of these emission lines , 300 of these lying in the region of the diagram associated with agn as defined by kewley et al . + in this section we outline the results of our age and metallicity determinations using lick indices . the distribution of ages and metallicities are presented in fig . the sample has been sub - divided into passive- and emission - galaxies as described in section [ meas_em_lines ] . emission - galaxies classified as agn ( section [ emlingal ] ) are shown as grey symbols , the remainder are plotted in black . before interpreting this plot , we must note the apparent clusterings and ` zones - of - avoidance ' , the clearest examples of which are evident in the emission galaxy plot . these are the inevitable result of linear interpolations in non - rectilinear spaces . returning attention to the actual results of our age and metallicity determinations ; similar trends of increasing [ z / h ] with decreasing log(age ) are evident in both of passive and emission galaxies albeit with emission galaxies tending to younger ages . galaxies classified as possessing agn in section [ emlingal ] also often possess young central ages . we investigate the effect of aperture size on derived parameter using fig . [ aperture ] . in these plots , the galaxies with the largest and smallest sizes ( [email protected] arcsec and [email protected] arcsec ) are shown as red and blue points respectively . the blue points therefore represent the galaxies in which the 6.7 arcsec aaomega aperture encompasses the largest fraction of galaxy light ( on average @xmath03.0 r@xmath15 ) , while the red points represent galaxies in which the smallest fraction of galaxy light ( @xmath00.5 r@xmath15 ) is encompassed by the fibre . the observed age metallicity and age velocity dispersion relations can be seen to be similar in these two extremes , with the most noticeable difference a slightly higher [ z / h ] ( by @xmath00.1 dex ) in large galaxies . we therefore find aperture - size effects to be generally small . due to the large , circular aperture of the 6dfgs , a direct comparison with the literature , which are mainly based on long - slit observations of galaxy centres , is not straight forward . nevertheless , we note that the observed trend is at least qualitatively in agreement with the trends observed in previous studies ( trager et al . 2000 ; proctor & sansom 2002 ; mehlert et al . 2003 ; gallazzi et al . 2005 ; collobert et al . the young central ages of agn are also consistent with studies in the literature showing that star - bursts are often found to accompany agn activity ( e.g. see cid fernandes et al . 2004 and references therein ) . an intriguing feature of fig . [ agez ] is the presence of a few old galaxies with extremely low [ z / h ] values , ( @xmath330.5 dex ) . these galaxies , present in both passive and emission samples , possess a range of velocity dispersions . visual inspection of both spectra and photometry generally reveals nothing exceptional about these galaxies . further analysis of these galaxies are beyond the scope of this paper , but we note that their low numbers result in them having no impact on the results and conclusions of the present paper . as in these literature studies , the slope in the age metallicity relation exhibited by our data is similar to the slope of the age metallicity degeneracy ( worthey 1994 ; ` the 3/2 rule ' ) . before comparing these results with the photometry , it is therefore clearly important to establish that the age metallicity degeneracy has been broken . to this end , in the following section we consider the errors in the derived parameters . as described in section [ errdps ] , errors in age and metallicity were characterised by considering the results of the 50 monte - carlo realisations of each of the 7000 galaxies in our sample . fits using both the index combinations shown in fig . [ chis ] ( with and without balmer lines ) were carried out for the purposes of this analysis ( a total of approximately one million realisations ) . the differences between the input - model values of log(age ) and [ z / h ] and those of each of the realisations were calculated . these were combined in a number of bins depending upon the age metallicity of the input models . the analysis was also carried out using the _ observed _ index values as the inputs to the monte - carlo realisations ( rather than the best - fit model values ) . no significant quantitative or qualitative differences were found . [ error ] shows the results of this analysis as 1-sigma confidence contours for three age metallicity bins . each contour in these plots is based on tens of thousands of individual realisations of the best - fit galaxy data . the contour levels were defined as the iso - densities corresponding to e@xmath12 of the peak values ( as expected for the 1-sigma contour of a two - dimensional gaussian distribution ) . the use of the iso - density peak in estimating the extent of 1-sigma confidence contours results in a small uncertainty in the estimates . however , we estimate these to be of order 2% and therefore insignificant . results of the error analysis for both passive galaxies ( in which balmer indices were included in fits ) and emission galaxies ( in which balmer indices were _ excluded _ from fits ) are shown separately . we recall that h@xmath29 was excluded from all the fits . [ error ] also shows the data divided into two signal - to - noise ratio regimes ( s / n@xmath2316 @xmath12 and s / n@xmath2516 @xmath12 ) . marginal distributions in log(age ) and [ z / h ] are shown at the edges of each main plot . we note that the positions of the contours in fig . [ error ] were chosen for clarity only . to get a full picture one must imagine the whole surface of the log(age)[z / h ] plane populated by such contours . [ error ] shows that , as might be expected , error estimates increase with decreasing signal - to - noise - ratio . for passive galaxies , the analysis reveals log(age ) and [ z / h ] errors of @xmath00.1 dex in galaxies with s / n above @xmath016 . this falls to @xmath00.2 dex in galaxies with s / n below @xmath016 . the errors of the emission - galaxies ( in which all balmer lines are excluded from the fits ) are somewhat larger , but are generally still of order 0.2 dex . it is therefore clear that while the age metallicity degeneracy is still present ( as evidenced by the sloped elliptical error contours in fig . [ error ] ) , the magnitude of the errors is sufficiently small for its effects to be negligible in our results - even in galaxies fit without balmer lines . this is the result , and main advantage , of using large numbers ( @xmath2510 ) indices in the determination of ages and metallicities . the good age resolution achieved is emphasised in the marginal distributions of fig . [ error ] , in which galaxies older than 10 gyr ( 1.0 dex ) , and galaxies younger than 1.5 gyr ( 0.15 dex ) ( delimited by dashed lines in the marginal distributions ) are largely uncontaminated by galaxies with ages of 3 gyr ( 0.5 dex ) . since the following analysis concentrates mainly on the very oldest and very youngest galaxies , these provide confidence that results do not suffer significantly from the effects of the age metallicity degeneracy . the final sample for which velocity dispersions and reliable age / metallicity estimates were measured consists of 4500 passive galaxies and 1000 emission galaxies . + the combination of 6dfgs , 2mass and supercosmos data permits the comparison for spectroscopic age and metallicity determinations to the b , r and k band photometric data for some 6000 galaxies . specifically , by combining the photometry with the velocity dispersion measures from 6dfgs spectroscopy , we are able to investigate trends with age and metallicity in b , r and k band mass - to - light - ratios . this analysis is carried out using a _ dynamical _ mass ( m@xmath51 ) calculated as : @xmath52 where @xmath32 is the central velocity dispersion , r@xmath53 is the half - light radius and the constant c has a value of 5.0 ( cappellari et al . unfortunately half - light radii are only available for 25% of our sample , while k band 20th magnitude isophotal radii ( r@xmath54 ) are available for the entire sample . a calibration was therefore carried out by using equation [ masseq ] to calculate masses for the galaxies in the 25% sub - sample ( which also possess much more accurate , aperture corrected , velocity dispersion estimates ) , and comparing them with mass estimates using r@xmath14 ( instead of r@xmath53 ) in the same equation . a plot of the comparison is shown in fig . [ mass_comp ] . the correlation has equation m@[email protected] with only 0.12 dex of scatter . given error estimates in m@xmath51 and m@xmath56 of 0.06 and 0.10 dex respectively , the correlation is clearly extremely good . masses quoted throughout the remainder of this work are therefore those based on r@xmath14 corrected as detailed . as well as the photometry of 2mass and supercosmos , we also use the photometric results from the ssp models of bruzual & charlot ( 2003 ) . some key values from the bruzual & charlot models are shown in table [ photometry ] . we use these models almost exclusively to estimate differential properties ; e.g. the rate of change of mass - to - light - ratio with log(age ) in populations of given [ z / h ] . it should be noted that the rates of change of the photometric properties are reasonably constant with metallicity and age , as long as age is expressed in logarithmic form . this clearly indicates that a differential approach is a robust use of the model values . we shall flag the one occasion in which the models are used in a non - differential manner . . examples of b , r and k magnitudes from bc03 for single stellar population models of varying age and metallicity are presented . note that these values are used almost exclusively in a differential manner ( see section [ photo ] ) . [ cols="^,^,^,^,^,^ " , ] we now turn our attention to the effects of the young central populations detected in our 6dfgs spectra on the global properties of galaxies as indicated by their photometry . there are 400 ` young ' ( @xmath231.5 gyr ) galaxies ( see section [ eamopp ] ) included in this analysis . plots of mass - to - light - ratio with mass for both young and old galaxies are shown in fig . mass - binned averages of young and old data are shown as blue and red solid lines respectively . also shown in these plots as dashed blue lines are model predictions of the effects of the differing ages and metallicities of the young and old populations in each mass bin . these are based on the comparison of luminosity values from the bc03 ssp models for the differences in age and metallicity of the young and old galaxies in each mass bin . a direct comparison of the data for young galaxies to these lines would therefore implicitly assume that the young central populations detected in the spectroscopy pervade the galaxies populations _ as a whole _ - i.e represents the same fraction of the galaxy population at all radii . young galaxies can be seen to be displaced with respect to the old galaxies in a sense consistent with their young ages . however , the displacements are small compared to the predicted values . they also become smaller towards higher masses . the data therefore show that , on average : ( i ) young populations do _ not _ pervade their host galaxies , but are instead centrally concentrated ; ii ) young populations must constitute relatively small fractions of total galaxy masses ; and ( iii ) the mass fraction in the form of a young stellar population must decrease with increasing galaxy mass . to quantify these conclusions , for each mass bin , the fraction of total galaxy stellar mass and luminosity involved in the recent star - burst ( f@xmath57 and f@xmath58 ) were estimated . these estimates were made by a differential comparison of the bc03 predictions for mass - to - light - ratios to the observed values . in each waveband , the difference between the observed mass - to - light - ratios of the young galaxies from those of the old galaxies in the same mass bin ( @xmath59log[m / l]@xmath60 ) was compared to the difference between model predictions for young and old populations of appropriate metallicity ( @xmath59log[m / l]@xmath61 ) . denoting these mass - to - light - ratio differences as s and r respectively , we derive : @xmath62_{obs}}-1}{10^{-\delta\log[m / l]_{model}}-1 } \ \ or \ \ \frac{10^{-s}-1}{10^{-r}-1}.\ ] ] the derivation of equation 6 is given in appendix a. we estimate the _ luminosity _ fraction in the recent burst by : @xmath63 finally , the total mass ( m@xmath64 , in solar masses ) is estimated by : @xmath65 where m@xmath66 is the total mass of the galaxy . the derivation of these expressions again assumes that the young populations at any given mass are seen against a background old population with a metallicity and mass - to - light - ratio appropriate to that mass ( as given by the mass - binned values for old galaxies in table [ mol ] ) . the young stellar population is nonetheless assumed to _ dominate the central regions _ ( such that issues related to degeneracies with burst - strength do not arise ) . the estimates of mass- and luminosity - fractions and burst - masses are plotted against galaxy mass in fig . [ mass_frac ] . despite the ( expected ) variation in luminosity fractions between wavebands , agreement between the values for the mass fractions derived from the three wavebands is extremely good . this suggests that both our assumption of the same underlying dynamical ` tilt ' in young galaxies as exhibited by old galaxies is correct , and that our age and metallicity corrections are accurate . the trend shows a decreasing mass fraction with increasing mass falling from @xmath010% at 10@xmath67 m@xmath27 to 2% at 10@xmath68 m@xmath27 . this is in agreement with the study of treu et al . ( 2005 ) , who find the stellar fraction formed in recent times varies from 20%-40% below 10@xmath69 m@xmath27 , to below 1% above 10@xmath68 m@xmath27 . also shown in fig . [ mass_frac ] are the estimates of average total mass involved in the recent burst of star formation . agreement between wavebands is again good . however , the absolute values given here must be treated with caution , as they are calculated by the only non - differential application of bc03 models used in this work . nevertheless , in a relative sense , the data exhibit a narrow range of burst - masses and an apparent _ upper limit _ on their size in high mass galaxies . as a final consideration we next investigate the implications of our findings for colour - magnitude relations . note that these relations are independent of dynamical effects . in fig . [ cmr ] we plot the colour - magnitude relations b k and r k against m@xmath70 . galaxies with both old ( grey points ) and young ( black points ) central populations are shown . averages ( binned by luminosity ) of the old and young populations are again shown as solid red and blue lines respectively . slopes in the old galaxies are in good accord with literature values ( thin black lines in each plot ) . as the sample includes late - type galaxies , particular amongst the young galaxies , these must lie close to the ` red sequence ' in our colour - magnitude diagrams . this suggests they therefore correspond to the ` dusty red - sequence ' of wolf , gray & meisenheimer ( 2005 ) . the lack of galaxies in the ` blue - cloud ' is almost certainly the result of a number of selection effects . these include the effects of template mismatches during velocity dispersion measurements and the removal of galaxies with poor fits or large errors during age / metallicity measurement . [ cmr ] shows that , for both colours , young galaxies tend to be bluer than the old galaxies . however , the offsets are extremely small . our results from the previous section indicate that this is largely the result of the low mass fractions formed in the recent star formation events . [ cmr ] therefore shows that galaxies lieing on , or close to , the red sequence may not be ` red and dead ' , but may still be forming modest numbers of new stars . indeed , we note that _ most _ of the galaxies with young central populations also exhibit on - going star formation ( fig . [ agez ] ) .

we also derive dynamical masses and dynamical mass - to - light ratios for these galaxies by combining the central velocity dispersions with global photometry in the b , r and k bands from supercosmos and 2mass . together , these data allow us to reduce the degeneracies between age , metallicity and star formation burst - strength that have limited previous studies . we find that the central regions of galaxies of all masses often exhibit young stellar populations . we conclude that the young stellar populations detected in spectroscopic studies are generally centrally concentrated , and that there is an upper limit on the mass of star - forming events in massive galaxies . these results have important ramifications for mass - to - light ratios estimated from photometric observations . latexl-.36em.3ex-.15em t-.1667em.7ex-.125emx [ section ] [ firstpage ] galaxies : general , galaxies : stellar content , galaxies : kinematics and dynamics

we present an analysis of the stellar populations as a function of mass in a sample of galaxies of all morphological and emission types from the 6df galaxy survey ( 6dfgs ) . we measure velocity dispersions and lick indices from the spectra of the central regions of these galaxies , deriving ages and metallicities from the lick indices using stellar population models . we also derive dynamical masses and dynamical mass - to - light ratios for these galaxies by combining the central velocity dispersions with global photometry in the b , r and k bands from supercosmos and 2mass . together , these data allow us to reduce the degeneracies between age , metallicity and star formation burst - strength that have limited previous studies . we find that old galaxies exhibit a mass - metallicity relation with logarithmic slope d[fe / h]/dlog , while young galaxies show slopes consistent with zero . when we account for the effects of the mass - metallicity relation , we obtain a single , consistent relation between mass - to - light ratio and mass for old galaxies in all passbands , . as we have accounted for stellar population effects , this remaining variation in the mass - to - light with mass ( the residual ` tilt ' of the fundamental plane ) must have a dynamical origin . however , we demonstrate that any simple trend between mass - to - light - ratio and mass or luminosity is inconsistent with the observations , and that a more complex relationship must exist . we find that the central regions of galaxies of all masses often exhibit young stellar populations . however it is only in the lowest - mass galaxies studied ( ) that these populations are evident in the global photometry . in higher - mass galaxies , young central populations have decreasing influence on the global photometry , with there being no discernible impact in galaxies more massive than . we conclude that the young stellar populations detected in spectroscopic studies are generally centrally concentrated , and that there is an upper limit on the mass of star - forming events in massive galaxies . these results have important ramifications for mass - to - light ratios estimated from photometric observations . latexl-.36em.3ex-.15em t-.1667em.7ex-.125emx [ section ] [ firstpage ] galaxies : general , galaxies : stellar content , galaxies : kinematics and dynamics

0803.4032

c

we have measured ages and metallicities from the 6dfgs spectra of @xmath07000 galaxies of all morphological and emission types . we have demonstrated that the age metallicity degeneracy has been broken in this study , particularly with regard to the very young and very old galaxies used in the subsequent analysis . these data for the _ central regions _ are compared with the _ global _ photometry from 2mass and supercosmos for the @xmath06000 galaxies for which such photometry is available . combining spectroscopy and photometry we are able to investigate trends in global b , r and k band mass - to - light - ratios with the central stellar populations of the galaxies . an age - metallicity trend is identified in passive galaxies that is consistent with the results of similar studies in the literature . the stellar populations of galaxies exhibiting emission lines ( arising from either star - bursts or agn ) follow a similar trend , but exhibit younger ages than their passive counterparts . to minimise the effects of complex mixtures of stellar populations , we confine our analysis to the data for very old or very young galaxies . we find a steep mass - metallicity relation in old galaxies ( logarithmic slope of @xmath00.25 ) , while in young galaxies the trend is consistent with zero slope . using only the old galaxies , we find a dynamical trend in mass - to - light - ratios with mass of logarithmic slope @xmath00.15 . this is in good agreement with values reported in the literature . however , there is no agreement in the literature as to the cause of this tilt . padmanabhan et al . ( 2004 ) concluded that most of the tilt is due to variations in dark matter content . from the modelling perspective , robertson et al . ( 2006 ) perform simulations of a series of mergers and investigate the scaling relations of their merger remnants . from mergers of gas - rich disk galaxies they derive a fp with @xmath28 = 1.55 and @xmath29 = 0.82 , in good agreement with our k band values ( see table 3 ) . in their models the tilt away from the virial scaling relation is caused by gas dissipative effects and not dark matter content . trujillo et al . ( 2004 ) on the other hand attribute the tilt to ` non - homology ' , deriving almost identical contributions to the tilt from dynamical and stellar population effects as those we find in this work . our findings , however , also show that whatever the cause of the tilt , it is not a simple in mass - to - light - ratio ( or its scatter ) with mass or luminosity . instead we find a that a more complex description will be required ; this will be a subject of a future paper . considering extremely young galaxies , our results show that the properties of the central stellar populations identified by the spectroscopy do _ not _ generally correlate with the global properties of their host galaxies . indeed , the data indicate that recent bursts of star formation are limited to a mass of @xmath3310@xmath24 m@xmath27 and are strongly centrally concentrated . consequently , in high mass galaxies exhibiting young central populations , recent star forming events have negligible effect on their global photometry . as a consequence we have shown that galaxies sitting on , or near , the red sequence may not be ` red and dead ' , but may still be forming modest numbers of new stars . we therefore conclude that the young ages found in studies of galaxy centres in this and other spectroscopic studies should not be considered representative of the age of their host galaxies as a whole . they may , however , be reasonably interpreted as indicators of recent assembly events . as a result of these considerations , the small scatter in photometric scaling relations , which initially motivated the monolithic collapse models , ceases to be inconsistent with either the hierarchical merging model or the results of the lick index studies indicating young central populations . in future work we will include the results from the completed 6df galaxy survey and report on the full chemical properties ( including @xmath28-abundance ratios ) of the central regions of galaxies . the effective radii , peculiar velocities and environmental parameters of the 6dfgs sample , currently being analysed by the 6dfgs team , will be included . this will allow not only more accurate determination of photometric properties and a more detailed analysis of the fundamental plane , but also the investigation of the strong _ quantitative _ predictions of trends in galaxy formation / assembly epochs and metals - content with galaxy mass and environment . + * acknowledgements * this publication makes use of data products from the two micron all sky survey ( 2mass ) , which is a joint project of the university of massachusetts and the infrared processing and analysis center / california institute of technology , funded by the national aeronautics and space administration and the national science foundation . the authors acknowledge the data analysis facilities provided by iraf , which is distributed by the national optical astronomy observatories and operated by aura , inc . , under cooperative agreement with the national science foundation . we thank the australian research council for funding that supported this work . the measuring of line values from 6dfgs spectra was begun in philip lah s honours thesis 2003 `` exploring the stellar population of early - type galaxies in the 6df galaxy survey '' which looked at the 6dfgs early data release , ( december 2002 ) . this honours thesis was supervised by matthew colless and heath jones . assistance with the software used in this work was provided by craig harrison . lachlan campbell provided additional assistance with the 6dfgs data particularly with the velocity dispersions and 6df arc spectra . the authors also thank chris blake for valuable input . * references * + baldry i.k . et al . , 2002 , apj , 569 , 582 + baum w.a . , 1959 , iau symp . 10 , the hertzsprung - russell diagram , ed . j. l. greenstein ( paris : iau ) , 23 + bernardi m. et al . , 2003a , aj , 125 , 1882 + bernardi m. et al . , 2003b , aj , 125 , 1866 + bernardi m. et al . , 2006 , aj , 131 , 2018 + bernardi m. , 2007 , aj , 133 , 1954 + blumenthal g.r . , faber s.m . , primack j.r . , rees m.j . , 1984 , nature , 311 , 517 + bower r.g . , lucy j.r . , ellis r.s . , 1992a , mnras , 254 , 589 + bower r.g . , lucy j.r . , ellis r.s . , 1992b , mnras , 254 , 601 + bruzual a.g . , charlot s. , 2003 , mnras , 344 , 1000 + caldwell n. , rose j.a . , concannon k.d . , 2003 , aj , 125 , 2891 + cappellari m. et al . , 2006 , mnras , 366 , 1126 + carlberg r.g . , 1984 , apj , 286 , 403 + chang r. , gallazzi a. , kauffmann g. , charlot s. , ivezi z. , brinchmann j. , heckman t.m . , 2006 , mnras , 366 , 717 + cid fernandes r. , gu q. , melnick j. , terlevich e. , terlevich r. , kunth d. , rodrigues lacerda r. , joguet b. , 2004 , mnras , 355 , 273 + cid fernandes r , mateus , a. , sodr l. , stasinska g. , gomes j.m . , 2005 , mnras , 358 , 363 + clemens m.s . , bressan a. , nikolic b. , alexander p. , annibali f. , rampazzo r. , 2006 , mnras , 370 , 702 + colless m. , saglia r.p . , burstein d. , davies r.l . , mcmahan jr r.k . , wegner g. , 2001a , mnras , 321 , 277 + colless m. et al . , 2001b , mnras , 328 , 1039 + collobert m. , sarzi m. , davies r.l . , kuntschner h. , colless m. , 2006 , mnras , 370 , 1213 + de lucia g. , springel v. , white s.d.m . , croton d. , kauffmann g. , 2006 , mnras , 366 , 499 + dressler a. , lynden - ball d. , burstein d. , faber s.m . , terlevich r. , wegner g. , 1987 , apj , 313 , 42 + eisenstein d.j . et al . , 2003 , apj , 585 , 694 + fernandes r.c . , mateus , a. , sodr l. , stasinska g. , gomes j.m . , 2005 , mnras , 358 , 363 + gallazzi a. , charlot s. , brinchmann j. , white s.d.m . , tremonti c. , 2005 , mnras , 362 , 41 + gebhardt , k. et al . , apj , 539 , l13 + gibbons r.a . , fruchter a.s . , bothun g.d . , 2001 , aj , 121 , 649 + gonzlez j. j. , 1993 , phd thesis , univ . california + hambly n.c . , irwin m.j . , macgillivray h.t . , 2001 , mnras , 326 , 1295 + hudson m.j . , lucey j.r . , smith r.j . , steel j. , 1997 , mnras , 291 , 488 + jarrett t.h . , chester t. , cutri r. , schneider s. , skrutskie m. , huchra j.p . , 2000a , aj , 119 , 2498 + jarrett t.h . , cheste , t. , cutri r. , schneider s. , rosenberg j. , huchra j.p . , mader j. , 2000b , aj , 120 , 298 + jones d.h . et al . , 2004 , mnras , 355,747 + jones d.h . , peterson b.a . , colless m. , saunders w. , 2006 , mnras , 369 , 25 + jrgensen i. , franx m. , kjrgaard p. , 1995 , mnras , 276 , 1341 + jrgensen i. , franx m. , kjrgaard p. , 1996 , mnras , 280 , 167 + kauffmann g. , 1996 , mnras , 281 , 487 + kauffmann g. et al . , 2003 , mnras , 346 , 1055 + kauffmann g. et al . , 2003a , mnras , 341 , 33 + kauffmann g. et al . , 2003b , mnras , 341 , 54 + kewley l.j . , dopita m.a . , sutherland r.s . , heisler c.a . , trevena j. , 2001 , apj , 556 , 121 + korn a.j . , maraston c. , thomas d. , 2005 , a&a , 438 , 685 + larson r.b . , 1974 , mnras , 166 , 585 + lewis i. et al . , 2002 , mnras , 334 , 673 + mehlert d. , thomas d. , saglia r.p , bender r. , wegner g. , 2003 , a&a , 407 , 423 + moore s.a.w . , 2001 , ph.d thesis , univ of durham + padmanabhan n. et al . , 2004 , new astron . , 9 , 329 + pahre m.a . , djorgovski s.g . , de carvalho r.r . , 1998 , aj , 116 , 1591 + proctor r.n . , sansom a.e . , 2002 , mnras , 333 , 517 + proctor r.n . , forbes d.a . , beasley m.a . , 2004a , mnras , 355 , 1327 + proctor r.n . , forbes d.a . , hau g.k.t . , beasley m.a . , de silva g.m . , contreras r. , terlevich a.i . , 2004b , mnras , 349 , 1381 + proctor r.n . , forbes d.a , forestell a. , gebhardt k. , 2005 , mnras , 362 , 857 + robertson b. , cox t.j . , hernquist l. , franx m. , hopkins p.f . , martini p. , springel v. , 2006 , apj , 641 , 21 + smith r.j . , hudson m.j . , lucey , j.r . , nelan , j.e . , wegner , g.a . , 2006 , mnras , 369 , 1419 + spergel d.n . et al . , 2003 , apjs , 148 , 175 + terlevich a.i . , forbes d.a . , 2002 , mnras , 330 , 547 + thomas d. , maraston c. , bender r. , 2003 , mnras , 339 , 897 + thomas d. , maraston c. , bender r. , mendes de oliveira c. , 2005 , apj , 621 , 673 + thomas d. , maraston c. , korn a. , 2004 , mnras , 351 , 19 + trager , s. , faber , s. , worthey , g. , gonzalez , j. , 2000 , aj , 120 , 165 + tremonti c. et al . , apj , 613 , 898 + treu t. , ellis r.s . , liao t.x . , van dokkum p.g . , tozzi p. , coli , a. , newman j. , cooper m.c . , davis m. , 2005 , apj , 633 , 174 + trujjillo i. , burkert a. , bell e.f . , 2004 , aj , 600 , l39 + white s.d.m . , rees m.j . , 1978 , mnras , 183 , 341 + wolf c. , gray m.e . , meisenheimer k. , 2005 , a&a , 443 , 435 + worthey g. , 1994 , apjs , 95,107 + worthey g. , ottaviani d.l . , 1997 , apjs , 111 , 377 + york d.g . et al . , 2000 , aj , 120 , 1579 + [ lastpage ]

we find that old galaxies exhibit a mass - metallicity relation with logarithmic slope d[fe / h]/dlog , while young galaxies show slopes consistent with zero . when we account for the effects of the mass - metallicity relation , we obtain a single , consistent relation between mass - to - light ratio and mass for old galaxies in all passbands , . as we have accounted for stellar population effects , this remaining variation in the mass - to - light with mass ( the residual ` tilt ' of the fundamental plane ) must have a dynamical origin . however , we demonstrate that any simple trend between mass - to - light - ratio and mass or luminosity is inconsistent with the observations , and that a more complex relationship must exist . however it is only in the lowest - mass galaxies studied ( ) that these populations are evident in the global photometry . in higher - mass galaxies , young central populations

we present an analysis of the stellar populations as a function of mass in a sample of galaxies of all morphological and emission types from the 6df galaxy survey ( 6dfgs ) . we measure velocity dispersions and lick indices from the spectra of the central regions of these galaxies , deriving ages and metallicities from the lick indices using stellar population models . we also derive dynamical masses and dynamical mass - to - light ratios for these galaxies by combining the central velocity dispersions with global photometry in the b , r and k bands from supercosmos and 2mass . together , these data allow us to reduce the degeneracies between age , metallicity and star formation burst - strength that have limited previous studies . we find that old galaxies exhibit a mass - metallicity relation with logarithmic slope d[fe / h]/dlog , while young galaxies show slopes consistent with zero . when we account for the effects of the mass - metallicity relation , we obtain a single , consistent relation between mass - to - light ratio and mass for old galaxies in all passbands , . as we have accounted for stellar population effects , this remaining variation in the mass - to - light with mass ( the residual ` tilt ' of the fundamental plane ) must have a dynamical origin . however , we demonstrate that any simple trend between mass - to - light - ratio and mass or luminosity is inconsistent with the observations , and that a more complex relationship must exist . we find that the central regions of galaxies of all masses often exhibit young stellar populations . however it is only in the lowest - mass galaxies studied ( ) that these populations are evident in the global photometry . in higher - mass galaxies , young central populations have decreasing influence on the global photometry , with there being no discernible impact in galaxies more massive than . we conclude that the young stellar populations detected in spectroscopic studies are generally centrally concentrated , and that there is an upper limit on the mass of star - forming events in massive galaxies . these results have important ramifications for mass - to - light ratios estimated from photometric observations . latexl-.36em.3ex-.15em t-.1667em.7ex-.125emx [ section ] [ firstpage ] galaxies : general , galaxies : stellar content , galaxies : kinematics and dynamics

astro-ph0509148

i

one of the key predictions of any cosmological structure formation model is the distribution of matter on large scales . we now know that the cold dark matter ( cdm ) paradigm can account for many of the observed characteristics of the galaxy distribution . in this picture , the universe contained tiny ( gaussian ) density fluctuations at the time the cosmic microwave background was last scattered . bound structures assembled themselves through gravitational instability around these perturbations in the relatively recent past . in the past two decades , with the advent of high - resolution numerical simulations , it has become possible to follow this picture through the formation of massive galaxies and clusters . the cdm model accurately describes the abundance and clustering of collapsed objects from dwarf galaxies to rich galaxy clusters over a wide range of redshifts ( although the details of galaxy formation itself remain somewhat mysterious ) as well as the distribution of neutral gas in the intergalactic medium . these systems all correspond to density peaks in the matter distribution ( with the exception of the lowest column density ly@xmath7 forest absorbers ) . for a variety of reasons , the other end of the density distribution underdense voids has received considerably less attention , despite their long observational history @xcite and their place as the most visually striking features of the galaxy distribution . this is largely because voids subtend enormous volumes and so require large surveys to garner representative samples . although voids have been found in every redshift survey @xcite , the first statistically significant sample came only with the 2df redshift survey @xcite . now , with the deep2 redshift survey and the sloan digital sky survey , it is even possible to constrain the evolution of voids over the redshift interval @xmath8@xmath9 @xcite . @xcite presented the most complete search for voids to date . they found that voids with characteristic radii @xmath10 fill @xmath11 of the universe . but their search illustrates a second difficulty in studying voids : how to define and identify them precisely ( and meaningfully ) . @xcite proposed the _ voidfinder _ algorithm based on separating the observed galaxies into `` void '' and `` wall '' populations and building voids around gaps in the wall population ( see also @xcite ) . while clearly defined for any given observational sample , the results can nevertheless be difficult to interpret in relation to the underlying physical quantities of interest . for example , the distribution of observed sizes depends on the galaxy sample ( intrinsically brighter galaxies yield larger voids ) as well as the search algorithm ( @xcite restrict their search to radii greater than @xmath12 , for example ) . smaller voids are difficult to pinpoint because of confusion with random fluctuations in the galaxy distribution . we will nevertheless follow this approach and define a void to be any coherent region where the galaxy density falls below some threshold . note that this differs from many other studies ( e.g. , @xcite ) that require a void to be _ completely _ empty of galaxies . these voids , which are much smaller than the voids we consider here , do not correspond to the voids detected by eye in the galaxy distribution . the void phenomenon is also relatively difficult to study theoretically . the simplest model is a spherical tophat underdensity . the early evolution of such a system is well - described by spherical expansion ( the analog of the well - known spherical collapse model ; @xcite ) . underdensities expand in comoving units , gradually deepening , until they reach `` shell - crossing , '' when the center is evacuated and individual mass shells cross paths . at this stage , the initial spherical expansion model breaks down . later , the voids continue to expand relatively slowly in a self - similar fashion @xcite . unfortunately , these dark matter models are idealizations in that real voids are observed only through the galaxy distribution . because galaxies are biased relative to the dark matter , we must ask what kind of physical systems observed voids actually represent . they are in fact nearly empty of galaxies but does that require shell - crossing , or can they be at an earlier evolutionary stage ? another difficulty is that the large size of voids restricts the usefulness of numerical simulations . early efforts focused on understanding the dynamics of individual voids , which did not require particularly high resolution @xcite . but placing voids in their proper cosmological context demands both large volumes and high mass resolution the latter because we must resolve the galaxies from whose absence we identify voids ( e.g. , @xcite ) . only recently have @xmath13-body simulations of the required dynamic range become practical . this has allowed the first systematic studies of the structure of large voids @xcite as well as of simulated voids in the dark matter distribution @xcite and in the galaxy distribution @xcite . nevertheless , there are still no full hydrodynamic simulations of the void phenomenon : galaxy properties are currently determined through semi - analytic models @xcite . there are a number of reasons to study the void phenomenon . early attention focused on using the observed voids to constrain cosmological parameters . @xcite and @xcite argued that the scale of typical voids would depend on the matter power spectrum , just as the abundance of clusters does . it was quickly realized that the void distribution seemed to have much more large - scale power than collapsed objects most obviously , although voids with sizes @xmath14 are not uncommon , collapsed objects do not reach the same mass scales . this is puzzling because an underdense region reaches shell - crossing only after the equivalent overdensity would virialize . @xcite pointed out that a solution might lie in a proper treatment of the galaxies used to define voids : galaxy bias could lower the required dark matter underdensity , allowing voids to be larger for a given amount of cdm power . they argued that the observed population of large voids preferred a @xmath15cdm model . recently , interest in voids has focused on their role in galaxy formation . @xcite argued that ( in the cdm model ) voids should be populated by small dark matter haloes but that the observed voids appear to lack faint galaxies as well as bright ones . is this a fundamental problem for the cdm paradigm , or does it indicate that galaxy formation proceeds differently in voids ? one popular explanation is that photoheating during reionization may have suppressed the formation of dwarf galaxies in low - density environments @xcite , helping to clear faint galaxies from the voids . to study these problems , recent attention has focused on the variation of the galaxy luminosity function with the large - scale environment . @xcite and @xcite ( see also @xcite ) found that the characteristic galaxy luminosity and the galaxy density both decrease significantly in low - density environments but that the faint - end slope remains nearly constant . the latter may be surprising given the expected variation in the halo mass function with environment . before we can answer any of these questions , however , we require a proper understanding of voids in cdm models . @xcite and @xcite made important strides forward with their studies of voids in @xmath13-body simulations . they claimed that current semi - analytic galaxy formation models predict void properties similar to those observed and argued that the conflict pointed out by @xcite was illusory . but such studies are still limited by their dynamic range and by the ( necessarily complex ) galaxy formation prescriptions imposed on the dark matter haloes . the goal of this paper is to produce a straightforward analytic model of the void distribution and of galaxy populations within voids . such a model will sharpen our understanding of voids in a cdm model and generate a baseline prediction with which we can contrast their observed properties . we will build on @xcite , the most compelling theoretical model of voids to date . they used the excursion set formalism , which reproduces the abundance of collapsed haloes extremely well , to predict the distribution of void sizes . however , they required voids to reach shell - crossing and defined their properties in terms of the dark - matter underdensity . as a result , they predicted characteristic void radii much smaller than those observed . our main goal is to modify their approach so as to describe voids in the _ galaxy _ distribution . along the way , we will also be able to predict the halo populations within voids and quantify the claimed discrepancy with observational results . the remainder of this paper is organized as follows . in [ exp ] , we examine the nonlinear evolution of the void density . then , in [ defn ] , we show how to compute the linearized underdensity of voids with a specified galaxy underdensity , and we briefly discuss the expected galaxy populations inside voids . in [ abundance ] , we show how to compute the cosmological abundance of large _ galaxy _ voids and compare our results to observations . finally , we conclude in [ disc ] . in our calculations , we assume a cosmology with @xmath16 , @xmath17 , @xmath18 , @xmath19 ( with @xmath20 ) , @xmath21 , and @xmath22 , consistent with the most recent measurements @xcite .

we find that , because of bias , galaxy voids are typically significantly larger than dark matter voids and should fill most of the universe . we show that voids selected from catalogs of luminous galaxies should be larger than those selected from faint galaxies : the characteristic radii range from for galaxies with absolute-band magnitudes to .

we present an analytic model for the sizes of voids in the galaxy distribution . peebles and others have recently emphasized the possibility that the observed characteristics of voids may point to a problem in galaxy formation models , but testing these claims has been difficult without any clear predictions for their properties . in order to address such questions , we build a model to describe the distribution of galaxy underdensities . our model is based on the `` excursion set formalism , '' the same technique used to predict the dark matter halo mass function . we find that , because of bias , galaxy voids are typically significantly larger than dark matter voids and should fill most of the universe . we show that voids selected from catalogs of luminous galaxies should be larger than those selected from faint galaxies : the characteristic radii range from for galaxies with absolute-band magnitudes to . these are reasonably close to , though somewhat smaller than , the observed sizes . the discrepancy may result from the void selection algorithm or from their internal structure . we also compute the halo populations inside voids . we expect small haloes ( ) to be up to a factor of two less underdense than the haloes of normal galaxies . within large voids , the mass function is nearly independent of the size of the underdensity , but finite - size effects play a significant role in small voids ( ) . cosmology : theory large - scale structure of the universe galaxies : luminosity functions

astro-ph0509148

c

we have described a simple analytic model for voids in the galaxy distribution . our model is based on @xcite , who showed how to apply the excursion set formalism to underdensities . its most important parameter is @xmath62 , the linearized underdensity of a void . those authors originally set @xmath62 to be the density corresponding to shell - crossing . however , this condition produces voids much smaller than the observed structures , whose masses can exceed those of galaxy clusters @xcite . our major contribution has been to show how to define @xmath62 through the _ galaxy _ underdensity . because galaxies are biased relative to the dark matter , voids can be nearer to the mean density than the observations naively indicate . this significantly increases the characteristic void size and the volume filling fraction . our model predicts voids with characteristic radii @xmath222@xmath223 . this is similar to , though somewhat smaller than , the voids found in galaxy redshift surveys @xcite . it is a closer match to the void population in semi - analytic galaxy formation models @xcite . because bright galaxies are more highly biased , our model also predicts that voids selected from shallow surveys should be characteristically larger , in agreement with observations @xcite . however , we also predict larger galaxy densities inside voids than observed . the significance of these discrepancies is not clear : we argued in [ obs ] that the algorithms used to identify voids in surveys , redshift - space distortions , and the internal structure of voids are all important in detailed comparisons to the observations . the last of these is especially crucial , because simulations show that galaxies inside voids tend to congregate near their edges @xcite , where they are difficult to separate from genuine `` wall '' galaxies . two other subtleties may also affect the comparison . the first is the nonlinear evolution of substructure within the void ; we have used linear theory to describe the halo population inside voids , which likely breaks down in some regimes because power can be transferred between scales . the second is poisson variation in the galaxy number counts @xcite . this effect makes direct detection of small voids impossible , forcing existing surveys to search only for voids with @xmath224 . thus it is difficult to make a precise comparison between our simple model and observed voids . however , the qualitiative agreement is reasonable . our analytic approach is the first to contain voids with characteristic radii @xmath209 and hence to be in even qualitative agreement with the observations . our model shows that the enormous extent of observed voids is not particularly surprising and should not be perceived as a `` crisis '' for the cdm model . by using the excursion set approach , we have also self - consistently predicted the halo population inside of voids . in agreement with naive expectations , small haloes should be less underdense than massive haloes . but the discrepancy is by no means large typically smaller than a factor of two over three decades in halo mass . including satellite galaxies in this calculation ( which we have not done ) will decrease the difference . thus , the predicted steepening of the mass function in low - density environments is only modest and is not ruled out by existing observations of the galaxy luminosity function @xcite . we also emphasize that these observations determined the environmental density on relatively small scales ( @xmath225 ) , where random fluctuations in the galaxy field may mimic true underdensities . if so , differences between mean and low - density environments will be further washed out . moreover , on such scales , finite - size effects become important in setting the characteristic luminosity @xmath122 . the test can be sharpened by focusing on the luminosity function within large , easily identified voids although , even there , separating void and wall galaxies may be difficult . thus , we find no compelling reason to believe that galaxy formation differs inside and outside of voids although we have no evidence against such a possibility , either . statistical properties of the galaxy distribution can also be used to test our predictions , especially for small voids that can not be identified unambiguously . the void probability function quantifies the probability that a sphere of a given size contains no galaxies @xcite , and the underdensity probability function quantifies the probability that a sphere is more underdense than some specified value @xcite . these measures are free from bias in the void selection process , but they are harder to predict directly from the excursion set formalism and are less useful in identifying galaxies that reside in voids . our model is in qualitative agreement with the observed trends in these statistics @xcite , but more work is needed to connect our formalism to them . r. , loeb a. , 2004 , apj , 609 , 474 a. j. , hoyle f. , torres f. , vogeley m. s. , 2003 , mnras , 340 , 160 e. , 1985 , apjs , 58 , 1 m. r. , et al . , 2003 , apj , 592 , 819 g. r. , da costa l. n. , goldwirth d. s. , lecar m. , piran t. , 1992 , apj , 388 , 234 j. r. , cole s. , efstathiou g. , kaiser n. , 1991 , apj , 379 , 440 r. , mo h. j. , sheth r. k. , boerner g. , 2002 , mnras , 333 , 730 r. , 2005 , submitted to apj ( astro - ph/0507014 ) j. m. , sheth r. k. , diaferio a. , gao l. , yoshida n. , 2005 , mnras , 360 , 216 s. , kaiser n. , 1989 , mnras , 237 , 1127 c. , et al . , 2005 , submitted to apj ( astro - ph/0508250 ) a. , 2005 , submitted to mnras ( astro - ph/0505421 ) d. j. , et al . , 2005 , mnras , 356 , 1155 v. , geller m. j. , huchra j. p. , 1986 , apj , 302 , l1 j. , da costa l. n. , goldwirth d. s. , lecar m. , piran t. , 1993 , apj , 410 , 458 j. , einasto m. , gramann m. , 1989 , mnras , 238 , 155 g. , & frenk c. s. , & white s. d. m. , & davis m. , 1988 , mnras , 235 , 715 h. , piran t. , 1997 , apj , 491 , 421 h. , piran t. , 2000 , mnras , 313 , 553 h. , piran t. , dacosta l. n. , 1997 , mnras , 287 , 790 j. a. , goldreich p. , 1984 , apj , 281 , 9 y. , piran t. , 2001 , apj , 548 , 1 s. r. , mcquinn m. , hernquist l. , 2005 , mnras , in press ( astro - ph/0507524 ) d. m. , jones t. d. , hoyle f. , rojas r. r. , vogeley m. s. , blanton m. r. , 2005 , apj , 621 , 643 d. m. , vogeley m. s. , 2004 , apj , 605 , 1 s. , okas e. l. , klypin a. , hoffman y. , 2003 , mnras , 344 , 715 s. a. , thompson l. a. , 1978 , apj , 222 , 784 d. j. , 1977 , mnras , 179 , 351 f. , rojas r. r. , vogeley m. s. , brinkmann j. , 2005 , apj , 620 , 618 f. , vogeley m. s. , 2002 , apj , 566 , 641 f. , vogeley m. s. , 2004 , apj , 607 , 751 a. , frenk c. s. , white s. d. m. , colberg j. m. , cole s. , evrard a. e. , couchman h. m. p. , yoshida n. , 2001 , mnras , 321 , 372 r. p. , oemler a. , schechter p. l. , shectman s. a. , 1981 , apj , 248 , l57 a. v. , berlind a. a. , wechsler r. h. , klypin a. a. , gottl " ober s. , allgood b. , primack j. r. , 2004 , apj , 609 , 35 c. , cole s. , 1993 , mnras , 262 , 627 h. white s. d. m. , 1993 , mnras , 282 , 347 v. , arbabi - bidgoli s. , einasto j. , tucker d. , 2000 , mnras , 318 , 280 h. , white s. d. m. , 2002 , mnras , 337 , 1193 m. , furlanetto s. r. , hernquist l. , zahn o. , zaldarriaga m. , 2005 , apj , in press ( astro - ph/0504189 ) p. j. e. , 1980 , the large - scale structure of the universe . princeton : princeton university press p. j. e. , 2001 , apj , 557 , 495 t. , lecar m. , goldwirth d. s. , da costa l. n. , blumenthal g. r. , 1993 , mnras , 265 , 681 w. h. , schechter p. , 1974 , apj , 187 , 425 r. k. , lemson g. , 1999 , mnras , 304 , 767 r. k. , tormen g. , 1999 , mnras , 308 , 119 r. k. , tormen g. , 2002 , mnras , 329 , 61 r. k. , van de weygaert r. , 2004 , mnras , 350 , 517 d. n. , et al . , 2003 , apjs , 148 , 175 y. , sato k. , sato h. , 1984 , progress of theoretical physics , 71 , 938 r. b. , somerville r. s. , trentham n. , verheijen m. a. w. , 2002 , apj , 569 , 573 r. , van kampen e. , 1993 , mnras , 263 , 481 m. s. , geller m. j. , huchra j. p. , 1989 , baas , 21 , 1171 m. s. , geller m. j. , park c. , huchra j. p. , 1994 , aj , 108 , 745 s. d. m. , 1979 , mnras , 186 , 145 j. , hui l. , 2005 , submitted to apj ( astro - ph/0508384 )

we present an analytic model for the sizes of voids in the galaxy distribution . peebles and others have recently emphasized the possibility that the observed characteristics of voids may point to a problem in galaxy formation models , but testing these claims has been difficult without any clear predictions for their properties . in order to address such questions , we build a model to describe the distribution of galaxy underdensities . our model is based on the `` excursion set formalism , '' the same technique used to predict the dark matter halo mass function . we also compute the halo populations inside voids . cosmology : theory large - scale structure of the universe galaxies : luminosity functions

we present an analytic model for the sizes of voids in the galaxy distribution . peebles and others have recently emphasized the possibility that the observed characteristics of voids may point to a problem in galaxy formation models , but testing these claims has been difficult without any clear predictions for their properties . in order to address such questions , we build a model to describe the distribution of galaxy underdensities . our model is based on the `` excursion set formalism , '' the same technique used to predict the dark matter halo mass function . we find that , because of bias , galaxy voids are typically significantly larger than dark matter voids and should fill most of the universe . we show that voids selected from catalogs of luminous galaxies should be larger than those selected from faint galaxies : the characteristic radii range from for galaxies with absolute-band magnitudes to . these are reasonably close to , though somewhat smaller than , the observed sizes . the discrepancy may result from the void selection algorithm or from their internal structure . we also compute the halo populations inside voids . we expect small haloes ( ) to be up to a factor of two less underdense than the haloes of normal galaxies . within large voids , the mass function is nearly independent of the size of the underdensity , but finite - size effects play a significant role in small voids ( ) . cosmology : theory large - scale structure of the universe galaxies : luminosity functions

astro-ph0009158

i

this paper presents the photometric redshift catalogs that have been used in previous papers to investigate different aspects of the evolution of galaxies in the high redshift universe , such as the history of the uv luminosity density and the number of massive galaxies already assembled at early epochs ( fontana et al 1999b ) and the evolution of galaxy sizes ( poli et al 1999 , giallongo et al 2000 ) . a new deep multicolor ( ubvrijk@xmath0 ) photometric catalog has been produced of the galaxies in the ntt deep field and in the slightly overlapping field centered on the @xmath1 quasar br1202 - 07 . this has been obtained by combining the existing bvri and jk@xmath0 images with new , deep , u band observations of both fields acquired with ntt - susi2 . we have further presented in this paper the photometric redshift catalog drawn from this galaxy sample , using a @xmath2 minimization technique based on the bruzual and charlot spectral library , with the addition of dust and intergalactic absorption . we have also presented the results of applying the same photometric redshift techniqe to public catalogs of the hdf n and hdf s , where a simlar optical - infrared coverage is available . the global redshift distribution of i selected galaxies shows a distinct peak at intermediate redshifts , @xmath5 at @xmath6 and @xmath7 at @xmath8 followed by a tail extending to @xmath9 . systematic differences exist amongst the fields , most notably the hdf s contains a much smaller number of galaxies at @xmath10 and at @xmath11 than the hdf we have also presented the redshift distribution of the total ir - selected sample , which may be useful to tailor the planned surveys with ir spectrographs at large telescopes that will target the redshift range @xmath201 . we find that the number density of galaxies in the redshift range is @xmath202arcmin@xmath16 at @xmath17 and @xmath203arcmin@xmath16 at @xmath13 . we have also discussed the different results from applying color selection criteria and photometric redshifts for detecting galaxies in the redshift range @xmath21 using the hdfs data sets . we find that photometric redshifts predict a two times larger number of high @xmath22 candidates in both the hdf n and hdf s and show that this is primarily due to the inclusion of slightly dusty ( @xmath204 with smc extinction law ) models that were discarded in the original color selection criteria conservatively applied by madau et al 1998 . in several cases , the selection of these objects is made possible by the additional constraints from the ir bands . this effect partially reflect the poor spectral sampling of the hdf filter set , and is not present in ground based observations where a @xmath23 color selection criteria may be applied . finally , we show that galactic m stars may mimic @xmath24 candidates in the hdf filter set and that the 4 brightest candidates at @xmath24 in the hdf - s are , indeed , most likely to be m stars . the estimates of the uv luminosity density @xmath198 at @xmath199 from these data , when the selection against halo stars is applied , show that @xmath198 changes by a factor of @xmath200 between the hdf n and the hdf s ( fontana et al 1999b ) . * acknowledgments * + the paper is based on observations made with : the eso new technology telescope at the la silla observatory ( some under the eis programs 59.a-9005(a ) , 60.a-9005(a ) ) , the nasa / esa hubble space telescope and the kitt peak national observatory . the ultraviolet observations of the nttdf were performed in susi-2 guaranteed time of the observatory of rome in the framework of the eso - rome observatory agreement for this instrument . arnouts , s. , cristiani , s. , moscardini , l. , matarrese , s.,lucchin , f. , fontana , a. , giallongo , e. , 1999a , 310 , 540 arnouts , dodorico , s. , cristiani , s. , zaggia , s. , fontana , a. , giallongo , e. , 1999b , a&a 341 , 641 baum , w.a . , 1962 , iau symp 15 , 390 benitez , n. , 1999 apj subm . , astro - ph/9811189 cohen , j. , hogg , d. , blandford , r. , cowie , l. , hu , e. , songaila , a. , shopbell , p. , richberg , k. , 2000 , apj in press , astro - ph/9912048 connolly , a. j. , szalay , a. s. , dickinson , m. , subbarao , m. u. , & brunner , r. j. 1997 , apj , 486 , l11 cristiani , s. , appenzeller , i. , arnouts , s. , nonino , m. , aragn salamanca , a. , benoist , c. , da costa , l. , dennefield , m. , rengelink , r. , renzini , a. , szeifert , t. , white , s. , 2000 , a&a in press , astro ph/0004213 fontana a. , dodorico s. , fosbury , r. , giallongo , e. , hook , i. , poli , f. , renzini , a. , viezzer r. , , 1999a , a&a 343 , l19 fontana , a. , menci , n. , dodorico , s. , giallongo , e. , poli , f. , cristiani , s. , moorwood , a. , saracco , p. , 1999b , mnras , 310 l27 giallongo , e. , dodorico , s. , fontana , a. , cristiani , s. , egami , e. , hu , e. , mcmahon , r. g. , 1998 , aj 115 , 2169 ( g98 ) lanzetta , k. m. , chen , h. , , fernandez soto , a. , pascarelle , s. , yahata , n. , yahil , a. , 2000 , in : the hy - redshift universe : galaxy formation and evolution at high redshift , eds bunker , a and van breugel , w.,j.,m .. , astro - ph/9910554 madau , p. , 1995 , apj , 441 , 18 madau , p. , pozzetti , l. , dickinson , m. , 1998 , apj , 498 , 106 moorwood , a. , cuby , j. g. , lidman , c. , 1998a , the messenger , 91 , 9 moorwood , a. , cuby , j. g. , devillard , n. , lidman , c. , saracco , p. , 1998b , http://www.hq.eso.org/science/sofi_deep/ pascarelle , s. m. , lanzetta , k. m. , fernandez soto , a. , 1998 , apj , 501 , l1 pei , y.c . , 1992 , apj , 395 , 130 pello , r , kneib , j. p. , le borgne , j. f. , bezecourt , j. , ebbels , t. m. , tijera , i. , bruzual , g. , miralles , j. m. , smail , i. , soucail , g. , bridges , t. j. , 1999 , a&a 346 , 359 pickles , a. j. , 1998 , pasp , 110 , 863 pei , y. c. 1992 , apj , 395 , 130 petitjean , p. , pcontal , e. valls - gabaud , d. , charlot , s. 1996 , nature , 380 , 411 pettini , m. , steidel , c. c. , adelberger , k. l. , kellogg , m. , dickinson , m. , giavalisco , m. , 1997 , in cosmic origins : evolution of galaxies , stars , planet and life , shull , woodward and thronson eds . , asp series poli , f. , giallongo , e. , menci , n. , dodorico , s. , fontana , 1999 , apj , 527 , 662 br1202 & 4 & 26.9 & 26.5 & 26 & 25.6 & 25 & 23.4 & & 21.7 + nttdf & 4.8 & 26.9 & 27.5 & 26.85 & 26.5 & 26.4 & 23.4 & & 21.7 + hdf - n & 4.2 & 28.4 & 29.2 & 29.5 & & 29 & 23.8 & 22.9 & 22.4 + hdf - s & 3.9 & 28.2 & 29 & 29.3 & & 28.7 & 24 & 22.1 & 22 + rcccccccllll id&@xmath74 [ h]&@xmath205 [ deg]&u@xmath206&b&v&r&i & j & k&@xmath120&@xmath207 + n0001&12:05:18.01&-7:44:40.98&24.74&24.28&23.33&22.85&22.63&21.92&21.01&0.34&0.25 + n0002&12:05:18.43&-7:44:40.43&24.72&24.46&23.83&23.06&22.12&20.93&19.43&0.80&0.63 + n0003&12:05:18.21&-7:44:39.53&23.35&23.20&22.42&21.79&21.61&21.14&20.28&0.26&0.66 + n0005&12:05:17.81&-7:43:33.34&25.19&24.86&24.27&23.96&23.54&22.28&21.76&0.10 & + n0006&12:05:17.79&-7:43:26.36&25.80&25.88&24.43&24.12&23.97&22.49&21.63&0.29 & + c|cc|cc|cc @xmath22 & gal & gal / arcmin@xmath130 & gal & gal / arcmin@xmath130 & gal & gal / arcmin@xmath130 + & & & + & 43 @xmath208 & 3.08 & 96 @xmath209 & 6.88 & 196 @xmath210 & 14.06 + @xmath211 & 1 @xmath212 & 0.07 & 12 @xmath213 & 0.86 & 38 @xmath214 & 2.73 + @xmath215 & 0 @xmath216 & 0 . & 1 @xmath217 & 0.07 & 13 @xmath218 & 0.93 + @xmath219 & 1 @xmath220 & 0.07 & 1 @xmath220 & 0.07 & 5 @xmath221 & 0.36 + @xmath102 & 0 @xmath216 & 0 & 0 @xmath216 & 0 & 2 @xmath222 & 0.14 + & & & + & 69 @xmath223 & 4.95 & 125 @xmath224 & 8.98 & 247 @xmath225 & 17.72 + @xmath211 & 6 @xmath226 & 0.43 & 24 @xmath227 & 1.72 & 48 @xmath228 & 3.44 + @xmath215 & 1 @xmath217 & 0.07 & 5 @xmath229 & 0.36 & 15 @xmath230 & 1.08 + @xmath219 & 1 @xmath220 & 0.07 & 5 @xmath231 & 0.36 & 17 @xmath232 & 1.21 + @xmath102 & 0 @xmath216 & 0 & 0 @xmath216 & 0 & 13 @xmath233 & 0.93 +

we present and compare in this paper new photometric redshift catalogs of the galaxies in three public fields : the ntt deep field , the hdf n and the hdf s . in the case of the ntt deep field , we present here a new photometric catalog , obtained by combining the existing bvri and jk with new deep u observations acquired with ntt - susi2 , and which includes also the contiguous field centered on the quasar br1202 - 07 . s ) , by adopting a minimization technique on a spectral library drawn from the bruzual and charlot synthesis models , with the addition of dust and intergalactic absorption . systematic differences exist among the fields , most notably the hdf s which contains a much smaller number of galaxies at and at than the hdf we also present for the first time the redshift distribution of the total ir - selected sample to faint limits ( and ) . the hdfs data sets are used to compare the different results from color selection criteria and photometric redshifts in detecting galaxies in the redshift range . photometric redshifts predict a number of high candidates in both the hdf n and hdf s that is nearly 2 times larger than color selection criteria , and it is shown that this is primarily due to the inclusion of dusty models that were discarded in the original color selection criteria by madau et al 1998 . in several cases , the selection of these objects is made possible by the additional constraints from the ir bands . this effect partially reflect the poor spectral sampling of the hdf filter set , and is not present in ground based observations where a color selection criteria may be applied . finally , it is shown that galactic m stars may mimic candidates in the hdf filter set and that the 4 brightest candidates at in the hdf - s are indeed most likely m stars .

we present and compare in this paper new photometric redshift catalogs of the galaxies in three public fields : the ntt deep field , the hdf n and the hdf s . in the case of the ntt deep field , we present here a new photometric catalog , obtained by combining the existing bvri and jk with new deep u observations acquired with ntt - susi2 , and which includes also the contiguous field centered on the quasar br1202 - 07 . photometric redshifts have been obtained for the whole sample ( nttdf + hdf n + hdf s ) , by adopting a minimization technique on a spectral library drawn from the bruzual and charlot synthesis models , with the addition of dust and intergalactic absorption . the accuracy , determined from 125 galaxies with known spectroscopic redshifts , is in the redshift intervals . the global redshift distribution of i selected galaxies shows a distinct peak at intermediate redshifts , at and at followed by a tail extending to . systematic differences exist among the fields , most notably the hdf s which contains a much smaller number of galaxies at and at than the hdf n . we also present for the first time the redshift distribution of the total ir - selected sample to faint limits ( and ) . it is found that the number density of galaxies at isarcmin at andarcmin at , and drops toarcmin ( at ) at . the hdfs data sets are used to compare the different results from color selection criteria and photometric redshifts in detecting galaxies in the redshift range . photometric redshifts predict a number of high candidates in both the hdf n and hdf s that is nearly 2 times larger than color selection criteria , and it is shown that this is primarily due to the inclusion of dusty models that were discarded in the original color selection criteria by madau et al 1998 . in several cases , the selection of these objects is made possible by the additional constraints from the ir bands . this effect partially reflect the poor spectral sampling of the hdf filter set , and is not present in ground based observations where a color selection criteria may be applied . finally , it is shown that galactic m stars may mimic candidates in the hdf filter set and that the 4 brightest candidates at in the hdf - s are indeed most likely m stars . the data and photometric redshift catalogs presented here are available on line at http://www.mporzio.astro.it/highz . 2cm hsd

astro-ph0201185

i

there is now a large and compelling body of observational evidence that suggests that most , if not all , galaxies contain supermassive black holes at their centers ( e.g. , * ? ? ? however , active galactic nuclei ( agn ) are only found in a small minority of all galaxies in the local universe ( e.g. , * ? ? ? * ; * ? ? ? what is it that makes some galaxies agn but most others quiescent ? one line of inquiry is to ask if the differences are to be found in their circumnuclear environments . in particular , is the difference simply a matter of whether or not the central black hole is being provided with interstellar gas to fuel the nuclear activity ? the problem of providing fuel to an agn from the vast reservoirs of interstellar gas found in the disks of spiral galaxies is how to remove the angular momentum from the gas so it can fall into the nucleus . the two classical mechanisms that are invoked are interactions @xcite and bars @xcite , including nuclear bars @xcite . there is an extensive literature devoted to demonstrating that both are theoretically viable mechanisms for fueling agn @xcite . however , neither interactions nor bars of either type are sufficiently common among agn compared to non - agn galaxies to be the fueling mechanism in all cases @xcite . previous investigations have used _ hst _ to study the circumnuclear environments of seyfert galaxies and search for differences between the seyferts with and without broad - line components . @xcite obtained pre - costar imaging of a large sample of seyfert and non - seyfert markarian galaxies to look for differences in the nuclear structure of these galaxies at higher angular resolution than is possible with ground - based imaging . they discovered that the nuclei of broad - line seyfert @xmath0 galaxies are dominated by strong point sources . in contrast , seyfert @xmath1 galaxies and other markarian galaxies that lack broad - line regions contained weak or no strong nuclear source superimposed on the underlying galaxy s surface brightness profile . this result is further borne out in the extensive _ hst _ snapshot program of @xcite . they also invariably find more strong central point sources in seyfert 1 galaxies than in the seyfert 2s . as this survey was carried out with the unaberrated wfpc2 pc camera , these investigators were also able to look for differences in the nuclear environments of these galaxies . they found seyfert 2 galaxies were more likely to possess dusty nuclear environments than seyfert 1 galaxies , lending support to unified models , which propose obscuration of the broad - line region in seyfert 2s by dust . their observations are evidence for the presence of dust on large scales in the nuclear region , not in a torus immediately outside the broad - line region . visible near - infrared color maps obtained with the _ hubble space telescope _ ( _ hst _ ) have shown that the circumnuclear ( @xmath2 pc ) regions of a large number of low - luminosity agn contain nuclear spiral dust lanes . these spirals are distinct from the spiral arms on kpc scales in the main galaxy disk . @xmath3 color maps of these galaxies show that these ` nuclear spirals ' extend from 100 s of pc scales into the unresolved nucleus @xcite . theoretical models for the formation of nuclear spiral structure suggest that it is dynamically distinct from the main disk spiral arms @xcite . @xcite showed that nuclear spirals in agn reside in nonself - gravitating disks and are therefore likely due to shocks in nuclear gaseous disks . they postulated that as shocks can dissipate energy and angular momentum , these nuclear spirals may be the signature of the fueling mechanism in these galaxies . nuclear spirals in agn have generally been seen mostly in seyfert 2s , although this could easily be a selection effect : the samples of @xcite and @xcite were mostly comprised of seyfert 2s . also , @xcite only observed seyfert 2s with _ nicmos _ as they have fainter nuclear psfs , and thus the circumnuclear environments of seyfert 2s are easier to study with _ hst _ than the circumnuclear environments of seyfert 1s . the question remains , however , if seyfert 1s and 2s contain nuclear spirals with the same relative frequency ; that is , nearly 100% as seen in seyfert 2s . as we only have near - infrared _ nicmos _ imaging of the seyfert 2s in the cfa sample , we need to use an alternate technique to look for nuclear spiral structure in these seyferts . in [ sec : procedure ] we discuss our data - processing , and introduce a technique for creating `` structure maps '' in [ sec : stmaps ] that are an excellent surrogate for color maps for detecting small - scale dust - extinction and emission - line features present in the visible - band images . we then use this technique to compare the circumnuclear environments of the seyfert 1s and 2s in [ sec : morph ] and [ sec : sey ] , connecting the nuclear structures seen to the larger host galaxies in most cases . in [ sec : conc ] we present a summary of our results , and discuss the implications for the fueling of the active nuclei .

we present archival _ hubble space telescope _ images of the nuclear regions of 43 of the 46 seyfert galaxies found in the volume - limited , spectroscopically complete cfa redshift survey sample . using an improved method of image contrast enhancement , we see no significant differences in the circumnuclear dust morphologies of seyfert 1s and 2s , and very few seyfert 2 nuclei are obscured by large - scale dust structures in the host galaxies . if seyfert 2s are obscured seyfert 1s , then the obscuration must occur on smaller scales than those probed by hst .

we present archival _ hubble space telescope _ images of the nuclear regions of 43 of the 46 seyfert galaxies found in the volume - limited , spectroscopically complete cfa redshift survey sample . using an improved method of image contrast enhancement , we create detailed high - quality `` structure maps '' that allow us to study the distributions of dust , star clusters , and emission - line gas in the circumnuclear regions ( 100 - 1000 pc scales ) and in the associated host galaxy . essentially all of these seyfert galaxies have circumnuclear dust structures with morphologies ranging from grand - design two - armed spirals to chaotic dusty disks . in most seyferts there is a clear physical connection between the nuclear dust spirals on hundreds of parsec scales and large - scale bars and spiral arms in the host galaxies proper . these connections are particularly striking in the interacting and barred galaxies . such structures are predicted by numerical simulations of gas flows in barred and interacting galaxies , and may be related to the fueling of agn by matter inflow from the host galaxy disks . we see no significant differences in the circumnuclear dust morphologies of seyfert 1s and 2s , and very few seyfert 2 nuclei are obscured by large - scale dust structures in the host galaxies . if seyfert 2s are obscured seyfert 1s , then the obscuration must occur on smaller scales than those probed by hst . accepted for publication in the apj , v569 ( 2002 apr 20 )

0912.1760

i

our _ ab initio _ studies of the electronic structure , magnetic moments , exchange interactions and curie temperatures in zb crx ( x = as , sb , s , se and te ) and cras@xmath0x@xmath0 ( x = sb , s , se and te ) reveal that half - metallicity in these alloys is maintained over a wide range of lattice parameters . the results for the exchange interaction and the curie temperature show that these alloys have relatively high curie temperatures , i.e. room temperature and above . the exceptions occur for the alloys involving s , se and te at some low values of lattice parameters , where significant inter - atomic antiferromagnetic exchange interactions indicate ground states to be either antiferromagnetic or of complex magnetic nature . a comparison of total energies for the fm , dlm , and two zb antiferromagnetic configurations ( afm[001 ] and afm[111 ] ) show the lowest energy configuration to be afm[111 ] for crs and crse for compressed lattice parameters ( table[table2 ] ) . the possibility of afm ground states for compressed lattice parameters for crs was noted by zhao and zunger@xcite and for crse by sasioglu @xcite . our search for the antiferromagnetic ground states is more thorough than what was reported in these two studies . an extensive study of several antiferromagnetic configurations as well as ferrimagnetic and more complex magnetic structures for crs , crse and crte is currently underway . the mixed pnictide - chalcogenide alloys cras@xmath0x@xmath0 ( x= s , se , te ) do not show any tendency to antiferromagnetic spin fluctuations for the entire range of the lattice parameter studied . presumably the pnictogens suppress antiferromagnetic tendencies . such alloys may play an important role in fabricating stable zb half - metallic materials , as the concentration of the pnictogens and the chalcogens may be varied to achieve lattice - matching with a given substrate . as long as the concentration of as or sb is higher than the chalcogen concentration , half - metallic ferromagnetic state can be achieved . there is a large variation in the curie temperature of these alloys ( fig . [ fig19 ] ) as the lattice parameter varies from the low ( @xmath4 5.4 ) to the mid ( @xmath4 6.1 ) range of the lattice parameters studied . this variation is much smaller for the isoelectronic alloys cras , crsb and cras@xmath0sb@xmath0 ( fig . [ fig17 ] ) over this range of lattice parameters . note that most ii - vi and iii - v zb semiconductors have lattice parameters in this range . large changes in @xmath1 can be brought about by changing the carrier concentrations . the pnictides in general have a higher @xmath1 than the chalcogenides . our results for the curie temperature , the lattice fourier transform of the exchange interactions , and the resulting stability analysis are based on the exchange interactions between the cr atoms only . for the fm reference states this causes some errors due to the neglect of the effects of the induced moments . the dlm results are free from such errors . it is expected that the present study will provide both qualitative and quantitative guidance to experimentalists in the field . 99 k. sato and h. katayama - yoshida , semicond . * 17 * , 367 ( 2002 ) . see k. sato , t. fukushima and h. katayama - yoshida , j. phys . : matter * 19 * , 365212 ( 2007 ) , and references therein . see b. belhadji , l. bergqvist , r. zeller , p.h . dederichs , k. sato and h. katayama - yoshida , j. phys . : condens . matter * 19 * , 436227 ( 2007 ) , and references therein . h. saito , v. zayets , s. yamagata , and k. ando , phys . lett * 90 * , 207202 - 1 ( 2003 ) . k. sato and h. katayama - yoshida , jpn . . phys . * 40 * , l651 ( 2001 ) . h. akinaga , t. manago , and m. shirai , jpn . j. appl . phys.*39 * , l1118 ( 2000 ) . s. li , j - g duh , f. bao , k - x liu , c - l kuo , x. wu , liya l , z. huang , and y du , j. phys . phys . * 41 * 175004 ( 2008 ) . m. shirai , j. appl . phys . * 93 * , 6844 ( 2003 ) . h. akinaga , m. mizuguchi , k. nagao , y. miura , and m. shirai in _ springer lecture notes in physics _ * 676 * , 293 - 311 ( springer - verlag , berlin 2005 ) . k. yamana , m. geshi , h. tsukamoto , i. uchida , m. shirai , k. kusakabe , and n. suzuki , j. phys . : condens . matter * 16 * , s5815 ( 2004 ) . l. kahal , a. zaoul , m. ferhat , j. appl . phys . * 101 * , 093912 ( 2007 ) . i. galanakis and p. mavropoulos , , 104417 ( 2003 ) ; see also i. galanakis , , 012406 ( 2002 ) . pask , l.h . yang , c.y . fong , w.e . pickett , and s. dag , , 224420 ( 2003 ) . t. ito , h. ido , and k. motizuki , j. mag . mag . mat . * 310 * , e558 ( 2007 ) . l - j shi and b - g liu , j. phys . : condens . matter * 17 * , 1209 ( 2005 ) . m. zhang , j. phys . : matter * 15 * , 5017 ( 2003 ) . j. kbler , , 220403(r ) ( 2003 ) . b. sanyal , l. bergqvist , and o. eriksson , , 054417 ( 2003 ) . xie , y - q . xu , b - g . liu , and d.g . pettifor , 037204 ( 2003 ) . zhao , f. matsukura , k. takamura , e. abe , d. chiba , and h. ohno , appl . lett . * 79 * , 2776 ( 2001 ) . k. ono , j. okabayashi , m. mizuguchi , m. oshima , a. fujimori , and h. akinaga , j. appl . phys . * 91 * , 8088 ( 2001 ) . zhao and a. zunger , , 132403 ( 2005 ) . deng , j.h . zhao , j.f . niu , f.h . yang , x.g . wu , and h.z . zheng , j. appl . phys . * 99 * , 093902 ( 2006 ) . j. kudrnovsk and v. drchal , , 7515 ( 1990 ) . i. turek , v. drchal , j. kudrnovsk , m. ob , and p. weinberger , _ electronic structure of disordered alloys , surfaces and interfaces _ ( kluwer , boston - london - dordrecht , 1997 ) . vosko , l. wilk , and m. nusair , can . j. phys . * 58 * , 1200 ( 1980 ) . savrasov , and d.yu . savrasov , , 12181 ( 1992 ) . v. heine , j.h . samson , and c.m.m . nex , j. phys . phys . * 11 * , 2645 ( 1981 ) . v. heine and j.h . samson , j. phys . * 13 * , 2155 ( 1983 ) . h. hasegawa , j. phys . jpn . * 46 * , 1504 ( 1979 ) . pettifor , j. magn . mater * 15 - 18 * , 847 ( 1980 ) . staunton , b.l . gyorffy , a.j . pindor , g.m . stocks , and h. winter , j. phys . f * 15 * , 1387 ( 1985 ) . pindor , j. staunton , g.m . stocks , h. winter , j. phys . f * 13 * , 979 ( 1983 ) . e. sasaio@xmath7lu , i. galanakis , l.m . sandratskii , and p. bruno , j. phys : condens . matter * 17 * 3915 ( 2005 ) . sandratskii , r. singer , and e. sasiolu , , 184406 ( 2007 ) . m. pajda , j. kudrnovsk , i. turek , v. drchal , and p. bruno , phys . b * 64 * , 174402 ( 2001 ) . liechtenstein , m.i . katsnelson and v.a . gubanov , j. phys.f : met.phys . * 14 * , l125 ( 1984 ) . a. i. liechtenstein , m. i. katsnelson , v. p. antropov , v. a. gubanov , j. magn . magn . mater . * 67 * , 65 ( 1987 ) . liechtenstein , m.i . katsnelson and v.a . gubanov , solid.state.commun . * 51 * , 1232 ( 1984 ) . a.i . liechtenstein , m.i . katsnelson , v.p . antropov and v.a . gubanov , j.magn.magn.mater . * 21 * , 35 ( 1988 ) . gubanov , a.i . liechtenstein , a.v . postnikov _ magnetism and the electronic structure of crystals _ , edited by m. cardona , p. fulde , k. von klitzing , h .- j . queisser ( springer , berlin , 1992 ) . see , e.g. , o.k . andersen , o. jepsen , and d. gltzel , in _ highlights of condensed matter theory _ , edited by f. bassani ( north - holland , amsterdam , 1985 ) , p.59 . v. heine , _ solid state physics _ * 35 * ( academic press , new york ) , 1 ( 1980 ) . a. oswald , j. phys . f * 15 * , 193 ( 1985 ) . f. ducastelle , `` order and phase stability in alloys '' ( north - holland , amsterdam , 1991 ) . m. sluiter , and p.e . a. turchi , phys . b * 40 * , 11215 ( 1989 ) . connolly and a.r . williams , phys . b * 27 * , 5169 ( 1983 ) . z.w . lu , s .- h . wei , a. zunger , s. frota - pessoa , and l.g . ferreira , phys . b * 44 * , 512 ( 1991 ) . l.m . sandratski , j. phys . : condens . matter 3 , 8565 ( 1991 ) . i. turek , j. kudrnovsk , v. drchal , and p. bruno , philos . mag . * 86 * , 1713 ( 2006 ) . wang , r.e . prange , and v. korenman , phys . b * 25 * , 5766 ( 1982 ) . j. rusz , l. bergqvist , j. kudrnovsk , and i. turek , phys . b * 73 * , 214412 ( 2006 ) . l. jiang , q. feng , y. yang , z. chen , and z. huang , sol . . comm . * 139 * , 40 ( 2006 ) . s. curtarolo , d. morgan and g. ceder , comp . phase diagrams and thermochemistry * 29 * , 163 ( 2005 ) .

we present calculations of the exchange interactions and curie temperatures in cr - based pnictides and chalcogenides of the form crx with x = as , sb , s , se and te , and the mixed alloys crasx with x = sb , s , se , and te . the disorder effect in the as - sublattice for crasx ( x = sb , s , se , te ) alloys is taken into account via the coherent potential approximation ( cpa ) .

we present calculations of the exchange interactions and curie temperatures in cr - based pnictides and chalcogenides of the form crx with x = as , sb , s , se and te , and the mixed alloys crasx with x = sb , s , se , and te . the calculations are performed for zinc blende ( zb ) structure for 12 values of the lattice parameter between 5.44 and 6.62 , appropriate for some typical ii - vi and iii - v semiconducting substrates . electronic structure is calculated via the linear muffin - tin - orbitals ( lmto ) method in the atomic sphere approximation ( asa ) , using empty spheres to optimize asa - related errors . whenever necessary , the results have been verified using the full - potential version of the method , fp - lmto . the disorder effect in the as - sublattice for crasx ( x = sb , s , se , te ) alloys is taken into account via the coherent potential approximation ( cpa ) . exchange interactions are calculated using the linear response method for the ferromagnetic ( fm ) reference states of the alloys , as well as the disordered local moments ( dlm ) states . these results are then used to estimate the curie temperature from the low and high temperature side of the ferromagnetic / paramagnetic transition . estimates of the curie temperature are provided , based on the mean field and the more accurate random phase approximations . dominant antiferromagnetic exchange interactions for some low values of the lattice parameter for the fm reference states in crs , crse and crte prompted us to look for antiferromagnetic ( afm ) configurations for these systems with energies lower than the corresponding fm and dlm values . results for a limited number of such afm calculations are discussed , identifying the afm[111 ] state as a likely candidate for the ground state for these cases .

0705.2736

c

we present sensitive observations of a complete sample of compact polarized radio sources , as part of a deep integration of the elais n1 region made with the synthesis telescope at the dominion radio astrophysical observatory . a total of 83 polarized sources was detected in the ten - field mosaic . the distribution of fractional polarization of faint polarized sources was investigated with a monte - carlo analysis that generates synthetic source lists with the same noise statistics and observational selection criteria as the data . maximum - likelihood fits of the synthetic source lists to the data in the @xmath134 - @xmath135 plane yielded a best fitting gauss - hermite function with @xmath136 , @xmath137 for the distribution of intrinsic fractional polarization . the data demonstrate a trend of increasing fractional polarization with decreasing flux density . polarized source counts from the elais n1 deep field are presented down to @xmath138 . we find that the euclidean - normalized polarized counts remain flat below @xmath117 . the distribution of fractional polarization derived from our monte carlo analysis is convolved with the total - intensity source counts to produce a prediction of the polarized source counts . the predicted euclidean - normalized polarized counts are nearly flat to @xmath139 , in good agreement with the data . however , the data at the faintest polarized flux densities suggests a continuing trend of increased polarization fraction with decreasing flux density . the near - infrared color - color diagram for host galaxies identified with the polarized sources in the elais n1 field shows that most of the host galaxies are ellipticals , or galaxies for which the near - infrared spectrum is dominated by stochastically heated very small grains , presumably from the vicinity of an agn . some of the host galaxies appear to have pah bands in their near - infrared spectrum , but the morphological resemblance with ellipticals , and the fact that some of these polarized sources are resolved radio galaxies in the first survey , indicates that these objects also harbour agn . we suggest that the higher degree of polarization indicates a difference between agn observed at a flux density of hundreds of mjy , and fainter agn .

survey north 1 region ( elais n1 ) as part of the dominion radio astrophysical observatory _ planck _ a few host galaxies have colors that suggests significant pah emission in the near - infrared . a small fraction , 12% , of the polarized sources are not detected in the swire data .

we present deep polarimetric observations at 1420 mhz of the european large area _ iso _ survey north 1 region ( elais n1 ) as part of the dominion radio astrophysical observatory _ planck _ deep fields project . by combining closely spaced apertures synthesis fields , we image a region of 7.43 square degrees to a maximum sensitivity in stokes and ofjy beam , and detect 786 compact sources in stokes . of these , 83 exhibit polarized emission . we find that the differential source counts ( ) for polarized sources are nearly constant down tojy , and that these faint polarized radio sources are more highly polarized than the strong source population . the median fractional polarization is% for polarized sources with stokes flux density between 1 and 30 mjy ; approximately three times larger than sources with mjy . the majority of the polarized sources have been identified with galaxies in the _ spitzer _ wide area infrared extragalactic survey ( swire ) image of elais n1 . most of the galaxies occupy regions in the irac 5.8/3.6 m vs. m color - color diagram associated with dusty agns , or with ellipticals with an aging stellar population . a few host galaxies have colors that suggests significant pah emission in the near - infrared . a small fraction , 12% , of the polarized sources are not detected in the swire data . none of the polarized sources in our sample appears to be associated with an actively star - forming galaxy .

astro-ph9912004

c

we have presented detailed models for the formation of disk galaxies in both a cdm and a mond universe . in the case of dm , the structure of the disk is governed by the mass and angular momentum of the proto - galaxy . for mond , however , the distribution of angular momenta of proto - galaxies is not known , and we have instead assigned scale lengths to the disks that are in agreement with observations . in addition to these recipes that determine the structure and dynamics of the disks , we include recipes that describe how ( part of ) the gas is transformed into stars over the lifetime of the galaxy . we take account of a stability related threshold density for star formation and feedback from supernovae . the models have been tuned to fit the observed @xmath16-band tf relation , and compared to numerous observations . both models do remarkably well in reproducing a wide variety of observations of disk galaxies that span several orders of magnitude in both luminosity and surface brightness . to more easily compare the different models a summary of the results is given in table [ tab : results ] , where we indicate , for each of the observational constraints discussed in this paper , whether the models are consistent with the data or not . in the case of dm ( model l5 ) , sn feedback is required to yield a tf slope as steep as observed . although the mond model without feedback ( m1 ) is in excellent agreement with the empirical tf relation , it does not reproduce the observed deficit of hsb dwarf galaxies with small mass - discrepancies . this can be remedied by introducing sn feedback ( model m2 ) , but at the cost of a tf relation which is too steep and reveals an amount of scatter that is only marginally consistent with the data . this is a serious problem for mond ; since pure dynamics already predicts a tf relation as steep as observed , it leaves virtually no room for any other galaxy characteristics to vary systematically with mass . the dm model is consistent with all observational constraints against which we have tested it . consequently , we strongly disagree with the picture that emerges from the literature ( for example , mb98a , mb98b , and mcgaugh 1998 ) , that the dm hypothesis suffers from numerous serious fine - tuning problems , which do not seem to have a clear - cut solution , whereas mond is free from such problems and capable of fitting virtually everything . we have shown here that once the dm model is tuned to fit the slope of the tf relation , it automatically passes all the tests devised by mb98a and mcgaugh ( 1998 ) to argue against it . in particular , our dm model reproduces * a close correlation between global mass - to - light ratio and surface brightness , such that hsb and lsb galaxies follow the same tf relation without a systematic offset . * the presence of a characteristic acceleration as observed . * a close relation between the characteristic acceleration , @xmath145 , and central surface brightness of the form @xmath176 . this is a remarkable result : there is no obvious reason why the dm model would reveal a characteristic acceleration , unlike in the case of mond , where it is integral to the theory . furthermore , it is encouraging that the same feedback parameters that yield the correct tf slope , result in an amount of scatter around the @xmath145@xmath140 relation that is in excellent agreement with observations , and in addition explains the observed absence of hsb dwarfs with small mass - discrepancies . of the list of observational facts about disk galaxies presented in [ sec : intro ] , the upper limit to the observed surface brightness of disk galaxies ( item 5 ) has not been addressed by our models . several studies have shown , however , that in the dm picture the presence of a maximum central surface brightness of disk galaxies is related to stability arguments ( dss97 ; mo et al . 1998 ; scorza & van den bosch 1998 ) . in [ sec : tfmond ] we have shown that the same stability argument also yields an upper limit on the central surface brightness of disks under the hypothesis of mond ( cf . milgrom 1989 ) . henceforth , both the dm and the mond models are consistent with the observational constraint of item 5 , as long as disk stability is taken into account . the main problem for cdm is related to rotation curve shapes . we have compared predicted rotation curve shapes within the context of both mond and cdm . only in the case of lsb systems do the two scenarios yield rcs that are significantly different . several studies have pointed out that dark halos with a steep cusp are inconsistent with the central rotation curves of lsb and dwarf galaxies . this problem is also evident from the fact that at accelerations of @xmath177 the dm model predicts mass - to - light ratios that are slightly too high . however , it is important to realize that most data on the more massive lsb disks is severely affected by beam - smearing . this tends to underestimate the central gradients of rotation velocities , especially in galaxies that have a central hole in their hi distribution . when beam smearing is taken into account , the hi rotation curves of massive lsb galaxies are consistent with dark halos that follow a nfw density profile ( van den bosch et al . further indications that beam smearing plays an important role comes from a comparison of the ratio @xmath170 . both the dm and the mond models predict ratios that are lower than observed , which is most likely due to beam smearing . therefore , in order to discriminate between mond and dm , high resolution rotation curves of lsb disk galaxies are required . such data is currently only available for dwarf galaxies , which , because of their relative proximity , have been observed with high spatial resolution . it has been demonstrated convincingly that the rotation curves of these low - mass systems are inconsistent with centrally cusped dark matter halos . it is currently still under debate whether mond can fit these rotation curves . if it can , this is where mond has a clear advantage over cdm , unless ( close to ) constant density cores can be produced in dark halos that form in a cdm universe ( see e.g. , navarro , eke & frenk 1996 , kravtsov et al . 1998 , and bullock et al . 1999 for possible solutions ) . one might argue that within the dm scenario one can fit basically anything as long as there are a sufficient number of free parameters . in that respect it is important to realize that model l5 has only a very limited number of truly free parameters . where possible , we have used parameters that have either empirically determined values , or that are otherwise constrained : the dark halo properties are taken from high resolution numerical simulations combined with the press - schechter formalism , @xmath19 is constrained by stellar population models , the star formation recipe uses values for @xmath7 and the schmidt law that have been determined empirically , and @xmath88 is constrained by numerical simulations . this leaves only @xmath133 and @xmath27 as real free parameters . tuning these two parameters to obtain a tf relation with the observed slope of @xmath178 , the model predicts gas mass fractions , characteristic accelerations , an @xmath128 - @xmath140 `` conspiracy '' , and global mass - to - light ratios which are in excellent agreement with observations , without additional tuning of the parameters . furthermore , we note that while mond is often presented as being nearly free of fine - tuning , in order to match the systematic properties of disk galaxies it is necessary to adjust the mond feedback parameters to the same degree as required for dm . probably the most amazing aspect of the models presented here , is that the dm and mond models are so very similar . however , both mond and dm were constructed to fit the rotation curves of disk galaxies . the fact that both theories correctly predict many other properties , which are themselves closely related to the internal dynamics ( i.e. , tf relation , gas mass fractions that are set by stability related threshold densities , sn feedback whose efficiency depends on the escape velocity , etc ) , should therefore not be seen as too remarkable . it has often been argued that even if mond turns out to not be correct , one should provide an explanation as to why it fits the properties of galaxies so well . the demonstrations in this paper suggest that _ any _ theory which yields stable disks and fits their rotation curves , would probably perform as well as any of the models presented here . this work has benefited greatly from discussions with george lake . we are grateful to stacey mcgaugh for sending us his data in electronic format , and to the anonymous referee for his suggestions that helped to improve the paper . fvdb was supported by nasa through hubble fellowship grant # hf-01102.11 - 97.a awarded by the space telescope science institute , which is operated by aura for nasa under contract nas 5 - 26555 .

the dm and mond models are almost indistinguishable . they both yield gas mass fractions and dynamical mass - to - light ratios which are in good agreement with observations . both models reproduce the narrow relation between global mass - to - light ratio and central surface brightness , and reveal a characteristic acceleration , contrary to claims that these relations are not predicted by dm models . both models require sn feedback in order to reproduce the lack of high surface brightness dwarf galaxies . however , the introduction of feedback to the mond models steepens the tf relation and increases the scatter , making mond only marginally consistent with observations . however , the dm rotation curves are only slightly steeper than those of mond , and are only marginally inconsistent with the poor resolution data on lsb galaxies .

we present detailed semi - analytical models for the formation of disk galaxies both in a universe dominated by dark matter ( dm ) , and in one for which the force law is given by modified newtonian dynamics ( mond ) . we tune the models to fit the observed near - infrared tully - fisher ( tf ) relation , and compare numerous predictions of the resulting models with observations . the dm and mond models are almost indistinguishable . they both yield gas mass fractions and dynamical mass - to - light ratios which are in good agreement with observations . both models reproduce the narrow relation between global mass - to - light ratio and central surface brightness , and reveal a characteristic acceleration , contrary to claims that these relations are not predicted by dm models . both models require sn feedback in order to reproduce the lack of high surface brightness dwarf galaxies . however , the introduction of feedback to the mond models steepens the tf relation and increases the scatter , making mond only marginally consistent with observations . the most serious problem for the dm models is their prediction of steep central rotation curves . however , the dm rotation curves are only slightly steeper than those of mond , and are only marginally inconsistent with the poor resolution data on lsb galaxies .

astro-ph0004044

r

table [ modelle ] contains the best fit model for each image . together with the galaxy name , the filter , and the referring image , with integration time , and run i d , we list the inclination , the best fitting function for the z - distribution , the calibration index ( ref . section [ calib ] ) , and the central surface brightness of the model , without correcting for inclination . according to the distance tabulated in table [ galaxies ] , the cut - off radius @xmath27 is given in kpc and arcsec as well as the scalelength @xmath25 and the vertical scaleheight @xmath26 which is normalised to the isothermal case , being two times an exponential scaleheight @xmath38 . for the seven galaxies with available images in more than one filter , we do not see any correlation of fitted parameters with different wavelength , although we find the same inclination angle for the best fitted disk within the range of the errors . appendix a shows the best fitting model as an overlay to selected radial profiles for each image . the subsequent analysis of the distribution for the different parameters concerning the formation and evolution of galaxies will be given in forthcoming papers . l c c [email protected] lcc [email protected] [email protected] [email protected]@.l [email protected]@.l galaxy&filter&image&&@xmath77&cali.&@xmath83 & & & & & + & & & & & & mag@xmath62 & & & & & & & & + ' '' '' & & & & & & & & & & & + + eso 112 - 004 & r&40e3 & 87&5 & sech & l & 21.14 & 45&0 & & & 22&7 & & & 2&8 & + eso 150 - 014 & r&20e3 & 90&0 & sech & l & 22.00 & 64&1 & 33&27 & & 23&4 & 12&14 & & 5&5 & 2&85 + ngc 585 & r&20l1 & 88&0 & sech & e & 21.76 & 70&0 & 24&49 & & 32&7 & 11&44 & & 10&6 & 3&71 + eso 244 - 048 & r&15e3 & 87&0 & @xmath84 & l & 20.54 & 45&4 & 19&17 & & 13&7 & 5&78 & & 7&2 & 3&04 + ngc 973 & r&10l1 & 89&5 & @xmath84 & e & 20.83 & 105&0 & 33&74 & & 51&5 & 16&55 & & 12&4 & 3&99 + ugc 3326 & r&30l1 & 88&0 & sech@xmath17&e & 21.50 & 101&5 & 28&45 & & 70&4 & 19&74 & & 4&7 & 1&32 + ugc 3425 & r&30l1 & 87&0 & sech@xmath17&e & 21.01 & 80&5 & 22&26 & & 29&2 & 8&08 & & 7&4 & 2&05 + ngc 2424 & r&15l1 & 86&5 & @xmath84 & e & 20.52 & 112&0 & 23&40 & & 31&7 & 6&62 & & 11&9 & 2&49 + ic 2207 & r&10l2 & 86&5 & @xmath84 & e & 21.04 & 56&0 & 17&67 & & 38&2 & 12&06 & & 6&7 & 2&12 + eso 564 - 027 & r&30e2 & 88&0 & sech & l & 21.11 & 140&4 & 18&33 & & 50&9 & 6&65 & & 6&6 & 0&86 + eso 436 - 034 & g&60e3 & 88&0 & sech & l & 20.96 & 82&1 & 18&32 & & 22&6 & 5&04 & & 6&8 & 1&52 + eso 319 - 026 & g&30e3 & 86&5 & sech & e & 21.51 & 60&1 & 13&21 & & 14&2 & 3&12 & & 3&1 & 0&68 + eso 319 - 026 & r&30e2 & 86&0 & @xmath84 & i & 21.14 & 63&0 & 13&84 & & 14&7 & 3&23 & & 3&2 & 0&70 + eso 319 - 026 & r&30e3 & 88&0 & sech@xmath17&i & 21.94 & 61&9 & 13&60 & & 13&2 & 2&90 & & 2&8 & 0&62 + eso 319 - 026 & i&30e3 & 88&0 & sech & l & 21.01 & 64&8 & 14&24 & & 14&3 & 3&14 & & 3&4 & 0&75 + eso 321 - 010 & g&30e3 & 88&0 & sech & l & 20.11 & 64&4 & 12&34 & & 19&9 & 3&81 & & 4&7 & 0&90 + eso 321 - 010 & r&30e3 & 88&0 & sech & l & 19.54 & 64&8 & 12&42 & & 21&6 & 4&14 & & 5&0 & 0&96 + ngc 4835a & r&40e3 & 85&5 & @xmath84 & l & 20.92 & 90&0 & 18&56 & & 40&4 & 8&33 & & 7&9 & 1&63 + eso 575 - 059 & r&15e2 & 87&0 & sech@xmath17&l & 20.78 & 60&5 & 17&53 & & 22&6 & 6&55 & & 6&4 & 1&86 + eso 578 - 025 & g&30e2 & 86&5 & @xmath84 & l & 21.58 & 50&7 & 20&63 & & 17&0 & 6&92 & & 7&5 & 3&05 + eso 578 - 025 & g&30e3 & 86&5 & @xmath84 & l & 21.69 & 50&7 & 20&63 & & 17&0 & 6&92 & & 7&5 & 3&05 + eso 578 - 025 & r&30e2 & 86&0 & @xmath84 & l & 21.02 & 50&4 & 20&51 & & 14&9 & 6&06 & & 7&4 & 3&01 + eso 578 - 025 & i&30e2 & 86&0 & @xmath84 & l & 20.27 & 47&5 & 19&33 & & 20&3 & 8&26 & & 7&6 & 3&09 + eso 446 - 018 & r&30e2 & 86&5 & sech & l & 20.45 & 75&6 & 22&78 & & 23&7 & 7&14 & & 3&9 & 1&18 + ic 4393 & r&30e2 & 87&0 & @xmath84 & l & 20.30 & 75&6 & 12&88 & & 29&8 & 5&08 & & 5&8 & 0&99 + eso 581 - 006 & r&30e3 & 86&5 & @xmath84 & l & 21.33 & 55&8 & 11&07 & & 18&2 & 3&61 & & 5&6 & 1&11 + eso 583 - 008 & r&30e3 & 87&0 & @xmath84 & i & 20.84 & 56&2 & 26&70 & & 14&0 & 6&65 & & 3&8 & 1&81 + ugc 10535 & r&25e2 & 88&0 & sech@xmath17&i & 21.50 & 41&4 & 20&57 & & 11&4 & 5&66 & & 4&3 & 2&14 + ngc 6722 & r&10e3 & 86&5 & sech & l & 19.47 & 86&0 & 30&77 & & 21&6 & 7&73 & & 6&2 & 2&22 + eso 461 - 006 & r&60e3 & 87&5 & @xmath84 & l & 20.96 & 61&6 & 23&37 & & 21&1 & 8&01 & & 3&8 & 1&44 + ic 4937 & g&20e1 & 88&5 & sech@xmath17&l & 22.73 & 74&9 & 10&41 & & 93&8 & 13&04 & & 5&6 & 0&78 + ic 4937 & r&30e3 & 88&0 & sech@xmath17&l & 21.29 & 83&5 & 11&61 & & 32&1 & 4&46 & & 5&9 & 0&82 + ic 4937 & i&20e1 & 88&5 & sech & e & 19.92 & 78&1 & 10&86 & & 38&6 & 5&37 & & 7&2 & 1&00 + eso 528 - 017 & g&30e3 & 86&5 & @xmath84 & l & 21.98 & 57&6 & 22&54 & & 20&5 & 8&02 & & 3&5 & 1&37 + eso 528 - 017 & r&60e3 & 86&5 & @xmath84 & l & 21.28 & 55&1 & 21&56 & & 19&9 & 7&79 & & 2&7 & 1&06 + eso 528 - 017 & i&30e3 & 86&5 & sech & l & 20.89 & 50&0 & 19&57 & & 22&6 & 8&85 & & 3&2 & 1&25 + eso 187 - 008 & r&30e3 & 85&5 & @xmath84 & l & 20.77 & 50&0 & 13&65 & & 15&1 & 4&12 & & 4&4 & 1&20 + eso 466 - 001 & i&40e3 & 87&0 & @xmath84 & e & 19.52 & 52&6 & 23&76 & & 13&0 & 5&87 & & 8&2 & 3&70 + eso 189 - 012 & g&60e3 & 87&0 & @xmath84 & l & 21.50 & 56&2 & 29&81 & & 26&8 & 14&21 & & 3&8 & 2&02 + eso 189 - 012 & r&30e3 & 86&5 & @xmath84 & l & 20.71 & 56&5 & 29&97 & & 20&6 & 10&93 & & 3&4 & 1&80 + eso 189 - 012 & i&20e3 & 87&0 & sech & l & 20.44 & 54&7 & 29&01 & & 22&7 & 12&04 & & 3&2 & 1&70 + eso 533 - 004 & r&20e1 & 88&0 & @xmath84 & l & 20.32 & 68&4 & 11&12 & & 33&1 & 5&38 & & 7&7 & 1&25 + ic 5199 & g&30e3 & 86&5 & @xmath84 & l & 21.55 & 64&1 & 20&44 & & 19&9 & 6&35 & & 5&6 & 1&79 + ic 5199 & i&30e3 & 86&5 & @xmath84 & l & 19.97 & 57&6 & 18&37 & & 19&9 & 6&35 & & 4&9 & 1&56 + eso 604 - 006 & r&30e3 & 90&0 & sech & l & 21.22 & 70&6 & 34&54 & & 27&9 & 13&65 & & 3&8 & 1&86 + [ modelle ] the different fitting methods were independently developed within two diploma theses ( ltticke 1996 , schwarzkopf 1996 ) . the quality of the data basis for each project was the same . from the sample presented here there were five objects in common . these are used to compare the two methods and determine the quantitative difference of the derived parameters . r l c c c c cl galaxy&@xmath77 & @xmath40 & @xmath18 & @xmath25 & @xmath27 + & & [ 1.0ex][@xmath85 & [ 1.0ex ] [ @xmath86 & [ 1.0ex ] [ @xmath86 & [ 1.0ex ] [ @xmath86 + + & sech&88.0&5.0&21.6&64.8 + [ 1.1ex]eso 321 - 010 r \{&@xmath84&88.0&5.2&25.9&64.1 + & @xmath84&87.5&3.8&21.1&61.6 + [ 1.1ex]eso 461 - 006 r \{&@xmath84&87.5&3.6&16.2&61.2 + & @xmath84&86.5&6.1&26.8&83.5 + [ 1.1ex]ic 4937 r \{&sech&89.5&7.0&27.4&75.6 + & @xmath84&86.5&3.2&19.3&56.2 + [ 1.1ex]eso 189 - 012 r \{&@xmath84&88.0&3.6&21.3&56.9 + & sech & 88.5&4.1&39.8&67.7 + [ 1.1ex]eso 604 - 006 r \{&sech@xmath17&89.5&3.2&24.5&73.4 + table [ verglrlus ] shows the results for the five images . the mean deviation in the determined inclination is @xmath87 and 12.4% for the scaleheight ( ranging from 5.0%-26.6% ) whereas for three images different functions for the z distribution were used . the mean difference for the radial scalelength is 20.6% ( 2.1%-47.2% ) and 4.2% for the determination of the cut - off radius . + a subsequent analysis shows that it is not possible to ascribe the sometimes quite large discrepancies to the quality of the individual method . it turns out , that the main problem is the non - uniform determination of the fitting area . the intrinsic asymmetric variations of a real galactic disk compared to the model enforce a more subjective restriction of the galaxy image to the fitting region , whereby for example one has to exclude the bulge area and the dust lane . + this finding is in agreement with the study of knapen & van der kruit ( @xcite ) who compared published values of the scalelength and find an average value of 23% for the discrepancy between different sources . as already mentioned by schombert & bothun ( @xcite ) the limiting factor for accuracy of the decomposition is not the typical s / n from the ccd - telescope combination nor the errors in the determination of the sky background , but the deviation of real galaxies from the standard model . in our former study ( paper i ) with an earlier method to adapt equation ( [ numcalhatl ] ) , 20 of our 45 galaxy images have already been used . we decided to re - use them in this study to get models for as many galaxies as possible in a homogeneous way . additionally , paper i only presents the best fit values for the isothermal model , and uses a different definition of the cut - off radius . + only three galaxies are in common with the sample of de grijs ( @xcite ) : eso 564 - 027 , eso 321 - 010 , and eso 446 - 018 . the mean difference for the scalelength is 10.4 % ( ranging from 1.2%-20.3% ) and for the scale height ( normalised to the isothermal case ) 4.0% ( 0.0%-8.2% ) . for the remaining galaxies there are no models in the literature . our model only represents a rather simple axisymmetric three - dimensional model for a galactic disk , consisting of an one component radial exponential disk with three different laws for the density distribution in the z - direction and a sharp outer truncation . therefore it does not include additional components , such as bulges , bars , thick disks , or rings , and can not deal with any asymmetries . features like spiral structure or warps are not included , whereas reshetnikov & combes ( @xcite ) multiply their exponential disk by a spiral function introducing an expression to characterise an intrinsic warp depending on the position angle outside of a critical radius . + the choice of our fitting area tries to avoid the dust lane , possible only for almost edge - on galaxies , as a first step to account for the dust influence ( cf . section [ dust ] ) . examples of models including a radiative transfer with an extinction coefficient @xmath88 can be found in xilouris et al . ( @xcite ) . however , introducing more and more new components and features automatically increases the amount of free parameters . therefore we restricted our model to the described six parameters , to obtain statistically meaningful characteristics for galactic disks . + in the following we demonstrate that a simple disk model omitting the bulge component and the dust lane give indeed reasonable parameters . we have studied the influence of the bulge for some of our objects including the earliest type galaxy in our sample ( eso 575 - 059 ) presented here . we have subtracted our derived disk model from the galaxy and then tried to find the best representation for the remaining bulge by a de vaucouleurs @xmath89 or an exponential model . taking the slope of the vertical profile at @xmath90 and a fixed axis ratio , we have constructed the 2-dimensional model of the bulge . in agreement with andredakis et al . ( @xcite ) we find , that bulges of early type galaxies are better fitted by an exponential profile than by a @xmath89 . figure [ bulge ] shows the resulting vertical and radial cuts for eso 575 - 059 together with the models . despite the deviation between @xmath91 and @xmath92 which could be attributed to an additional component ( inner disk or bar ) , we do not find any evidence for changing our disk model due to the influence of the bulge . + therefore we conclude , that it is possible to nearly avoid any influence of the central component by fitting outside the clearly visible bulge region . dust disturbs the light profile by a combination of absorption and extinction and the net effect has to be calculated by radiation transfer models . therefore it is not obvious that outside the `` visible '' dust lane , which is excluded for the fitting area as a first step , the dust will not play a major role in shaping the light distribution . xilouris et al . ( @xcite ) , bianchi et al . ( @xcite ) , de jong ( @xcite ) , and byun et al . ( @xcite ) have recently addressed this problem in more detail . although they investigated the influence of the dust on the light distribution by quoting best fit structural parameters for the star - disk as well as the dust - disk , they did not quantify the influence on the star - disk parameters derived by standard fitting methods without dust . even kylafis & bahcall ( @xcite ) state within their fundamental paper on finding the dust distribution for ngc 891 that `` in order to avoid duplication of previous work , we will take ... ( the values for the star - distribution estimated by standard fitting methods ) '' . + we checked the influence of the dust on our determined parameter set , by studying simulated galaxy images with three different dust distributions . these _ dusty galaxies _ were kindly provided by simone bianchi who calculated images with known input parameters for the star and dust distribution with his monte carlo radiative transfer method ( bianchi et al . we have defined a _ worst _ , _ best _ , and _ transparent _ case , according to the dust distributions presented by xilouris et al . ( @xcite ) , of our mean stellar disk ( @xmath93 , @xmath94 , and @xmath95 ) . worst _ case is calculated with @xmath96 , @xmath97 , and @xmath98 , the _ best _ case with @xmath99 , @xmath100 , and @xmath101 , and a _ transparent _ case without dust . to be comparable we used the same method for selecting the fitting region by masking the `` visible '' dust lane and reserving a typical area for a possible bulge component using mean values for the _ transparent _ case . in contrast to our standard procedure we do not restrict the inclination range from the appearance of the dust lane , but specify the best @xmath72 model in the range @xmath102 in table [ dustvgl ] . for the models marked with a @xmath10 we pretend the correct input inclination . table [ dustvgl ] demonstrates , that even for the worst case we are able to reproduce the input parameters within the range of the typical 20% error discussed in section [ comp ] . it should be mentioned that for each case we overestimate the input scalelength and -height , whereas the determination of @xmath27 does not depend on the dust distribution . the implication on the distribution of the ratio @xmath55 will be discussed in a forthcoming paper . .comparison of the results for the stellar disk ( input : @xmath103 , @xmath104 , and @xmath105exp ) from our fitting procedure for the three different dust distributions . [ cols="<,^,^,^,^,^ , > , > " , ] trying to adapt a simple , perfect , and exact symmetric model to real galaxies always implies a compromise between the degree of any deviation and the final model ( section [ pec ] ) . the following list will provide some typical caveats found during the fit procedure which will characterise the quality of the specified model for individual galaxies . + * eso 112 - 004 * : warped , asymmetric , central part slightly tilted compared to disk , after fitting still remaining residuals + * eso 150 - 014 * : slightly warped , minor flatfield problems + * ngc 585 * : remaining residuals + * eso 244 - 048 * : possible two component system , slope of inner radial profile significantly higher than of an outer one , final model fits the inner parts + * ngc 973 * : one side disturbed by stray light of nearby star , seems to be radially asymmetric , remaining residuals + * ugc 3425 * : superimposed star on one edge + * ngc 2424 * : model does not fit very well without obvious reason + * eso 436 - 034 * : strong bulge component , possibly barred , hard to pinpoint final model , remaining residuals + * eso 319 - 026 * : outer parts show u - shaped behaviour , remaining residuals , therefore large ( @xmath106 ) difference in inclination angle + * eso 321 - 010 * : u - shaped , no clear major axis visible , therefore uncertain rotation angle , bar visible , bulge rotated against disk + * ngc 4835a * : strong residuals + * eso 446 - 018 * : the different sides of the disk are asymmetric visible in radial profiles and on the contour plot + * ic 4393 * : similar to ngc 4835a + * eso 581 - 006 * : galaxy shows typical late type profile , @xmath27 questionable , but nevertheless final model seems to fit well + * eso 583 - 008 * : disturbed by superimposed star , shows warp feature and a bar structure , @xmath27 questionable , remaining residuals + * ugc 10535 * : one side slightly extended + * ngc 6722 * : only one side observed , bulge rotated against disk , barred , strongly disturbed by dust absorption , radial extension visible , therefore @xmath27 should be treated with caution + * eso 461 - 006 * : minor flatfield problem seems to cause asymmetry , although model looks fine + * ic 4937 * : similar to ngc 6722 , dominating bulge , small disk , model significantly different compared to the i and r image , model possibly hampered by strong dust lane + * eso 578 - 025 * : bar visible + * eso 466 - 001 * : maybe two components , final model represents only inner part , outer part clearly different from normal disk component + * eso 189 - 012 * : slightly warped + * eso 533 - 004 * : similar to ngc 4835a , model fits the whole galaxy , leaving more or less no bulge component + * ic 5199 * : slightly radial asymmetric + * eso 604 - 006 * : only one side observed , bar structure visible the model limitations described above constrain the application of our fitting process . therefore we had to exclude about 20 galaxies from our original sample . they all show significant deviations from the simple geometry and an inclusion of their parameters obtained by forcing the model to fit the data will spoil the resulting parameter distribution . + one larger group classified mainly as s0 galaxies ( e.g. ngc 2549 , eso 376 - 009 , ngc 7332 , eso 383 - 085 , eso 506 - 033 ) shows a completely different behaviour of the luminosity distribution in the outer parts compared to the other galaxies . they all show an additional component , mainly characterised as an elliptical envelope . this is already visible in the contour plot , but becomes even more evident in a radial cut parallel to the major axis . in these cases the usual common curved decline of the profile ( e.g. eso 578 - 025 ) is missing , and is replaced by a more or less straight decline into the noise level , sometimes even by an upwards curved profile . fitting these luminosity distribution by our one component exponential disk with cut - off , will therefore naturally provide parameters qualitatively different compared to late type disks . this will be discussed in detail in a forthcoming paper . + another group consists of galaxies dominated mainly by their bulges , whereas the disk is only an underlying component , partly characterized as having thick boxy bulges ( dettmar & ltticke @xcite ) , e.g. ic 4745 , eso 383 - 005 , although there are also pure elliptical bulges ( e.g. eso 445 - 049 , ngc 6948 ) . + in the case of eso 383 - 048 and eso 510 - 074 the radial profiles clearly indicate that a more complex model will be needed to fit these kind of multicomponent galaxies . galaxies like ugc 7170 or eso 113 - 006 were excluded due to their strong warps , which made it impossible to fit the model in a consistent way . mainly late type galaxies such as eso 385 - 008 , ic4871 , ugc 1281 , or eso 376 - 023 show a too patchy and asymmetric light distribution , that any attempt to fit the profiles will give only very crude , low quality parameters . ugc 11859 and ugc 12423 were rejected due to their thin faint disks , which will maybe overcome by taking new images with longer integration times to get a higher signal - to - noise ratio , whereas ngc 5193a is completely embedded into the surface brightness distribution of its near companion . this work was supported by the _ deutsche forschungsgemeinschaft , dfg_. this research has made use of the nasa / ipac extragalactic database ( ned ) which is operated by the jet propulsion laboratory , california institute of technology , under contract with the national aeronautics and space administration . we have made use of the leda database ( www-obs.univ-lyon1.fr ) . the authors wish to thank simone bianchi , who kindly provided us _ dusty - galaxy images _ produced with his radiative transfer code . andredakis y.c . , peletier r. , balcells m. , 1995 , mnras 275 , 874 barteldrees a. , dettmar r .- j . , 1989 , in : dynamics and interaction of galaxies , ed . wielen , springer - verlag , heidelberg , p.348 barteldrees a. , dettmar r .- j . , 1994 , a&as 103 , 475 ( paper i ) bianchi s. , ferrara a. , giovanardi c. , 1996 , apj 465 , 127 bottinelli l. , gouguenheim l. , paturel g. , de vaucouleurs g. , 1983 , a&a 118 , 4 byun y. i. , freeman k. c. , kylafis n. d. , 1994 , apj 432 , 114 camm g.l . , 1950 , mnras 110 , 305 courteau s. , 1996 , apjs 103 , 363 dalcanton j.j . , spergel d.n . , summers f.j . , 1997 , apj 482 , 659 de grijs r. , peletier r.f . , van der kruit p.c . , 1997 , a&a 327 , 966 de grijs r. , 1998 , mnras 299 , 595 de jong r.s . , 1996 , a&a 313 , 377 de jong r.s . , 1996 , a&as 118 , 557 de vaucouleurs g. , 1959 , in : handbuch der physik liii , ed . flgge s. , springer - verlag berlin , p. 275 de vaucouleurs g. , de vaucouleurs a. , corwin jr . buta r.j . , paturel g. , fouqu p. , 1991 , third reference catalogue of bright galaxies , springer - verlag new york dettmar r .- j . , ltticke r. , 1999 , in : asp conference series volume 165 , eds . gibson b.k . , axelrod t.s . , putman m.e . , p. 95 fasano g. , filippi m. , 1998 , a&as 129 , 583 firmani c. , hernandez x. , gallagher j. , 1996 , a&a 308 , 403 guthrie b.n.g . , 1992 , a&as 93 , 255 freeman k.c . , 1970 , apj 160 , 811 fuchs b. , wielen r. , 1987 , in : the galaxy , eds . gilmore g. , carswell , d. reidel publishing co. dordrecht , p.375 hernquist l. , mihos j.c . , 1995 apj 448 , 41 kennicutt r.c . , 1989 , apj 344 , 685 knapen j.h . , van der kruit p.c . , 1991 , a&a 248 , 57 kylafis n.d . , bahcall j.n . , 1987 , apj 317 , 637 lauberts a. , 1982 , the eso / uppsala survey of the eso(b ) atlas , eso lauberts a. , valentijn e.a . , 1989 , the surface photometry catalogue of the eso - uppsala galaxies , european southern observatory , garching lin d.n.c . , pringle j.e . , 1987 , apj 320 , l87 ltticke r. , 1996 , diploma thesis , ruhr - universitt bochum marleau f.r . , simard l. 1998 , apj 507 , 585 mo h.j . , mao s. , white s.d.m . , 1998 , mnras 295 , 319 nelder j.a . , mead r. , 1965 , computer journal , vol . 7 , p. 308 nilson p. , 1973 , uppsala general catalogue of galaxies , uppsala pohlen , m. , dettmar , r .- j . , ltticke r. , 2000 , a&a 357 , l1 press w.h . , fannery b.p . , teukolsky s.a . , vetterling w.t . , 1988 , in : numerical recipes in c , cambridge university press reshetnikov v. , combes f. , 1998 , a&a 337 , 9 saio h. , yoshii y. , 1990 , apj 363 , 40 schombert j.m . , bothun g. , 1987 , aj 92 , 60 schwarzkopf u. , 1996 , diploma thesis , ruhr - universitt bochum seiden p. e. , schulman l. s. , elmegreen b. g. , 1984 , apj 282 , 95 shaw m.a . , 1993 , mnras 261 , 718 shaw m.a . , gilmore g. , 1989 , mnras 237 , 903 shaw m.a . , gilmore g. , 1990 , mnras 242 , 59 spitzer l. , 1942 , apj 95 , 329 struck - marcell c. , 1991 , apj 368 , 348 syer d. , mao s. , mo h.j . , 1999 , mnras 305 , 357 takamiya m. , 1999 , apjs 122 , 109 thuan t.x . , gunn j.e . , 1976 , pasp 88 , 543 toth g. , ostriker j.p . , 1992 apj 389 , 5 van der kruit p.c . , 1979 , a&as 38 , 15 van der kruit p.c . , 1987 , a&a 173 , 59 van der kruit p.c . , 1988 , a&a 192 , 117 van der kruit p.c . , searle , l. , 1981a , a&a 95 , 105 van der kruit p.c . , searle , l. , 1981b , a&a 95 , 116 van der kruit p.c . , searle , l. , 1982a , a&a 110 , 61 van der kruit p.c . , searle , l. , 1982b , a&a 110 , 79 wainscoat r.j . , 1986 , in : ph.d . thesis , australian national university wainscoat r.j . , freeman k.c . , hyland a.r , 1989 , apj 337 , 163 wyse r.f.g . , gilmore g. , franx m. , 1997 , ara&a 35 , 637 xilouris e.m . , byun y.i . , kylafis n.d . , paleologou e.v . , papamastorakis p. , 1999 a&a 344 , 868

these differences are not due to the individual method itself , but rather to the selected fitting region , which masks the bulge component , the dust lane , or present foreground stars . other serious limitations are small but appreciable intrinsic deviations of real disks compared to the simple input model .

we present detailed three - dimensional modelling of the stellar luminosity distribution for the disks of 31 relatively nearby ( 110 mpc ) edge - on spiral galaxies . in contrast to most of the standard methods available in the literature we take into account the full three - dimensional information of the disk . we minimize the difference between the observed 2d - image and an image of our 3d - disk model integrated along the line of sight . thereby we specify the inclination , the fitting function for the z - distribution of the disk , and the best values for the structural parameters such as scalelength , scaleheight , central surface brightness , and a disk cut - off radius . from a comparison of two independently developed methods we conclude , that the discrepancies e.g. for the scaleheights and scalelengths are of the order of . these differences are not due to the individual method itself , but rather to the selected fitting region , which masks the bulge component , the dust lane , or present foreground stars . other serious limitations are small but appreciable intrinsic deviations of real disks compared to the simple input model . + in this paper we describe the methods and present contour plots as well as radial profiles for all galaxies without previously published surface photometry . resulting parameters are given for the complete sample .

1110.4226

i

the observed global properties of galaxies obey a diverse set of scaling relations , which are fundamental tools to constrain models of galaxy formation and evolution . particularly interesting are the correlations between dynamics and luminosity / stellar mass or size , such as the tully - fisher ( tf ; * ? ? ? * ) relation for spirals , and the faber - jackson ( fj ; * ? ? ? * ) , @xmath4-@xmath5 @xcite , or fundamental plane ( fp ; * ? ? ? * ; * ? ? ? * ) relations for spheroids , because they link the luminous to the total mass of galaxies , thus providing insights into the interplay between their baryonic and dark matter components . historically , the importance of the tight tf and fp relations as secondary distance indicators has driven the community to assemble galaxy samples that meet strict selection criteria , in order to minimize systematic errors and scatter . while fp studies targeted elliptical and s0 galaxies , the samples used for tf applications have typically been restricted to late - type spirals with inclinations to the line - of - sight larger than 30 - 40 , preferably observed in red or infrared photometric bands to minimize extinction effects ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? these data sets are not ideal for characterizing the statistical properties of galaxies in general , because the excessive pruning means that they are not fair samples of the local universe . this issue is particularly problematic for the comparison with theoretical studies and numerical simulations of galaxy formation and evolution , which should be based on representative samples . for instance , tf studies of different classes of objects , such as s0s and early - type spirals ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? * and references therein ) , polar ring galaxies @xcite , barred spirals @xcite , and gas - rich dwarfs @xcite , have sometimes found disagreement with the tf relation of late - type spirals . most notably , gas - rich dwarfs lie systematically below the tf relation defined by bright galaxies . during the last decades , the interest in tf and fj - like scaling relations has shifted from cosmic flow applications to constraining galaxy formation models , and samples with broader morphological properties have been constructed specifically for this purpose ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? however , current large and homogeneous data sets include either spirals ( e.g. , in addition to the works mentioned above , @xcite and @xcite , hereafter ss11 ) or early - type galaxies ( e.g. , * ? ? ? * ; * ? ? ? . there have been attempts to move beyond the spiral / elliptical dichotomy , and uncover relations between stellar content and dynamics that hold for all galaxies , independent of morphology . @xcite found that all galaxies , from disks to spheroids and from dwarf spheroidals to giant ellipticals , lie on a two - dimensional surface defined by surface brightness , half - light radius , internal velocity and mass - to - light ratio . as a measure of internal velocity @xmath6 , they adopt either the rotational velocity @xmath7 for disks or the velocity dispersion @xmath5 for spheroids , thus the two types of systems are still treated separately ( especially because the sample is a large but heterogeneous collection of published data sets , for which either @xmath7 or @xmath5 is available ) . it is also important to point out that the search for scaling relations that are valid for all types of galaxies should make use of _ baryonic _ masses ( i.e. , the sum of stellar and gas masses ) instead of luminosities or stellar masses , because they could be more fundamental quantities . this is demonstrated by the fact that baryonic scaling relations hold for subsets of galaxies that do not follow the corresponding stellar relations . for example , the offset of the gas - rich dwarf galaxies from the stellar tf relation disappears when their gas mass is taken into account @xcite . the baryonic tf relation ( btf ; * ? ? ? * ) is linear over 5 orders of magnitude in ( stellar + gas ) mass , suggesting that the tf is fundamentally a relation between baryonic ( rather than luminous ) and total mass of the galaxy . intriguingly , although supported by limited statistics , there is some evidence that giant and dwarf ellipticals might lie on the same btf as the spirals @xcite . unfortunately , because they require estimates of both stellar and gas masses , btf samples ( e.g. , * ? ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ) are significantly smaller than tf ones . to summarize , it is still unclear if disk - dominated galaxies and spheroids obey the same dynamical scaling relations , mainly due to the lack of well - defined , representative samples of galaxies for which both rotation and stellar dispersion are measured . certainly , ellipticals might have no gas and no detectable rotation , and pure disks might have negligible stellar dispersions , but a significant fraction of local galaxies ( especially massive ones ) have a disk _ and _ a bulge ( e.g. * ? ? ? * ) , and the dynamical scaling relations should account for the smooth transition across galaxies with different bulge - to - disk ratios . as @xcite point out , a connection between the scaling relations of early - type and late - type galaxies is expected on the grounds that early - type systems are generally assumed to form through mergers of late - type ones . this is especially true at the high stellar mass end , where the blue sequence of star - forming disks merges onto the red sequence of passively - evolving , bulge - dominated galaxies ( e.g. * ? ? ? * ) , and the systems typically host a bulge and a disk . in this paper , we investigate dynamical scaling relations for a representative sample of @xmath0500 massive galaxies that are selected _ only _ by stellar mass ( @xmath8 m@xmath2 ) and redshift ( @xmath3 ) , as part of the ongoing galex arecibo sdss survey ( gass ; * ? ? ? * hereafter paper i ) . for these galaxies , we have homogeneous measurements of structural parameters and velocity dispersions from the sloan digital sky survey ( sdss ; * ? ? ? * ) , nuv@xmath9 colours from galex @xcite and sdss imaging , and hi masses and rotational velocities for the subset of objects detected at 21 cm with the arecibo radio telescope . this unique sample , which includes massive galaxies of all morphological types , allows us to investigate how objects that are typically not included in tf or fj / fp data sets scatter around those relations . in the spirit of works like those of @xcite and @xcite , we wish to establish if there is a _ fundamental _ correlation between baryonic mass and dynamics that is obeyed by the complete galaxy population , regardless of morphology . we show that , at least for the massive galaxies in our sample , such a relation does exist , and has a scatter smaller than 0.1 dex , comparable to that of the tf and fj relations applied to their respective pruned subsets . this paper is organized as follows . we summarize sample selection and measurements of relevant quantities in [ s_sample ] . we present the baryonic mass - velocity relations , starting with the tf and fj , in [ s_bvrel ] , and the velocity - size relations in [ s_vsize ] . we discuss our findings and conclude in [ s_disc ] . all the distance - dependent quantities in this work are computed assuming @xmath10 , @xmath11 and @xmath12 km s@xmath13 mpc@xmath13 .

we present dynamical scaling relations for a homogeneous and representative sample of massive galaxies , selected only by stellar mass ( m ) and redshift ( ) as part of the ongoing galex arecibo sdss survey . [ firstpage ] galaxies : kinematics and dynamics galaxies : evolution galaxies : fundamental parameters radio lines : galaxies

we present dynamical scaling relations for a homogeneous and representative sample of massive galaxies , selected only by stellar mass ( m ) and redshift ( ) as part of the ongoing galex arecibo sdss survey . we compare baryonic tully - fisher ( btf ) and faber - jackson ( bfj ) relations for this sample , and investigate how galaxies scatter around the best fits obtained for pruned subsets of disk - dominated and bulge - dominated systems . the bfj relation is significantly less scattered than the btf when the relations are applied to their maximum samples ( for the btf , only galaxies with hi detections ) , and is not affected by the inclination problems that plague the btf . disk - dominated , gas - rich galaxies systematically deviate from the bfj relation defined by the spheroids . we demonstrate that by applying a simple correction to the stellar velocity dispersions that depends only on the concentration index of the galaxy , we are able to bring disks and spheroids onto the same dynamical relation in other words , we obtain a generalized bfj relation that holds for all the galaxies in our sample , regardless of morphology , inclination or gas content , and has a scatter smaller than 0.1 dex . we compare the velocity - size relation for the three dynamical indicators used in this work , i.e. , rotational velocity , observed and concentration - corrected stellar dispersion . we find that disks and spheroids are offset in the stellar dispersion - size relation , and that the offset is removed when corrected dispersions are used instead . the generalized bfj relation represents a fundamental correlation between the global dark matter and baryonic content of galaxies , which is obeyed by all ( massive ) systems regardless of morphology . [ firstpage ] galaxies : kinematics and dynamics galaxies : evolution galaxies : fundamental parameters radio lines : galaxies

1110.4226

c

the main result of this work is the existence of a tight relation between baryonic mass and velocity that is independent of galaxy morphology . to first order , we can remove the dependence of dynamical scaling relations on the internal structure of the galaxy by applying a simple correction to the measured stellar velocity dispersions , which depends only on the concentration index @xmath57(eq . [ eq_modsigma ] ) . the correlation between baryonic mass and corrected @xmath5 thus obtained has a scatter of only 0.08 dex ( fig . [ bfj_corr ] and table [ t_fits ] ) . it is encouraging that our correction removes the offset between disks and spheroids also in the size - velocity relation ( expressed in terms of @xmath17 versus stellar dispersion ) . we tested if an even tighter baryonic mass - velocity relation could be obtained by using the @xmath91 parameter , which combines rotational velocity and stellar dispersion . we found no improvement , because the @xmath91 relation still suffers from the problem of low inclination galaxies that affects the btf , to a lesser extent . despite being measured at a smaller spatial scale than the hi width , the stellar velocity dispersion turns out to be a better tracer of mass , at least for the massive galaxies in our sample . this is surprising , as one would expect the rotational velocity to provide a more reliable measurement of mass for galaxies with hi detections , the vast majority of which are rotation dominated ( i.e. , 216/228 galaxies have @xmath101 ) , regardless of the presence of a bulge . and yet , the same result was obtained by @xcite for a sample of s0 galaxies with inclinations @xmath035@xmath7160 , which are also , overall , rotation - dominated : the central stellar velocity dispersion is a better predictor of _ luminosity than the circular speed at 2 - 3 exponential disk scale lengths , which they carefully measured from long - slit optical absorption - line spectra . as we already pointed out , the main limitation with rotational velocities is observational , not intrinsic . we are simply not able to reliably measure circular speeds from line - of - sight velocities , except for well - selected samples of undisturbed late - type , inclined spirals ( i.e. , the typical tf samples ) , for which the deprojection to edge - on view does not introduce very large uncertainties . as for the velocity dispersions , which are measured through 3-diameter fibers , there might be a concern about contamination from disk stars in circular orbits , especially for the most edge - on , disk - dominated galaxies . however , we have argued that this effect is negligible ( [ s_bfj ] ) . this agrees with the conclusion reached by @xcite , who reported that the contamination is small ( @xmath102 , based on simulations ) for milky way - type galaxies with pressure - supported bulges . we argued that the reason why disk - dominated galaxies are simply offset from the spheroids on the bfj plane ( which is the key point that allows us to bring all massive galaxies onto the same baryonic mass - velocity relation ) is that @xmath5 is proportional to @xmath7 , and their ratio is a function of galaxy morphology . although initial work indicated that spirals and ellipticals follow the same tight @xmath7-@xmath5 correlation for @xmath103 km s@xmath13(e.g . * ; * ? ? ? * ; * ? ? ? * ) , more recent analyses demonstrate that the relation between @xmath7 and @xmath5 is not universal , but depends on morphology . as mentioned in [ s_mbfj ] , @xcite show that the brightest , bulge - dominated galaxies lie on the @xmath86 relation expected for isothermal stellar systems , whereas later - type spirals and dwarfs are offset by an amount that depends on morphology or total light concentration . we have shown that correcting the stellar dispersions according to equation [ eq_modsigma ] effectively translates into accounting for this departure from isothermality for disk - dominated galaxies . @xcite noted that , despite the fact that a detailed understanding of what sets the relation between @xmath7 , @xmath5 and concentration index is still missing , one could use this relation to empirically reduce the scatter of scaling relations that involve dynamical parameters , such as the tf or fj . we showed that the existence of such relation allows us to do more than that we obtained a generalized bfj relation that holds for all the massive galaxies in our sample , with a scatter ( 0.079 dex ) that is as small as that of the btf and bfj relations applied to their pruned subsets ( i.e. 0.076 and 0.074 dex , respectively ) . for comparison , @xcite and @xcite report a scatter of 0.06 dex for the btf and 0.071 dex for the stellar fj relations , respectively . the implications of our generalized bfj relation for extragalactic studies are very promising . because it holds for all massive galaxies in our sample _ regardless of morphology _ , this relation appears to provide a more fundamental link between dark matter halo mass and baryonic content than that obtained using rotational velocities or velocity dispersions . as such , it gives more fundamental constraints to galaxy formation models than the tf or fj / fp relations . also this is the reference baryonic mass - internal velocity relation that higher redshift studies , which naturally target the most massive galaxies , should compare with . this relation is more resilient to systematic effects than the tf , and , contrary to both tf and fj / fp relations , does not require any morphological pruning a significant advantage since accurate morphological classifications are difficult to obtain for large samples beyond the very local universe . as pointed out throughout this work , our results are based on a representative sample of galaxies with stellar masses larger than @xmath14 m@xmath2 . it remains to be established how far down in stellar mass these results can be extrapolated . it would be beneficial to investigate whether a bfj relation for disk galaxies holds down to low baryonic masses ( where the gas contribution is more important ) , and with similarly low scatter , based on a representative and homogeneous sample such as gass .

disk - dominated , gas - rich galaxies systematically deviate from the bfj relation defined by the spheroids . we demonstrate that by applying a simple correction to the stellar velocity dispersions that depends only on the concentration index of the galaxy , we are able to bring disks and spheroids onto the same dynamical relation in other words , we obtain a generalized bfj relation that holds for all the galaxies in our sample , regardless of morphology , inclination or gas content , and has a scatter smaller than 0.1 dex .

we present dynamical scaling relations for a homogeneous and representative sample of massive galaxies , selected only by stellar mass ( m ) and redshift ( ) as part of the ongoing galex arecibo sdss survey . we compare baryonic tully - fisher ( btf ) and faber - jackson ( bfj ) relations for this sample , and investigate how galaxies scatter around the best fits obtained for pruned subsets of disk - dominated and bulge - dominated systems . the bfj relation is significantly less scattered than the btf when the relations are applied to their maximum samples ( for the btf , only galaxies with hi detections ) , and is not affected by the inclination problems that plague the btf . disk - dominated , gas - rich galaxies systematically deviate from the bfj relation defined by the spheroids . we demonstrate that by applying a simple correction to the stellar velocity dispersions that depends only on the concentration index of the galaxy , we are able to bring disks and spheroids onto the same dynamical relation in other words , we obtain a generalized bfj relation that holds for all the galaxies in our sample , regardless of morphology , inclination or gas content , and has a scatter smaller than 0.1 dex . we compare the velocity - size relation for the three dynamical indicators used in this work , i.e. , rotational velocity , observed and concentration - corrected stellar dispersion . we find that disks and spheroids are offset in the stellar dispersion - size relation , and that the offset is removed when corrected dispersions are used instead . the generalized bfj relation represents a fundamental correlation between the global dark matter and baryonic content of galaxies , which is obeyed by all ( massive ) systems regardless of morphology . [ firstpage ] galaxies : kinematics and dynamics galaxies : evolution galaxies : fundamental parameters radio lines : galaxies

1302.6894

i

in this paper we discuss efficient ways to model optically pumped atoms in a regime where velocity - selective optical pumping ( vsop ) is possible , but where collisional rates with buffer gases are too high to permit the use of models for cooling and trapping in the near absence of collisions . this is the regime of sodium guidestar atoms . these naturally - occurring layers of sodium atoms at altitudes of 90100 km above the earth s surface @xcite are illuminated by ground - based lasers , and the returning photons are used to measure the relative retardation of wave fronts across an optical aperture . this retardation information can be used with a deformable mirror to correct for the aberrations from atmospheric turbulence and to allow the receiving optics to produce a more nearly diffraction - limited image of astronomical objects . the performance of guidestar systems is limited by the loss of atoms from the most strongly backscattering spin sublevels and velocity groups . the most important reasons for these losses are : collisions with the residual atmospheric gases , which transfer atoms from strongly absorbing to weakly absorbing spin sublevels , or which shift the atoms into velocity groups that are not in resonance with the pumping light ; larmor precession of the spins away from strongly absorbing orientations if the geomagnetic field is not parallel to the direction of the laser beam ; and unwanted optical pumping into weakly absorbing sublevels . the powerful modeling methods discussed here make it easier to explore the parameter space of these processes and to optimize the performance of guidestar systems . these methods also provide a more realistic and numerically convenient way to model laboratory experiments with vsop of atoms in low - pressure buffer gases . in contrast to previous work on this topic , for example references @xcite , we account for the full hyperfine structure of real alkali - metal atoms , we show how to use spin - relaxation modes @xcite to incorporate into the model the complicated spin relaxation of na guidestar atoms due to collisions with paramagnetic oxygen atoms , and we use recently developed cusp kernels @xcite to realistically and efficiently model velocity - changing collisions . this article is organized as follows . in section [ denmat ] , we introduce the liouville space of the coupled spin and velocity distributions of the atoms . in section [ spin ] we introduce spin - relaxation modes @xmath0 @xcite to describe the spin distributions , and we introduce the concept of conjugate spin - mode indices , @xmath1 and @xmath2 . the amplitudes @xmath3 of the spin modes in velocity space provide a complete description of the spin and velocity polarization of the atoms . the little - known liouville conjugate operation , denoted by the superscript @xmath4 , and the transposition operator @xmath5 for liouville - space operators are discussed in section [ spinb ] . using liouville conjugates reduces the numerical computing requirements by nearly a factor of 2 . in section [ velocity ] we show how to describe velocity distributions with velocity - relaxation modes @xmath6 @xcite . we use the velocity modes to show that a simple transformation of the widely used keilson - storer kernels @xcite leads to much more realistic and useful cusp kernels @xcite for describing velocity - changing collisions . in section [ pressure ] we present a simple model for the transition from collision - free , ballistic flight to any container walls at very low buffer - gas pressure to diffusional wall losses at higher pressure . we sketch how the relative sizes of the laser beam and the cell affect these processes . in an appendix we show how to deduce the rate of velocity - changing collisions , @xmath7 from the spatial diffusion coefficient @xmath8 and the smallest - nonzero eigenvalue @xmath9 of the collision operator @xmath10 with the little - known formula ( [ sd80 ] ) . in ( [ rid16 ] ) of section [ dark ] , we show that spin - changing and velocity - changing collisions cause the spin - mode amplitudes @xmath3 to relax exponentially in time at the rate @xmath11 , where @xmath11 is a kernel in velocity space . in section [ pumping ] , we introduce a velocity - dependent optical pumping operator @xmath12 , which we write as the sum of poles in the complex - velocity plane at locations determined by the laser frequency and the optical bohr frequencies . the pole expansion facilitates velocity averages in terms of faddeeva functions ( voigt profiles ) @xcite . the poles have `` residue matrices '' that are independent of the laser frequency and atomic velocity . in section viii , we show that the steady - state mode amplitudes @xmath3 generated by the combined effects of optical pumping , spin relaxation , and velocity relaxation can be written in terms of green s functions , @xmath13 . we show that if the kernel @xmath11 is a cusp kernel or linear combination thereof , @xmath14 is also cusp kernel or linear combination thereof . being able to invert cusp kernels in closed form greatly simplifies the numerical evaluation of the mode amplitudes @xmath3 . this simplification is not possible with keilson - storer kernels or any other collision kernel that we know of . we present an explicit , series solution ( [ ss2 ] ) for the mode amplitudes @xmath3 in powers of the pumping light intensity , with particular emphasis on the first - order solution ( [ sr10 ] ) . finally , in section ix , we use ( [ sr10 ] ) to demonstrate how the methods we present are in excellent agreement with existing vsop experiments . we also describe a new type of magnetic - depolarization experiment that can be carried out under laboratory conditions and readily interpreted with the powerful modeling methods described in this paper . such experiments would provide much more detailed experimental information about the nature of velocity - changing collisions .

we present efficient theoretical tools for describing the optical pumping of atoms by light propagating at arbitrary directions with respect to an external magnetic field , at buffer - gas pressures that are small enough for velocity - selective optical pumping ( vsop ) but large enough to cause substantial collisional relaxation of the velocities and the spin . we use spin and velocity relaxation modes to describe the distribution of atoms in spin space ( including both populations and coherences ) and velocity space . cusp kernels are used to describe velocity - changing collisions . optical pumping operators are represented as a sum of poles in the complex velocity plane . signals simulated with these methods are in excellent agreement with previous experiments and with preliminary experiments in our laboratory .

we present efficient theoretical tools for describing the optical pumping of atoms by light propagating at arbitrary directions with respect to an external magnetic field , at buffer - gas pressures that are small enough for velocity - selective optical pumping ( vsop ) but large enough to cause substantial collisional relaxation of the velocities and the spin . these are the conditions for the sodium atoms at an altitude of about 100 km that are used as guidestars for adaptive optics in modern ground - based telescopes to correct for aberrations due to atmospheric turbulence . we use spin and velocity relaxation modes to describe the distribution of atoms in spin space ( including both populations and coherences ) and velocity space . cusp kernels are used to describe velocity - changing collisions . optical pumping operators are represented as a sum of poles in the complex velocity plane . signals simulated with these methods are in excellent agreement with previous experiments and with preliminary experiments in our laboratory .

1112.0063

c

we present the empirical formulae of intrinsic reddening as a function of intrinsic h@xmath0 luminosity or surface brightness , metallicity and disc inclination ( equation [ eq : powfit ] , [ eq : powfit_br ] and table [ tab : lha ] and [ tab : br ] ) for a large sample of @xmath122 000 well - defined star - forming disc galaxies selected from the sdss . with the empirical formulae , the reddening parameterized by @xmath2 could be predicted within 1@xmath122 uncertainty of 0.07 mag . the @xmath20 , defined as the observed reddening estimated from the balmer decrement minus predicted value , does not correlate with the three parameters used in the formulae . we also find that @xmath20 is independent of the stellar mass , 4000 break strength or electron density within the galaxy . the observed trend between the dust reddening and h@xmath0 surface brightness could be reproduced by a plane - parallel slab toy model , in which the dust is scaled with gas , and all the distributions of the dust , gas and star - formation regions are smooth and follow exponential laws in the disc . we find from the comparisons of the models to the observation , the relative vertical scale of dust distribution to hii regions distribution may vary with h@xmath0 surface brightness or metallicity . the higher intensity of star formation could drive the dust to distribute in a thicker layer , or the dust disc is thicker than the disc of the hiiregions when the metallicity is higher . however , the real galaxies may be not fully described by this simple toy model . for example , when the spiral arms are present , or the sf regions are clumpy , the assumptions of smooth distributions will be improper . in addition , our result implies a different scaling law of dust - to - gas ratio as a function of metallicity @xmath22 from the linearly relation ( equation [ eq : rhod1 ] ) . our empirical relations , which suggest dust reddening is partially dependent on metallicity ( @xmath123 ) , can be an observational constraints on the output of radiative - transfer model calculations , and hence provides constraints on the assumed dust - to - gas ratio in the model . since metallicity - dependence has been introduced , we believe that our empirical formulae should be applicable to the early evolutionary stage of a galaxy , when the gas metallicity might be low . in the starburst galaxies at high redshift , dust surrounding star - forming region has been proven to be prone to sub - milimeter and far - infrared observations . but most of those galaxies are relatively dim in optical , implying severe extinction and reddening to these galaxies . the observed galaxies at high redshifts are affected by reddening , in the way that the galaxies with heavy attenuation and reddening may escape from the detection . the importance of the reddening correction depends on the wavelength band that is used . for example , uv continuum luminosity is usually used to infer sfr in high - redshift galaxies . it is very sensitive to the attenuation / reddening . calzetti et al . ( 2000 ) find that the extinction to emission lines and to stellar continuum in optical band are different , and the continuum reddening is only about 40 percent of that for emission lines . however , the extinction to uv continuum is likely similar to the hii region because uv continuum is also emitted by young hot stars , which are likely surrounded by hiiregions . as a result , our results may be used to estimate the reddening for these high - redshift galaxies with caution , that the metallicity should be well - calibrated and be within the range covered by our sample . following equation [ eq : powfit ] and the discussions in section [ sec : relation ] , sfr in luminous star formation or high metallicity galaxies might have been under - estimated ( e.g. heckman et al . 1998 ; hopkins et al . 2001 ; panuzzo et al . 2007 ) , or those galaxies are even completely lost in optical surveys . the quantitative relationship of reddening with the metallicity , luminosity or surface brightness , and the inclination for disc galaxies can also be incorporated into the current semi - analytic models of galaxy formation and evolution ( e.g. croton et al . 2006 ; de lucia et al . 2004 ; kang et al . 2005 ) . in those models , a simple chemical enrichment scheme has been included , and the metal abundance has been actually predicted . with the known metallicity and the star formation rate ( intrinsic luminosity ) provided by the models , the reddening can be estimated accurately with our formulae , and then used to shape the spectral energy distribution and predict the emerging luminosity more realistically .

we present empirical relations between the global dust reddening and other physical galaxy properties including the h luminosity , h surface brightness , metallicity and axial ratio for star - forming disc galaxies . the study is based on a large sample of 22 000 well - defined star - forming galaxies selected from the sloan digital sky survey ( sdss ) . the empirical formulae can be incorporated into semi - analytical models of galaxy formation and evolution to estimate the dust reddening and enable comparison with observations more practically .

we present empirical relations between the global dust reddening and other physical galaxy properties including the h luminosity , h surface brightness , metallicity and axial ratio for star - forming disc galaxies . the study is based on a large sample of 22 000 well - defined star - forming galaxies selected from the sloan digital sky survey ( sdss ) . the reddening parameterized by color excess is derived from the balmer decrement . besides the dependency of reddening on hluminosity / surface brightness and gas phase metallicity , it is also correlated with the galaxy inclination , in the sense that edge - on galaxies are more attenuated than face - on galaxies at a give intrinsic luminosity . in light of these correlations , we present the empirical formulae of as a function of these galaxy properties , with a scatter of only 0.07 mag . the empirical relation can be reproduced if most dust attenuation to the hii region is due to diffuse interstellar dust distributing in a disc thicker than that of hii regions . the empirical formulae can be incorporated into semi - analytical models of galaxy formation and evolution to estimate the dust reddening and enable comparison with observations more practically . [ firstpage ] galaxies : ism galaxies : abundance hii regions dust , extinction .

1101.4966

c

we searched for nebular emission in a sample of 77 bcgs from the 160sd x - ray cluster survey . no or emission stronger than @xmath1 and @xmath2 , respectively , was detected in any bcg . a comparison between our sample and the bcs clusters of similar x - ray luminosity suggests that we should have detected roughly one dozen emission - line bcgs stronger than our detection threshold . 160sd clusters of a given x - ray luminosity lie at systematically higher redshifts compared to bcs clusters of the same luminosity . this may indicate an increase in the number density of _ strong _ cooling flows between @xmath97 and today . because of our relatively high detection thresholds , our result does not imply that cooling flows do nt exist at the redshifts surveyed in the 160sd survey ( see bauer et al . 2005 ) . however , we find a significant decrease in the number of _ strong _ cooling flows like the perseus cluster . similar conclusions were reached using independent methods by santos et al . ( 2008 ) , vikhlinin et al . ( 2007 ) and by voevodkin et al . ( 2010 ) , who also found fewer cooling flow clusters at higher redshifts . the decline may be related to the destruction of cooling flows by mergers , which occurred more frequently than they do today ( santos et al . weaker mergers may also disturb x - ray atmospheres and miscenter the bcg , which could prevent a strong cooling flow from forming . agn feedback is thought to be responsible for quenching cooling flows in nearby clusters ( peterson & fabian 2006 ; mcnamara & nulsen 2007 ) , but this is less clear in higher redshift clusters lacking the signatures of cooling flows . however , in a study of radio agn located in the 400 square degree x - ray cluster survey ( burenin et al . 2007 ) , ma et al . ( 2011 , in preparation ) found that the power injected by radio sources in galaxies lying within 250 kpc of cluster centers is significant compared to the power radiated by the hot atmospheres . the selection function of the 400 square degree survey is nearly identical to the 160 sd survey , but it covers a much larger area of the sky . the ma et al . study suggests that agn are able to deposit enough energy in these clusters to delay or prevent a massive cooling flow from forming at redshifts of a few tenths or so . mcnamara and samuele would like to acknowledge ohio university s astrophysical institute for its hospitality while much of the data analysis was performed in pursuit of samuele s masters thesis . mcnamara acknowledges conversations with alastair edge , megan donahue , paul nulsen , andy fabian , and a generous grant from canada s natural sciences and engineering research council . balogh , m. , bower , r.g . , smail , i. , ziegler , b.l . , davies , r.l . , gaztelu , a. , fritz , a. , 2002 mnras , 337 , 256 bauer , f. e. , fabian , a. c. , sanders , j. s. , allen , s. w. , & johnstone , r. m. 2005 , , 359 , 1481 bernardi , m. , et al . 2003 , , 125 , 1882 burenin , r. a. , vikhlinin , a. , hornstrup , a. , ebeling , h. , quintana , h. , & mescheryakov , a. 2007 , , 172 , 561 cardelli , j.a . , clayton , g.c . , mathis , j.s . , 1989 , apj , 345 , 245 cavagnolo , k. w. , donahue , m. , voit , g. m. , & sun , m. 2008 , , 683 , l107 crawford , c.s . , edge , a.c . , fabian , a.c . , allen , s.w . , bhringer , h. , ebeling , h. , mcmahon , r.g . , voges , w. , 1995 , mnras , 274 , 75 crawford , c.s . , allen , s.w . , ebeling , h. , edge , a.c . , fabian , a.c . , 1999 , mnras , 306 , 857 davies , r.l . , sadler , e.m . , peletier , r.f . , 1993 , mnras , 262 , 650 donahue , m. , mack , j. , voit , g. m. , sparks , w. , elston , r. , & maloney , p. r. 2000 , , 545 , 670 donahue , m. , et al . 2010 , , 715 , 881 edge , a. c. 2001 , , 328 , 762 edwards , l. o. v. , hudson , m. j. , balogh , m. l. , & smith , r. j. 2007 , , 379 , 100 ebeling , h. , edge , a.c . , bhringer , h. , allen , s.w . , crawford , c.s . , fabian , a.c . , voges , w. , huchra , j.p . , 1998 , mnras , 301 , 881 faber , s.m . , friel , e.d . , burstein , d. , gaskell , c.m . , 1985 , apj , 57 , 711 fabian , a. c. 1994 , , 32 , 277 fabian , a.c . , 1992 , clusters and superclusters of galaxies , kluwer , 151 heckman , t.m . , baum , s.a . , van breugel , w.j.m . , mccarthy , p. , 1989 , apj , 338 , 48 hudson , d. s. , mittal , r. , reiprich , t. h. , nulsen , p. e. j. , andernach , h. , & sarazin , c. l. 2010 , , 513 , a37 ma , c.j , mcnamara , b.r . , nulsen , p.e.j . , vikhlinin , a. , r. schaffer , 2011 , in preparation . mcnamara , b. r. , vikhlinin , a. , hornstrup , a. , quintana , h. , whitman , k. , forman , w. , & jones , c. 2001 , , 558 , 590 mcnamara , b.r . , oconnell , r.w . , 1989 , a.j . , 98 , 6 mullis , c.r . , mcnamara , b.r . , quintana , h. , vikhlinin , a. , henry , j.p . , gioia , i.m . , hornstrup , a. , forman , w. , jones , c. 2003 , apj , 594 , 154 oconnell , r.w . , 1973 , a.j . , 78 , 10 rafferty , d. a. , mcnamara , b. r. , & nulsen , p. e. j. 2008 , , 687 , 899 odea , k. p. , et al . 2010 , arxiv:1006.3796 peterson , j. r. , & fabian , a. c. 2006 , , 427 , 1 rosati , p. , borgani , s. , & norman , c. 2002 , , 40 , 539 salom , p. , & combes , f. 2003 , , 412 , 657 santos , j. s. , rosati , p. , tozzi , p. , bhringer , h. , ettori , s. , & bignamini , a. 2008 , , 483 , 35 vikhlinin , a. , burenin , r. , forman , w. r. , jones , c. , hornstrup , a. , murray , s. s. , & quintana , h. 2007 , heating versus cooling in galaxies and clusters of galaxies , 48 vikhlinin , a. , mcnamara , b.r . , forman , w. , jones , c. , quintana , h. , hornstrup , a. , 1998 , apj , 502 , 558 voevodkin , a. et al . , in preparation cccc [ oii ] & 3727 & 3721.003732.00 & 3692.003722.00 , 3732.003762.00 + h@xmath98 & 6562 & 6557.006568.00 & 6527.006557.00 , 6567.006597.00 + h@xmath99 & 4861 & 4849.504877.00 & 4829.504848.25 , 4878.254892.00 + mgib & 5175 & 5162.005193.25 & 5144.505162.00 , 5193.255207.00 + mg@xmath100 & & 5156.005197.25 & 4897.004958.25 , 5303.005366.75 + [ table : widths ] rxj0041.12339 & 0.06 & 0.112 & ... & ... & -0.34 @xmath77 0.31 & 0.85 + rxj0050.90929 & 0.64 & 0.200 & -2.82 @xmath77 2.28 & 3.97 & -0.11 @xmath77 0.48 & 2.40 + rxj0056.92213 & 0.16 & 0.116 & ... & ... & 0.47 @xmath77 0.66 & ... + rxj0110.3 + 1938 & 0.36 & 0.317 & 1.25 @xmath77 2.72 & ... & 0.83 @xmath77 1.29 & ... + rxj0122.52832 & 0.78 & 0.256 & 2.24 @xmath77 1.91 & ... & 1.27 @xmath77 0.73 & ... + rxj0124.5 + 0400 & 0.34 & 0.316 & -1.85 @xmath77 2.26 & 3.59 & -1.36 @xmath77 1.65 & 8.36 + rxj0142.8 + 2025 & 0.84 & 0.271 & 3.68 @xmath77 2.47 & ... & ... & ... + rxj0144.4 + 0212 & 0.13 & 0.166 & -2.70 @xmath77 13.02 & 0.77 & ... & ... + rxj0159.3 + 0030 & 2.11 & 0.386 & 1.95 @xmath77 1.23 & ... & ... & ... + rxj0206.3 + 1511 & 0.36 & 0.248 & 7.56 @xmath77 2.66 & ... & 0.17 @xmath77 1.29 & ... + rxj0206.81309 & 1.18 & 0.320 & 2.02 @xmath77 2.94 & ... & 0.78 @xmath77 1.24 & ... + rxj0258.7 + 0012 & 0.33 & 0.259 & 2.61 @xmath77 4.16 & ... & 1.87 @xmath77 1.69 & ... + rxj0259.5 + 0013 & 0.54 & 0.194 & 1.79 @xmath77 8.36 & ... & 0.79 @xmath77 0.84 & ... + rxj0351.63649 & 0.56 & 0.372 & 1.22 @xmath77 2.47 & ... & ... & ... + rxj0506.02840 & 0.16 & 0.136 & ... & ... & 0.88 @xmath77 1.01 & ... + rxj0521.12530 & 2.57 & 0.581 & -0.29 @xmath77 3.11 & 0.60 & ... & ... + rxj0522.23625 & 1.79 & 0.472 & 4.83 @xmath77 5.85 & ... & ... & ... + rxj0826.4 + 3125 & 0.22 & 0.209 & -4.98 @xmath77 9.83 & 2.97 & 0.90 @xmath77 0.71 & ... + rxj0841.1 + 6422 & 1.49 & 0.342 & 4.91 @xmath77 10.11 & ... & ... & ... + rxj0842.8 + 5023 & 0.52 & 0.423 & 3.40 @xmath77 2.84 & ... & ... & ... + rxj0852.5 + 1618 & 0.16 & 0.098 & -2.33 @xmath77 3.12 & ... & 1.17 @xmath77 0.79 & ... + rxj0858.4 + 1357 & 0.70 & 0.488 & -2.60 @xmath77 0.97 & 11.28 & 1.29 @xmath77 1.14 & ... + rxj0907.2 + 3330 & 0.48 & 0.483 & 1.58 @xmath77 4.76 & ... & ... & ... + rxj0921.2 + 4528 & 1.05 & 0.315 & -1.74 @xmath77 6.75 & 0.38 & -2.37 @xmath77 4.65 & 0.78 + rxj0926.6 + 1242 & 1.75 & 0.489 & 3.81 @xmath77 3.64 & ... & ... & ... + rxj0943.7 + 1644 & 0.31 & 0.180 & -0.36 @xmath77 2.13 & ... & 2.48 @xmath77 1.11 & ... + rxj0958.2 + 5516 & 0.97 & 0.214 & ... & ... & 1.47 @xmath77 2.11 & ... + rxj1013.6 + 4933 & 0.36 & 0.133 & ... & ... & 0.24 @xmath77 0.56 & ... + rxj1015.1 + 4931 & 0.71 & 0.383 & 3.55 @xmath77 2.66 & ... & ... & ... + rxj1036.1 + 5713 & 0.35 & 0.203 & ... & ... & 3.45 @xmath77 1.89 & ... + rxj1049.0 + 5424 & 0.26 & 0.251 & 3.25 @xmath77 3.28 & ... & 0.33 @xmath77 0.78 & ... + rxj1117.2 + 1744 & 0.50 & 0.305 & 1.30 @xmath77 2.46 & ... & -3.16 @xmath77 3.33 & 1.24 + rxj1117.5 + 1744 & 1.90 & 0.548 & 0.48 @xmath77 1.14 & ... & ... & ... + rxj1120.9 + 2326 & 2.89 & 0.562 & -0.48 @xmath77 1.02 & 0.98 & ... & ... + rxj1123.1 + 1409 & 0.93 & 0.340 & 2.72 @xmath77 3.54 & ... & -0.98 @xmath77 3.70 & 2.08 + rxj1124.01700 & 0.81 & 0.407 & -2.21 @xmath77 3.32 & 4.58 & ... & ... + rxj1124.6 + 4155 & 0.67 & 0.195 & ... & ... & -0.07 @xmath77 0.87 & 0.59 + rxj1135.9 + 2131 & 0.14 & 0.133 & 2.26 @xmath77 5.48 & ... & 0.23 @xmath77 0.51 & ... + rxj1142.0 + 2144 & 0.35 & 0.131 & ... & ... & 0.76 @xmath77 0.82 & ... + rxj1146.4 + 2854 & 0.38 & 0.149 & 1.27 @xmath77 10.24 & ... & 0.51 @xmath77 0.60 & ... + rxj1158.1 + 5521 & 0.04 & 0.135 & -2.74 @xmath77 14.32 & 0.64 & 0.47 @xmath77 0.37 & ... + rxj1200.90327 & 1.28 & 0.396 & 1.05 @xmath77 3.23 & ... & ... & ... + rxj1206.50744 & 0.26 & 0.068 & ... & ... & 0.20 @xmath77 0.34 & ... + rxj1213.5 + 0253 & 1.07 & 0.409 & 4.21 @xmath77 8.94 & ... & ... & ... + rxj1218.4 + 3011 & 0.33 & 0.368 & -4.16 @xmath77 3.98 & 8.61 & 8.25 @xmath77 3.41 & ... + rxj1221.4 + 4918 & 4.27 & 0.700 & 0.45 @xmath77 0.50 & ... & ... & ... + rxj1237.6 + 2632 & 0.25 & 0.278 & -5.23 @xmath77 3.76 & 7.38 & -0.53 @xmath77 1.78 & 1.95 + rxj1254.6 + 2545 & 0.17 & 0.193 & -0.87 @xmath77 2.92 & ... & 3.40 @xmath77 1.02 & ... + rxj1254.8 + 2550 & 0.32 & 0.233 & ... & ... & 0.71 @xmath77 0.36 & ... + rxj1256.0 + 2556 & 0.24 & 0.232 & 1.08 @xmath77 4.54 & ... & 0.16 @xmath77 0.81 & ... + rxj1301.7 + 1059 & 0.66 & 0.231 & ... & ... & 2.16 @xmath77 1.27 & ... + rxj1309.9 + 3222 & 0.34 & 0.290 & -0.62 @xmath77 2.95 & ... & -0.55 @xmath77 2.43 & ... + rxj1337.8 + 3854 & 0.41 & 0.252 & ... & ... & 0.13 @xmath77 1.22 & ... + rxj1342.8 + 4028 & 1.62 & 0.699 & 1.03 @xmath77 1.10 & ... & ... & ... + rxj1343.4 + 5547 & 0.04 & 0.069 & ... & ... & 0.06 @xmath77 0.25 & ... + rxj1354.8 + 6917 & 0.13 & 0.207 & 9.35 @xmath77 4.39 & ... & -2.02 @xmath77 0.92 & 3.42 + rxj1438.9 + 6423 & 0.25 & 0.146 & ... & ... & -1.38 @xmath77 0.49 & 3.03 + rxj1500.0 + 2233 & 0.35 & 0.230 & -2.66 @xmath77 4.30 & 3.69 & 0.08 @xmath77 0.74 & ... + rxj1515.5 + 4346 & 0.29 & 0.137 & ... & ... & -0.04 @xmath77 0.50 & 0.21 + rxj1515.6 + 4350 & 0.28 & 0.243 & 7.76 @xmath77 7.80 & ... & -0.48 @xmath77 0.71 & 13.99 + rxj1524.6 + 0957 & 3.45 & 0.516 & -0.72 @xmath77 1.36 & 1.60 & ... & ... + rxj1537.7 + 1200 & 0.21 & 0.134 & -4.22 @xmath77 10.94 & 2.08 & 0.62 @xmath77 0.32 & ... + rxj1540.8 + 1445 & 0.68 & 0.441 & 7.24 @xmath77 3.52 & ... & ... & ... + rxj1547.3 + 2056 & 0.79 & 0.266 & 1.84 @xmath77 4.29 & ... & 3.01 @xmath77 0.84 & ... + rxj1552.2 + 2013 & 0.40 & 0.136 & 1.94 @xmath77 2.08 & ... & 1.32 @xmath77 0.70 & ... + rxj1630.2 + 2434 & 0.33 & 0.066 & ... & ... & 1.19 @xmath77 0.40 & ... + rxj1642.6 + 3935 & 0.58 & 0.355 & -6.54 @xmath77 8.92 & 22.03 & ... & ... + rxj1659.7 + 3410 & 0.52 & 0.341 & -5.57 @xmath77 4.15 & 10.64 & 1.04 @xmath77 4.64 & ... + rxj1722.8 + 4105 & 1.23 & 0.309 & 6.08 @xmath77 2.80 & ... & 0.03 @xmath77 3.10 & ... + rxj1729.0 + 7440 & 0.35 & 0.213 & -1.02 @xmath77 11.77 & 0.84 & 1.94 @xmath77 1.88 & ... + rxj1746.4 + 6848 & 0.47 & 0.217 & -6.83 @xmath77 4.53 & ... & 0.53 @xmath77 1.58 & ... + rxj2202.71902 & 0.58 & 0.438 & -6.93 @xmath77 8.21 & 22.77 & ... & ... + rxj2212.61713 & 0.04 & 0.134 & -0.28 @xmath77 4.04 & 0.24 & 0.64 @xmath77 0.22 & ... + rxj2213.51656 & 0.71 & 0.297 & -13.55 @xmath77 9.81 & 11.18 & -2.82 @xmath77 2.86 & 10.26 + rxj2257.8 + 2056 & 0.44 & 0.297 & -4.48 @xmath77 4.62 & 13.26 & 4.03 @xmath77 1.72 & ... + rxj2328.8 + 1453 & 0.85 & 0.497 & 4.47 @xmath77 2.07 & ... & ... & ... + rxj2348.83117 & 0.49 & 0.184 & 8.86 @xmath77 10.67 & ... & 1.16 @xmath77 0.85 & ... + rxj0041.12339 & 2.24 @xmath77 0.64 & 3.30 @xmath77 0.94 & 0.26 @xmath77 0.11 + rxj0050.90929 & 1.07 @xmath77 0.42 & 2.97 @xmath77 0.73 & 0.31 @xmath77 0.07 + rxj0056.92213 & 2.32 @xmath77 0.90 & 3.79 @xmath77 0.72 & 0.29 @xmath77 0.13 + rxj0110.3 + 1938 & 1.90 @xmath77 0.85 & ... & 0.27 @xmath77 0.10 + rxj0122.52832 & 2.32 @xmath77 0.41 & 3.95 @xmath77 0.60 & 0.29 @xmath77 0.11 + rxj0124.5 + 0400 & 2.17 @xmath77 0.92 & 4.19 @xmath77 1.03 & 0.27 @xmath77 0.14 + rxj0142.8 + 2025 & 2.24 @xmath77 0.46 & 3.09 @xmath77 0.91 & 0.24 @xmath77 0.12 + rxj0144.4 + 0212 & 0.61 @xmath77 0.47 & -0.54 @xmath77 0.78 & 0.21 @xmath77 0.06 + rxj0159.3 + 0030 & 1.63 @xmath77 1.00 & 0.05 @xmath77 0.69 & 0.14 @xmath77 0.08 + rxj0206.3 + 1511 & 1.11 @xmath77 1.02 & 3.86 @xmath77 1.21 & 0.25 @xmath77 0.16 + rxj0206.81309 & 0.51 @xmath77 1.43 & ... & 0.28 @xmath77 0.11 + rxj0258.7 + 0012 & 2.32 @xmath77 1.33 & 1.93 @xmath77 1.30 & 0.25 @xmath77 0.16 + rxj0259.5 + 0013 & 2.38 @xmath77 0.76 & 5.00 @xmath77 1.11 & 0.25 @xmath77 0.14 + rxj0351.63649 & 2.29 @xmath77 0.90 & 4.78 @xmath77 0.99 & 0.22 @xmath77 0.18 + rxj0506.02840 & 2.37 @xmath77 1.13 & 0.13 @xmath77 1.79 & 0.22 @xmath77 0.12 + rxj0521.12530 & 1.13 @xmath77 1.77 & 5.99 @xmath77 3.26 & 0.38 @xmath77 0.41 + rxj0522.23625 & 1.63 @xmath77 1.34 & ... & 0.60 @xmath77 0.60 + rxj0826.4 + 3125 & 2.54 @xmath77 0.99 & 3.25 @xmath77 0.75 & 0.24 @xmath77 0.11 + rxj0841.1 + 6422 & -0.99 @xmath77 1.64 & 5.74 @xmath77 2.51 & 0.25 @xmath77 0.28 + rxj0842.8 + 5023 & 1.98 @xmath77 1.75 & 3.20 @xmath77 2.04 & 0.20 @xmath77 0.29 + rxj0852.5 + 1618 & 1.57 @xmath77 1.38 & 4.56 @xmath77 1.41 & 0.34 @xmath77 0.24 + rxj0858.4 + 1357 & 1.68 @xmath77 0.45 & 2.55 @xmath77 0.79 & 0.19 @xmath77 0.11 + rxj0907.2 + 3330 & 2.41 @xmath77 0.92 & 1.90 @xmath77 1.56 & 0.24 @xmath77 0.25 + rxj0921.2 + 4528 & 0.82 @xmath77 1.22 & 1.35 @xmath77 0.84 & 0.16 @xmath77 0.19 + rxj0926.6 + 1242 & 3.71 @xmath77 2.61 & ... & 0.18 @xmath77 0.51 + rxj0943.7 + 1644 & 2.59 @xmath77 0.76 & 4.97 @xmath77 1.58 & 0.31 @xmath77 0.25 + rxj0958.2 + 5516 & -0.33 @xmath77 1.84 & 7.49 @xmath77 1.70 & 0.34 @xmath77 0.28 + rxj1013.6 + 4933 & 1.49 @xmath77 1.38 & 3.99 @xmath77 1.03 & 0.32 @xmath77 0.13 + rxj1015.1 + 4931 & 1.77 @xmath77 1.48 & -0.15 @xmath77 0.95 & 0.22 @xmath77 0.17 + rxj1036.1 + 5713 & 6.68 @xmath77 1.44 & 1.55 @xmath77 1.31 & 0.08 @xmath77 0.13 + rxj1049.0 + 5424 & 2.64 @xmath77 0.36 & 1.59 @xmath77 0.67 & 0.16 @xmath77 0.10 + rxj1117.2 + 1744 & 2.24 @xmath77 1.43 & 1.42 @xmath77 1.66 & 0.03 @xmath77 0.12 + rxj1117.5 + 1744 & ... & 2.88 @xmath77 0.68 & 0.24 @xmath77 0.16 + rxj1120.9 + 2326 & ... & 1.17 @xmath77 0.77 & 0.28 @xmath77 0.14 + rxj1123.1 + 1409 & 2.09 @xmath77 0.93 & 1.38 @xmath77 1.58 & 0.23 @xmath77 0.14 + rxj1124.01700 & -0.32 @xmath77 1.65 & 0.43 @xmath77 1.24 & 0.17 @xmath77 0.23 + rxj1124.6 + 4155 & 3.88 @xmath77 1.19 & 2.72 @xmath77 0.92 & 0.20 @xmath77 0.11 + rxj1135.9 + 2131 & 0.34 @xmath77 0.54 & 2.67 @xmath77 0.48 & 0.24 @xmath77 0.11 + rxj1142.0 + 2144 & 4.21 @xmath77 1.50 & 3.61 @xmath77 1.31 & 0.41 @xmath77 0.16 + rxj1146.4 + 2854 & 0.80 @xmath77 0.79 & 1.25 @xmath77 0.42 & 0.23 @xmath77 0.14 + rxj1158.1 + 5521 & 1.56 @xmath77 0.80 & 2.48 @xmath77 0.65 & 0.26 @xmath77 0.10 + rxj1200.90327 & 0.58 @xmath77 1.33 & 2.53 @xmath77 1.71 & 0.22 @xmath77 0.17 + rxj1206.50744 & ... & ... & 0.05 @xmath77 0.05 + rxj1213.5 + 0253 & 4.27 @xmath77 1.39 & 3.05 @xmath77 1.87 & 0.19 @xmath77 0.16 + rxj1218.4 + 3011 & 2.05 @xmath77 1.08 & 4.73 @xmath77 1.60 & 0.22 @xmath770.21 + rxj1221.4 + 4918 & -0.18 @xmath77 0.44 & 6.87 @xmath77 0.99 & 0.13 @xmath770.28 + rxj1237.6 + 2632 & -2.36 @xmath77 0.82 & 3.81 @xmath77 1.22 & 0.16 @xmath77 0.11 + rxj1254.6 + 2545 & -0.28 @xmath77 1.13 & 4.47 @xmath77 1.33 & 0.23 @xmath77 0.25 + rxj1254.8 + 2550 & 1.69 @xmath77 0.89 & 3.26 @xmath77 0.73 & 0.24 @xmath77 0.08 + rxj1256.0 + 2556 & 1.00 @xmath77 0.69 & 3.40 @xmath77 0.84 & 0.24 @xmath77 0.13 + rxj1301.7 + 1059 & 1.88 @xmath77 2.13 & 4.50 @xmath77 1.89 & 0.30 @xmath77 0.16 + rxj1309.9 + 3222 & 2.71 @xmath77 1.45 & 4.64 @xmath77 2.24 & 0.21 @xmath77 0.36 + rxj1337.8 + 3854 & -1.68 @xmath77 1.36 & 0.60 @xmath77 1.08 & 0.24 @xmath77 0.11 + rxj1342.8 + 4028 & 0.50 @xmath77 0.95 & 3.69 @xmath77 1.72 & 0.28 @xmath77 0.36 + rxj1343.4 + 5547 & 0.90 @xmath77 0.38 & 3.29 @xmath77 0.81 & 0.27 @xmath77 0.10 + rxj1354.8 + 6917 & -1.15 @xmath77 0.75 & 3.73 @xmath77 1.33 & 0.40 @xmath77 0.17 + rxj1438.9 + 6423 & 3.32 @xmath77 0.69 & 2.30 @xmath77 0.90 & 0.25 @xmath77 0.15 + rxj1500.0 + 2233 & 1.93 @xmath77 0.97 & 1.22 @xmath77 1.20 & 0.15 @xmath77 0.11 + rxj1515.5 + 4346 & 2.42 @xmath77 1.85 & 3.63 @xmath77 1.43 & 0.25 @xmath77 0.10 + rxj1515.6 + 4350 & 1.08 @xmath77 0.58 & 2.96 @xmath77 0.74 & 0.21 @xmath77 0.10 + rxj1524.6 + 0957 & 2.18 @xmath77 0.64 & 3.09 @xmath77 1.17 & 0.17 @xmath77 0.17 + rxj1537.7 + 1200 & 1.84 @xmath77 0.49 & 2.07 @xmath77 0.72 & 0.25 @xmath77 0.06 + rxj1540.8 + 1445 & 1.61 @xmath77 1.96 & 6.07 @xmath77 2.14 & 0.10 @xmath77 0.38 + rxj1547.3 + 2056 & 0.79 @xmath77 0.72 & 1.01 @xmath77 1.15 & 0.26 @xmath77 0.12 + rxj1552.2 + 2013 & 1.88 @xmath77 1.18 & 5.70 @xmath77 1.22 & 0.32 @xmath77 0.20 + rxj1630.2 + 2434 & 1.41 @xmath77 0.86 & 4.42 @xmath77 0.49 & 0.38 @xmath77 0.17 + rxj1642.6 + 3935 & -0.41 @xmath77 2.74 & -0.90 @xmath77 1.71 & 0.03 @xmath77 0.29 + rxj1659.7 + 3410 & 1.34 @xmath77 0.86 & 4.24 @xmath77 0.88 & 0.27 @xmath77 0.16 + rxj1722.8 + 4105 & -0.26 @xmath77 1.36 & 2.47 @xmath77 1.90 & 0.32 @xmath77 0.22 + rxj1729.0 + 7440 & 5.10 @xmath77 2.00 & -0.44 @xmath77 2.57 & 0.23 @xmath77 0.36 + rxj1746.4 + 6848 & 0.04 @xmath77 2.03 & 2.22 @xmath77 1.84 & 0.23 @xmath77 0.25 + rxj2202.71902 & 2.19 @xmath77 1.11 & 3.43 @xmath77 1.24 & 0.22 @xmath77 0.20 + rxj2212.61713 & 1.45 @xmath77 0.57 & 3.06 @xmath77 0.44 & 0.27 @xmath77 0.09 + rxj2213.51656 & 2.15 @xmath77 1.35 & 4.26 @xmath77 1.06 & 0.26 @xmath77 0.15 + rxj2257.8 + 2056 & 2.62 @xmath77 0.66 & 4.04 @xmath77 1.04 & 0.24 @xmath77 0.10 + rxj2328.8 + 1453 & -1.44 @xmath77 1.15 & 3.94 @xmath77 2.30 & 0.30 @xmath77 0.24 + rxj2348.83117 & 1.05 @xmath77 0.38 & 3.15 @xmath77 0.92 & 0.23 @xmath77 0.14 + [ tab : tabab ]

we find no or emission stronger than or , respectively , in any bcg . a comparison to the detection frequency of nebular emission in bcgs at drawn from the brightest cluster survey indicates that we should have detected roughly one dozen emission - line galaxies , assuming the two surveys are selecting similar clusters in the x - ray luminosity range to . the absence of luminous nebular emission ( ie . , perseus - like systems ) in our sample is consistent with an increase in the number density of _ strong _ cooling flow ( cooling core ) clusters between and today .

we present equivalent widths of the and nebular emission lines for 77 brightest cluster galaxies ( bcgs ) selected from the 160 square degree x - ray survey . we find no or emission stronger than or , respectively , in any bcg . the corresponding emission line luminosities lie below , which is a factor of 30 below that of ngc1275 in the perseus cluster . a comparison to the detection frequency of nebular emission in bcgs at drawn from the brightest cluster survey indicates that we should have detected roughly one dozen emission - line galaxies , assuming the two surveys are selecting similar clusters in the x - ray luminosity range to . the absence of luminous nebular emission ( ie . , perseus - like systems ) in our sample is consistent with an increase in the number density of _ strong _ cooling flow ( cooling core ) clusters between and today . the decline in their numbers at higher redshift could be due to cluster mergers and agn heating .

astro-ph0210398

c

the fuv spectrum of ux uma is everywhere characterized by line features . the fuse and archival hst spectra show that from 910 1600 , the time - averaged spectrum is composed of prominent emission lines and a complex blend of weaker features . the emission lines arise from the usual transitions seen in cv spectra in the uv : resonance lines of neutral hydrogen and ionized metals and prominent excited state transitions . in eclipse , these lines are largely unocculted , while the weaker features are completely eclipsed . the spectrum of the eclipsed light suggests that the weaker line features are primarily in absorption . because of the severe line blending , it is difficult to determine where to place the continuum in the fuv and uv spectra , and some of the eclipsed line region may be in emission as well as in absorption . at mid - eclipse , the spectrum in between the prominent resonance emission lines is smooth and appears to consist only of continuum emission . as a result , the emission line component of the eclipsed light , if present , must have substantial asymmetry or originate very close to the wd to explain why it is not seen during eclipse , when much of the back of the accretion disk remains unocculted . a possible source for asymmetric emission is the mass accretion stream . @xcite created spatially resolved spectra of the accretion disk in ux uma from uv and optical hst / fos observations . they found that the quadrant of the disk containing the mass accretion stream is a source of uv emission , including line emission , down to 0.1r@xmath26 , which they attributed to overflow of the mass stream above the disk . the introduction of an asymmetry in the line emission by the mass stream is also consistent with the phase lag in the orbital radial velocity variations in the fuv spectrum with respect to the photometric ephemeris . the origin of the line absorption also has important implications for the structure of the accretion disk and the wind . @xcite modeled the @xmath41548,1552 line profile in ux uma in and out of eclipse with a kinematic model of a rotating , biconical disk wind . they found that to match the eclipse of the narrow absorption features at the rest velocity of the doublet , their model had to include a dense , slowly outflowing transition region between the fast wind and the disk , which they dubbed an accretion disk chromosphere . some hydrodynamical models of radiatively - driven disk winds have predicted regions of chaotic , low velocity motion that may be the source of the chromospheric features seen in ux uma @xcite . the results of knigge & drew were based on only one epoch of observation of ux uma and only one line , however , which left open the possibility that secular variability and/or non - wind contributions to could be responsible for the narrow absorption components of the line . an alternate explanation for the absorption components is absorption by low - velocity material in the outer disk ; narrow , low - velocity absorption lines from the outer disk are present in the fuv spectrum of the dwarf nova u gem in outburst , for example @xcite . @xcite showed that the narrow line absorption in ux uma also occurred in @xmath41393,1402 and that the absorption line ews decreased during a broadband light curve dip that they attributed to asymmetric vertical extension in the outer disk . the ew decline suggested that the line absorption components were interior to the outer edge of the disk , which made the disk wind / chromosphere model of @xcite more attractive . the fuse observations , obtained 5 years after the hst spectra , show that the line absorption components are present throughout the uv and fuv in a wide variety of species and ionization states and are persistent . the ews of the absorption lines in the fuv spectrum also decrease during the broadband light curve dip . if the line absorption occurred in the outer disk edge , the line ews would increase , not decrease , when the amount of absorbing material increased around phase 0.6 . the low velocities and narrow profiles of the absorption lines are consistent with origin in the dense , slowly outflowing transition region between the accretion disk and the fast wind . nevertheless , both the hst and fuse observations sample a small number of orbits , so it remains possible that the orbital spectral changes are being confused with variability caused by flickering ; additional broadband uv observations spread over the binary orbit of ux uma will be necessary to properly average out the flickering effects . another ongoing question in ux uma and in other cvs concerns the continuum source in the ultraviolet . in ux uma , as in other nls , the accretion disk is believed to be in a high mass - accretion rate steady - state , and we expect it to be the dominant continuum source in the optical and uv . previous studies have not been successful in modeling the continuum emission in ux uma with steady - state disk models , however . @xcite explored model disk fits to broadband hst spectra ( 1250 9000 ) of ux uma . their models were constructed from sums of appropriately weighted stellar atmosphere spectra at the temperature and gravity of each accretion disk annulus . they demonstrated that the model disk spectra provide a poor fit to the overall spectral shape . at optical wavelengths , the models overestimate the amplitude of the balmer jump , while in the uv , the models are too blue to match the observed spectrum . the discrepancy at blue wavelengths between model and observed spectra becomes worse when the fuv spectrum is also taken into account : the model disk spectrum of @xcite would exceed the observed 2001 march fuv flux by @xmath2750% . we examined the fuv fluxes predicted by accretion disk models for the geometry of ux uma ( testing the geometries of both baptista et al . 1995 and smak 1994 ) . optical and uv eclipse mapping analyses have given mass accretion rates for ux uma from to @xcite , where the range reflects a real variation in the observed flux and a likely variation in the mass flow rate through the disk . the distance to ux uma has been estimated to lie from 215 345 pc ( baptista et al . 1995 and sources therein ) , and the reddening is e(b v ) = 0 , or at most e(b v ) @xmath14 0.04 @xcite . we compared the fluxes of fuv and uv spectra of ux uma to model disk spectra constructed from summed , area - weighted , doppler - broadened spectra of stellar atmospheres set to the appropriate temperature and gravity for a given disk annulus , assuming the standard , steady - state temperature distribution in the disk ( see froning et al . 2001 for a full description of the models ) . there is a continued discrepancy between the models and the observed fluxes at blue wavelengths . if the accretion disk is assumed to be the dominant fuv continuum source which corresponds to scaling the model to something like the minima in the fuv spectrum away from the strong lines : at 965 and at 1095 the predicted mass accretion rate is , assuming the upper range to the distance ( increasing the reddening to the upper limit gives @xmath28 ) . although the 1996 november and 2001 march hst and fuse spectra have the same out of eclipse fluxes in their region of overlap ( 1149 1182 ) , the same accretion disk model will not fit both spectra . the @xmath28 model that scales to the flux in the fuv only provides half of the flux observed in 1996 november at 1250 and one third of the observed flux at 1400 . increasing the reddening to e(b v ) = 0.04 has a negligible effect on the slope of the spectrum in the uv . the unknown degree of line blanketing affecting the fuv and uv limits our ability to assess the discrepancy between the observed continuum and the disk models , but it is clear that a single accretion disk model can not fit the continuum over the full fuv+uv range . although the discrepancy between the disk models and observed spectra in high state cvs could result from some problem with the physics of the models , a more likely cause is an incomplete description of the three - dimensional structure of the accretion disk . the presence of substantial vertical extent in the disk , as indicated by the absorption components in the spectral lines and the orbital dip at phase 0.6 , points to a more complex accretion disk structure than current disk models assume . @xcite discussed several ways in which to bring model continuum spectra in closer agreement with the observed spectrum in ux uma . one possibility , first discussed by @xcite with respect to discrepant disk model fits to the spectrum of the nl , ix vel , is that the hottest , inner annuli of the disk are not present , causing the disk spectrum to be less blue than predicted by a standard disk model . the presence of 29 s oscillations in the optical and uv light in ux uma is an inner disk phenomenon and may be indicative of disrupted accretion at the innermost annuli @xcite , but the oscillations are not always present and we do not detect them our observations . another way to bring model spectra into agreement with observed spectra is to posit a second continuum source : @xcite showed that the addition of optically thin recombination emission can flatten both the slope of the continuum from red to blue and fill in the balmer jump . the optically thin continuum emission may originate from the accretion disk chromosphere ; further modeling will be needed to determine if this picture is consistent with theoretical understanding of the properties of the disk - wind transition region . the wd may also be a significant source of uv continuum emission . if so , its presence will worsen the discrepancy between the models and the observed continuum in ux uma @xcite . @xcite concluded from their analysis of uv light curves that the wd is seen in the uv . they identified a signal in the derivative of the eclipse light curve from which they determined contact points for the wd eclipse ingress and egress . they also determined the flux of the wd at 1600 , 4.8 mjy ( ) , which was 25% of the continuum flux at the time of their observation . from their identification of the wd in the uv eclipse and a distance estimate of 345 pc obtained from their optical and uv eclipse maps , @xcite determined the wd diameter and a temperature : 52,000 k if the full wd is visible , or 70,000 k if the lower hemisphere of the wd is obscured by the disk . they combined their measure of the duration of wd eclipse with an empirical zams mass - radius relation and limits on @xmath5 to determine the binary parameters for ux uma , including the mass ratio , inclination , binary separation , and masses and radial velocities of the primary and the secondary . their derived parameters have become the standard adopted values for ux uma . a simple extension of the spectrum of a wd with the flux and temperature found by @xcite to fuv wavelengths indicates a problem with their analysis , however . this is illustrated in figure [ fig_wd ] , in which we compare model wd spectra to the 2001 fuv and ( for reference ) the 1996 november uv spectra of ux uma . the da wd models shown were created using tlusty and synspec @xcite , and have @xmath29 . a wd with a 1600 flux of 4.8 mjy ( ) and a temperature of 52,000 k exceeds the observed flux at nearly every wavelength in the fuv ; only in regions of strong fuv line emission does the model even fall below the observed flux . at its peak near 1000 , the model wd spectrum exceeds the observed flux by 60% . a 70,000 k wd ( the baptista et al . temperature if half the wd is obscured by the disk ) will deviate even more severely from the observed fuv spectrum . we examined the possibility that the wd flux found by @xcite is correct but that their temperature is too high . the second wd model in figure [ fig_wd ] shows that a cooler wd with a 4.8 mjy flux at 1600 is equally unlikely , however . the model shown is for a 20,000 k wd , which is the hottest wd of that uv flux consistent with the observed fuv spectrum . the problem with a cool wd with substantial uv flux is evident : its spectral shape is too variable with wavelength to be consistent with the observed spectrum . moreover , such a wd would need to have a radius of ( for the 216 345 pc distance range to ux uma ) to emit the flux shown , substantially larger than the radius found by baptista et al . we conclude , therefore , that the wd in ux uma is unlikely to have the high flux value determined by baptista et al . a strong change in the wd flux over time is also unlikely . the 1600 flux in 1993 ( the time of the baptista et al . observations ) is roughly the same as the flux in the 1996 hst spectrum . a comparison of the baptista et al . wd model to that of the 1996 spectrum in figure [ fig_wd ] shows that the model begins to overshoot the observed flux around 1200 , even before an accretion disk contribution to the fuv flux is taken into account . this inconsistency between the derived wd parameters and observed fuv spectra calls into question whether the features identified by baptista et al . in the uv eclipse were actually the ingress and egress of the wd . the s / n of the uv eclipse light curve used by baptista et al . to identify the wd was not high , and we consider their identification of the wd contact points in the derivative light curve as insecure . as all subsequent binary parameters found by @xcite depend on the accurate identification of the wd in the eclipse , we also consider their determination of the wd temperature , wd radius , primary and secondary masses , and inclination unreliable . the derived parameters themselves indicate a problem : a mass ratio of 1 , as found by baptista et al . , is theoretically inconsistent with stable mass transfer , for which there is ample and unambiguous evidence in ux uma . although the wd does not have the large flux contribution found by baptista et al . , it may still be an important source of flux in the fuv . we do not see obvious wd eclipse features in the fuv high time resolution light curve , however . figure [ fig_eclc ] shows a slight notch in eclipse egress near phase 0.02 , but that feature is only present in one of the two eclipses we observed ( the individual eclipses in broadband fuv light can be seen in figure [ fig_contlc ] ) . we can set limits on the wd contribution based on the spectrum of the eclipsed light , shown in the bottom panel of figure [ fig_eclipse ] . since the wd is fully eclipsed , the flux in the eclipsed light spectrum gives an upper limit to the wd contribution . if we assume a standard wd mass @xmath30 , a corresponding radius of order @xmath31 cm ( using the mass - radius relation from nauenberg 1972 ) , and a maximum distance of 345 pc , we find that the maximum temperature of the wd can not exceed 40,000 k if the wd flux is to fall below the observed fuv eclipsed flux . at this temperature and projected area , the wd would be responsible for all of the eclipsed flux , which is unlikely , considering that most of the inner accretion disk and the chromosphere are also eclipsed , and no clear - cut wd eclipse features can be seen in the fuv light curve . if we assume therefore that the wd is only responsible for 50% of the fuv eclipsed flux , its temperature would drop to 25,000 k , and a 20% wd contribution would predict a cool wd of 20,000 k. theoretically , half of the energy of mass accretion is emitted in the boundary layer at or near the wd surface ; for the typical accretion rates of ux uma , this indicates that the surface of the wd ( or some part thereof ) should be quite hot ( @xmath11100,000 k ) . if the wd is hot , it must have a small emitting area to be consistent with the eclipsed fuv flux . if fully visible , a 50,000 k wd must have a radius equal to or smaller than , assuming the minimum distance to ux uma of 216 pc . a 70,000 k wd would have to have a radius @xmath32 cm for its flux to be smaller than the observed fuv eclipsed flux . both of these values correspond to a massive wd ( m@xmath33 ) if a standard wd mass - radius relation holds for for ux uma . observationally , a hot bl / wd has not been observed in ux uma , whose x - ray luminosity ( @xmath34 ergs s@xmath1 for a 340 pc distance ) is too low to be consistent with the theoretically predicted bl emission @xcite . this may indicate that we are not seeing the wd in ux uma at all due to shrouding or occultation by the outer disk . occultation of the wd has been demonstrated in the the nl dw uma ( @xmath35 ) when a low state in the disk exposed the normally hidden wd @xcite . a clear detection of the eclipse of the wd would rule out the occultation hypothesis . ideally , a search for the contact points of the wd would be conducted using multiple ( at least 10 ) uv observations ( with fuse or hst ) of the eclipse to decrease the effects of flickering and variability in the shape of the eclipse .

we discuss the implications of our results in the context of models of an accretion disk with a chromosphere between the disk and the extended wind . finally , we note that the observed fuv flux is too low to be consistent with the temperature and radius of the wd derived by baptista et al . ( 1995 ) , suggesting that their remaining binary parameters , including a mass ratio of 1 , ought to be viewed with skepticism .

we present far - ultraviolet ( 905 1182 ) , time - series spectroscopy of the eclipsing , novalike cataclysmic variable , ux uma , acquired with fuse . the time - averaged spectrum is complex and is dominated by overlapping spectral features . the most prominent features are emission lines of , , , and . they are broad ( fwhm 1800 km s ) and double - peaked with a central absorption at zero velocity . during eclipse , the spectrum is simpler : the emission lines remain bright , but the absorption components of the lines and the weaker features between the emission lines disappear entirely , leaving a flat continuum . this behavior is also evident in ghrs ( 1149 1660 ) spectra that we retrieved from the hst archive . the fuv spectra show flickering on time scales of several minutes . the flickering is seen primarily in the continuum and/or the weaker lines rather than in the prominent emission lines . the orbital light curve has a dip in the fuv flux between orbital phases 0.45 0.65 , similar to a pre - eclipse dip detected in hst observations . the ews of the line absorption features decrease during the dip . we have detected a systematic wavelength shift of spectral features on the orbital period , but with a phase lag of , a phenomenon that has been reported at optical wavelengths . we discuss the implications of our results in the context of models of an accretion disk with a chromosphere between the disk and the extended wind . finally , we note that the observed fuv flux is too low to be consistent with the temperature and radius of the wd derived by baptista et al . ( 1995 ) , suggesting that their remaining binary parameters , including a mass ratio of 1 , ought to be viewed with skepticism .

astro-ph0210398

c

in this manuscript , we presented the observational characteristics of the first high resolution time series fuv spectroscopy of the eclipsing nl , ux uma . our main results are as follows : 1 . the time - averaged fuv spectrum of ux uma is dominated by line emission from numerous , overlapping transitions of , , , , , , , , and , and additional unidentified , blended features . there is no region in the time - averaged spectrum in which the continuum alone can be discerned . 2 . the emission lines have double - peaked profiles , with a narrow ( fwhm @xmath19 600 km s@xmath1 ) central absorption reversal at the rest wavelength of each line . in the strongest lines @xmath13977 , @xmath13992 , the blend at 920924 , and the doublet lines of @xmath41032,1038 the red emission peak is enhanced relative to the blue peak . 3 . the mid - eclipse spectrum is much less complex than the time - averaged spectrum . most of the weaker spectral features are occulted in eclipse , exposing a flat continuum with the prominent emission lines superposed . the central absorption components of these lines are eclipsed . the spectrum of the eclipsed light is a mix of continuum and line emission , with the line features primarily in absorption . several lines appear in emission only at mid - eclipse , when the strong absorption components of the lines are occulted . archival hst spectra of ux uma show similar morphology and behavior during eclipse . 4 . over the full fuv range ( 9101182 ) , the mid - eclipse flux is 40% below the uneclipsed level . for regions of the spectrum away from the prominent emission lines , the mid - eclipse flux is on average 82% below the uneclipsed level . the eclipse depth does not change significantly at longer uv wavelengths : the mid - eclipse flux is 77% of the mean out of eclipse flux in the continuum - dominated regions of the hst ( 1149 1660 ) spectra . the full width at half depth of the eclipse away from the prominent emission lines is roughly the same over the fuv and uv as well . the high time resolution ( 1 s ) light curve of the fuv observation shows strong flickering on time scales of several minutes . we do not detect a periodicity associated with the flickers and we do not detect the 29 s oscillations seen in some previous optical and uv observations . spectra of the flickering peaks and minima show that the flickering occurs in regions near the plane of the disk ( i.e. , the regions that are eclipsed ) ; the continuum fluctuates but the prominent emission lines do not flicker . there is an orbital variation in the fuv light curve : the phase - binned flux around orbital phase 0.55 dips to a minimum 25% below the flux at other phases outside of eclipse . the light curve dip occurs at roughly the same orbital phase as but is shallower than the dip seen in hst observations of ux uma . spectra in the light curve dip , averaged over two orbits , show lower ews for the line absorption features than spectra taken outside the dip . the same behavior was seen in the 1996 hst data . we have detected a wavelength shift in the fuv spectrum phased on the orbital period . the amplitude of the shift is 70 km s@xmath1 over broad regions of the fuv but may range as high as 140 200 km s@xmath1 for the absorption components of the lines . the wavelength shifts lag the phasing of the photometric ephemeris by @xmath36 . the wd parameters found by @xcite from hst uv observations requires a fuv flux from the wd in excess of the observed fuv flux . a wd with a radius ( corresponding to a mass range 0.7 0.47 @xmath37 ) must be cooler than 40,000 k to remain below the observed fuv flux . alternatively , the wd can be hotter but smaller in radius . . a curve of growth analysis of interstellar absorption lines in the fuv spectrum of ux uma ( see appendix [ sec_cog ] ) gives a range of @xmath20 for the column density along the line of sight to ux uma .

we present far - ultraviolet ( 905 1182 ) , time - series spectroscopy of the eclipsing , novalike cataclysmic variable , ux uma , acquired with fuse . the time - averaged spectrum is complex and is dominated by overlapping spectral features . the fuv spectra show flickering on time scales of several minutes . we have detected a systematic wavelength shift of spectral features on the orbital period , but with a phase lag of , a phenomenon that has been reported at optical wavelengths .

we present far - ultraviolet ( 905 1182 ) , time - series spectroscopy of the eclipsing , novalike cataclysmic variable , ux uma , acquired with fuse . the time - averaged spectrum is complex and is dominated by overlapping spectral features . the most prominent features are emission lines of , , , and . they are broad ( fwhm 1800 km s ) and double - peaked with a central absorption at zero velocity . during eclipse , the spectrum is simpler : the emission lines remain bright , but the absorption components of the lines and the weaker features between the emission lines disappear entirely , leaving a flat continuum . this behavior is also evident in ghrs ( 1149 1660 ) spectra that we retrieved from the hst archive . the fuv spectra show flickering on time scales of several minutes . the flickering is seen primarily in the continuum and/or the weaker lines rather than in the prominent emission lines . the orbital light curve has a dip in the fuv flux between orbital phases 0.45 0.65 , similar to a pre - eclipse dip detected in hst observations . the ews of the line absorption features decrease during the dip . we have detected a systematic wavelength shift of spectral features on the orbital period , but with a phase lag of , a phenomenon that has been reported at optical wavelengths . we discuss the implications of our results in the context of models of an accretion disk with a chromosphere between the disk and the extended wind . finally , we note that the observed fuv flux is too low to be consistent with the temperature and radius of the wd derived by baptista et al . ( 1995 ) , suggesting that their remaining binary parameters , including a mass ratio of 1 , ought to be viewed with skepticism .

1102.1983

c

all of our models of the co and co@xmath0 gas were done using a single temperature slab model . we fit the @xmath13co r(16 - 19 ) transitions with the assumption the transitions are optically thin and found a temperature of 2887 k with column density @xmath18 ( see population diagram in fig . 7 ) . if we fit the high - j @xmath11co p(25 - 27 , 30 ) lines to an optically thin model as well , we find the gas would be at a high temperature of @xmath19 with @xmath20 . however , this would indicate a @xmath11co/@xmath13co ratio of only 2.2 , much lower than the mean ratio found in the local ism ( wilson 1999 ) . this indicates all of the observed @xmath11co lines are likely optically thick . to constrain the temperature of the absorbing gas , we modeled @xmath11co gas with a column density @xmath21 , estimated using the n(@xmath13co ) and mean isotopic ratio of 69@xmath226 from wilson ( 1999 ) , and varied the turbulence and temperature . we could not model all the @xmath11co transitions to a single temperature , with the low - j transitions always showing higher optical depths than permitted by fits to the high - j transitions . this is likely due to multiple velocity components being viewed in the low - j transitions , which would account for the difference in velocity shift and width from the high - j transitions . we thus fit only the high - j p(25 - 27 , 30 ) transitions , and the goodness - of - fit , shown in fig . 7 , indicates the gas is likely @xmath23 with a doppler b - value @xmath24 . if the @xmath11co is self - shielding and the @xmath11co/@xmath13co is raised to 100 , similar to the ratios found in s09 , these values would not change much as the temperature would still be @xmath25 and @xmath26 . we find the model for the co@xmath0 gas has degenerate fits by changing the microturbulence and temperature , so we use the doppler b - value of @xmath27 found with the co gas , and find @xmath28 and @xmath29 . we further find upper limits for c@xmath0h@xmath0 and hcn , two other organic molecules detected toward irs 46 ( lahuis et al . 2006 ) , to compare molecular abundance ratios . to get upper limits on column densities , we note that temperatures of both the co and co@xmath0 gases toward dg tau b are similar to that found toward irs 46 ( 400 and @xmath30 , respectively ) , so we start with the assumption that if c@xmath0h@xmath0 and hcn are present , they would have similar temperatures to irs 46 ( 700 and 400 k , respectively ) . to derive the upper limits , we individually modeled the absorption spectrum for the molecules at those temperatures . we then used the model - generated profiles to measure feature strengths at various wavelengths in the dg tau b spectrum , giving us a distribution of feature strengths we would expect to measure given only the spectral noise . we then set the upper limit as the column density that would provide a 3@xmath31 integrated intensity . we find the upper limit on the c@xmath0h@xmath0/co@xmath0 ratio ( < 0.19 ) is not very different from the ratio measured for irs 46 ( 0.3 ) , suggesting the non - detection of c@xmath0h@xmath0 in our data may be the result of the lower absorbing column in this source . the upper limit on the hcn / co@xmath0 ratio ( < 0.22 ) is a factor of two smaller than was measured for irs 46 ( 0.5 ) . modeling the gas at higher temperatures with similar column densities decreases the integrated intensity of the absorption features , so higher gas temperatures would result in higher upper limits on the column densities . a temperature increase of @xmath32 would raise the column density upper limit by @xmath215% for both molecules . if we model the molecules at @xmath30 , like the co@xmath0 gas , the upper limits of the c@xmath0h@xmath0/co@xmath0 and hcn / co@xmath0 ratios would be 0.06 and 0.18 , respectively . the [ neii ] emission , as seen toward dg tau b , is a feature commonly found by spitzer irs toward ttss ( lahuis et al . gdel et al . ( 2010 ) showed that among tts the [ neii ] emission strength increases with the jet mass loss rate , the latter as inferred from the high velocity component of the [ oi ] emission . mundt & fried ( 1983 ) found that hh 159 , a jet associated with dg tau b , has a mass loss rate of @xmath33 . if we assume dg tau b is @xmath34 away , our measured [ neii ] luminosity of @xmath35 is similar to that of tts with comparable mass loss rates ( cf . 5 , gdel et al . as we do not see c@xmath0h@xmath0 and hcn in the disk , a simple explanation is they have low abundances in the disk . this could possibly be due to photodissociation . the [ neii ] emission may be an indication the gas is irradiated by x - rays ( glassgold et al . 2007 ) , and/or uv irradiation ( hollenbach & gorti 2009 ) , which in turn could reduce the c@xmath0h@xmath0 and hcn in the disk . however , many t tauri stars with [ neii ] emission similar to or stronger than that of dg tau b show emission in c@xmath0h@xmath0 and hcn ( carr & najita 2011 , in prep ) . alternatively , as c@xmath0h@xmath0 and hcn are commonly found in circumstellar disks ( pontoppidan et al . 2010 ) , these species could be in the disk but simply not in our line of sight . this may be implied by the most striking difference between dg tau b and irs 46 : while the column density of co toward dg tau b is only 10 times that toward irs 46 , the gas ratios to co in dg tau b are 1%-2% of that found toward irs 46 ( see table 6 ) . also , the limits on the ratios n(c@xmath0h@xmath0)/n(co ) and n(hcn)/n(co ) are much smaller for dg tau b than the measured ratios toward gv tau n ( gibb et al . 2007 ) . while the large co column toward dg tau b implies we are probing a large amount of disk material , we are not detecting the large columns of other gaseous species . the disk chemistry model in agndez 2008 ) may provide further insight into the disk conditions probed by our observations . they focus their models on the photon dominated region of the disk , and use photochemistry to estimate the steady state abundances of simple organic molecules . they show the c@xmath0h@xmath0 and hcn have low abundances for disk radii @xmath36 , and the co / co@xmath0 is large for radii @xmath37 . while the temperatures we measured for co and co@xmath0 are cooler than expected for radii < 1 au , the vertical structure of these molecules is not available . one possible scenario is the line of sight passes below the main temperature inversion layer . doppmann et al . ( 2008 ) note gv tau n has a larger @xmath38 than other class i ysos in taurus ( typically @xmath39 ; furlan et al . 2008 ) , and suggest this could happen if a majority of the probed silicate is residing in a warm region of the disk atmosphere . we calculate this ratio for dg tau b and irs 46 for comparison . to do this , we used a polynomial fit across the features to determine the initial continua . for the silicate feature , we fit to the continuum in the 13 - 14.8 mm region , and to the continuum between the water and `` methanol '' ice at wavelengths < 7.8 mm . for the co@xmath0 feature , we assume the co@xmath0 ice is located at larger disk radii than the silicate , as found by watson et al . ( 2004 ) , so the ice is absorbing from a continuum already including the silicate absorption features . we thus use a continuum for dg tau b that is modified from that used in watson et al . ( 2004 ) , found by fitting the spectrum at wavelengths 13 - 14.8 mm and > 15.6 mm . the continuum fits are shown in fig . we then used the relation @xmath40 , where i@xmath41 is the observed flux and i@xmath42 is the emitted continuum , to find @xmath43 and @xmath44 . lahuis et al . ( 2006 ) found irs 46 to have @xmath45 , and pontoppidan et al . ( 2008 ) shows @xmath46 . while both irs 46 and dg tau b have higher ratios than found for class i objects in taurus , both gv tau n and irs 46 have higher ratios ( @xmath47 and 8.6 , respectively ) than dg tau b ( [email protected] ) . additionally , using the co@xmath0 ice column density measurements for irs 46 and dg tau b found in pontoppidan et al . ( 2008 ) , we find irs 46 has a higher co@xmath0 gas / ice ratio , about 8 times that of dg tau b. this suggests the line of sight toward irs 46 is predominantly probing warmer regions than toward dg tau b , resulting in the detection of c@xmath0h@xmath0 and hcn . alternatively ( or in addition ) , the ratios of co@xmath48 gas , ice and silicate could simply be a factor of disk inclination . if the lines of sight for irs 46 and gv tau n pass higher in the disk atmosphere than for dg tau b , raising the silicate to ice ratio , it could also be probing a moderately different region of the disk where the organic molecules reside . if we are only viewing the cold to moderate temperature ( @xmath49 ) regions deeper in the disk in dg tau b , the c@xmath0h@xmath0 and hcn production would be lower in these regions , thus lowering the steady state abundances ( agndez et al . 2008 ) and resulting in their non - detection . if this is the case , the detection of pre - biotic molecules in absorption may be highly sensitive to the line of sight , even for edge - on disks . this could possibly further explain the rarity of molecular absorption among spitzer irs spectra of young stars . our phoenix spectra of the vv cra system shows @xmath11co transitions in absorption toward both components . however , the optically thick lines for the irc and few lines for the primary make it difficult to extract accurate temperature and column density measurements . we therefore use the co abundances reported in s09 for the irc . to get the temperature and column density of the cold @xmath11co toward the primary , we obtained the crires spectra used in s09 from the eso archive . we reduced the data using a similar procedure to that used for the phoenix and nirspec , although the wavelength calibration and telluric division were performed using a telluric model fitter ( seifahrt et al . 2010 ) . we find the cold gas toward the primary has a column density of @xmath50 and temperature @xmath51 ( population diagram is shown in fig . while doppler velocity shifts of the @xmath11co v=(2 - 0 ) absorption lines reported for the irc in s09 are @xmath52 , the crires spectrum shows a smaller shift for the v=(1 - 0 ) lines ( @xmath53 ) , similar to the shift we found with the phoenix spectrum ( @xmath54 ) . one change between the phoenix and crires spectra is the crires spectrum does not show the broad , shallow lines , but rather displays other @xmath11co lines in the -40 to -@xmath55 range ( see fig . 9 ) . thus , these broad lines are in a turbulent , varying region , possibly located in an outflow . while the near - ir spectrum of vv cra is quite complicated , the spitzer irs spectrum of the vv cra binary shows only amorphous silicate and co@xmath0 ice absorption . to determine the optical depths of the silicate and co@xmath0 ice , we used the broadband fluxes found with the t - recs images to estimate continua for the two components . the t - recs images indicate almost all of the silicate absorption is toward the irc , so this feature is not associated with the ism . to get the optical depth of the silicate and co@xmath0 ice , we first subtracted the primary continuum from the spitzer irs spectrum , then we used the same technique used to get the optical depths for dg tau b. the peak optical depth for the 9.7 mm feature seen toward the irc is @xmath56 . using the relation @xmath57 from rieke & lebofsky ( 1985 ) for grains in the ism , we estimate the extinction to be @xmath58 , similar to the estimate by koresko et al . ( 1997 ) of @xmath59 . for the co@xmath0 ice feature , if all of the absorption for this feature is associated with the irc , the optical depth would be @xmath60 , which would give a ratio @xmath61 . at first glance , this high ratio could imply an edge - on disk geometry when comparing to the class i ysos in taurus ( furlan et al . however , the @xmath62 ratio found for the irc is typical of values found by gibb et al . ( 2004 ) for embedded , massive ysos , and found by alexander et al . ( 2003 ) for low- and intermediate - mass ysos . further , the selection in alexandar et al . ( 2003 ) included ysos in the cra cloud , so an edge - on geometry may not be necessary for the high @xmath62 ratio seen toward the irc . we want to characterize the disk structures of dg tau b and vv cra in order to have a context for the gas absorption , so we used the online sed fitting tool described in robitaille et al . ( 2006 ) to get estimates on the disk structure . this tool finds a best fit to an input spectral energy distribution ( sed ) using a database of 200,000 pre - computed radiation transfer models of circumstellar disks , with 14 physical parameters varied at 10 different disk viewing angles . this modeler also estimates the interstellar and circumstellar extinction separately . the user is able to input spectral points , along with the apertures used in the measurements , at wavelengths ranging from the uv to radio . the user can also use upper and lower limits with a confidence percentage that rates how much to penalize models that exceed the limit . the output disk models can be limited to those within a distance and interstellar extinction range . the models are returned in order of increasing @xmath63 value , and while a @xmath63 cut - off is arbitrary and impossible to make formal when fitting well behaved , symmetric models to less behaved ysos ( robitaille et al . 2007 ) , we find for our disks the sed modeler provides good fits up to @xmath64 . while stark et al . ( 2006 ) created synthetic images to model the envelope and disk structure seen in hst near - ir images , giving them the advantage of directly modeling images that show disk structure , we can compare and contrast to the model estimates from the sed fitter . this can provide insight into the model results for vv cra . if we compare the model extinction for vv cra to our extinction estimates using the optical depth @xmath65 , we may be able to constrain the location of the dust , whether in the disk or in the interstellar medium , and test the two disk scenarios presented in s09 . we constrain the models for dg tau b to be between 120 and @xmath66 away ( gottlieb & upson 1969 ; elias 1978 ; kenyon et al . 1994 ) with interstellar extinction up to @xmath67 ( whittet et al . our spitzer irs spectrum was included by inputting the flux at every 1 mm with a 10% error . the sed model fits can be seen in fig . 10 , and the parameter results for the best fit model are listed in table 7 . the models predict a disk that is not seen edge - on ( @xmath68 ) , which is not surprising as robitaille et al . ( 2007 ) found this sed fitting tool does not determine the disk inclination well for most models . robitaille et al . ( 2007 ) used the sed fitting tool to investigate dg tau b and found an edge - on orientation ( inclination=@xmath69 ) for those models with no envelope accretion . however , when envelope accretion - as has been found for dg tau b by stark et al . ( 2006 ) - was added , all inclinations were allowed . we thus exclude disk inclinations from the reported disk parameters . stark et al . ( 2006 ) modeled dg tau b with a central star with @xmath70 , @xmath71 , and minimum disk radius @xmath72 . they calculated an envelope infall rate of @xmath73 , which would be @xmath74 after correcting for our model stellar mass of @xmath75 , similar to the sed model prediction of @xmath76 . however , the cavity opening angle and disk mass for our models are consistently smaller than those found by stark et al . ( 2006 ) . if we again use the ratio @xmath57 from rieke & lebofsky ( 1985 ) , we can use our measured optical depth @xmath43 to estimate the extinction toward dg tau b to be @xmath77 . watson et al . ( 2004 ) also used this absorption feature to estimate an extinction of 30 mag , which they associate with the circumstellar disk . the best fit sed model shows a somewhat higher circumstellar extinction of @xmath78 , but the circumstellar extinction ranges from 26.6 to @xmath79 for the disks within the @xmath63 cut - off , comparable to our extinction estimate of @xmath80 . while the model results for cavity opening angle and disk mass are small , the other parameters appear to be comparable to previous results and estimates . although the sed fitter returns circumstellar extinctions that can be higher than previous estimates , the estimates are within range of those disks in the @xmath63 cut - off . while the irc dominated the system flux at wavelengths > 2.2 mm in the past , it has been progressively fading and is now fainter than the primary at all wavelengths ( see table 8) . we can thus only use data taken within a relatively short period of time for modeling the irc . the primary , on the other hand , has remained constant . regular observations by the the all sky automated survey ( pojmanski 2002 ) show the visually dominated primary has a mean vmag=@xmath81 from the beginning of 2001 through the end of 2009 without any trend . for the primary fluxes at optical wavelengths , we use the ubvri magnitude ratios reported in marraco & rydgren ( 1981 ) . reipurth and zinnecker ( 1993 ) could not extract a gunn z magnitude ratio for vv cra because the primary was saturated . as the largest ratio they quote is 2% , and because the irc is even weaker toward optical wavelengths relative to the primary ( koresko et al . 1997 ) , we constrain the irc flux to be less than 2% of the primary for all wavelengths less than 0.91 mm . for the near - infrared magnitudes of the primary , we use flux values reported by prato et al . ( 2003 ) with observations taken in 1996 and 1997 . again assuming the primary does not vary , we then use relative ratios found with crires spectra reported in s09 to get a k - band irc flux . to get far - infrared information , we use iso observations which were taken when the irc dominated the system flux at those wavelengths . because the irc has faded while the primary has remained constant , we use the far - infrared magnitudes as upper limits for both the primary and irc . for all models of vv cra , we constrained the distance to be 120 - 160 pc away ( casey et al . 1998 ; de zeeuw et al . 1999 ) . vv cra has been classified as a late k - type ( appenzeller et al . 1986 ; prato et al . 2003 ) , so we also constrained ourselves to those models with stellar temperatures @xmath825,000 k. in order to model the two disk system scenarios proposed by s09 , we modeled the irc twice with two interstellar extinction ranges for the two cases . for the case a structure of vv cra , where the gas absorption is in the primary disk , the primary disk will appear as extinction in the line of sight to the irc disk rather than a part of the irc disk . the sed model fitter uses the kmh model of dust grain sizes described in kim et al . ( 1994 ) to model the disk and interstellar extinction , with a modification at wavelengths 1.25 - 8 mm for the interstellar extinction model ( indebetouw et al . because the sed of the irc is not highly constrained in this wavelength region , we approximate the primary disk as having a grain size distribution similar to the ism , and thus model it as part of the interstellar extinction up to an estimated 45 mag . for this sed fitting , the number of models returned is very large for our @xmath63 cut - off , so we lower the cut - off to @xmath83 so we are not overestimating the parameters ( robitaille et al . the result for the best fit seen in fig . 11 is given in table 9 . the extinction that we would associate with the primary disk ( the interstellar extinction ) is 28.6 mag while the other models within the @xmath63 cut - off have a range 24.3 to 30.2 mag . our estimate of @xmath84 is at the lower end of the model extinction range , similar to what was found for dg tau b. for modeling s09 case b , if the co gas is only in the irc disk , we want to use the interstellar extinction seen toward the primary as an estimate for the irc , so we first model the primary ( see fig . 11 ) . the returned models have interstellar extinction @xmath85<1 mag , with the best fit model showing 0 mag . we compare this extinction to the column density measured with the n(@xmath11co)=@xmath50 measured toward the primary . the average n(co)/@xmath86 for the ism ( rettig et al . 2006 ) would suggest a small @xmath87 , showing this interstellar extinction range is accurate . we therefore constrain the models for the irc to have an interstellar extinction a@xmath88<1 mag . the resulting best fit seen in fig . 11 has circumstellar extinction of 43.4 mag , and the lowest extinction in the @xmath63 cut - off is 35.9 mag , much greater than the estimate of @xmath84 . thus , modeling the sed of the irc for case a , where the primary disk is the main source of extinction for the irc , results in extinction estimates that are more constistent with extinction estimates using the optical depth @xmath65 , suggesting this is a more likely scenario . the gas and dust in disks originate from the surrounding ism , so if we compare the dust extinction and co abundance in the disk to the ism , we can investigate dust settling ( brittain et al . a change in dust - to - gas ratio in the line of sight to a central emitting source is largely due to the fact that the two components are not well mixed throughout the disk . models in rettig et al . ( 2006 ) show that turbulence in the disk will push dust away from the midplane , and the settling velocity increases with grain size due to friction with the gas . this results in a stratification with larger grains closer to the midplane , and more grain growth leads to a higher gas / dust ratio above the midplane . further , the large gas / dust ratios require grain growth as well as dust settling . we adopt the ratio used by rettig et al . ( 2006 ) @xmath89_{disk}/[\mbox{n(co)/a}_{v}]_{interstellar}$ ] , where @xmath90_{interstellar}=1.4\times10^{17}\mbox { cm}^{-2}\mbox { mag}^{-1}$ ] . rettig et al . ( 2006 ) probed the gas / dust ratio in the line of sight to the central emitting source of four disks and found the @xmath3-ratio ranges from 1.8 for edge - on disks to @xmath91 for face - on disks , and we use these values as a starting point of comparison for dg tau b and vv cra . rettig et al . ( 2006 ) uses a@xmath88 estimates from jhk - band photometry , so the gas and dust are being probed at similar wavelengths . our extinction estimates are taken from these wavelengths as well as the rest of the sed . however , while the near - infrared spectrum of dg tau b and the vv cra primary are well - constrained , this is not the case for the vv cra irc . thus , this should be taken as a first estimate for the irc . we use our @xmath11co column density of @xmath92 and circumstellar extinction , estimated using the @xmath65 , of @xmath80 to find @xmath93 . this is significantly larger than the ratios found toward other edge - on disks in rettig et al . ( 2006 ) , which were @xmath21.8 . if we use the extinction found with the sed modeling , @xmath94 , the @xmath3-ratio would still be 4.8 . further , we are using the column density estimated from the @xmath13co column density and the mean local ism isotopic ratio ( wilson 1999 ) . if the @xmath11co is self - shielding , as was found in the disks in s09 , the n(@xmath11co)/n(@xmath13co ) would increase from the ism , and the @xmath3-ratio would only be larger . this indicates a larger grain growth and settling in dg tau b than seen toward the disks in rettig et al . ( 2006 ) . for comparison , lahuis et al . ( 2006 ) found @xmath45 ( a@xmath95 ) and n(co)=@xmath96 , giving @xmath97 for irs 46 . for gv tau n , bowey & adamson ( 2001 ) found @xmath98 and gibb et al . ( 2007 ) found n(co)=@xmath99 , giving @xmath100 . irs 46 and gv tau n thus have gas / dust ratios more similar to the ism , as expected for edge - on disks , further emphasizing the high @xmath3-ratio found in dg tau b. this also suggests the line of sight in dg tau b is not probing deeper in the disk atmosphere than irs 46 and gv tau n , nearer where the dust is settled , as is suggested by the @xmath62 ratio . rather , the higher @xmath62 ratios in irs 46 and gv tau n may just be due to having lines of sight that view more silicate in the warmer regions of the disk , where the co@xmath48 ice has sublimated , than for dg tau b. the @xmath3-ratios in rettig et al . ( 2006 ) are for probing the gas / dust ratio in the radial structure toward the central emitting source . this method can therefore not be used to test the case a structure for vv cra , as the geometry in this scenario would be probing the vertical structure of the primary disk . for case b , the vv cra irc was found by s09 to have a hot @xmath11co gas component with column density @xmath101 and cold @xmath11co gas of @xmath102 . by adding these column densities , and using the a@xmath88 estimate of @xmath103 found with the silicate absorption , we find @xmath1042.74 . comparing to the other sources in rettig et al . ( 2006 ) , this is comparable to a disk with inclination @xmath105 . the warm gases toward irs 46 and gv tau n show velocity shifts that suggest they are not in simple keplerian rotation in the disk . irs 46 shows @xmath11co and hcn blueshifted @xmath106 from the quiescent cloud ( lahuis et al . 2006 ) , and gv tau n shows warm hcn redshifted @xmath107 from the stellar velocity ( doppmann et al . similarly , dg tau b shows a shift that raises the question of where the gas resides . the molecular cloud surrounding dg tau b has a velocity @xmath108 ( mitchell et al . 1994 ) , so we would expect a similar velocity for dg tau b ( herbig 1977 ) . however , the @xmath13co and high - j @xmath11co gas seen toward dg tau b have an excess redshift of @xmath109 , so the gas does not appear to be in simple keplerian rotation . the gas is not likely associated with a disk wind , which would result in a blue shift , or with the optical jets as they have velocities @xmath110 and @xmath111 relative to the cloud velocity ( eislffel & mundt 1998 ) . a molecular outflow was found to have a blue ( v@xmath112=-3.8 to 5.1 @xmath113 ) and red ( 7.5 to 10.9@xmath114 ; mitchell et al . 1997 ) component , but our observed shift is outside these bounds . while the co velocity shift is near the extreme end of the red outflow , jhk - band images show the main source of continuum is southeast of the dark lane ( padgett et al . 1999 ) , in the direction of the blue outflow component . the molecular outflow does encase the surrounding cloud velocity , indicating the dg tau b systemic velocity is likely similar to the surrounding cloud . one explanation for the observed shift could be infalling gas . assuming the stellar mass of @xmath115 from our sed model fitting , the velocity of the infalling gas @xmath116 would be @xmath117 at @xmath118 . an envelope infall rate of @xmath76 would imply a density @xmath119 . assuming a co / h ratio of @xmath120 and a column length of 5 au , we would estimate a column density of @xmath121 , similar our measured @xmath122 . at this radius , the accretion radiation from an infalling envelope has been modeled to heat the gas to @xmath123 for l@xmath124 ( see fig . 4 in ceccarelli et al . 1996 ) . however , the total luminosity of dg tau b is quite small compared to those models , found previously to range between 0.2 and 2.5 l@xmath125 ( stark et al . 2006 ; furlan et al . 2008 ) , and between 0.7 and 1.5 l@xmath125 by our sed model fits . if we were to increase the stellar mass to @xmath126 , this would push a @xmath117 infall out to 25 au where even an accretion radiation of 65 l@xmath125 would result in gas temperatures @xmath127 . also , if the co gas is associated with an infalling envelope , we would expect the gas / dust ratio to be near the ism value , contrary to our findings . the absorption is thus unlikely to be associated with an infalling envelope . another possible explanation for the observed velocity shift could be the co is in the disk of a non - axisymmetric system . the main m - band continuum source is likely the disk which could be non - axisymmetric , possibly due to having a close stellar or planetary companion , hot spots in the disk , or some similar phenomenon . also , the absorbing gas may be non - axisymmetric , as would be the case for an asymmetric disk atmosphere . in either case , the absorption would vary in velocity but average out over time to the systemic velocity . although observations that show such velocity variation would be needed to support this hypothesis , it currently appears to be the most likely scenario . while the spitzer infrared spectrograph was an excellent instrument to detect molecular signatures in circumstellar disks , the detection of organic molecules in absorption was rare . dg tau b is one disk that showed molecular gas in absorption . while we were able to directly detect co@xmath0 , we only found upper limits for c@xmath0h@xmath0 and hcn , the other species found toward irs 46 with the spitzer irs ( lahuis et al . this could simply be a function of lower gas abundances in the disk , or this could be due to probing a slightly different region of the disk where c@xmath0h@xmath0 and hcn have lower abundances , as is suggested by the gas column densities and gas / ice ratios . although we are viewing an edge - on disk like irs 46 and gv tau n , this could imply molecular absorption is highly sensitive to even moderate changes in the line of sight in the disk , which may account for the rarity of such detections even in edge - on disks . further , the large co column density to extinction ratio indicates the disk is likely showing a large amount of grain growth and dust settling when compared to other disks ( rettig et al . finally , while the velocity shift of the absorbing co gas is difficult to explain , the most likely scenario currently appears to be the co is in the disk while the main continuum emitting source or the absorbing gas is non - axisymmetric . while the spitzer irs spectrum of vv cra shows only silicate and co@xmath0 ice absorption , high - resolution spectroscopy reveals the disk has a complicated and changing near - ir spectrum . smith et al . ( 2009 ) proposed the absorbing co gas seen toward the irc is located ( a ) within the primary disk in the line of sight to the irc , or ( b ) within the disk around the irc . to test these cases , we used t - recs images to separate the fluxes of the vv cra components in the irs spectrum , allowing us to determine a dust extinction estimate for the irc . we then modeled the sed for the irc with two interstellar extinction ranges , simulating the two cases proposed by s09 , and found the disk models for case a return circumstellar extinction magnitudes more consistent with our extinction estimate . this supports the view by smith et al . ( 2009 ) that this is a more likely case . alexander , r. d. , casali , m. m. , andr , p. , persi , p. , & eiroa , c. 2003 , a&a , 401 , 613 andrews , s. m. , & williams , j. p. 2005 , apj , 631 , 1134 appenzeller , i. , jankovics , i. , & jetter , r. 1986 , a&as , 64 , 65 beckwith , s. v. w. , & sargent , a. i. 1991 , apj , 381 , 250 bergin , e. a. , aikawa , y. , blake , g. a. , & van dishoeck , e. f. 2007 , in protostars and planets v , ed . b. reipurth , d. jewitt , & k. keil ( tucson , az : univ . arizona press ) , 751 bowey , j. e. & adamson , a. j. 2001 , mnras , 320 , 131 brittain , s. d. , simon , t. , najita , j. r. , & rettig , t. w. 2007 , apj , 659 , 685 brittain , s. d. , rettig , t. w. , simon , t. , & kulesa , c. 2005 , apj , 626 , 283 carr , j. s. , & najita , j. r. 2008 , science , 319 , 1504 casey , b. w. , mathieu , r. d. , vaz , l. p. r. , anderson , j. , & suntze , n. b. 1998 , aj , 115 , 1617 ceccarelli , c. , hollenbach , d. j. , & tielens , a. g. g. m. 1996 , apj , 471 , 400 chelli , a. , cruz - gonzalez , i. , & reipurth , b. 1995 , a&as , 114 , 135 correia , s. , ratzka , t. , duchene , g. , & zinnecker , h. 2007 , in iau symp . 240 , binary stars as critical tools and tests in contemporary astrophysics , ed . w. i. hartkopf , e. f. guinan , & p. harmanec ( cambridge : cambridge univ . press ) , 213 de zeeuw , p. t. , hoogerwerf , r. , de bruijne , j. h. j. , brown , a. g. a. , & blaauw , a. 1999 , aj , 117 , 354 doppmann , g. w. , najita , j. r. & carr , j. s. 2008 , apj , 685 , 298 eislffel , j. , & mundt , r. 1998 , aj , 115 , 1554 elias , j. h. 1978 , apj , 224 , 857 evans , n. j. , et al . 2003 , pasp , 115 , 965 furlan , e. , et al . 2008 , apjs , 176 , 184 gibb , e. l. , van brunt , k. a. , brittain , s. d. , & rettig , t. w. 2007 , apj , 660 , 1572 gibb , e. l. , whittet , d. c. b. , boogert , a. c. a. , & tielens , a. g. g. m. 2004 , apjs , 151 , 35 glassgold , a. e. , meijerink , r. , & najita , j. r. 2009 , 701 , 142 glassgold , a. e. , najita , j. r. & igea , j. 2007 , apj , 656 , 515 gottlieb , d. m. , & upson , w. l. 1969 , apj , 157 , 611 gdel , m. , et al . 2010 , a&a , 519 , a113 herbig , g. h. 1977 , apj , 214 , 747 hinkle , k. h. , cuberly , r. w. , gaughan , n. a. , heynssens , j. b. , joyce , r. r. , ridgway , s. t. , schmitt , p. , & simmons , j. e. 1998 , spie , 3354 , 810 hollenbach , d. & gorti , u. 2009 , apj , 703 , 1203 houck , j. , et al . 2004 , apjs , 154 , 18 indebetouw , r. , et al . 2005 , apj , 619 , 931 kenyon , s. j. , dobrzycka , d. , & hartmann , l. 1994 , aj , 108 , 1872 kim , s. h. , martin , p. g. , & hendry , p. d. 1994 , apj , 422 , 164 khler , r. , neuhuser , r. , krmer , s. , leinert , ch . , ott , t. , & eckart , a. 2008 , a&a , 488 , 997 koresko , c. d. , herbst , t. m. , & leinert , ch . 1997 , apj , 480 , 741 lahuis , f. , van dishoeck , e. f. , blake , g. a. , evans , n. j. , ii , kessler - silacci , j. e. , & pontoppidan , k. m. 2007 , apj , 665 , 492 lahuis , f. , van dishoeck , e. f. , boogert , a. c. a. , et al . 2006 , apj , 636 , l145 luhman , k. l. , allen , p. r. , espaillat , c. , hartmann , l. , & calvet , n. 2010 , apjs , 167 , 256 markwick , a. j. , ilgner , m. , millar , t. j. , & henning , th . 2002 , a&a , 385 , 632 marraco , h. g. , & rydgren , a. e. 1981 , aj , 86 , 62 mclean et al . 1998 , spie , 3354 , 566 mitchell , g. f. , hasegawa , t. i. , dent , w. r. f. , & matthews , h. e. 1994 , apj , 436 , l177 mitchell , g. f. , sargent , a. i. , & mannings , v. 1997 , apj , 483 , l127 mundt , r. , & fried , j. w. 1983 , apj , 274 , l83 najita , j. , carr , j. s. , & mathieu , r. d. 2003 , apj , 589 , 931 padgett , d. l. , brandner , w. , stapelfeldt , k. r. , strom , s. e. , tereby , s. , & koerner , d. 1999 , aj , 117 , 1490 pojmanski , g. 2002 , acta astronomica , 52 , 397 pontoppidan , k. m. , salyk , c. , blake , g. a. , meijerink , r. , carr , j. s. , & najita , j. 2010 , apj , 720 , 887 pontoppidan , k. m. , boogert , a. c. a. , fraser , h. j. , et al . 2008 , apj , 678 , 1005 prato , l. , greene , t. p. , & simon , m. 2003 , apj , 584 , 853 ratzka , th . , leinert , ch . , przygodda , f. , & wolf , s. 2008 , the power of optical / ir interferometry : recent scientific results and 2nd generation instrumentation , ed . a. richichi , f. delplancke , f. paresce , & a. chelli , 431 reipurth , b. , & zinnecker , h. 1993 , a&a , 278 , 81 rettig , t. , brittain , s. , simon , t. , gibb , e. , balsara , d. s. , tilley , d. a. , & kulesa , c. 2006 , apj , 646 , 342 rieke , g. h. , & lebofsky , m. j. 1985 , apj , 288 , 618 robitaille , t. p. , whitney , b. a. , indebetouw , r. , & wood , k. 2007 , apjs , 169 , 328 robitaille , t. p. , whitney , b. a. , indebetouw , r. , wood , k. , & denzmore , p. 2006 , apjs , 167 , 256 rothman , l. s. , et al . 1998 , j. quant . . transfer , 60 , 665 salyk , c. , pontoppidan , k. m. , blake , g. a. , lahuis , f. , van dishoeck , e. f. , & evans , n. j. , ii . 2008 , apj , 676 , l49 seifahrt , a. , kufl , h. u. , zngl , g. , bean , j. l. , richter , m. j. , & siebenmorgen , r. 2010 , a&a , 524 , a11 smith , r. l. , pontoppidan , k. m. , young , e. d. , morris , m. r. , & van dishoeck , e. f. 2009 , apj , 701 , 163 stapelfeldt , k. , burrows , c. j. , krist , j. e. , & the wfpc2 science team . 1997 , in iau symp . 182 , herbig - haro flows and the birth of stars , ed . b. reipurth & c. bertout182 , ( dordrecht : kluwer academic publishers ) , 355 stark , d. p. , whitney , b. a. , stassun , k. , & wood , k. 2006 apj , 649 , 900 telesco , c. m. , pina , r. k. , hanna , k. t. , julian , j. a. , hon , d. b. , & kisko , t. m. 1998 , proc . spie , 3354 , 534 watson , d. m. , et al . 2004 , apjs , 154 , 391 weaver , w. b. , & jones , g. 1992 , apjs , 78 , 239 whittet , d. c. b. , gerakines , p. a. , hough , j. h. , & shenoy , s. s. 2001 , apj , 547 , 872 wilson , t. l. 1999 , rep . phys . , 62 , 143 woods , p. m. & willacy , k. 2009 , apj , 693 , 1360 zasowski , g. , kemper , f. , watson , d. m. , furlan , e. , bohac , c. j. , hull , c. , & green , j. d. 2009 , apj , 694 , 459 ice ( 15.2 mm ) absorption . overplotted are the continua used to find the silicate and co@xmath0 ice peak optical depths ( dashed lines ; see text ) . the vv cra spectrum also shows the summed fluxes of the binary components as found with t - recs ( diamonds ) , with each point summing 70% of the primary flux and 100% of the irc flux . ] ice ( 15.2 mm ) absorption . overplotted are the continua used to find the silicate and co@xmath0 ice peak optical depths ( dashed lines ; see text ) . the vv cra spectrum also shows the summed fluxes of the binary components as found with t - recs ( diamonds ) , with each point summing 70% of the primary flux and 100% of the irc flux . ] co line profile for the vv cra irc showing an optically thick , low - j transitions ( dotted line ; v@xmath128 , fwhm=14.8@xmath114 ) and a broad , optically thin line ( dashed line ; v@xmath129=-24@xmath114 , fwhm=32.9@xmath114 ) . the solid line is the combined fit . ] . the 5@xmath31 contours , in jy / pixel , correspond to : @xmath130=0.080 , @xmath131=0.020 , @xmath132=0.023 , @xmath133=0.038.,title="fig : " ] . the 5@xmath31 contours , in jy / pixel , correspond to : @xmath130=0.080 , @xmath131=0.020 , @xmath132=0.023 , @xmath133=0.038.,title="fig : " ] . the 5@xmath31 contours , in jy / pixel , correspond to : @xmath130=0.080 , @xmath131=0.020 , @xmath132=0.023 , @xmath133=0.038.,title="fig : " ] . the 5@xmath31 contours , in jy / pixel , correspond to : @xmath130=0.080 , @xmath131=0.020 , @xmath132=0.023 , @xmath133=0.038.,title="fig : " ] co and @xmath13co absorption as seen toward dg tau b. the @xmath13co ( diamonds ) is modeled as being optically thin with temperature @xmath134 . the @xmath11co line transitions are optically thick ( shown with triangles , indicating these are lower limits ) and we model the gas ( crosses ) with column density n(@xmath11co)=@xmath21 , t@xmath135 and v@xmath136 . @xmath137 goodness - of - fit contours for the temperature and turbulence velocity of the @xmath11co p(25 - 27 , 30 ) absorption lines seen toward dg tau b assuming a column density @xmath21 . the contours show the 68 , 95 , and 99% confidence levels . ] co line profile for the vv cra irc found with crires , using the p(3 , 6 , 7 , 14 ) transitions as examples . the broad absorption found with phoenix is absent , replaced by other @xmath11co absorption lines . ]

we present findings for dg tau b and vv cra , two of the objects observed in our spitzer irs project to search for molecular absorption in edge - on disks , along with near - ir spectroscopy of the co fundamental transitions and mid - ir imaging . while the only gas absorption seen in the spitzer irs spectrum toward dg tau b is co , we use gas abundances and gas / ice ratios to argue that we are probing regions of the disk that have low organic molecule abundances . this implies the rarity of detecting molecular absorption toward even edge - on disks with spitzer irs is a result of high dependence on the line of sight . we also argue the disk around dg tau b shows high amounts of grain growth and settling . for vv cra , we use the silicate absorption feature to estimate a dust extinction , and model the disk with a spectral energy distribution fitting tool to give evidence in support of the disk geometry presented by smith et al . ( 2009 ) where the primary disk is the main source of extinction toward the infrared companion .

we present findings for dg tau b and vv cra , two of the objects observed in our spitzer irs project to search for molecular absorption in edge - on disks , along with near - ir spectroscopy of the co fundamental transitions and mid - ir imaging . while the only gas absorption seen in the spitzer irs spectrum toward dg tau b is co , we use gas abundances and gas / ice ratios to argue that we are probing regions of the disk that have low organic molecule abundances . this implies the rarity of detecting molecular absorption toward even edge - on disks with spitzer irs is a result of high dependence on the line of sight . we also argue the disk around dg tau b shows high amounts of grain growth and settling . for vv cra , we use the silicate absorption feature to estimate a dust extinction , and model the disk with a spectral energy distribution fitting tool to give evidence in support of the disk geometry presented by smith et al . ( 2009 ) where the primary disk is the main source of extinction toward the infrared companion .

astro-ph0510836

i

our study of the near - ir spectrum of p cygni s nebula has led to the main conclusions below : \1 . the near - ir spectrum of p cygni is dominated by collisionally - excited lines of [ fe ii ] , mainly from the three lowest terms . it closely resembles the near - ir spectrum of the homunculus nebula around @xmath5 car , except that p cygni s nebula lacks the molecular hydrogen seen there ( smith 2002b ) . \2 . using 17 near - ir lines of [ fe ii ] we have determined empirical values for the transition ratios and einstein a coefficients for lines between the lowest 3 terms of fe@xmath78 ( @xmath9 , @xmath8 , and @xmath7 ) . these empirical values are in better agreement with the calculations of ns88 than with those of q96 , but in some cases our values differ substantially from both . the ratio @xmath212567/@xmath216435 is of particular importance , as the two bright lines share a common upper level and are often used to deduce the reddening . we find an experimentally - determined value for this intrinsic line ratio of 1.49 ( adopting @xmath79=1.89 ; lamers et al . 1983 ) ; higher than previous values . \4 . from high - resolution spectra of [ fe ii ] , the geometric parameters of p cygni s inner shell are as follows : the shell radius is 10@xmath8010@xmath81 cm , the average expansion velocity is 136 km s@xmath16 , the dynamic age is about 530 yr , and the geometric filling factor for [ fe ii]-emitting filaments appears to be roughly 0.2 . \5 . based on the present expansion speed compared to the speed of the stellar wind , the dynamic age of the nebula , and the lack of any bright nebulosity inside the main shell , we conclude that the 10-radius [ fe ii]-emitting shell around p cygni is the major product of the 1600 a.d . outburst . \7 . based on models for the observed [ n ii ] @xmath26583/[n i ] @xmath25200 line ratio , we find a likely hydrogen ionization fraction of 0.80.95 for a wide range of possible excitation scenarios , while he is mostly neutral . the total mass of the nebula is about 0.1 @xmath82 , subject to errors in the filling factor and the electron density . this , combined with the weak thermal - ir emission from dust , suggests that the gas : dust mass ratio exceeds the canonical value of 100 by an order of magnitude . destruction of dust grains that released fe into the gas phase may help explain the very bright [ fe ii ] emission in p cygni s nebula . we are grateful to richard green for granting a half - night of director s discretionary time to obtain optical spectra of p cygni s nebula at kpno , and to robert d. gehrz for assistance during the irtf / spex observing run . we thank paco najarro and john hillier for sharing results prior to publication regarding their recent work on radiative transfer modeling of p cygni s atmosphere and wind . n.s.was supported by nasa through grant hf-01166.01a from the space telescope science institute , which is operated by the association of universities for research in astronomy , inc . , under nasa contract nas5 - 26555 . 1.03199 & he i & 3.4@xmath1010@xmath83 & 1.0@xmath1010@xmath83 + 1.03931 & [ n i ] & 4.9@xmath1010@xmath83 & 1.5@xmath1010@xmath83 + 1.04612 & [ ni ii ] & 4.4@xmath1010@xmath83 & 2.0@xmath1010@xmath83 + 1.05428 & n i & 7.4@xmath1010@xmath83 & 1.9@xmath1010@xmath83 + 1.07112 & [ ni ii ] & 4.2@xmath1010@xmath83 & 1.4@xmath1010@xmath83 + 1.08277 & he i & 2.17@xmath1010@xmath84 & 1.3@xmath1010@xmath83 + 1.09123 & he i+[ni ii ] & 5.7@xmath1010@xmath83 & 1.3@xmath1010@xmath83 + 1.09392 & pa@xmath85 & 1.54@xmath1010@xmath84 & 2.0@xmath1010@xmath83 + 1.10603 & ? & 6.4@xmath1010@xmath83 & 1.7@xmath1010@xmath83 + 1.11644 & ? & 4.9@xmath1010@xmath83 & 1.5@xmath1010@xmath83 + 1.12262 & ? & 4.7@xmath1010@xmath83 & 1.5@xmath1010@xmath83 + 1.12890 & o i & 5.9@xmath1010@xmath83 & 1.7@xmath1010@xmath83 + 1.13477 & [ fe ii ] & 1.38@xmath1010@xmath84 & 2.0@xmath1010@xmath83 + 1.23249 & n i & 2.9@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.24819 & [ fe ii ] & 4.7@xmath1010@xmath83 & 0.8@xmath1010@xmath83 + 1.25642 & [ fe ii ] ( @xmath86 ) & 1.904@xmath1010@xmath87&0.8@xmath1010@xmath83 + 1.27005 & [ fe ii ] ( @xmath88 ) & 5.4@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.27855 & [ fe ii ] ( @xmath89 ) & 1.12@xmath1010@xmath84 & 0.9@xmath1010@xmath83 + 1.28177 & pa@xmath90 & 2.58@xmath1010@xmath84 & 1.0@xmath1010@xmath83 + 1.28521 & fe ii & 2.1@xmath1010@xmath83 & 0.4@xmath1010@xmath83 + 1.29396 & [ fe ii ] ( @xmath91 ) & 2.64@xmath1010@xmath84 & 0.6@xmath1010@xmath83 + 1.29749 & [ fe ii ] ( @xmath92 ) & 4.2@xmath1010@xmath83 & 0.7@xmath1010@xmath83 + 1.30014 & [ fe ii ] & 1.4@xmath1010@xmath83 & 0.4@xmath1010@xmath83 + 1.31615 & o i & 1.51@xmath1010@xmath84 & 1.0@xmath1010@xmath83 + 1.32028 & [ fe ii ] ( @xmath93 ) & 5.27@xmath1010@xmath84 & 1.4@xmath1010@xmath83 + 1.32746 & [ fe ii ] ( @xmath94 ) & 1.78@xmath1010@xmath84 & 1.0@xmath1010@xmath83 + 1.53309 & [ fe ii ] ( @xmath95 ) & 2.23@xmath1010@xmath84 & 0.2@xmath1010@xmath83 + 1.59905 & [ fe ii ] ( @xmath96 ) & 8.6@xmath1010@xmath83 & 0.3@xmath1010@xmath83 + 1.63735 & he i & 1.3@xmath1010@xmath83 & 0.3@xmath1010@xmath83 + 1.64319 & [ fe ii ] ( @xmath97 ) & 1.519@xmath1010@xmath87&0.9@xmath1010@xmath83 + 1.65759 & he i & 2.6@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.66369 & [ fe ii ] ( @xmath98 ) & 4.5@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.67664 & [ fe ii ] ( @xmath99 ) & 2.00@xmath1010@xmath84 & 0.6@xmath1010@xmath83 + 1.68121 & br11 & 2.7@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.68713 & fe ii & 5.9@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.69731 & he i & 2.2@xmath1010@xmath83 & 0.3@xmath1010@xmath83 + 1.70099 & he i & 4.0@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.71101 & [ fe ii ] ( @xmath100 ) & 2.7@xmath1010@xmath83 & 0.3@xmath1010@xmath83 + 1.73674 & br10 & 4.2@xmath1010@xmath83 & 0.5@xmath1010@xmath83 + 1.74497 & [ fe ii ] ( @xmath101 ) & 3.9@xmath1010@xmath83 & 0.6@xmath1010@xmath83 + 1.79736 & [ fe ii ] ( @xmath102 ) & 3.9@xmath1010@xmath83 & 0.9@xmath1010@xmath83 + 1.80033 & [ fe ii ] ( @xmath103 ) & 9.6@xmath1010@xmath83 & 0.8@xmath1010@xmath83 + 1.80963 & [ fe ii ] ( @xmath104 ) & 3.62@xmath1010@xmath84 & 1.2@xmath1010@xmath83 + 2.05835 & he i & 7.3@xmath1010@xmath83 & 1.0@xmath1010@xmath83 + 2.14273 & mg ii ( ? ) & 3.6@xmath1010@xmath83 & 0.7@xmath1010@xmath83 + 2.16547 & br@xmath85 & 8.2@xmath1010@xmath83 & 1.1@xmath1010@xmath83 + 2.22333 & [ fe ii ] & 1.8@xmath1010@xmath83 & 0.6@xmath1010@xmath83 + 2.30786 & [ ni ii ] ( @xmath20 ) & 2.7@xmath1010@xmath83 & 0.7@xmath1010@xmath83 + 2.36907 & [ ni ii ] ( @xmath20 ) & 2.3@xmath1010@xmath83 & 0.7@xmath1010@xmath83 + 2.40651 & ? & 4.1@xmath1010@xmath83 & 1.2@xmath1010@xmath83 + 10 - 1/10 - 2 1.25668/1.32055 & 3.80 & 3.82 & 3.66 & 3.68 & 3.62 & 3.61 & 0.10 + 10 - 1/10 - 6 1.25668/1.64355 & 1.04 & 1.36 & 0.87 & 1.15 & 1.25 & 1.25 & 0.01 + 10 - 1/10 - 7 1.25668/1.80939 & 5.17 & 6.75 & 4.17 & 5.44 & 5.24 & 5.25 & 0.18 + 10 - 2/10 - 6 1.32055/1.64355 & 0.27 & 0.36 & 0.24 & 0.31 & 0.35 & 0.35 & 0.01 + 10 - 2/10 - 7 1.32055/1.80939 & 1.36 & 1.77 & 1.14 & 1.48 & 1.45 & 1.46 & 0.06 + 10 - 6/10 - 7 1.64355/1.80939 & 4.99 & 4.97 & 4.77 & 4.75 & 4.19 & 4.20 & 0.14 + 11 - 2/11 - 3 1.24819/1.29427&&&&&0.35 & 0.35 & 0.06 + 11 - 3/11 - 4 1.29427/1.32778 & 1.74 & 1.64 & 1.70 & 1.61 & 1.48 & 1.48 & 0.09 + 11 - 3/11 - 6 1.29427/1.53347 & 0.75 & 0.94 & 0.67 & 0.84 & 1.18 & 1.18 & 0.03 + 11 - 3/11 - 7 1.29427/1.67688 & 1.03 & 1.30 & 0.88 & 1.11 & 1.32 & 1.32 & 0.05 + 11 - 3/11 - 8 1.29427/1.80002 & 1.51 & 1.89 & 1.25 & 1.56 & 2.75 & 2.75 & 0.24 + 11 - 4/11 - 6 1.32778/1.53347 & 0.43 & 0.57 & 0.39 & 0.52 & 0.80 & 0.80 & 0.05 + 11 - 4/11 - 7 1.32778/1.67688 & 0.59 & 0.79 & 0.52 & 0.69 & 0.89 & 0.89 & 0.06 + 11 - 4/11 - 8 1.32778/1.80002 & 0.87 & 1.15 & 0.73 & 0.97 & 1.86 & 1.85 & 0.19 + 11 - 6/11 - 7 1.53347/1.67688 & 1.37 & 1.38 & 1.31 & 1.31 & 1.11 & 1.12 & 0.03 + 11 - 6/11 - 8 1.53347/1.80002 & 2.01 & 2.00 & 1.86 & 1.85 & 2.33 & 2.32 & 0.19 + 11 - 7/11 - 8 1.67688/1.80002 & 1.47 & 1.46 & 1.42 & 1.41 & 2.09 & 2.08 & 0.18 + 12 - 4/12 - 5 1.27878/1.29777 & 2.30 & 2.28 & 2.28 & 2.26 & 2.67 & 2.67 & 0.49 + 12 - 4/12 - 7 1.27878/1.59947 & 0.73 & 0.86 & 0.64 & 0.74 & 1.31 & 1.30 & 0.11 + 12 - 4/12 - 8 1.27878/1.71113 & 2.78 & 3.27 & 2.33 & 2.74 & 4.14 & 4.15 & 0.57 + 12 - 4/12 - 9 1.27878/1.79710 & 1.62 & 1.89 & 1.33 & 1.55 & 2.89 & 2.87 & 0.70 + 12 - 5/12 - 7 1.29777/1.59947 & 0.32 & 0.38 & 0.28 & 0.33 & 0.49 & 0.49 & 0.08 + 12 - 5/12 - 8 1.29777/1.71113 & 1.21 & 1.43 & 1.02 & 1.21 & 1.55 & 1.56 & 0.31 + 12 - 5/12 - 9 1.29777/1.79710 & 0.71 & 0.83 & 0.58 & 0.69 & 1.08 & 1.08 & 0.31 + 12 - 7/12 - 8 1.59947/1.71113 & 3.79 & 3.81 & 3.66 & 3.69 & 3.16 & 3.19 & 0.37 + 12 - 7/12 - 9 1.59947/1.79710 & 2.22 & 2.21 & 2.09 & 2.09 & 2.20 & 2.21 & 0.51 + 12 - 8/12 - 9 1.71113/1.79710 & 0.58 & 0.58 & 0.57 & 0.57 & 0.70 & 0.69 & 0.18 + 13 - 5/13 - 8 1.27035/1.66377 & 0.92 & 1.02 & 0.77 & 0.86 & 1.20 & 1.20 & 0.17 + 13 - 5/13 - 9 1.27035/1.74493 & 1.85 & 2.05 & 1.52 & 1.69 & 1.38 & 1.38 & 0.25 + 13 - 8/13 - 9 1.66377/1.74493 & 2.02 & 2.01 & 1.97 & 1.96 & 1.15 & 1.15 & 0.22 + lccc 10 - 1 12566.80 & 4.74@xmath1010@xmath64 & 4.83@xmath1010@xmath64 & 4.83@xmath1010@xmath64 + 10 - 2 13205.54 & 1.31@xmath1010@xmath64 & 1.33@xmath1010@xmath64 & 1.35@xmath1010@xmath64 + 10 - 6 16435.50 & 5.98@xmath1010@xmath64 & 4.65@xmath1010@xmath64 & 4.26@xmath1010@xmath64 + 10 - 7 18093.95 & 1.32@xmath1010@xmath64 & 1.03@xmath1010@xmath64 & 1.07@xmath1010@xmath64 + 11 - 2 12481.9 & & & 0.35@xmath1010@xmath64 + 11 - 3 12942.68 & 1.98@xmath1010@xmath64 & 1.94@xmath1010@xmath64 & 1.94@xmath1010@xmath64 + 11 - 4 13277.76 & 1.17@xmath1010@xmath64 & 1.21@xmath1010@xmath64 & 1.32@xmath1010@xmath64 + 11 - 6 15334.71 & 3.12@xmath1010@xmath64 & 2.44@xmath1010@xmath64 & 1.74@xmath1010@xmath64 + 11 - 7 16768.76 & 2.49@xmath1010@xmath64 & 1.94@xmath1010@xmath64 & 1.63@xmath1010@xmath64 + 11 - 8 18000.16 & 1.82@xmath1010@xmath64 & 1.43@xmath1010@xmath64 & 0.81@xmath1010@xmath64 + 12 - 4 12787.76 & 2.45@xmath1010@xmath64 & 2.25@xmath1010@xmath64 & 2.66@xmath1010@xmath64 + 12 - 5 12977.73 & 1.08@xmath1010@xmath64 & 1.00@xmath1010@xmath64 & 1.00@xmath1010@xmath64 + 12 - 7 15994.73 & 4.18@xmath1010@xmath64 & 3.28@xmath1010@xmath64 & 2.20@xmath1010@xmath64 + 12 - 8 17111.29 & 1.18@xmath1010@xmath64 & 0.92@xmath1010@xmath64 & 0.72@xmath1010@xmath64 + 12 - 9 17971.04 & 2.12@xmath1010@xmath64 & 1.67@xmath1010@xmath64 & 1.06@xmath1010@xmath64 + 13 - 5 12703.46 & 3.32@xmath1010@xmath64 & 2.91@xmath1010@xmath64 & 2.91@xmath1010@xmath64 + 13 - 8 16637.66 & 4.75@xmath1010@xmath64 & 3.73@xmath1010@xmath64 & 2.68@xmath1010@xmath64 + 13 - 9 17449.34 & 2.47@xmath1010@xmath64 & 1.95@xmath1010@xmath64 & 2.39@xmath1010@xmath64 + llccc v@xmath105 ( heliocentric ) & km s@xmath16 & & & 38 + radial expansion velocity & km s@xmath16 & 120 & 152 & 136 + radius ( @xmath28 and @xmath29 ) & arcseconds & 7.8 & 9.7 & 8.8 + radius ( @xmath28 and @xmath29 ) & log@xmath106 cm & 17.3 & 17.4 & 17.35 + dynamic age ( r / v ) & years & 531 & 524 & 527 + filling factor & & & & 0.2 + avg . electron density & @xmath43 & & & 6000 + h ionization fraction & & & & 0.8 - 0.95 + mass & m@xmath3 & & & 0.1 + kinetic energy & log@xmath106 ergs&&&46.3 + ke / lt & & & & @xmath480.01 +

high - dispersion spectra of [ fe ii ] constrain the geometry , detailed structure , and kinematics of p cygni s nebula , which is the major product of p cygni s outburst in 1600 a.d . we propose that the dust which obscured the star after the outburst has since been largely destroyed , releasing fe into the gas phase to produce the bright [ fe ii ] emission .

we present moderate and high - dispersion 12.5 spectra of the-radius nebula around p cygni , dominated by bright emission lines of [ fe ii ] . observed [ fe ii ] line ratios disagree with theoretical transition rates in the literature , so we use the spectrum of p cygni s nebula to constrain the atomic data for low - lying levels of [ fe ii ] . of particular interest is the ratio [ fe ii ]/ , often used as a reddening indicator , for which we empirically derive an intrinsic value of 1.49 , which is 1040% higher than previous estimates . high - dispersion spectra of [ fe ii ] constrain the geometry , detailed structure , and kinematics of p cygni s nebula , which is the major product of p cygni s outburst in 1600 a.d . we use the [ n ii]/[n i ] line ratio to conclude that the nebula is mostly ionized with a total mass of.1 m ; more than the mass lost by the stellar wind since the eruption . for this mass , we would expect a larger infrared excess than observed . we propose that the dust which obscured the star after the outburst has since been largely destroyed , releasing fe into the gas phase to produce the bright [ fe ii ] emission . the kinetic energy of this shell is ergs , far less than the kinetic energy released during the giant eruption of car in the 1840s , but close to the value for car s smaller 1890 outburst . in this respect , it is interesting that the infrared spectrum of p cygni s nebula resembles that of the `` little homunculus '' around car , ejected in that star s 1890 eruption . the mass and kinetic energy in the nebulae of car and p cygni give insight to the range of parameters expected for extragalactic car - like eruptions .

0901.0710

i

in recent years dedicated searches and serendipitous discoveries have uncovered several optical transients with luminosities intermediate between the well - studied classes of nova eruptions ( @xmath7 mag ) and supernova ( sn ) explosions ( @xmath8 mag ) . as can be expected from the wide gap in luminosity , these intermediate luminosity optical transients ( hereafter , ilots ) appear to be diverse in their properties and origins . some have been classified as low luminosity core - collapse sne , with inferred energies and @xmath9ni masses that are at least an order of magnitude below typical sn events ( e.g. , @xcite ) . others , sometimes initially classified as type iin sne , have been subsequently tagged with the catch - all designation of `` sn impostors '' ( e.g. , @xcite ) , that includes events resembling luminous blue variable ( lbv ) eruptions . finally , a small number of somewhat dimmer events have been recently grouped under the proposed name `` luminous red novae '' ( lrne ; @xcite ) , based on their red optical colors ; the origin of lrne and their relation to each other remain unclear . overall , the various groupings have been subject to debate , and there is no clear agreement about the nature of individual objects or the degree of overlap between the various designations . equally important , the connection between the different types of events and different classes of progenitors remains unclear . for example , some lbv eruptions have been initially classified as sne ( e.g. , sne 1961v and 1954j ; @xcite ) . similarly , the progenitor of the recent event sn2008s , which was detected in archival _ spitzer space telescope _ images , has been argued to be an extreme agb star @xcite or an lbv @xcite . the nature of the event itself is also unclear , with claims of a low mass electron - capture sn @xcite and an lbv - like outburst @xcite . the recent event 85 , and the possibly related eruptive object v838 mon , have been speculated to possibly result from stellar mergers @xcite , but other possibilities have been proposed such as a low - luminosity sn for 85@xcite , and an agb pulse , nova eruption from an embedded common envelope white dwarf , or planet capture for v838 mon ( e.g. , @xcite and references therein ) . regardless of the exact interpretation , it is clear that the various proposed eruption / explosion scenarios mark important phases in the evolution of massive stars , possibly on the path to diverse types of sne . therefore , a mapping between ilots and their progenitors will provide important constraints on the initial conditions ( i.e. , progenitors and circumstellar properties ) of sn explosions . the present uncertainty in this mapping highlights the need for detailed observations that can uncover the conditions both prior to and during the ilot events . these include the identification of progenitors , a detailed study of the circumstellar environment , and measurements of the event properties across the electromagnetic spectrum . such observations will also allow us to determine the connection between the various proposed categories of ilots , and to map the overall diversity of this seemingly heterogeneous class . in this paper we present such a combination of observations for an ilot discovered in the nearby galaxy ngc300 on 2008 april 24 ut ( hereafter , ) . for the first time for such an event , we combine multi - wavelength observations ( uv , optical , radio , and x - rays ) , high - resolution echelle spectroscopy , and archival _ hubble space telescope _ and _ spitzer space telescope _ data to show that was the result of an eruption from a @xmath10 m@xmath5 dust - obscured and relatively compact star , possibly in a binary system . our high - resolution spectra , the first for this event , resemble the yellow hypergiant star irc+10420 ( see also @xcite ) , suggesting that may mark a transition to a similar evolutionary stage . a comparison to the recent events sn2008s and 85 shows that all three events may have originated from the same phenomenon , but with a scatter in the energy release and circumstellar properties , that may point to a range of progenitor masses .

we present multi - epoch high - resolution optical spectroscopy , uv / radio / x - ray imaging , and archival _ hubble _ and _ spitzer _ observations of an intermediate luminosity optical transient recently discovered in the nearby galaxy ngc300 . our high - resolution spectra , the first for this event , are dominated by intermediate velocity ( km s ) hydrogen balmer lines and emission and absorption lines that point to a complex circumstellar environment , reminiscent of the yellow hypergiant irc+10420 . hubble _ data .

we present multi - epoch high - resolution optical spectroscopy , uv / radio / x - ray imaging , and archival _ hubble _ and _ spitzer _ observations of an intermediate luminosity optical transient recently discovered in the nearby galaxy ngc300 . we find that the transient ( ngc300ot2008 - 1 ) has a peak absolute magnitude of mag , intermediate between novae and supernovae , and similar to the recent events m85ot2006 - 1 and sn2008s . our high - resolution spectra , the first for this event , are dominated by intermediate velocity ( km s ) hydrogen balmer lines and emission and absorption lines that point to a complex circumstellar environment , reminiscent of the yellow hypergiant irc+10420 . in particular , we detect broad h&k absorption with an asymmetric _ red _ wing extending to km s , indicative of gas infall onto a massive and relatively compact star ( blue supergiant or wolf - rayet star ) ; an extended red supergiant progenitor is unlikely . the origin of the inflowing gas may be a previous ejection from the progenitor or the wind of a massive binary companion . the low luminosity , intermediate velocities , and overall similarity to a known eruptive star indicate that the event did not result in a complete disruption of the progenitor . we identify the progenitor in archival _ spitzer _ observations , with deep upper limits from _ hubble _ data . the spectral energy distribution points to a dust - enshrouded star with a luminosity of about l , indicative of a m progenitor ( or binary system ) . this conclusion is in good agreement with our interpretation of the outburst and circumstellar properties . the lack of significant extinction in the transient spectrum indicates that the dust surrounding the progenitor was cleared by the outburst . we thus predict that the progenitor should be eventually visible with _ hubble _ if the transient event marks an evolutionary transition to a dust - free state , or with _ spitzer _ if the event marks a cyclical process of dust formation .

1409.0026

i

a quantitative understanding of radiation transport in the cool molecule - dominated regions in planetary and ultracool dwarf atmospheres is essential to many aspects of understanding the temperature structure , thermal evolution , and formation of these objects . in freedman08 ( hereafter , f08 ) we presented a detailed discussion of atomic and molecular line opacities used by our group in modeling the atmospheres of brown dwarfs and giant planets [ e.g.][]marley02,fortney05,saumon08,fortney08a , marley10 . since mean opacities can also be widely used in many contexts , in f08 we also computed rosseland and planck mean opacities from 75 to 4000 k , 0.3 mbar to 300 bar , at metallicities of [ m / h ] = -0.3 , 0.0 , and + 0.3 . these tabulations have since found wide use in the communities working to understand giant planet formation [ e.g.,][]mordasini12a , bodenheimer13 , the temperature structure of giant planet atmospheres [ e.g.][]paxton13 , and planetary and ultracool dwarf thermal evolution [ e.g.,][]batygin11,valencia13,paxton13 . here we extend our previous work in a number of important aspects . in addition to updates in opacities of particular molecules ( described below ) , these new calculations are performed over a much larger phase space . mean opacity calculations at pressures from 1 @xmath0bar to 300 bars and 75 to 4000 k are presented , which is a much larger range in pressure , in particular at higher temperatures and lower pressures , compared to f08 . the overall temperature resolution of this compilation is also much finer . furthermore , we present calculations over a wide range of metallicities , from solar up to [ m / h]=+1.7 ( @xmath150@xmath2 solar ) , in several increments . these high metallicities may well - approximate the metal - rich atmospheres of giant planets , up to levels of the solar system s ice giant planets , uranus and neptune guillot07 . finally , for use in models of irradiated planetary atmospheres , we calculate mean opacities where the temperature in the weighting function is not the local temperature , but rather stellar blackbody temperatures from 3000 - 7000 k , to simulate mean `` incident flux '' or `` visible '' opacities to understand the absorption of incident stellar flux in planetary atmospheres .

we present new calculations of rosseland and planck gaseous mean opacities relevant to the atmospheres of giant planets and ultracool dwarfs . such calculations are used in modeling the atmospheres , interiors , formation , and evolution of these objects . our calculations are an expansion of those presented in freedman et al . ( 2008 ) to include lower pressures , finer temperature resolution , and also the higher metallicities most relevant for giant planet atmospheres . span 1bar to 300 bar , and 75 k to 4000 k , in a nearly square grid . opacities at metallicities from solar to 50 times solar abundances are calculated . in addition to computing mean opacities at these local temperatures , we also calculate them with weighting functions up to 7000 k , to simulate the mean opacities for incident stellar intensities , rather than locally thermally emitted intensities . we provide our extensive opacity tables for public use .

we present new calculations of rosseland and planck gaseous mean opacities relevant to the atmospheres of giant planets and ultracool dwarfs . such calculations are used in modeling the atmospheres , interiors , formation , and evolution of these objects . our calculations are an expansion of those presented in freedman et al . ( 2008 ) to include lower pressures , finer temperature resolution , and also the higher metallicities most relevant for giant planet atmospheres . span 1bar to 300 bar , and 75 k to 4000 k , in a nearly square grid . opacities at metallicities from solar to 50 times solar abundances are calculated . we also provide an analytic fit to the rosseland mean opacities over the grid in pressure , temperature , and metallicity . in addition to computing mean opacities at these local temperatures , we also calculate them with weighting functions up to 7000 k , to simulate the mean opacities for incident stellar intensities , rather than locally thermally emitted intensities . the chemical equilibrium calculations account for the settling of condensates in a gravitational field and are applicable to cloud - free giant planet and ultracool dwarf atmospheres , but not circumstellar disks . we provide our extensive opacity tables for public use .

1107.5065

i

understanding the formation and evolution of the first galaxies and the supermassive black holes ( smbhs ) they host is a key goal of both observational and theoretical astronomy . the first sample of more than twenty bright quasars at z@xmath06 were selected from the imaging data of the sloan digital sky survey ( hereafter sdss main survey , e.g. , fan et al . 2000 , 2006 ) . these objects are very bright in quasar uv and optical emission ( sdss magnitudes of @xmath9 ) and have smbhs with masses of @xmath10 a few @xmath11 ( jiang et al . 2007 ; kurk et al . 2007 ) . the dust and molecular gas properties of this bright quasar sample at z@xmath06 have been studied at millimeter wavelengths , using the max planck millimeter bolometer array ( mambo ) , the iram plateau de bure interferometer ( pdbi ) , and the very large array ( vla , bertoldi et al . 2003a ; 2003b ; petric et al . 2003 ; walter et al . 2003 , 2004 , 2009 ; carilli et al . 2004 , 2007 ; riechers et al . 2009 ; wang et al . 2007 , 2008b , 2010 ) . about 30% of them have strong continuum emission at 250 ghz from 40-to-60 k warm dust ( bertoldi et al . 2003a ; petric et al . 2003 ; wang et al . 2007 , 2008b ) . the fir luminosities of the mambo detections are a few @xmath12 to @xmath13 , and the estimated dust masses are a few @xmath6 ( beelen et al . 2006 ; bertoldi et al . 2003a ; wang et al . most of the 250 ghz - detected sources have also been detected in molecular co line emission with pdbi , yielding molecular gas masses of @xmath14 ( walter et al . 2003 ; carilli et al . 2007 ; riechers et al . 2009 ; wang et al . 2010 , 2011 ) . the fir - co and fir - radio luminosity ratios of these objects are consistent with the trend defined by dusty starburst systems at low and high redshifts ( riechers et al . 2006 ; beelen et al . 2006 ; wang et al . the results provide direct evidence for massive star formation at a rate of @xmath15 , indicating that bulges are being formed at the same time that the smbh is rapidly accreting in at least 1/3 of the optically bright quasar systems at z@xmath06 ( wang et al . 2008b ) . molecular co ( 3 - 2 ) and ( 7 - 6 ) line emission was resolved in one of the fir and co luminous quasars , sdss j114816.64 + 525150.3 at z=6.42 , revealing a source size of @xmath16 ( walter et al . 2004 , 2009 ) . the dynamical mass estimated with the size and line width of the co emission suggest a smbh - spheroidal bulge mass ratio an order of magnitude higher than the typical present - day value ( walter et al . 2004 ) . z@xmath06 quasars that are fainter at uv and optical wavelengths have been discovered from deep optical - near infrared surveys ( the sdss southern deep imaging survey , @xmath17 , jiang et al . 2008 , 2009 , the canada - france high redshift quasar survey [ cfhqs ] , @xmath18 , willott et al . 2007 , 2009a , 2009b ) . near - infrared spectroscopic observations of some of these faint quasars reveal that they are likely to be accreting at the eddington rate , and have less massive black holes ( i.e. , 10@xmath19 to @xmath20 ) than the sdss main survey z@xmath06 quasar sample ( kurk et al . 2009 ; willott et al . the steepness of the quasar luminosity function at high redshifts is such that lower - luminosity / mass objects are more common . thus , these faint z@xmath06 quasars provide an important sample to understand smbh - host galaxy formation at the earliest epoch . the millimeter dust continuum emission from the faint z@xmath06 quasars was first studied with a small sample of four objects discovered by cfhqs ( willott et al . . one of them was detected at 250 ghz with a flux density of @xmath21 , indicating a fir continuum luminosity from dust of @xmath22 . the mean 250 ghz flux density of the four quasars when stacked is @xmath23 , which is about a factor of two smaller than the average emission of the z@xmath06 luminous quasar sample ( willott et al . 2007 ; wang et al . we subsequently detected 250 ghz dust continuum and molecular co ( 6 - 5 ) line emission in another two z@xmath06 quasars with @xmath24 ( sdss j205406.42 - 000514.8 and ndwfs j142516.30 + 325409.0 , jiang et al . 2008 ; cool et al . 2006 ; wang et al . 2008b , 2010 ) and co ( 2 - 1 ) line emission from one cfhqs quasar at z=6.2 ( cfhqs j142952 + 544717 , willott et al . 2010a ; wang et al . these observations revealed a few @xmath25 of 40-to-60 k warm dust and @xmath7 of molecular gas in the quasar host galaxies , i.e. , masses comparable to those found with the bright z@xmath06 quasars . in this paper , we report new millimeter and radio observations of nine z@xmath06 quasars . we then study the fir and radio emission with a sample of 18 z@xmath06 quasars that are faint at uv and optical wavelengths , and investigate the dust and molecular gas emission properties and star forming activity in the host galaxies of the millimeter - detected objects . we summarize the current quasar sample at z@xmath06 and describe the new millimeter and radio observations in section 2 , present the results in section 3 , analyse the continuum emission properties of the faint z@xmath06 quasar sample in section 4 , discuss the star formation and quasar - host evolution in section 5 , and conclude in section 6 . a @xmath26-cdm cosmology with @xmath27 , @xmath28 and @xmath29 is adopted throughout this paper ( spergel et al .

we present new millimeter and radio observations of nine z quasars discovered in deep optical and near - infrared surveys . combining with previous millimeter and radio observations , we study the fir and radio emission and quasar - host galaxy evolution with a sample of 18 z quasars that are faint at uv and optical wavelengths ( rest - frame 1450 magnitudes of ) . the average fir - to - agn uv luminosity ratio of this faint quasar sample is about two times higher than that of the bright quasars at z ( ) . a fit to the average fir and agn bolometric luminosities of both the uv / optically faint and bright z quasars , and the average luminosities of samples of submillimeter

we present new millimeter and radio observations of nine z quasars discovered in deep optical and near - infrared surveys . we observed the 250 ghz continuum in eight of the nine objects and detected three of them . new 1.4 ghz radio continuum data have been obtained for four sources , and one has been detected . we searched for molecular co ( 6 - 5 ) line emission in the three 250 ghz detections and detected two of them . combining with previous millimeter and radio observations , we study the fir and radio emission and quasar - host galaxy evolution with a sample of 18 z quasars that are faint at uv and optical wavelengths ( rest - frame 1450 magnitudes of ) . the average fir - to - agn uv luminosity ratio of this faint quasar sample is about two times higher than that of the bright quasars at z ( ) . a fit to the average fir and agn bolometric luminosities of both the uv / optically faint and bright z quasars , and the average luminosities of samples of submillimeter / millimeter - observed quasars at z to 5 , yields a relationship of . five of the 18 faint z quasars have been detected at 250 ghz . these 250 ghz detections , as well as most of the millimeter - detected optically bright z quasars , follow a shallower trend of defined by the starburst - agn systems in local and high - z universe . the millimeter continuum detections in the five objects and molecular co detections in three of them reveal a few of fir - emitting warm dust and of molecular gas in the quasar host galaxies . all these results argue for massive star formation in the quasar host galaxies , with estimated star formation rates of a few hundred . additionally , the higher fir - to - agn luminosity ratio found in these 250 ghz - detected faint quasars also suggests a higher ratio between star formation rate and supermassive black hole accretion rate than the uv / optically most luminous quasars at z .

1107.5065

c

we present observations of millimeter and radio continuum and co(6 - 5 ) line emission from the host galaxies of quasars at z@xmath06 . the new observations complete our mambo dust continuum survey of all the uv / optically faint z@xmath06 quasars ( @xmath2 ) discovered from sdss . combining with previous data , we calculate the average fir and radio emission for a sample of 18 z@xmath06 quasars with @xmath2 . the mean fir - to - agn uv luminosity ratio of this faint quasar sample is about two times higher than that of the bright z@xmath06 quasars ( @xmath3 ) . a fit to the average fir and agn bolometric luminosities of both the uv / optically faint and bright z@xmath06 quasars , and the average luminosities of samples of submillimeter / millimeter - observed quasars at z@xmath02 to 5 , yields a relationship of @xmath4 . the millimeter observations have detected the most fir luminous objects among these faint z@xmath06 quasars . the 250 ghz dust continuum detections in five of the 18 faint quasars and co ( 6 - 5 ) line detections in three of them ( see table 1 and 2 ) reveal dust masses of a few @xmath25 and molecular gas masses of @xmath7 in the quasar host galaxies , which are comparable to the dust and gas masses found in the bright z@xmath06 quasars ( beelen et al . 2006 ; bertoldi et al . 2003b ; carilli et al . 2007 ; wang et al . 2008a , 2010 ) . their fir luminosities are estimated to be @xmath157 , and the fir and agn bolometric luminosities follow a shallower relationship of @xmath5 , in agreement with the starburst - agn systems at lower redshifts . all these results suggest active star formation in the host galaxies of the millimeter - detected and uv / optically faint quasars at z@xmath06 with sfrs of a few hundred @xmath8 . moreover , the higher average fir - to agn uv luminosity ratios ( table 4 ) found with these objects and the shallow luminosity relationship suggest higher sfr - to - agn accretion rate ratios than that of the more luminous / massive z@xmath06 quasars . further high - resolution imaging of the dust and molecular gas components ( e.g. , with alma ) in the quasar host galaxies will be critical to constrain the host galaxy dynamical mass , black hole - spheroidal host mass ratio , gas and star formation rate surface density of these objects . this work is based on observations carried out with the max planck millimeter bolometer array ( mambo ) on the iram 30 m telescope , the plateau de bure interferometer , and the very large array ( nrao ) . iram is supported by insu / cnrs ( france ) , mpg ( germany ) and ign ( spain ) . the national radio astronomy observatory ( nrao ) is a facility of the national science foundation operated under cooperative agreement by associated universities , inc . we acknowledge support from the max - planck society and the alexander von humboldt foundation through the max - planck - forschungspreis 2005 . dominik a. riechers acknowledges support from from nasa through hubble fellowship grant hst - hf-51235.01 awarded by the space telescope science institute , which is operated by the association of universities for research in astronomy , inc . , for nasa , under contract nas 5 - 26555 . m. a. strauss thanks the support of nsf grant ast-0707266 . beelen , a. 2004 , phd thesis , observatoire de paris beelen , a. , cox , p. , benford , d. j. , dowell , c. d. , kov@xmath158cs , a. , bertoldi , f. , omont , a. , & carilli , c. l. , 2006 , apj , 642 , 694 bertoldi , f. , carilli , c. l. , cox , p. , fan , x. , strauss , m. a. , beelen , a. , omont , a. , & zylka , r. , 2003 , a&a , 406 , l55 bertoldi , f. et al . 2003b , a&a , 409 , l47 carilli , c. l. et al . 2001 , apj , 555 , 625 carilli , c. l. et al . 2004 , aj , 128 , 997 carilli , c. l. et al . 2007 , apj , 666 , l9 condon , j. j. 1992 , ara & a , 30 , 575 cool , r. j. et al . , 2006 , aj , 132 , 823 dwek , e. , galliano , f. , & jones , a. p. 2007 , apj , 662 , 927 downes , d. , & solomon , p. m. 1998 , apj , 507 , 615 fan , x. et al . 2000 , aj , 120 , 1167 fan , x. et al . 2001 , aj , 122 , 2833 fan , x. et al . 2003 , aj , 125 , 1649 fan , x. et al . 2004 , aj , 128 , 515 fan , x. et al . 2006 , aj , 131 , 1203 fomalont , e. b. , kellermann , k. i. , cowie , l. l. , capak , p. , barger , a. j. , partridge , r. b , windhorst , r. a. , & richards , e. a. 2006 , apjs , 167 , 103 guilloteau , s. , & lucas , r. 2000 , aspc , 217 , 299 greve , t. r. et al . 2005 , mnras , 359 , 1165 hao , c. n. , xia , x. y. , mao , s. et al . 2005 , apj , 625 , 78 hao , c. n. , xia , x. y. , mao , s. , deng , z. g. , & wu , h. 2008 , chjaa , 8 , 12 hildebrand , r. h. 1983 , qjras , 24 , 267 isobe , t , feigelson , e. d. , akritas , m. g. & babu , g. j. 1990 , apj , 364 , 104 jiang , l. et al . 2007 , aj , 134 , 1150 jiang , l. et al . 2008 , aj , 135 , 1057 jiang , l. et al . 2009 , aj , 138 , 305 kellermann k. i. , sramek , r. , schmidt , m. et al . 1989 , aj , 98,1195 kennicutt , r. c. 1998 , ara&a , 36 , 189 kreysa , e. et al . 1998 , spie , 3357 , 319 kurk , j. d. et al . 2007 , 669 , 32 kurk , j. d. , walter , f. , fan , x. , jiang , l. , jester , s. , rix , h .- w . , & riechers , d. a. , 2009 , apj , 702 , 833 lutz , d. et al . 2010 , apj , 712 , 1287 maiolino , r. , schneider , r. , oliva , e. , bianchi , s. , ferrara , a. , mannucci , f. , pedani , m. , roca sogorb , m. 2004 , nature , 431 , 533 mcgreer , i. d. , becker , r. h. , helfand , d. j. , & white , r.l . 2006 , apj , 652 , 157 momjian , e. , carilli , c. l. , & mcgreer , i. d. 2008 , aj , 136 , 344 morgan , h. l. , & edmunds , m. g. , 2003 , mnras , 343 , 427 mortlock , d. j. et al . 2009 , a&a , 509 , 97 netzer , h. et al . 2007 , apj , 666 , 806 omont , a. , cox , p. , bertoldi , f. , cox , p. , carilli , c. l. , priddey , r. s. , mcmahon , r. c. , & isaak , k. g. 2001 , a & a , 374 , 371 omont , a. , beelen , a. , bertoldi , f. , mcmahon , r. g. , carilli , c. l. , & isaak , k. g. 2003,a & a , 398 , 857 petric , a. o. , carilli , c. l. , bertoldi , f. , fan , x. , cox , p. , strauss , m. a. , omont , a. , & schneider , d. p. 2003 , aj , 126 , 15 priddey , r. s. , isaak , k. g. , mcmahon , r. g. , & omont , a. 2003a , mnras , 339 , 1183 priddey , r. s. , isaak , k. g. , mcmahon , r. g. , roboson , e. i. , & pearson , c. p. 2003b , mnras , 344 , l74 richards , g. t. , vanden berk , d. e. , reichard , t. a. , hall , p. b. , schneider , d. p. , subbarao , m. , thakar , a. r. , & york , d. g. 2002 , aj , 124 , 1 richards , g. t. et al . 2006 , apjs , 166 , 470 riechers , d. a. et al . 2006 , apj , 650 , 604 riechers , d. a. , et al . 2009 , apj , 703 , 1338 seaquist , e. r. , ivison , r. j. , & hall , p. j. 1995 , mnras , 276 , 867 serjeant , s. & hatziminaoglou , e. 2009 , mnras , 397 , 265 solomon , p. m. , downes , d. , radford , s. j. e. , & barrett , j. w. 1997 , apj , 478 , 144 solomon , p. m. , & vanden bout , p. a. 2005 , ara&a , 43 , 677 , * sv05 * spergel , d. n. et al . 2007 , apjs , 170 , 377 venemans , b. p. , mcmahon , r. g. , warren , s. j. , gonzalez - solares , e. a. , hewett , p. c. , mortlock , d. j. , dye , s. , & sharp , r. g. 2007 , mnras , 376 , l76 walter , f. et al . 2003 , nature , 424 , 406 walter , f. , carilli , c. l. , bertoldi , f. , menten , k. , cox , p. , lo , k. y. , fan , x. , & strauss , m. a. 2004 , apj , 615 , l17 walter , f. , riechers , d. , cox , p. , neri , r. , carilli , c. , bertoldi , f. , wei@xmath159 , a. , & maiolino , r. 2009 , nature , 457 , 699 wang , r. et al . 2007 , aj , 134 , 617 wang , r. et al . 2008a , aj , 135 , 1201 wang , r. et al . 2008b , apj , 687 , 848 wang , r. et al . 2010 , apj , 714 , 699 wang , r. et al . 2011 , apjl , in press ( astro - ph/1105.4199 ) willott , c. j. et al . 2007 , aj , 134 , 2435 willott , c. j. et al . 2009 , aj , 137 , 3541 willott , c. j. et al . 2010a , aj , 139 , 906 willott , c. j. et al . 2010b , aj , 140 , 546 yun , m. s. , reddy , n. a. , & condon , j. j. 2001 , apj , 554 , 803 zheng , x. z. xia , x. y. mao , s. et al . 2002 , aj , 124 , 18 zylka , r. 1998 , mopsi users manual , ( iram : grenoble ) lccccccc name & redshift & @xmath34 & @xmath160 & @xmath161 & @xmath115 & fwhm & @xmath162 + & & & mjy & @xmath163 & @xmath116 & @xmath83 & mjy + ( 1 ) & ( 2 ) & ( 3 ) & ( 4 ) & ( 5 ) & ( 6 ) & ( 7 ) & ( 8) + sdss [email protected] & [email protected] & 22.28 & * [email protected] * & & [email protected] & 283@xmath8587 & [email protected] + sdss j020332.39 + 001229.3 & 5.72 & 20.97 & * [email protected] * & * 195@xmath8522@xmath164 * & @xmath1650.32 & & @xmath1650.12 + sdss [email protected] & 5.82 & 22.28 & [email protected] & & & & + ulas j131911.29 + 095051.4 & [email protected] & 19.65 & * [email protected] * & * 64@xmath8517 * & [email protected] & 537@xmath85123 & [email protected] + cfhqs [email protected] & 6.12 & 19.82 & [email protected]@xmath166 & 23@xmath8518 & & & + sdss j205321.77 + 004706.8 & 5.92 & 21.20 & [email protected] & & & & + sdss j214755.40 + 010755.0 & 5.81 & 21.65 & @[email protected] & @xmath3828@xmath8518 & & & + sdss j230735.40 + 003149.0 & 5.87 & 21.73 & @[email protected] & @xmath3821@xmath8517 & & & + sdss j235651.58 + 002333.3 & 6.00 & 21.77 & [email protected] & & & & + note column ( 1 ) , name ; column ( 2 ) , redshift . co redshifts are presented for the two sources detected by the pdbi , and we list the redshifts measured with the uv quasar emission lines from the discovery papers for the other six sources ( jiang et al . 2008 , 2009 ) ; column ( 3 ) , rest - frame 1450 @xmath1 magnitudes from the discovery papers ( mortlock et al . 2008 , jiang et al . 2008 , 2009 ) . column ( 4 ) , mambo measurements of the 250 ghz dust continuum ; column ( 5 ) , vla measurements of the 1.4 ghz radio continuum ; column ( 6 ) and ( 7 ) , line flux and fwhm of the co ( 6 - 5 ) line emission ; column ( 8) , measurement of the dust continuum at the co ( 6 - 5 ) line frequency . the 250 ghz and 1.4 ghz detections are marked as boldface . + @xmath164wang et al . ( 2008b ) ; @xmath166willott et al . 2007 . lcccccccc name&redshift&@xmath34&@xmath160&@xmath161&@xmath167&@xmath168&@xmath169&reference + & & & mjy & @xmath170jy & @xmath171 & @xmath25 & @xmath7 & + ( 1 ) & ( 2 ) & ( 3 ) & ( 4 ) & ( 5 ) & ( 6 ) & ( 7 ) & ( 8) & ( 9 ) + sdss [email protected]&5.85 & 20.83&[email protected]&40@xmath85130&@xmath1653.4&&&1,2 + cfhqs [email protected]&6.13&21.78@xmath164&*[email protected]*&@xmath3827@xmath8519&[email protected]&1.5&&2,3 + sdss [email protected]&6.08&21.28&[email protected]&@xmath3885@xmath8562&@xmath1653.5&&&2 + sdss j035349.76 + 010405.4&6.05&20.22&[email protected]&17@xmath8519&@xmath1653.2&&&2 + ndwfs j142516.30 + 325409.0&5.89&20.62@xmath164&*[email protected]*&20@xmath8520&[email protected]&3.0&2.0&2,4 + first j142738.59 + 331242.0&6.12&20.33@xmath164&[email protected]&*1730@xmath85131*&@xmath1654.6&&&2,5 + sdss j163033.90 + 401209.6&6.05&20.64&[email protected]&14@xmath8515&@xmath1654.2&&&2,6 + cfhqs j164121.64 + 375520.5&6.05&21.30@xmath164&[email protected]&@xmath3830@xmath8532&@xmath1653.2&&&2,3 + sdss [email protected]&6.04&20.60&*[email protected]*&17@xmath8523&[email protected]&3.1&1.2&2,4 + sdss [email protected]&6.12&21.34&[email protected]&31@xmath8516&@xmath1654.1&&&2 + cfhqs [email protected]&6.42&21.65@xmath164&[email protected]&14@xmath8522&@xmath1653.4&&&2 + + note reference:(1 ) wang et al . 2007 ; ( 2 ) wang et al . 2008b ; ( 3 ) willott et al . 2007 ; ( 4 ) wang et al . 2010 ; ( 5 ) mcgreer et al . 2006 ; ( 6 ) bertoldi et al . + @xmath164the values are calculated from the absolute 1450 @xmath135 magnitudes in the discovery papers ( cool et al . 2006 ; mcgreer et al . 2006 ; willott et al . 2007 ) . the 250 ghz and 1.4 ghz detections are marked as boldface . @xmath166the upper limits of @xmath142 are derived with the 3@xmath40 upper limits of the 250 ghz flux densities . lcccccc name & @xmath142 & @xmath168 & @xmath172 & @xmath173 & @xmath174 & @xmath169 + & @xmath171 & @xmath25 & @xmath175 & @xmath176 & @xmath176 & @xmath7 + ( 1 ) & ( 2 ) & ( 3 ) & ( 4 ) & ( 5 ) & ( 6 ) & ( 7 ) + j0129 - 0035 & [email protected] & 2.9 & & [email protected] & [email protected] & 1.2 + j0203 + 0012 & [email protected] & 2.5 & [email protected] & @xmath1651.0 & @xmath1651.3 & @xmath1651.0 + j0239 - 0045 & @xmath1653.9 & & & & & + j1319 + 0950 & [email protected] & 5.7 & [email protected] & [email protected] & [email protected] & 1.5 + j1509@xmath381749 & @xmath1653.2 & & @xmath1650.038 & & & + j2053 + 0047 & @xmath1654.4 & & & & & + j2147 + 0107 & @xmath1654.4 & & @xmath1650.034 & & & + j2307 + 0031 & @xmath1653.8 & & @xmath1650.032 & & & + j2356 + 0023 & @xmath1652.8 & & & & & + lccccccccc group & number@xmath164 & @xmath177 & @xmath178 & @xmath179 & @xmath180 & number@xmath181 & @xmath182 & @xmath183 & q + & & @xmath171 & mjy & @xmath171 & & & @xmath163 & @xmath175 & + ( 1 ) & ( 2 ) & ( 3 ) & ( 4 ) & ( 5 ) & ( 6 ) & ( 7 ) & ( 8) & ( 9 ) & ( 10 ) + + all objects & 18 & 4.3 & [email protected] & [email protected] & 0.46 & 14 & 19@xmath856 & [email protected] & [email protected] + ( radio quiet ) & 16 & 4.0 & [email protected] & [email protected] & 0.46 & 12 & 2@xmath856 & @xmath1650.013 & @xmath1841.57 + 250 ghz detections & 5 & 4.3 & [email protected] & [email protected] & 1.01 & 4 & 45@xmath8510 & [email protected] & [email protected] + ( radio quiet ) & 4 & 4.2 & [email protected] & [email protected] & 1.03 & 3 & 1@xmath8512 & @xmath1650.024 & @xmath1841.68 + 250 ghz non - detections & 13 & 4.3 & [email protected] & @xmath1651.8 & @xmath1650.42 & 10 & 7@xmath857 & @xmath1650.019 & + ( radio quiet ) & 12 & 3.9 & [email protected] & @xmath1651.8 & @xmath1650.46 & 9 & 2@xmath857 & @xmath1650.015 & + + all objects & 22 & 17.6 & [email protected] & [email protected] & 0.22 & 22 & 50@xmath853 & [email protected] & [email protected] + ( radio quiet ) & 21 & 16.9 & [email protected] & [email protected] & 0.23 & 21 & 38@xmath853 & [email protected] & [email protected] + 250 ghz detections & 8 & 19.1 & [email protected] & [email protected] & 0.33 & 8 & 43@xmath854 & [email protected] & [email protected] + 250 ghz non - detections & 14 & 16.8 & [email protected] & [email protected] & 0.08 & 14 & 59@xmath855 & [email protected] & [email protected] + ( radio quiet ) & 13&15.5 & [email protected] & [email protected] & 0.09 & 13 & 32@xmath855 & [email protected] & [email protected] + note we calculate the average flux densities using @xmath185 , @xmath186 , and @xmath187 ( omont et al . 2003 ) , where @xmath188 and @xmath189 are the measurement and 1@xmath40 rms for each object at 250 ghz or 1.4 ghz . for results with @xmath190 less than 3@xmath191 , we take the stacking average value plus 3@xmath192 as the upper limits of the average flux densities and calculate the corresponding upper limits for the luminosities . + @xmath164 number of sources observed at 250 ghz . @xmath166 average luminosity at rest - frame 1450 @xmath1 , derived from @xmath34 . @xmath181 number of sources observed at 1.4 ghz . @xmath193 a radio spectral index of -0.75 ( condon 1992 ) is adopted here to calculate the rest frame 1.4 ghz lunimosity . , which is plotted on the bottom left , and the contour levels are ( -3 , -2 , 2 , 3 , 4)@xmath19415 @xmath195 pdbi 97 ghz continuum image of the quasar averaged over the line - free channels . the synthesized beam size ( fwhm ) is 3.5@xmath196 , and the contour levels are ( -2 , -1 , 1 , 2 , 3)@xmath1940.08 @xmath197 . the cross in both plots denotes the optical position of the quasar ( mortlock et al . 2008).,title="fig:",height=249 ] -2.8 in 3.5 in 6 quasars observed with the pdbi . the continuum emission underlying the co spectra is removed for the two detections , j1319 + 0950 and j0129@xmath380035 . the synthesized beam size ( fwhm ) of the co velocity - integrated image of j1319 + 0950 is @xmath03.5@xmath196 , and @xmath198 for j0129@xmath380035 and j0203 + 0012 . the top abscissae of the co spectra give the redshift range of the spectral windows , and zero velocity in each panel corresponds to the optical / near - infrared redshifts of z=6.127 , 5.780 , 5.720 from the discovery papers ( mortlock et al . 2008 ; jiang et al . 2009 ) . the solid lines in the spectra show gaussian fits to the line emission . the co peak positions are consistent with the optical quasars for the two detections . , title="fig:",width=576,height=192 ] 6 quasars observed with the pdbi . the continuum emission underlying the co spectra is removed for the two detections , j1319 + 0950 and j0129@xmath380035 . the synthesized beam size ( fwhm ) of the co velocity - integrated image of j1319 + 0950 is @xmath03.5@xmath196 , and @xmath198 for j0129@xmath380035 and j0203 + 0012 . the top abscissae of the co spectra give the redshift range of the spectral windows , and zero velocity in each panel corresponds to the optical / near - infrared redshifts of z=6.127 , 5.780 , 5.720 from the discovery papers ( mortlock et al . 2008 ; jiang et al . the solid lines in the spectra show gaussian fits to the line emission . the co peak positions are consistent with the optical quasars for the two detections . , title="fig:",width=576,height=192 ]

we observed the 250 ghz continuum in eight of the nine objects and detected three of them . / millimeter - observed quasars at z to 5 , yields a relationship of . the millimeter continuum detections in the five objects and molecular co detections in three of them reveal a few of fir - emitting warm dust and of molecular gas in the quasar host galaxies . additionally , the higher fir - to - agn luminosity ratio found in these 250 ghz - detected faint quasars also suggests a higher ratio between star formation rate and supermassive black hole accretion rate than the uv / optically most luminous quasars at z .

we present new millimeter and radio observations of nine z quasars discovered in deep optical and near - infrared surveys . we observed the 250 ghz continuum in eight of the nine objects and detected three of them . new 1.4 ghz radio continuum data have been obtained for four sources , and one has been detected . we searched for molecular co ( 6 - 5 ) line emission in the three 250 ghz detections and detected two of them . combining with previous millimeter and radio observations , we study the fir and radio emission and quasar - host galaxy evolution with a sample of 18 z quasars that are faint at uv and optical wavelengths ( rest - frame 1450 magnitudes of ) . the average fir - to - agn uv luminosity ratio of this faint quasar sample is about two times higher than that of the bright quasars at z ( ) . a fit to the average fir and agn bolometric luminosities of both the uv / optically faint and bright z quasars , and the average luminosities of samples of submillimeter / millimeter - observed quasars at z to 5 , yields a relationship of . five of the 18 faint z quasars have been detected at 250 ghz . these 250 ghz detections , as well as most of the millimeter - detected optically bright z quasars , follow a shallower trend of defined by the starburst - agn systems in local and high - z universe . the millimeter continuum detections in the five objects and molecular co detections in three of them reveal a few of fir - emitting warm dust and of molecular gas in the quasar host galaxies . all these results argue for massive star formation in the quasar host galaxies , with estimated star formation rates of a few hundred . additionally , the higher fir - to - agn luminosity ratio found in these 250 ghz - detected faint quasars also suggests a higher ratio between star formation rate and supermassive black hole accretion rate than the uv / optically most luminous quasars at z .

1611.05259

i

luminous blue variables ( lbvs ) are evolved massive stars ( @xmath3 ) , intrinsically bright ( @xmath4 ) and hot ( o , b spectral type ) . they are unstable and exhibit spectroscopic and photometric variability . during the lbv variability cycle they can resemble a cooler supergiant of spectral type a or f , and show visual magnitude variations over a wide range of amplitudes and timescales ( as discussed and reviewed by * ? ? ? * ; * ? ? ? because of their instability , they suffer mass - loss at high rate ( @xmath5 ) and form circumstellar nebulae . the mechanism that causes this instability is still poorly understood . to explain the common `` s doradus type '' outbursts ( with visual magnitude variations of 1 - 2 mag on timescales of years ) , changes of the photospheric physical conditions have been invoked . this variation of the photospheric physical conditions is caused by a change of the wind efficiency due to variation of the ionisation of fe , which is the main carrier of line - driven stellar winds . this mechanism is known as the `` bi - stability jump '' ( explained by * ? ? ? * ; * ? ? ? * ) , a predicted effect of which is mass - loss variability @xcite . the observational mass - loss rates estimated from different indicators ( e.g. uv and optical emission lines , radio free - free emission ) have often been discrepant , most of the time depending on whether clumped or unclumped wind models were assumed . for example , @xcite found that mass - loss rates estimated from lines in clumpy stellar winds of o stars are systematically smaller than those obtained from squared electron density diagnostics ( e.g. h@xmath6 and radio free - free emission ) with unclumped wind models , resulting in empirical mass - loss rates overestimated by a factor 10 or more . the implication is that line - driven stellar winds are not sufficient to strip off quickly the h envelope , before they evolve to wolf - rayet ( wr ) stars @xcite . enhanced mass - loss was therefore proposed to reduce the stellar mass , possibly through short - duration eruptions or explosions . subsequently , @xcite showed that if macro - clumping ( instead of optically thin , micro - clumping ) is taken into account , lines become significantly weaker and lead to underestimation of the mass - loss rate . finally , showed that for moderate clumping ( factor up to 10 ) and reasonable mass - loss rate reductions ( of a factor of 3 ) the empirical mass - loss rates agree with the observational rates and , more importantly , with the model - independent transition mass - loss rate , which is independent of any clumping effects . the implication of this is that eruptive events are not needed to make wr stars . the mechanism that triggered the `` giant eruptions '' ( with visual magnitude changes larger than 2 mag ) witnessed in the 17th ( p cygni ) and in the 19th century ( @xmath7 carinae ) in our galaxy is still unknown , but some scenarios involving hydrodynamic ( sub - photospheric ) instabilities , rapid rotation and close binarity have been proposed ( e.g. * ? ? ? * and ref . therein ) . the presence of nebulae in most of the known objects ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? * ) suggests that these are a common aspect of the lbv behaviour @xcite . given the short duration of the lbv phase ( @xmath8 ) , combined with the rapid evolution of massive stars , lbvs are rare : only a few tens of objects in our galaxy and in the magellanic clouds ( mcs ) @xcite satisfy the variability criteria coupled with high mass - loss rates @xcite . nevertheless , based on the discovery of dusty ring nebulae surrounding luminous stars , the number of galactic candidate lbvs ( clbvs ) has increased recently to 55 @xcite . a few tens of confirmed lbvs have been discovered in farther galaxies ( e.g. m31 , m33 , ngc2403 * ? ? ? * and ref . therein ) . lbv ejecta are the fingerprints of the mass - loss phenomenon suffered by the star . the lbv nebulae ( lbvne ) observed in our galaxy usually consist of both gas and dust . previous studies of known galactic lbvs at radio wavelengths , which trace the ionised component , estimated the masses of the nebulae and their current mass - loss rates ( e.g. * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? on the other hand , ir observations revealed that the dust is often distributed outside of the ionised region , indicative of mass - loss episodes of different epochs and/or that the nebulae are ionisation - bounded ( e.g. , [email protected] , [email protected] , wray 15 - 751 , ag car , * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ; * ? ? ? * ) . these studies show that multi - wavelength , high spatial resolution observations are needed to determine the mass - loss history and the geometry associated with massive stars near the end of their lives @xcite . this information is fundamental to test evolutionary models . however , some of the parameters associated with the mass - loss still have large uncertainties , partly due to imprecise distance estimates , but also due to arbitrary assumptions about the nebula geometry . to understand the importance of eruptive mass - loss in different metallicity environments , we observed at radio - wavelengths a sub - sample of lbvs in the large magellanic cloud ( lmc ) , that has a lower metal content ( @xmath10 ) than the milky way . we selected this sub - sample based on the presence of an optical nebula @xcite . in @xcite ( hereafter paper i ) we presented for the first time radio observations , performed with atca at 5.5 and . we detected the radio emission associated with lbvs rmc@xmath2127 , rmc@xmath2143and candidates lbvs ( clbvs ) s61 and s119 . in this work we present the most recent observations of clbv s61 , covering a larger spectral domain and including australia telescope compact array ( atca ) , atacama large millimeter / submillimeter array ( alma ) , and very large telescope ( vlt ) visir data . the goals of this work are : _ ( i ) _ to introduce a quantifiable and objective method for determining the nebular mass via bayesian estimation of geometrical nebula parameters ; _ ( ii ) _ to derive the mass - loss history with high temporal resolution ; _ ( iii ) _ to compare the nebular properties of s61 with similar galactic lbvne , with respect to the nebular mass , kinematical age of the nebula and dust production . s61 ( also named sk@xmath2 - 67@xmath2266 and al@xmath2418 ) is only a candidate lbv because , since its first observations @xcite , it has not shown both spectroscopic and photometric variability . the star was classified as luminous supergiant ( ia ) spectral type o8fpe . originally rmc@xmath2127also belonged to this class , until it entered a state of outburst ( between 19781980 , * ? ? ? * ) , during which the of features disappeared and the spectrum evolved through an intermediate b - type to a peculiar supergiant a - type . in the meantime , of - type emission was discovered during a visual minimum of the lbv ag car @xcite , the galactic twin of rmc@xmath2127 . all these findings suggested that ofpe stars and lbvs are physically related ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? * ) , and ofpe supergiant stars are now considered quiescent lbvs . in this paper we will focus our attention on s61 for which @xcite derived the following stellar parameters : @xmath11 , @xmath12 , @xmath13 and @xmath14 . the paper is organised as follows : in section [ sec : data ] we present the new observations and data reduction ; we describe the nebula around s61 , its morphology , flux densities and spectral index ( section [ sec : analysis ] ) ; we present our new public code rhocube @xcite to model 3d density distributions , and derive via bayesian inference the geometrical nebula parameters ( section [ sec : model ] ) . from the marginalized posteriors of all parameters obtained from fitting the and maps of s61 , we estimate the posterior pdf of the ionised mass contained in the nebula . in section [ sec : discuss ] we also show a method to derive the mass - loss history with high temporal resolution and we compare it with s61 s empirical mass - loss rate . we discuss the derived extinction maps and interpret them together with the mid - ir and alma non - detections . finally , in section [ sec : summary ] we summarise our results .

we present new observations of the nebula around the magellanic candidate luminous blue variable s61 . these comprise high - resolution data acquired with the australia telescope compact array ( atca ) , the atacama large millimetre / submillimetre array ( alma ) , and visir at the very large telescope ( vlt ) . we present our new public code rhocube to model 3d density distributions , and determine via bayesian inference the nebula s geometric parameters . we show how the nebula models can be used to derive the mass - loss history with high - temporal resolution . [ firstpage ] stars : individual ( sk - 67 ) stars : mass - loss stars : circumstellar matter stars : massive radio continuum : stars methods : statistical

we present new observations of the nebula around the magellanic candidate luminous blue variable s61 . these comprise high - resolution data acquired with the australia telescope compact array ( atca ) , the atacama large millimetre / submillimetre array ( alma ) , and visir at the very large telescope ( vlt ) . the nebula was detected only in the radio , up to 17 ghz . the 17 ghz atca map , with 0.8 arcsec resolution , allowed a morphological comparison with the _ hubble space telescope _ image . the radio nebula resembles a spherical shell , as in the optical . the spectral index map indicates that the radio emission is due to free - free transitions in the ionised , optically thin gas , but there are hints of inhomogeneities . we present our new public code rhocube to model 3d density distributions , and determine via bayesian inference the nebula s geometric parameters . we applied the code to model the electron density distribution in the s61 nebula . we found that different distributions fit the data , but all of them converge to the same ionised mass , , which is an order of magnitude smaller than previous estimates . we show how the nebula models can be used to derive the mass - loss history with high - temporal resolution . the nebula was probably formed through stellar winds , rather than eruptions . from the alma and visir non - detections , plus the derived extinction map , we deduce that the infrared emission observed by space telescopes must arise from extended , diffuse dust within the ionised region . [ firstpage ] stars : individual ( sk - 67 ) stars : mass - loss stars : circumstellar matter stars : massive radio continuum : stars methods : statistical

1611.05259

c

starting from our best model in section [ sec : model ] , we now derive the mass - loss history of s61 with high temporal resolution , keeping in mind that we could not constrain the electron density distribution . if we know the expansion velocity of this nebula , each voxel of our datacube corresponds to a kinematical age . for instance , in the case of s61 the expansion velocity is 27 @xmath69 @xcite and each voxel in the model of the 17 ghz data has a 1-d size equal to 0.1 arcsec , which , at the assumed distance of the lmc ( 48.5 kpc ) corresponds to 7.3@xmath70 and therefore to a kinematical interval of @xmath71 . we know the mass in each voxel and therefore we can derive the average mass - loss rate @xmath72 in intervals of @xmath73 . according to the shell geometry adopted for the s61 nebula , we can assume that the star has lost mass isotropically . we can finally integrate @xmath72 over shells of radius @xmath16 and thickness @xmath74 , where @xmath16 can vary between 0 and @xmath75 ( @xmath76 is the number of pixels in each dimension of the cube ) . the resulting mass - loss rates for s61 are shown in fig . [ fig : s61massloss ] . in this particular example , the peak of the mass - loss has occurred at epoch @xmath77 yr with a rate of @xmath78 . however , we note that the finite resolution due to the synthesised beam may mean that the real distribution is less smooth . the mass - loss rates derived in fig . [ fig : s61massloss ] are consistent with our non - detection of the stellar wind in the radio maps . if we assume a spherical mass - loss for the star and then the model in @xcite , with a terminal velocity of 250 @xmath69 @xcite , an electron temperature of 6120 k @xcite and a flux density equal to 3 times the noise in the map at , we derive a 3@xmath79 upper limit of @xmath80 @xmath81 , for a fully ionised wind with solar abundances . this upper limit is consistent with the mass - loss rate of @xmath82 @xmath81 derived by @xcite , which would be within the distribution in fig . [ fig : s61massloss ] . the value by @xcite of @xmath83 @xmath81 derived from h emission lines seems inconsistent with the radio non - detection . we now compare the mass - loss history with the empirical mass - loss rate , as predicted for o , b normal supergiants following the procedure described in @xcite . we first assume the stellar parameters by @xcite ( @xmath84 , @xmath85 and @xmath86 ) , a metallicity of @xmath87 and an initial stellar mass of @xmath88 ( according to the evolutionary tracks by * ? ? ? the small nebular mass derived in the previous section suggests that the star has a stellar mass similar to its initial value and then the mass - loss rate is comparable to the one relative to the o , b main sequence stars ( the reduced stellar mass of lbvs causes a strong increase in the mass - loss rate with respect to normal o , b supergiants , as showed by * ? ? ? the empirical mass - loss rate derived with the mentioned parameters is @xmath89 , which is close to the average value of the distribution in fig . [ fig : s61massloss ] . the consequence of this result is either that the mass - loss occurred with a constant wind ( constant density model ) or the mass - loss rates varied due to excursion through the bi - stability jump ( power - law electron density model ) . in this latter case the bi - stability jump for s61 would occur at @xmath90 , with a mass - loss rate of @xmath91 , which is consistent with the peak in fig . [ fig : s61massloss ] within a factor of @xmath92 . in both cases we can probably exclude that s61 lost mass through eruptions , as normal stellar winds perfectly explain the observations . if we instead use the stellar parameters by @xcite ( @xmath93 , @xmath94 ) , the derived empirical mass - loss rate is @xmath95 , which is far higher than our observational values derived in the previous section . we have derived the extinction map of s61 by comparing pixel - by - pixel the highest - resolution radio image ( ) with the _ hst _ @xmath0 image , as the two emissions trace the same gas ( paper i ) . according to @xcite , if the optical @xmath0 emission is due to the de - excitation of the recombined h atom and the radio continuum emission to free - free encounters , one can determine the extinction of the optical line by comparing the two brightnesses , as @xmath96 \label{predicted}\ ] ] where @xmath97 is the electron temperature of the nebula in units of k , @xmath98 is the radio frequency in ghz , y is a factor incorporating the ionised he / h ratios ( assumed to be 1 , as in paper i ) and @xmath99 for the theoretical balmer decrement . , as derived by comparing the @xmath53 recombination line and the centimetre ( ) emission above @xmath100 . the central star was masked with a circular aperture in the optical image . the black contours are 5 , 7 and 9@xmath79 levels of the radio emission and the black ellipse is the resolution . the black cross represents the position of the star , according to the simbad database . ] we re - gridded the _ hst _ image to the same grid of the radio map and converted it to jy pixel@xmath101 unit . we convolved the optical image with the radio beam ( elliptical gaussian with hpbw as in table [ tab : atca2maps ] ) . adopting as electron temperature @xmath102 @xcite , we derive the expected radio map from the @xmath0 recombination - line emission . keeping in mind that we want to estimate the expected free - free emission from the optical line in the nebula , we masked the @xmath53 emission from the star . finally , the extinction map in @xmath53 was derived as @xmath103 in every common pixel with brightness above 5@xmath79 , where @xmath79 was computed by summing in quadrature the noise in the maps and calibration uncertainties ( however negligible ) . as a result of this procedure , we obtained the extinction map illustrated in fig . [ fig : ext1 ] . . the 3@xmath79 upper limit spatially integrated flux densities at the visir and alma observing wavelengths ( and , respectively ) are also shown . ] small extinction due to dust is evident across the whole region . the range of values for a@xmath104 across the nebula is between @xmath300.1 and 1.09 . the maximum value for a@xmath105 is 1.8 and corresponds to the spur - object in the northern ( upper ) part . to derive a range of possible characteristic temperatures for the dust that extinguishes the optical emission , we fitted the flux density distribution from the mid- to the far - ir . we consulted the ir catalogues with the vizier tool @xcite at the position of s61 , and we extracted the flux densities in the _ wise _ bands w3 and w4 @xcite , _ akari _ l18w at 18 @xmath106 @xcite , _ spitzer _ mips at 24 and 70 @xmath106 @xcite and _ herschel _ pacs at 100 @xmath106 @xcite . we fitted a single - temperature greybody with power - law opacity index @xmath107 at longer wavelengths and constant opacity at shorter wavelengths ( e.g. * ? ? ? * ) , @xmath108 we found a range of characteristic temperatures between @xmath109 to @xmath110 by varying the parameter @xmath107 ( we explored the cases for @xmath111 and @xmath112 ) and @xmath113 ( between 18 and 25 @xmath106 ) . the modified black - body that best fits the data is the one represented with a dark continuum line in the figure , with @xmath114 . note that @xmath115 implies either the existence of relatively large grains or different dust components ( of different temperatures ) . the latter is usually observed in galactic lbvne and the temperature of the dust decreases with increasing distance from the star ( e.g. * ? ? ? * ; * ? ? ? according to the flux density distributions in fig . [ fig : sed ] , the expected flux densities @xmath116 at the visir pah2@xmath1172 and q1 central wavelengths ( @xmath118 and @xmath119 ) are : @xmath120 and @xmath121 , respectively . if this emission arises from a point - like source close to the star as observed in other candidate lbvs ( e.g. g79.29 + 0.46 , * ? ? ? * ) , we would have detected it with visir . we deduce then that the dust is spread out over the nebula , at angular scales that our observations were not sensitive to . similarly , we have derived the expected alma flux density at : @xmath122 . even for the most favourable case for the dust ( @xmath123 , corresponding to optically thick large grains ) , the sensitivity achieved with only one execution was not sufficient to detect the thermal emission at sub - mm wavelengths . in fact , with an expected flux density of @xmath124 ( case @xmath125 ) , spread across 16.7 alma synthesised beams , the average brightness would be @xmath126 and a sensitivity of @xmath127 was needed for a @xmath128 detection . note that the extrapolated flux density at the alma frequency is consistent with the upper limit derived from the map : the rms - noise ( @xmath129 ) , integrated over the area corresponding the ionised nebula , yields a 3@xmath79 upper limit of @xmath130 . using the flux density extracted from the best fit ( case @xmath114 ) at the alma frequency ( see fig . [ fig : sed ] ) , we derived a dust mass @xmath131 to @xmath132 , considering that @xmath133 , and assuming @xmath134 . this means a low gas - to - dust ratio for the lmc . it suggests that the sub - mm emission might be even lower than that computed from the flux density distributions based on the mid- and far - ir data . the extinction map resembles the dusty nebula around the galactic lbv iras 18576@xmath9034 @xcite . this was also observed with visir in the filters pah2@xmath1172 and q1 . they derived for the mid - ir nebula a dust component of temperature ranging from 130 to 160 k. iras 18576@xmath9034 has a mid - ir nebula of 7 arcsec diameter , corresponding to 0.35 pc at the distance of 10 kpc . it has a physical size which is about half that of s61 , but in the sky the two sources have similar angular size . we rescaled the iras 18576@xmath9034 visir maps to the distance of the lmc , and we derived from the maps a mean value of @xmath135 in the pah2@xmath1172 filter and @xmath136 in the q1 filter . this means that with our visir observations , we would have detected at @xmath137 in the pah2@xmath1172 filter image a nebula like iras 18576@xmath9034 . the sensitivity reached in band q1 would have not been sufficient to detect the nebula . @xcite derived for iras 18576@xmath9034 a dust mass of @xmath138 ( depending on the assumed dust composition ) and @xcite derived a mass of @xmath139 for the ionised gas . this suggests the dust content in the nebula around s61 is similar in mass to that estimated in iras 18576@xmath9034 . conversely , the ionised mass in s61 is only a small fraction ( 1/20th ) of the mass in the iras 18576@xmath9034 nebula , despite s61 s nebula diameter ( @xmath140 ) being about 3.5 times bigger than iras 18576@xmath9034 ( @xmath141 ) . we recall , however , that iras 18576@xmath9034 has an estimated bolometric luminosity higher than s61 ( @xmath142 , * ? ? ? * ) and the mass , through the mass - loss quadratic dependence on luminosity , has a stronger effect than the metallicity . the inner shell around wray 15 - 751 , which has a luminosity similar to s61 , extends up to @xmath143 , similar to s61 and has gas and dust masses of @xmath144 and @xmath145 , respectively @xcite , more massive than the s61 nebula . this may suggest that s61 s mass - loss has been less efficient over time than the mentioned galactic lbvs . the dust production does not seem significantly different . however , a potential issue for this comparison is the larger uncertainty of the galactic lbvs distances than those of the magellanic objects .

the 17 ghz atca map , with 0.8 arcsec resolution , allowed a morphological comparison with the _ hubble space telescope _ the spectral index map indicates that the radio emission is due to free - free transitions in the ionised , optically thin gas , but there are hints of inhomogeneities .

we present new observations of the nebula around the magellanic candidate luminous blue variable s61 . these comprise high - resolution data acquired with the australia telescope compact array ( atca ) , the atacama large millimetre / submillimetre array ( alma ) , and visir at the very large telescope ( vlt ) . the nebula was detected only in the radio , up to 17 ghz . the 17 ghz atca map , with 0.8 arcsec resolution , allowed a morphological comparison with the _ hubble space telescope _ image . the radio nebula resembles a spherical shell , as in the optical . the spectral index map indicates that the radio emission is due to free - free transitions in the ionised , optically thin gas , but there are hints of inhomogeneities . we present our new public code rhocube to model 3d density distributions , and determine via bayesian inference the nebula s geometric parameters . we applied the code to model the electron density distribution in the s61 nebula . we found that different distributions fit the data , but all of them converge to the same ionised mass , , which is an order of magnitude smaller than previous estimates . we show how the nebula models can be used to derive the mass - loss history with high - temporal resolution . the nebula was probably formed through stellar winds , rather than eruptions . from the alma and visir non - detections , plus the derived extinction map , we deduce that the infrared emission observed by space telescopes must arise from extended , diffuse dust within the ionised region . [ firstpage ] stars : individual ( sk - 67 ) stars : mass - loss stars : circumstellar matter stars : massive radio continuum : stars methods : statistical

1611.05259

c

in this work we presented high spatial resolution observations from the radio to the mid - ir of the nebula associated with the candidate lbv s61 . it was detected only in the centimetre band . the nebula has a morphology resembling a shell , as in the optical , but in the radio there is more sub - structure . the emission mechanism is optically thin free - free , as evidenced by the spectral index map , although there are regions that suggest self - absorption . we developed and made publicly available a code in python that permits to model the 3-d electron density distribution and to derive the mass in the nebula . we tried different geometries for the shell ( truncated gaussian , constant density and power - law @xmath146 and @xmath147 ) and we found that at least three of these geometries give similar - quality results . for all the well - fitting models , the derived ionised mass is always about 0.1 m@xmath148 , which is an order of magnitude smaller than previous estimates and also a few factors smaller than the mass of similar galactic objects . the nebula is very likely density bounded , meaning that part of the stellar uv flux escapes from the nebula . as an application of our modelled electron density distribution , we also show how to derive the mass - loss history with high - temporal resolution ( @xmath30 850 yr ) . the derived mass - loss rates are consistent with the empirical mass - loss rate for s61 , implying that the nebula was likely formed by stellar winds , rather than eruptive phenomena . the present - day mass - loss is @xmath149 @xmath81 . based on the extinction map derived from the radio map and the @xmath150 _ hst _ image , we have explored the possibility that the nebular regions with higher spectral index are dusty , by means of high - resolution mid - ir and sub - mm observations . we did not detect any point - like source , or compact regions associated with the clumps , neither with visir nor with alma . the fit of the ir flux distribution from space telescope observations suggest the presence of dust with a range of characteristic temperatures of @xmath151 and dust mass @xmath152 of @xmath153 to @xmath154 . based on the observations with visir and alma , we exclude that the ir emission arises from a point - like source . the dust producing the infrared emission observed by space telescopes must be searched for within the angular scales of the ionised gas ( @xmath3015 arcsec ) . the dust is distributed in an optically thin configuration over the radio nebula , but not uniformly , as shown in the extinction map . the visir observations did not reach the required sensitivity to detect such extended thermal emission . with the alma observations , we obtained better sensitivity to study thermal emission , but still the nebula was not detected . we estimate that the thermal emission could be detected by deeper alma observations in the future , including 7 m antennas to enhance sensitivity on larger angular scales .

the nebula was detected only in the radio , up to 17 ghz . the radio nebula resembles a spherical shell , as in the optical . we applied the code to model the electron density distribution in the s61 nebula . we found that different distributions fit the data , but all of them converge to the same ionised mass , , which is an order of magnitude smaller than previous estimates . the nebula was probably formed through stellar winds , rather than eruptions . from the alma and visir non - detections , plus the derived extinction map , we deduce that the infrared emission observed by space telescopes must arise from extended , diffuse dust within the ionised region .

we present new observations of the nebula around the magellanic candidate luminous blue variable s61 . these comprise high - resolution data acquired with the australia telescope compact array ( atca ) , the atacama large millimetre / submillimetre array ( alma ) , and visir at the very large telescope ( vlt ) . the nebula was detected only in the radio , up to 17 ghz . the 17 ghz atca map , with 0.8 arcsec resolution , allowed a morphological comparison with the _ hubble space telescope _ image . the radio nebula resembles a spherical shell , as in the optical . the spectral index map indicates that the radio emission is due to free - free transitions in the ionised , optically thin gas , but there are hints of inhomogeneities . we present our new public code rhocube to model 3d density distributions , and determine via bayesian inference the nebula s geometric parameters . we applied the code to model the electron density distribution in the s61 nebula . we found that different distributions fit the data , but all of them converge to the same ionised mass , , which is an order of magnitude smaller than previous estimates . we show how the nebula models can be used to derive the mass - loss history with high - temporal resolution . the nebula was probably formed through stellar winds , rather than eruptions . from the alma and visir non - detections , plus the derived extinction map , we deduce that the infrared emission observed by space telescopes must arise from extended , diffuse dust within the ionised region . [ firstpage ] stars : individual ( sk - 67 ) stars : mass - loss stars : circumstellar matter stars : massive radio continuum : stars methods : statistical

astro-ph9909298

i

within the standard picture for galaxy formation , in which galaxies are thought to form out of the gas which dissipatively collapses within the potential wells provided by virialized dark halos ( white & rees 1978 ) , disk galaxies take a special place . their flatness and rotational support suggest a relatively smooth formation history without violent non - linear processes . the structure and dynamics of disk galaxies are thus expected to be strongly related to the properties of the dark halos in which they are embedded . this principle is in fact the key idea behind the standard model for disk formation which was set out by fall & efstathiou ( 1980 ) , and which has since been extended upon by numerous investigators ( faber 1982 ; fall 1983 ; van der kruit 1987 ; dalcanton , spergel & summers 1997 ; mo , mao & white 1998 ; van den bosch 1998 ) . in this picture , the gas radiates its binding energy but retains its angular momentum , thus settling into a rotationally supported disk , the scale length of which is proportional to both the size and the angular momentum of the dark halo . understanding galaxy formation is intimately linked to understanding the origin of the fundamental scaling relations of galaxies . in this paper we use new semi - analytical models for the formation of disk galaxies to investigate the origin of the tully - fisher ( hereafter tf ) relation ( tully & fisher 1977 ) , which has the form @xmath2 or , in terms of a linear law in a log - log plane : @xmath3 where @xmath4 is the absolute magnitude of the disk , and @xmath5 is called the slope of the relation . in the light of the strong coupling between disk and dark halo mentioned above , a scaling relation like this is central to theories of galaxy formation . understanding its origin can put constraints on theories of galaxy formation , and in particular on the detailed physics that cause the baryonic component of a dark halo to be transformed into a luminous galaxy . as such , any successful theory of ( disk ) galaxy formation should be able to explain the slope , zero - point , and small amount of scatter of this fundamental scaling relation . numerous studies in the past have addressed the origin of the tf relation . however , no consensus has been reached , and it is currently still under debate whether the origin of the tf relation is mainly governed by initial cosmological conditions ( e.g. , eisenstein & loeb 1996 ; avila - reese , firmani & hernndez 1998 ; firmani & avila - reese 1998a , b ) , or by the detailed processes governing star formation ( silk 1997 ; heavens & jiminez 1999 ) and/or feedback ( e.g. , kauffmann , white & guiderdoni 1993 ; cole et al . 1994 ; elizondo et al . 1999 ; natarajan 1999 ) . one of the main reasons for this lack of consistency is the discordant use of luminosity and rotation measures in tf relations . the slope and scatter of the tf relation depend strongly on the photometric band used to measure the luminosities : values of @xmath6 increase from @xmath7 in the @xmath8-band to @xmath9 in the infrared ( e.g. , aaronson , huchra & mould 1979 ; visvanathan 1981 ; tully , mould & aaronson 1982 ; wyse 1982 ; pierce & tully 1988 ; gavazzi 1993 ; verheijen 1997 ) . for the rotation velocity a large number of different measures have been used : hi line widths from single - dish observations , velocities measured from h@xmath10 rotation curves , rotation velocities at a fixed number of disk scale lengths , maximum observed rotation velocities , rotation velocities at the last measured point , and the velocity of the flat part of the rotation curve . all these different measures yield wildly different tf relations ( see courteau 1997 and verheijen 1997 for detailed discussions ) . therefore it is crucial that one extracts the same quantities from the models as the ones used in the empirical relation whose origin one seeks to explain . most previous studies have ignored the subtleties associated with these different measures , and allowed themselves the `` freedom '' to compare their model predictions with the best fitting tf relation available in the literature . the other reason why we still lack consensus on the origin of the tf relation regards the modeling techniques that have been used . several investigators have used numerical simulations ( e.g. , evrard , summers & davis 1994 ; navarro & white 1994 ; steinmetz & mller 1994 ; tissera , lambas & abadi 1997 ; domnguez - tenreiro , tissera & siz 1998 ; steinmetz & navarro 1999 ; elizondo et al . a general problem , however , is that the disks that form in these simulations have specific angular momenta that are more than an order of magnitude lower than those of observed disks . in addition , most of these simulations ignore star formation and feedback , making a direct comparison between simulated and observed disks a treacherous enterprise ( see evrard 1997 for a review ) . semi - analytical modeling ( sam ) of galaxy formation has been more successful ( white & frenk 1991 ; cole 1991 ; kauffmann , white & guiderdoni 1993 ; lacey et al . 1993 ; cole et al . 1994 ; heyl et al . 1995 ; somerville & primack 1998 ) . despite reasonable success in reproducing certain tf relations , these models are hampered by several shortcomings ( see discussion in [ sec : model ] ) . more recently , avila - reese , firmani & hernndez ( 1998 ; hereafter afh98 ) and firmani & avila - reese ( 1998a , b ; hereafter fa98 ) presented new sams that focus only on the formation of disk galaxies . these models improve upon previous sams in several important ways . from a detailed comparison with observed tf relations , these authors conclude that the tf relation represents a fossil of the primordial density fluctuation field . in this paper we present new semi - analytical models for the formation of disk galaxies . we improve upon previous studies by ( i ) making a well - motivated choice for the luminosity and velocity measures that define the tf relation whose origin we seek to understand , ( ii ) by extracting the same measures from the models , and ( iii ) by using new semi - analytical models that focus on the formation of disk galaxies and that alleviate several shortcomings of previous models . the main aim of this paper is to investigate the prevailing mechanism that is responsible for the slope , zero - point and small amount of scatter of the tf relation . although our models are , in many ways , complementary to those of fa98 , we reach a different conclusion .

we present new semi - analytical models for the formation of disk galaxies with the purpose of investigating the origin of the near - infrared tully - fisher ( tf ) relation . we emphasize the importance of extracting the proper luminosity and velocity measures from the models , something that has often been ignored in the past . finally we show that our models provide a natural explanation for the small amount of scatter that makes the tf relation useful as a cosmological distance indicator .

we present new semi - analytical models for the formation of disk galaxies with the purpose of investigating the origin of the near - infrared tully - fisher ( tf ) relation . the models assume that disks are formed by cooling of the baryons inside dark halos with realistic density profiles , and that the baryons conserve their specific angular momentum . adiabatic contraction of the dark halo is taken into account , as well as a recipe for bulge formation based on a self - regulating mechanism that ensures disks to be stable . only gas with densities above the critical density given by toomre s stability criterion is considered eligible for star formation . a schmidt law is assumed to prescribe the rate at which this gas is transformed into stars . the introduction of the star formation threshold density proves an essential ingredient of our models , and yields gas mass fractions that are in excellent agreement with observations . finally , a simple recipe for supernovae feedback is included . we emphasize the importance of extracting the proper luminosity and velocity measures from the models , something that has often been ignored in the past . we use the zero - point of the-band tf relation to place stringent constraints on cosmological parameters . in particular , we rule out a standard cold dark matter universe , in which disk galaxies are too faint to be consistent with observations . the tf zero - point , in combination with nucleosynthesis constraints on the baryon density , and with constraints on the normalization of the power spectrum , requires a matter density . the observed-band tf relation has a slope that is steeper than simple predictions based on dynamical arguments suggest . taking the stability related star formation threshold densities into account steepens the tf relation and decreases its scatter . however , in order for the slope to be as steep as observed , further physics are required . we argue that the characteristics of the observed near - infrared tf relation do not reflect systematic variations in stellar populations , or cosmological initial conditions . in fact , feedback seems an essential ingredient in order to explain the observed slope of the-band tf relation . finally we show that our models provide a natural explanation for the small amount of scatter that makes the tf relation useful as a cosmological distance indicator .

astro-ph9909298

i

we have used new semi - analytical models to investigate the origin of the near - infrared tf relation . by focusing on the near - infrared we are less susceptible to uncertainties related to extinction by dust and stellar populations . as empirical constraint we have used the @xmath0-band tf relation of v97 , which is the only near - infrared tf relation based on ccd imaging ( rather than aperture photometry ) combined with full hi rotation curves ( rather than hi line widths ) . an important improvement over previous studies aimed at understanding the tf relation concerns our special care in extracting the same luminosity and rotation measures from our model , as the ones used in the empirical tf relation . since dark halos have density profiles that at large radii fall off more rapidly than an isothermal ( i.e. , the nfw profiles used here fall off as @xmath165 ) , the circular velocities of the disk at the flat part of the rotation curve are generally larger than the circular velocities at the virial radius . this has important consequences for the tf zero - point , as it allows us to put constraints on cosmological parameters . in a scdm universe , dark halos are too centrally concentrated , and the baryon fraction is too low , to be able to fit the @xmath0-band tf relation . the observed zero - point puts an upper limit on the normalization of the power spectrum of @xmath268 , in clear contradiction with constraints from cobe and the abundances of rich clusters of galaxies . we have shown that the observed tf zero - point ( combined with the nucleosynthesis constraints on the baryon density ) favor a universe with @xmath269 . previous studies of the tf relation based on sams were only able to rule against scdm , _ if the models were normalized to fit the luminosity function of galaxies_. the reason why we can rule out scdm without this additional constraint is our more sophisticated treatment of the relationship between @xmath171 and @xmath118 . whereas simple dynamics predict a tf relation with a slope of @xmath125 , the @xmath0-band tf relation reveals a slope of @xmath15 . this implies that the physics regulating star formation and feedback , coupled with the mass dependence of halo densities and stellar populations , has to tilt the tf relation to its observed slope . the introduction of a stability - related star formation threshold density increases the slope of the tf relation , reduces its scatter , and yields gas mass fractions that are in excellent agreement with observations . setting toomre s @xmath51 parameter to its empirically determined value tilts the tf relation to a slope of @xmath270 , and additional physics are thus required to obtain a tf relation as steep as observed . we have presented four different physical mechanisms that all yield good fits to the observed @xmath0-band tf relation : lowering @xmath51 to @xmath271 , systematic variations in the @xmath0-band mass - to - light ratios , a power spectrum of initial density fluctuations that differs from that for cdm , and feedback . except for the latter , each of these results in inconsistencies with other observations , and can be ruled out as the prevailing mechanism for the observed characteristics of the near - infrared tf relation . the feedback efficiency required to tilt the tf relation to its observed slope has to be higher in lower mass systems , and depends strongly on @xmath145 and @xmath39 ( if the baryon fraction is taken to be constrained by the nucleosynthesis results ) . in particular , in a flat universe with @xmath272 and @xmath161 ( a cosmological model that yields low density halos in agreement with observed rotation curves ) , @xmath226 percent of the available sn energy is required to fit the observed tf zero - point . although in principle unphysical , we do not wish to draw too strong conclusions from our oversimplified model for feedback . we do emphasize , however , that some amount of feedback is required to yield tf relations with a slope as steep as observed . more sophisticated models of sn feedback , such as in the work of mac low & ferrara ( 1999 ) , are required to investigate whether our inferred feedback efficiencies are realistic . our best fitting models with feedback reveal a small amount of curvature in the tf relation , in qualitative agreement with recent observations of dwarf galaxies ( mathews , van driel & gallagher 1998 ; stil & israel 1998 ) . further constraints on the models come from the gas mass fractions , which require that virtually all the gas mass eligible for star formation ( i.e. , with densities above the critical density ) is transformed into stars over the lifetime of the galaxy . for the fiducial value of @xmath51 the scatter in halo spin parameters affects both the luminosities and the observed rotation velocities , but to such a degree that galaxies are scattered along the tf relation , rather than perpendicular to it . the introduction of star formation threshold densities thus yield a natural explanation for the small amount of scatter observed . taking account of a realistic spread in halo concentrations , consistent with what one expects from the spread in mass aggregation histories , our model that best fits the zero - point and slope of the @xmath0-band tf relation predicts a scatter of @xmath273 mag only , in excellent agreement with observations . in a follow - up paper ( van den bosch & dalcanton 1999 ) we compare the models presented here to numerous other observational constraints . this work has benefited greatly from discussions with marcella longhetti , claudia maraston , jeroen stil , and marc verheijen . i am indebted to julianne dalcanton and george lake for advice and critical assessments of an earlier draft of the paper , to stphane charlot for providing results from his stellar population models , and to the anonymous referee for his comments that helped improve the presentation of the paper . support for this work was provided by nasa through hubble fellowship grant # hf-01102.11 - 97.a awarded by the space telescope science institute , which is operated by aura for nasa under contract nas 5 - 26555 . lcccllccccc scdm & @xmath274 & @xmath275 & @xmath276 & @xmath245 & @xmath277 & @xmath278 & y & n & n & n + @xmath174cdm3 & @xmath279 & @xmath280 & @xmath280 & @xmath274 & @xmath281 & @xmath282 & y & y & n & y + @xmath174cdm2 & @xmath283 & @xmath284 & @xmath276 & @xmath285 & @xmath286 & @xmath287 & n & y & y & y + llrcclcrcrc s1 & scdm & @xmath170 & @xmath279 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath290 & @xmath291 + l0 & @xmath174cdm3 & @xmath292 & @xmath249 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath125 & @xmath293 + l1 & @xmath174cdm3 & @xmath170 & @xmath249 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath294 & @xmath295 + l2 & @xmath174cdm3 & @xmath284 & @xmath249 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath296 & @xmath297 + l3 & @xmath174cdm3 & @xmath170 & @xmath298 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath15 & @xmath251 + l4 & @xmath174cdm3 & @xmath170 & @xmath249 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath299 & @xmath15 & @xmath300 + l5 & @xmath174cdm3 & @xmath170 & @xmath249 & @xmath288 & @xmath286 & @xmath277 & @xmath301 & @xmath289 & @xmath15 & @xmath302 + l6 & @xmath174cdm2 & @xmath170 & @xmath249 & @xmath288 & @xmath286 & @xmath275 & @xmath275 & @xmath289 & @xmath303 & @xmath304 + l7 & @xmath174cdm2 & @xmath170 & @xmath249 & @xmath288 & @xmath286 & @xmath305 & @xmath275 & @xmath289 & @xmath306 & @xmath307 + l8 & @xmath174cdm2 & @xmath170 & @xmath249 & @xmath288 & @xmath308 & @xmath277 & @xmath309 & @xmath289 & @xmath310 & @xmath307 +

the introduction of the star formation threshold density proves an essential ingredient of our models , and yields gas mass fractions that are in excellent agreement with observations . we use the zero - point of the-band tf relation to place stringent constraints on cosmological parameters . the tf zero - point , in combination with nucleosynthesis constraints on the baryon density , and with constraints on the normalization of the power spectrum , requires a matter density . the observed-band tf relation has a slope that is steeper than simple predictions based on dynamical arguments suggest . taking the stability related star formation threshold densities into account steepens the tf relation and decreases its scatter . in fact , feedback seems an essential ingredient in order to explain the observed slope of the-band tf relation .

we present new semi - analytical models for the formation of disk galaxies with the purpose of investigating the origin of the near - infrared tully - fisher ( tf ) relation . the models assume that disks are formed by cooling of the baryons inside dark halos with realistic density profiles , and that the baryons conserve their specific angular momentum . adiabatic contraction of the dark halo is taken into account , as well as a recipe for bulge formation based on a self - regulating mechanism that ensures disks to be stable . only gas with densities above the critical density given by toomre s stability criterion is considered eligible for star formation . a schmidt law is assumed to prescribe the rate at which this gas is transformed into stars . the introduction of the star formation threshold density proves an essential ingredient of our models , and yields gas mass fractions that are in excellent agreement with observations . finally , a simple recipe for supernovae feedback is included . we emphasize the importance of extracting the proper luminosity and velocity measures from the models , something that has often been ignored in the past . we use the zero - point of the-band tf relation to place stringent constraints on cosmological parameters . in particular , we rule out a standard cold dark matter universe , in which disk galaxies are too faint to be consistent with observations . the tf zero - point , in combination with nucleosynthesis constraints on the baryon density , and with constraints on the normalization of the power spectrum , requires a matter density . the observed-band tf relation has a slope that is steeper than simple predictions based on dynamical arguments suggest . taking the stability related star formation threshold densities into account steepens the tf relation and decreases its scatter . however , in order for the slope to be as steep as observed , further physics are required . we argue that the characteristics of the observed near - infrared tf relation do not reflect systematic variations in stellar populations , or cosmological initial conditions . in fact , feedback seems an essential ingredient in order to explain the observed slope of the-band tf relation . finally we show that our models provide a natural explanation for the small amount of scatter that makes the tf relation useful as a cosmological distance indicator .

astro-ph0611092

i

the launch of the _ swift _ satellite @xcite has revolutionized the study of the interstellar medium of gamma - ray burst ( grb ) host galaxies and fostered new studies on the intervening intergalactic medium @xcite . while previous missions alerted the community to the existence of very bright grb afterglows ( e.g. * ? ? ? * ) , only a few events were identified in time to allow echelle observations @xcite . the rapid localization by _ swift _ of the hard x - ray emission to several arcminutes @xcite and the soft x - ray component to a few arcseconds @xcite has enabled observers to obtain high signal - to - noise , high - resolution spectroscopy of grbs with 10m - class telescopes . in this paper , we present a uniform dataset of echelle spectra of four grb afterglows by our grb afterglows as probes ( graasp ) collaboration . among other applications , the data presented here can be used to derive constraints on the physical conditions of the ism of grb host galaxies : the metallicity , dust - to - gas ratio , chemical abundance patterns , ionization state , the distance of the grb to the absorbing gas , etc . future papers in this series will examine these properties in greater detail and provide comparison with analogous observations along quasar sightlines . in addition to the data presented here , we have collected several low - resolution spectra of these and other grbs and will present that data in a companion paper @xcite along with ism abundance measurements obtained by different analysis techniques more suitable to lower resolution data @xcite . we also present measurements of the strong systems in the grb afterglow spectra . in @xmath0 [ sec : obs ] we summarize the observations and data reduction procedures . the methodology is briefly described in @xmath0 [ sec : colm ] and the line - profiles and column density values are given in @xmath0 [ sec : grb ] and @xmath0 [ sec : mgii ] .

we present optical echelle spectra of four gamma - ray burst ( grb ) afterglows ( grb 050730 , grb 050820 , grb 051111 , and grb 060418 ) discovered during the first 1.5 years of operation of the _ swift _ satellite and localized by either the _ swift _ telescope or follow - up ground - based imaging . we analyze the spectra to derive accurate column density measurements for the transitions arising in the interstellar medium ( ism ) of the grb host galaxies . these measurements can be used to constrain the physical properties of the ism including the metallicity , dust - to - gas ratio , ionization state , and chemical abundances of the gas . we also present measurements of the strong systems in the grb afterglow spectra . with the publication of this paper , we provide the first data release of echelle afterglow spectra by the graasp collaboration to the general community .

we present optical echelle spectra of four gamma - ray burst ( grb ) afterglows ( grb 050730 , grb 050820 , grb 051111 , and grb 060418 ) discovered during the first 1.5 years of operation of the _ swift _ satellite and localized by either the _ swift _ telescope or follow - up ground - based imaging . we analyze the spectra to derive accurate column density measurements for the transitions arising in the interstellar medium ( ism ) of the grb host galaxies . these measurements can be used to constrain the physical properties of the ism including the metallicity , dust - to - gas ratio , ionization state , and chemical abundances of the gas . we also present measurements of the strong systems in the grb afterglow spectra . with the publication of this paper , we provide the first data release of echelle afterglow spectra by the graasp collaboration to the general community .

1107.2121

c

our spectroscopic fitting clearly indicates that the compt and bb+bb models yield superior fits to the other possibilities . the compt model , with its power - law shape curtailed by an exponential turnover , is intended to mock - up the classic unsaturated comptonization spectrum realized in models of accretion disks such as in cyg x-1 or in active galactic nuclei ( see chapter 7 of rybicki & lightman 1979 for a summary of its development as a solution of the kompaneets equation ) . these models use hot , thermal electrons in a corona to repeatedly scatter low - energy photons , heating them gradually up to an energy @xmath158 consistent with the electron temperature @xmath159 , at which point a quasi - exponential spectral turnover emerges as further heating becomes impossible . the power law marks the scale - independence of the compton upscattering , and its slope depends only on the mean energy gain per collision ( @xmath160 for non - relativistic electrons ) and the probability of loss of photons from the scattering zone , being expressed as @xmath161 with the compton @xmath162-parameter in the domain @xmath163 . this parameter is the product of the average fractional energy change per scattering and the mean number of scatterings , and is @xmath164 for a non - relativistic situation for repeated compton upscattering . here @xmath165 is the scattering optical depth . in the context of magnetars , such comptonization can also ensue , but the strong magnetic field now plays an important role . ephemeral coronae of hot electrons could be expected in a dynamic inner magnetosphere . for example , this could be due to intense dissipation of magnetic energy in the closed field line region via field line twisting , i.e. , transient departures from poloidal geometry , as in the considerations of thomson , lyutikov & kulkarni ( 2002 ) , thompson & beloborodov ( 2005 ) , beloborodov & thompson ( 2007 ) , and nobili , turolla & zane ( 2008 ) . such a pumping of energy into electrons in low altitude regions would then be subject to irradiation by the intense bath of surface x - ray emission . the electron coronae would mimic those of their black hole counterparts discussed above , and serve as a comptonization target for the x - rays . temporally , the coronae could be quite variable , resulting in varying or chaotic time profiles such as are observed . if the deposition of energy in hot electrons is persistent over many light - crossing timescales , so also can the upscattered hard x - ray emission be . in this magnetic case , the spectrum would again naturally be a power law truncated by an exponential tail , whose energy is pinned by @xmath159 . the slope does differ somewhat in value from the non - magnetic case because the presence of the intense field anisotropizes the environment . this alters the average fractional energy change per scattering from the isotropic @xmath166 form in rybicki & lightman ( 1979 ) . we note that this picture is similar to that of lyutikov & gavriil ( 2006 ; see also rea et al . 2008 ) , who modeled the steep x - ray tails below 10 kev in quiescent magnetar emission using a resonant cyclotron scattering picture . since the electrons will move along the field lines in the zeroth landau level , the exact kinematics that impact the determination of @xmath167 depend on the colatitude and altitude of the collisions therefore , the effective _ magnetic compton @xmath162-parameter _ , @xmath168 , that could be substituted in eq . ( [ comp_index ] ) , would take a somewhat different value from its non - magnetic cousin , but one anticipates that the range of spectral realizations would be similar to the @xmath169 case . since head - on collisions generally yield greater heating of photons , and these are more likely at lower altitudes , it is expected that interaction zones close to the stellar surface would spawn larger @xmath168 and therefore flatter photon spectra . if there were a coronal radius expansion , this would then translate to an evolutionary steepening of the spectra with time . note that the photon retention probabilities in the coronae will need to be higher than in the lyutikov & gavriil ( 2006 ) scenario to generate the requisite moderate optical depths @xmath165 to match the flat outburst spectra discussed in this paper . the scatterings will take place below the cyclotron resonance unless their altitude is above around 10 stellar radii ( e.g. lyutikov & gavriil 2006 ) . below resonance , the cross section is far inferior to the thomson value ( e.g. see herold 1979 ) . then , the coronal electron density must accordingly be much higher than when @xmath169 in order to establish a sizable optical depth . this provides a possible distinction between the two classes of magnetar emission : the steady @xmath170kev signal may originate at higher altitudes where the resonance is accessed and the local electron density is low , whereas the bursts reported here may be triggered closer to the surface in higher density zones that initially precipitate scattering below the cyclotron resonance . as photons get energized above 20 kev , the altitude where the resonance is accessed is lowered , increasing the cross section for scattering and therefore @xmath168 . note also , that to a considerable extent , polarization mode - switching ( e.g. miller 1995 ) between the two photon polarizations can help increase the opacity , adding a further nuance to consider when modeling the scattering environment . observationally , it is difficult to discern unambiguously the presence of the intense field using a comptonization model scenario in the energy range of data presented in this paper : details of both the field strength and the geometry are subsumed in a single parameter @xmath168 . since the emission in this scenario should be strongly polarized , and the degree of polarization should depend on the interaction geometry , a hard x - ray polarimeter would provide insightful probes into the presence of a strongly anisotropizing super - critical field . the dual blackbody can also be envisaged as a viable alternative from a theoretical standpoint . the moderate thomson depths required to generate flat comptonization spectra could easily be higher , thereby pushing the electron - photon interactions more towards an equilibration . the saturation temperature is then controlled by the largely uncertain total energy dissipated in the inner magnetosphere per hot electron present . since any equilibration will be non - uniform over a coronal volume , there should be a modest temperature gradient throughout , smearing out the continuum . the range of temperatures will not be great because the system is not gravitationally hydrostatic in character . moreover , radiative transfer effects impact the spectral shape and further modify it . accordingly , pure blackbody shapes are not expected . it is quite conceivable that a two - component blackbody fit may well approximate the emergent continuum that is a superposition of distorted blackbodies spanning a small range of temperatures . most probably , due to general thermodynamic considerations , the base of the coronae ( e.g. centered near the source of magnetic dissipation ) should be hotter than the outer layers ( see also lyubarsky 2002 and thompson & duncan 1995 ) . interestingly , the fits here generate a smaller volume associated with the hotter blackbody contribution , consistent with the expectation for coronal structure . yet , somehow , we must obtain a view of the hotter zone , therefore indicating a strongly aspherical coronal geometry . modeling this semi - equilibration is a challenging task for theorists considering the influences of the field , the twisted magnetospheric geometry , and the inherent anisotropy and polarization - dependence of the scattering process . at present , it is not possible to distinguish between this thermal scenario , and a comptonization one . since the compt model fits all events , we used it to derive fluences ( time - integrated and time - resolved ) , and determined using the time - resolved values that the hardness - fluence correlation can be described by a broken power law , with a minimum at a fluence of @xmath171 erg @xmath6 . we have used @xmath172 values to characterize the hardness of the events in our sample . earlier studies of this correlation used hardness ratios to bypass proper spectral fits due to low quality or insufficient data . moreover , these studies @xcite had a very narrow overlap in fluence space . @xcite used 95 sgr@xmath9 events detected with the interplanetary sun earth explorer-3 ( isee-3 ) with fluences ranging between @xmath173 erg @xmath6 , and found a slightly positive correlation between the two quantities ( hardness increasing with fluence ) . on the other hand , @xcite analyzed 159 and 385 events from sgrs @xmath9 and @xmath79 , respectively , observed with _ rxte _ , with fluences ranging between @xmath174 erg @xmath6 , and found the opposite trend ( an anticorrelation between hardness ratios and fluences ) . since the gbm data sample covers both fluence ranges ( @xmath175 erg @xmath6 ) we were able to establish that both trends are indeed correct , and to define the turning point in the @xmath172 - fluence diagram . only two magnetar candidates were observed with gbm to emit a multitude of bursts , thus allowing us to construct their @xmath172 - flux diagrams . we find for the time - resolved data of , that @xmath172 reaches a minimum of @xmath4kev at a flux value of @xmath176 erg @xmath6 s@xmath7 . the second source ( sgrj@xmath8 ; van der horst et al . 2011 in preparation ) has flux values of @xmath177 erg @xmath6 s@xmath7 at a similar @xmath172 minimum . in addition we used the hardness ratio - count rate relationship of sgr@xmath9 found in the _ integral _ data @xcite to derive an approximate flux value of @xmath178 erg @xmath6 s@xmath7 ( using their conversion from count rate to flux ) at the hardness ratio minimum ( figure 3 in * ? ? ? * ) . we then converted these values into isotropic source luminosity , @xmath10 ( see table [ liso ] ) ; we note that ( with the caveat of small number statistics , and uncertainties in the distance measurements ) these values are comparable [ @xmath179 erg s@xmath7 ] , although the @xmath181fields and fluxes at @xmath172 minima vary by approximately a factor of ten between the two gbm sources and sgr@xmath9 . whether these differences reflect intrinsic source properties or instrumental effects is not yet clear . the @xmath172 trend is , however , clearly established in at least two different instrument data sets . the physical interpretation of this trend is beyond the scope of this paper . part of our sample ( 18 bursts ) was fit with the bb+bb spectral function applied to time - integrated intervals ; the remaining 11 events could not be fit due to poor statistics . in addition for five bright events we performed time - resolved spectral analysis . we found that the temperatures , emission areas and fluences ( luminosities ) of the two thermal components exhibit very similar properties in both cases ( time - integrated and -resolved ) . our results are consistent with those presented for intermediate and short bursts from sgr @xmath79 by @xcite and @xcite ; the burst emission areas and temperatures fall into the same region with those of the short bursts of sgr @xmath79 shown in figure 5 of @xcite . we also see a similar trend in the correlation between the luminosities of the two components as described by @xcite , namely that both luminosities increase in tandem . this behavior indicates that the hot and cool bb components may come from two separate emission regions , as also pointed out by @xcite : a smaller but hotter one from the surface of the magnetar , and a larger , cooler one from the star s magnetosphere ( but see also the discussion in section 5.1 ) . we note here that the time - integrated hot bb emission area of 13 bursts ( i.e. , not including the 5 brightest events ; figure [ r2_kt ] ) is similar to the emission area ( @xmath182 km@xmath183 ) of the bb component found in the persistent emission of sgr j1550@xmath155418 during one active bursting episode in january 2009 , believed to originate from a hot spot on the neutron star surface @xcite . finally , we performed a detailed temporal analysis of all 29 bursts and estimated their durations ( @xmath0/@xmath30 ) for the first time in count _ and _ in photon space . both estimates agree within statistics and thus validated all earlier ( count space ) duration estimates of magnetar candidate bursts . the durations ( and for four sgrs also the emission times and duty cycles ) of five magnetar candidates follow a log - normal or normal distribution . we find that events are very similar in average duration with four more magnetar candidates ( three sgrs and one axp ) . however , the @xmath0 distribution of the bursts from axp1e@xmath81 has a factor of two larger dispersion ( @xmath184 , * ? ? ? * ) than those of at least three other sgrs ( @xmath185 , * ? ? ? * , present work ) . bursts have an average duty cycle ( @xmath186 ) larger than other sgrs ( 0.45 , 0.46 for sgrs @xmath79 , @xmath9 , respectively ; * ? ? ? the differences in the intrinsic properties of the sources may be due to differences in the size of the active region responsible for the burst emission . we discuss below the burst energetics and its evolution . similarly to other magnetars , we also do not find a correlation between the pulse phase of the source and the burst peak times . the overall behavior of during its 13-day active period ( 2008 august 22 to september 3 ) is also very interesting . although about half of the bursts ( 16/29 ) occurred on one day ( 2008 august 23 ) , most of the burst energy was emitted from the source during august @xmath187 ( see also figure [ energetic ] ) . the average fluence of the beginning and the later part of the active period was constant and at a lower level . during these three days , gbm detected seven bursts , five of which are also distinguished in that the emission areas from both bb components are the largest . further , one of these five events ( bn080824.054 ) shows a double - peaked structure reminiscent of some bright thermonuclear type i x - ray bursts from accreting neutron stars , which has been interpreted as due to photospheric radius expansion ( pre ) . @xcite studied the effects of pre in high magnetic fields using this event and find that the predicted flux from pre theory is consistent with the one observed , opening the way to determining fundamental parameters of neutron stars , such as their equation of state . this publication is part of the gbm / magnetar key project ( nasa grant nnh07zda001-glast , pi : c. kouveliotou ) . m.g.b . acknowledges support from nasa through grant nnx10ac59a . e.g. and y.k . acknowledge the support from the scientific and technological research council of turkey ( tbitak ) through grant 109t755 . a.v.k . was supported by the bundesministeriums fr wirtschaft und technologie ( bmwi ) through dlr grant 50 og 1101 . alw acknowledges support from a netherlands organisation for scientific research ( nwo ) vidi grant . ramjw acknowledges support from the european research council via advanced investigator grant no . 247295 . 999 aptekar r. l. et al . , 2009 , , 698 , l82 barthelmy s. d. , et al , 2008 , gcn circ . , 8113 , 1 beloborodov , a. m. , & thompson , c. 2007 , , 657 , 967 bibby j. l. et al . , 2008 , , 386 , l23 bissaldi e. et al . , 2011 , arxiv11013325 camilo f. et al . , 2007 , , 666 , l93 enoto t. et al . , 2009 , , 693 , l122 evans p. a. et al . , 2009 , , 397 , 1177 fenimore e. e. , laros j. g. , ulmer a. , 1994 , , 432 , 742 feroci m. et al . , 2004 , , 612 , 408 gavriil f. p. , kaspi v. m. , woods p. m. , 2004 , , 607 , 959 g e. et al . , 1999 , , 526 , l93 g e. et al . , 2000 , , 532 , l121 g e. et al , 2001 , , 558 , 228 g e. , woods p. , & kouveliotou c. 2008 , grb coordinates network # 8118 g e. et al . , 2010 , , 722 , 899 gtz d. et al . , 2004 , a&a , 417 , l45 herold h. 1979 , phys . d , 19 , 2868 holland s. t. , sato g. , 2008 , gcn rep . , 160 , 1 israel g. l. et al . , 2008 , , 685 , 1114 kaneko y. et al . , 2006 , , 166 , 298 kaneko y. et al . , 2010 , , 710 , 1335 koshut t. m. , 1996 , ph . d. thesis , an investigation of temporal characteristics of gamma - ray bursts observed with batse koshut t. m. et al . , 1996 , , 463 , 570 kouveliotou c. et al . , 1993 , , 413 , l101 kumar h. s. 2010 , , 716 , 97 laros j. g. et al . , 1986 , , 322 , 152 lyubarsky , y. e. 2002 , , 332 , 199 lyutikov , m. & gavriil , f. p. 2006 , mnras , 368 , 690 meegan c. et al . , 2009 , , 702 , 791 mereghetti s. et al . , 2005 , , 628 , 938 mereghetti s. , 2008 , , 15 , 225 miller m. c. , 1995 , apj , 448 , l29 mitrofanov i. g. et al . , 1999 , , 522 , 1069 nobili l. , turolla r. , & zane , s. , 2008 , , 386 , 1527 olive j .- f . et al . , 2004 , , 616 , 1148 paciesas w. et al . , 2011 , in preparation rea n. et al . , 2008 , apj , 686 , 1245 rea n. et al . , 2009 , , 396 , 2419 rots , a. h. , jahoda , k. , & lyne , a.g , 2004 , , 605 , l129 rybicki g. b. & lightman a. p. 1979 _ radiative processes in astrophysics ( wiley , new york ) scargle , j. , 1998 , , 504 , 405 scargle , j. et al . , 2011 , in preparation starling r. l. c. et al , 2008 , , 672 , 433 strohmayer t. e. & ibrahim a. , 1998 , aipc , 428 , 947 thompson , c. , & duncan , r. c. 1995 , , 275 , 255 thompson c. & duncan r. c. , 2001 , , 561 , 980 thompson c. , lyutikov m. & kulkarni s. r. , 2002 , apj , 574 , 332 thompson c. & beloborodov a. m. , 2005 , apj , 634 , 565 tiengo a. et al . , 2010 , , 710 , 227 van der horst a. j. et al . , 2011 , in preparation watts a. l. et al . , 2010 , , 719 , 190 woods p. m. , thompson c. , 2006 , in compact stellar x - ray sources , eds . lewin @xmath188 m. van der klis , cambridge astrophysics series , 39 , p.547 woods p. m. et al . , 2008 , atel , 1691 xu y. et al . , 2006 , science , 311 , 54 _ clllccccccc 1 & bn080822.529 & 12:41:56.914 & 8 , 7 , 4 & @xmath189 & @xmath190 & @xmath191 & @xmath192 & @xmath193 & @xmath194 & @xmath195 + 2 & bn080822.647 & 15:36:35.200 & 9 , 10 & @xmath196 & @xmath197 & @xmath198 & @xmath199 & @xmath200 & @xmath201 & @xmath202 + 3 & bn080822.981 & 23:32:57.746 & 2 & @xmath203 & @xmath204 & @xmath205 & @xmath206 & @xmath207 & @xmath208 & @xmath209 + 4 & bn080823.020 & 00:28:09.904 & 3 , 4 & @xmath210 & @xmath211 & @xmath212 & @xmath213 & @xmath214 & @xmath215 & @xmath216 + 5 & bn080823.091 & 02:11:36.630 & 10 , 11 & @xmath217 & @xmath218 & @xmath219 & @xmath220 & @xmath221 & @xmath222 & @xmath223 + 6 & bn080823.174 & 04:10:19.280 & 0 , 1 & @xmath224 & @xmath225 & @xmath226 & @xmath227 & @xmath228 & @xmath229 & @xmath230 + 7 & bn080823.248 & 05:56:31.529 & 2 & @xmath231 & @xmath232 & @xmath233 & @xmath234 & @xmath235 & @xmath236 & @xmath237 + 8 & bn080823.293 & 07:01:09.967 & 3 , 0 , 1 , 5 & @xmath238 & @xmath239 & @xmath240 & @xmath241 & @xmath242 & @xmath243 & @xmath244 + 9 & bn080823.293 & 07:04:22.610 & 3 , 0 , 1 , 5 & @xmath245 & @xmath246 & @xmath247 & @xmath248 & @xmath249 & @xmath250 & @xmath251 + 10 & bn080823.319 & 07:39:32.257 & 9 , 10 & @xmath252 & @xmath253 & @xmath254 & @xmath255 & @xmath256 & @xmath257 & @xmath258 + 11 & bn080823.330 & 07:55:45.690 & 4 , 3 , 8 , 7 & @xmath259 & @xmath260 & @xmath261 & @xmath262 & @xmath263 & @xmath264 & @xmath265 + 12 & bn080823.354 & 08:30:01.633 & 11 , 8 & @xmath266 & @xmath267 & @xmath268 & @xmath269 & @xmath270 & @xmath271 & @xmath272 + 13 & bn080823.429 & 10:18:13.891 & 0 , 1 , 3 , 5 & @xmath273 & @xmath274 & @xmath275 & @xmath276 & @xmath277 & @xmath278 & @xmath279 + 14 & bn080823.478 & 11:27:32.306 & 8 , 4 & @xmath280 & @xmath281 & @xmath282 & @xmath283 & @xmath284 & @xmath285 & @xmath286 + 15 & bn080823.623 & 14:56:23.563 & 10 , 11 & @xmath287 & @xmath288 & @xmath289 & @xmath290 & @xmath291 & @xmath292 & @xmath293 + 16 & bn080823.714 & 17:08:49.038 & 9 , 10 & @xmath294 & @xmath295 & @xmath296 & @xmath297 & @xmath298 & @xmath299 & @xmath300 + 17 & bn080823.847 & 20:19:30.659 & 9 , 10 & @xmath301 & @xmath302 & @xmath303 & @xmath304 & @xmath305 & @xmath306 & @xmath307 + 18 & bn080823.847 & 20:23:42.822 & 9 , 10 & @xmath308 & @xmath309 & @xmath310 & @xmath311 & @xmath312 & @xmath313 & @xmath314 + 19 & bn080823.986 & 23:39:24.472 & 9 , 11 , 7 , 6 & @xmath315 & @xmath316 & @xmath317 & @xmath318 & @xmath319 & @xmath320 & @xmath321 + 20 & bn080824.054 & 01:17:55.394 & 2 , 5 & @xmath322 & @xmath323 & @xmath324 & @xmath325 & @xmath326 & @xmath327 & @xmath328 + 21 & bn080824.346 & 08:18:24.418 & 3 , 4 & @xmath329 & @xmath330 & @xmath331 & @xmath332 & @xmath333 & @xmath334 & @xmath335 + 22 & bn080824.828 & 19:52:51.264 & 2 , 10 & @xmath336 & @xmath337 & @xmath338 & @xmath339 & @xmath340 & @xmath341 & @xmath342 + 23 & bn080825.200 & 04:48:27.405 & 4 & @xmath343 & @xmath344 & @xmath345 & @xmath346 & @xmath347 & @xmath348 & @xmath349 + 24 & bn080825.401 & 09:37:42.158 & 4 , 3 , 8 & @xmath350 & @xmath351 & @xmath352 & @xmath353 & @xmath354 & @xmath355 & @xmath356 + 25 & bn080826.136 & 03:16:14.773 & 8 & @xmath357 & @xmath358 & @xmath359 & @xmath360 & @xmath361 & @xmath362 & @xmath363 + 26 & bn080826.236 & 05:40:19.425 & 9 , 10 & @xmath364 & @xmath365 & @xmath366 & @xmath367 & @xmath368 & @xmath369 & @xmath370 + 27 & bn080828.875 & 20:59:39.966 & 1 , 0 , 5 , 3 & @xmath371 & @xmath372 & @xmath373 & @xmath374 & @xmath375 & @xmath376 & @xmath377 + 28 & bn080903.421 & 10:06:35.329 & 4 , 5 & @xmath378 & @xmath379 & @xmath380 & @xmath381 & @xmath382 & @xmath383 & @xmath384 + 29 & bn080903.787 & 18:53:48.775 & 2 , 10 & @xmath385 & @xmath386 & @xmath387 & @xmath388 & @xmath389 & @xmath390 & @xmath391 + lcccccccc mean & @xmath392 & @xmath393 & @xmath394 & @xmath395 & @xmath396 & @xmath397 & @xmath398 & @xmath399 + @xmath22 & @xmath400 & @xmath401 & @xmath402 & @xmath403 & @xmath404 & @xmath405 & @xmath406 & @xmath407 + weighted mean & @xmath408 & @xmath409 & @xmath410 & @xmath411 & & & & +

we also estimate for the first time event durations of soft gamma repeater ( sgr ) bursts in photon space ( i.e. , using their deconvolved spectra ) and find that these are very similar to thes estimated in count space ( following a log - normal distribution with a mean value ofms ) . we expand on the physical interpretation of these two models and we compare their parameters and discuss their evolution . we show that the time - integrated and time - resolved spectra reveal that decreases with energy flux ( and fluence ) to a minimum ofkev at erg s , increasing steadily afterwards . the isotropic luminosity , , corresponding to these flux values is roughly similar for all sources ( erg s ) .

we present our temporal and spectral analyses of 29 bursts from , detected with the gamma - ray burst monitor onboard the _ fermi _ gamma - ray space telescope during the 13 days of the source activation in 2008 ( august 22 to september 3 ) . we find that the durations of the bursts can be fit with a log - normal distribution with a mean value ofms . we also estimate for the first time event durations of soft gamma repeater ( sgr ) bursts in photon space ( i.e. , using their deconvolved spectra ) and find that these are very similar to thes estimated in count space ( following a log - normal distribution with a mean value ofms ) . we fit the time - integrated spectra for each burst and the time - resolved spectra of the five brightest bursts with several models . we find that a single power law with an exponential cutoff model fits all 29 bursts well , while 18 of the events can also be fit with two black body functions . we expand on the physical interpretation of these two models and we compare their parameters and discuss their evolution . we show that the time - integrated and time - resolved spectra reveal that decreases with energy flux ( and fluence ) to a minimum ofkev at erg s , increasing steadily afterwards . two more sources exhibit a similar trend : sgrs j and . the isotropic luminosity , , corresponding to these flux values is roughly similar for all sources ( erg s ) .

astro-ph0702107

i

the absolute magnitude of the lmc @xmath0 scuti star # 28114 in df05 catalog was determined from the @xmath1 relation of the galactic @xmath0 scuti stars ( equation 1 ) and from the theoretical modeling of the star light curves ( section 4 ) , deriving values of @xmath26=1.23@xmath137 and 1.21@xmath138 mag , respectively . combined with the star apparent magnitude @xmath15=19.94 @xmath3 0.09 mag and assuming for the reddening @xmath4=0.08 @xmath69 ( which is the average of the pulsational " reddening : @xmath4=0.070 , the color excess inferred from the star intrinsic color : @xmath4=0.076 , and the lmc foreground reddening : @xmath4=0.075 , according to schlegel et al . 1998 maps ) , the corresponding distance moduli for the lmc are : [email protected] and [email protected] mag respectively . the star was used along with 24 @xmath0 scuti candidates discovered in the ogle ii survey of the lmc , and 7 @xmath0 scuti stars discovered in the lmc by kaluzny et al . ( 2003 , 2006 ) , to define the @xmath1 relation of the lmc @xmath0 scuti stars : @xmath139 , which is tied to df05 photometric zero - point . for @[email protected] mag this relation leads to @xmath140 = 18.50 @xmath3 0.22 at the period and metallicity of # 28114 . the final distance modulus of the lmc derived from the @xmath0 scuti stars is then : [email protected] , where the error is the standard deviations of the weighted average of the three above solutions . thanks to its composite stellar population the lmc hosts different types of pulsating stars which are standard candles and that , once distortions and depth effects due to the lmc geometry are properly taken into account , can provide the opportunity of a direct and easy cross check of the predicted distances . in the following we will compare our results from the lmc @xmath0 scuti variables with the distance inferred from the other two major types of pulsating variables found in the lmc , namely : the rr lyrae stars , and the classical cepheids . c03 found that the average @xmath120 magnitude of the rr lyrae stars in two regions of the lmc containing a total number of 108 of these variables , for stars with [ fe / h ] = @xmath141.5 is : @xmath120 = [email protected] mag . individual reddenings of the rr lyrae stars in each region were found that led to this accurate @xmath120 magnitude . cacciari & clementini ( 2003 ) suggest the best @xmath141 value at [ fe / h ] = @xmath141.5 is , @xmath141 = 0.59 @xmath3 0.03 . thus the distance modulus is @xmath140 = 18.46 @xmath3 0.07 mag . soszyski et al . ( 2003 ) give @xmath142 = [email protected] mag and @xmath143 = [email protected] mag from the ogle ii sample of rr lyrae stars in the lmc . these values transformed to df05 photometric zero point give 19.42 and 19.37 for _ ab- _ and _ c- _ type rr lyrae respectively , and an average of @[email protected] ( standard deviation of the average ) . adopting c03 reddening scale for the lmc rr lyrae stars that , on average , is 0.03 mag lower than ogle s , and cacciari & clementini ( 2003 ) @xmath141 , again we find @xmath145=19.05 and @xmath140 = 18.46 @xmath30.07 mag for the distance modulus of the lmc , in perfect agreement with c03 solution . marconi & clementini ( 2005 ) derive @xmath140 = 18.54 @xmath30.09 mag from the theoretical modeling of the light curves of lmc rr lyrae stars . we will thus adopt for the distance modulus from the rr lyrae stars the weighted average of the above solutions , namely : 18.49 @xmath30.06 ( standard deviation of the weighted average ) . next we turn to the cepheid variables . we will make our computation at @xmath79 = 0.55 which is approximately at the centroid of where the largest number of these variables are found , and will use udalski et al . ( 1999b ) @xmath1 relation for cepheids ( equation ( 11 ) ) , since it is based on a very large number of stars and is the relation most frequently used in the recent literature . the @xmath76 magnitude at @xmath79 = 0.55 , given by equation ( 11 ) is : @xmath76 = 15.524 , that corrected back to an original @xmath73 magnitude by adding @xmath58 = 0.476 mag and transformed to df05 photometric zero point , becomes : @xmath60 = 16.035 mag . at @xmath79 = @xmath140.75 and [ fe / h]=@xmath105 the @xmath0 scuti stars give @xmath73 = 19.764 @xmath3 0.131 mag ( equation 7 ) and @xmath26 = 0.981 @xmath3 0.16 mag ( equation 1 ) . combining differences , on the assumption of the same apparent distance modulus ( @xmath136=18.783 @xmath30.207 ) for the lmc @xmath0 scuti stars and cepheids , yields @xmath26 = @xmath146 mag for the cepheids . again we assume the color excess is 0.08@xmath3 0.02 mag and the absorption in @xmath15 is @xmath58 = 0.248 mag . we find a distance modulus of @xmath147=@xmath136-@xmath58= 18.535 @xmath148 mag . in this solution we have utilized the @xmath0 scuti stars to infer the luminosities of the cepheids . the above distance modulus is in excellent agreement with that obtained by an independent cepheid - based method , namely the theoretical fitting of the observed light curves of the lmc cepheids : [email protected] ( keller & wood 2002 , 2006 ) ; [email protected] ( bono , castellani , marconi 2002 ) . we adopt for the cepheids the weighted average of the above three solutions , namely : 18.53 @xmath30.02 ( standard deviation of the weighted average ) . the very good agreement found among results from the model - fitting of the observed light curves of the lmc cepheids , rr lyrae and @xmath0 scuti stars ( see sect . 4 of the present paper ) , demonstrates the reliability of this method , and its marginal dependence on the adopted pulsation code , as well as on physical and numerical assumptions in the model calculations . in summary , our best values for the distance modulus of the lmc from the three different types of pulsating variables are : 18.48 @xmath3 0.02 mag from the @xmath0 scuti stars ; 18.49 @xmath3 0.06 mag from the rr lyrae stars ; 18.53 @xmath3 0.02 mag from the cepheids . here the errors are the standard deviations of the averages of independent solutions . systematic errors related to reddening uncertainties and , for the delta scuti stars considered in this paper , uncertanties in the pulsation mode definition , dependence of the @xmath1 relation on chemical composition , pulsation and/or model atmosphere limitations , are difficult to quantify and can be significantly large ( up to a couple of tenths of magnitude ) . new data for the lmc delta scuti stars from the supermacho survey of the lmc ( cook 2006 , private communication ) enabling a direct definition of periods and pulsation modes , and reddening determinations from the star s intrinsic colors should in the future allow to cut down significantly these error contributions , thus streghtening the delta scuti star solution . however , and notwhitstanding the uncertainties still present , our finding that the three groups of pulsating stars give consistent distance moduli for the lmc already reprents a remarkable result , probably traceable to the use of the proper color excess and to the application of independent techinques to the various types of variables . in particular , all three groups of pulsating stars give distance moduli that are consistent with the long " astronomical distance scale for the lmc , in very good agreement with the value of 19.52 @xmath3 0.10 mag derived by c03 using a variety of different distance indicators of both population i and ii and a careful treatment of the error sources . in conclusion , let us emphasize the desirability of identifying other @xmath5 scuti stars in the magellanic clouds . in the smc they should be found between @xmath149 mag at log @xmath150 to @xmath151 mag at log @xmath152 . these variables can be also identified in other local group galaxies , provided that deep enough observations are obtained . for instance , a large number of sx phe stars have been successfully identified in the fornax dwarf spheroidal galaxy at 22 @xmath153 24 mag ( clementini et al . 2004 , poretti et al . 2006a , b ) . not only do the variables in these galaxies allow to improve our detailed knowledge of the properties of the variables , but distances to the galaxies as well . note that the lmc , the smc , fornax , as well as other galaxies close - by provide the opportunity to directly and easily compare luminosities of the three pulsating variables : @xmath5 scuti , rr lyrae and cepheids , that occupy the instability strip . the present paper is a first attempt in this direction . we thank mario mateo and ennio poretti for informations on the dwarf cepheid statistics outside the milky way and an anonimous referee for his / her comments and suggestions . financial support for this study was provided by miur , under the scientific project 2004020323 , ( p.i . : massimo capaccioli ) . alibert , y. , baraffe , i. , hauschild , p. , & allard , f. 1999 , a&a , 344 , 551 bono , g. , caputo , f. , cassisi , s. , castellani , v. , marconi , m. , & stellingwerf , r.f . 1997 , apj , 477 , 346 bono , g. , castellani , v. , & marconi , m. 2000 , apj , 532 , l129 bono , g. , castellani , v. , & marconi , m. 2002 , apj , 565 , l83 breger , m. 2000 , in delta scuti and related stars , eds . m. breger , & m.h . montgomery , asp conference series , 210,3 burnstein , d. , & heiles , c. 1982 , , 87 , 1165 buzasi , d.l . , bruntt , h. , bedding , r. , retter , a. , kieldsen , h. , preston , h.l . , mandeville , w.j . , suarez , j.c . , catanzarite , j. , conrow , t. & laher , r. 2005 , , 619 , 1072 cacciari , c. , & clementini , g. 2003 , in stellar candles for the extragalactic distance scale , eds . d. alloin , & w. gieren , lecture notes in physics , 635 , 105 caldwell , j.a.r . , & coulson , i.m . 1986 , mnras , 218 , 223 caputo , f. , marconi , m. , & musella , i. 2000 , a&a , 354 , 610 caputo , f. , marconi , m. , & musella , i. 2002 , , 566 , 833 castellani , v. , deglinnocenti , s. , & marconi , m. 2002 , in omega centauri , a unique window into astrophysics , eds . f. van leeuwen , j. d. hughes , and g. piotto , asp conference series , 265 , 193 clementini , g. , et al . 2004 , in variable stars in the local group , eds . kurtz , & k.r . pollard , asp conference series , 310 , 60 clementini , g. , gratton , r.g . , bragaglia , a. , carretta , e. , di fabrizio , l. , & maio , m. 2003 , aj , 125 , 1309 ( c03 ) cook , k.h . 2006 , private communication cousins , a. w. j. , & caldwell , j. a. r. 1985 , observatory , 105 , 134 di fabrizio , l. , et al . 2002 , mnras , 336 , 841 di fabrizio , l. , clementini , g. , maio , m. , bragaglia , a. , carretta , e. , gratton , r. , montegriffo , p. , & zoccali , m. 2005 , a&a , 430 , 603 ( df05 ) fernie , j. d. 1992 , aj , 103 , 1647 gieren , w.p . , fouqu , p. , & gomez , m. 1998 , , 496 , 17 hoyle , f. , shanks , t. , & tanvir , n.r . 2003 , mnras , 345 , 269 jurcsik , j. , szeidl , b. , vradi , m. , henden , a. , hurta , zs . , lakatos , b. , posztobnyi , k. , klagyivik , p. , & sdor , . 2006 , a&a , 445 , 617 kaluzny , j. , mochnacki , s. , & rucinski , s.m . 2006 , , 131 , 407 kaluzny , j. , & rucinski , s.m . 2003 , , 126 , 237 keller , s.c . , & wood , p.r . 2002 , , 578 , 144 keller , s.c . , & wood , p.r . 2006 , , 642 , 834 kervella , p. , nardetto , n. , bersier , d. , mourard , d. , & coud du foresto , v. 2004 , a&a , 416 , 941 laney , c.d . , joner , m. , & schwendiman , l. 2002 , in asp conf . 259 , radial and nonradial pulsations as probes of stellar physics , eds . c. aerts , t.r . bedding and j. christensen - dalsgaard , p. 112 laney , c.d . , & stobie , r.s . 1994 , mnras , 266 , 441 luck , r. e. , moffett , t. j. , barnes , t. g. iii , & gieren , w. p. 1998 , , 115 , 605 lutz , t.e . , & kelker , d.h . 1973 , pasp , 85 , 573 marconi , m. , & clementini , g. 2005 , , 129 , 2257 mateo , m. , hurley - keller , d. , & nemec , j. 1998 , , 115 , 1856 mcnamara , d.h . 1997 , pasp , 109,1221 mcnamara , d.h . 2000 , pasp , 112,1096 mcnamara , d. h. , rose , m.b . , brown , p.j . , ketcheson , d.i . , maxwell , j.e . , smith , k.m . , & wodey , r.c . , 2004 , in variable stars in the local group , eds . d. w. kurtz , and k. r. pollard , asp conference series , 310 , 525 nemec , j.m . , linnell nemec , a.f . , & lutz , t.e . 1994 , , 108 , 222 petersen , j.o . , & christiansen - dalsgaard , v. 1996 , a&a , 312,463 petersen , j.o . , & christiansen - dalsgaard , v. 1999 , a&a , 352,547 petersen , j.o . , & frehammer , l.m . 2002 , asp conf . 259 , radial and nonradial pulsations as probes of stellar physics , eds . c. aerts , t.r bedding and j. christensen - dalsgaard , p. 136 poretti , e. 1999 , , 343 , 385 poretti et al . 2006a , in workshop on stellar pulsation and evolution , eds . a. walker and g. bono , mem.s.a.it . 77 , 219 poretti et al . 2006b , letters , in preparation rodrguez , e. , & lpez - gonzlez , m. j. 2000 , a&a , 359 , 597 rodrguez , e. , lpez - gonzlez m.j . , & lpez de coca , p. 2000 , a&as , 144 , 469 russell , s.c . , & dopita , m.a . 1990 , , 74 , 93 santolamazza , p. , marconi , m. , bong , g. , caputo , f. , cassisi , s. , & gilliland , r.l . 2001 , , 554 , 1124 schlegel , d. j. , finkbeiner , d. p. , & davis , m. 1998 , apj , 500 , 525 soszyski , i. et al . 2003 , acta astronomica , 53 , 93 tammann , g.a . , sandage , a. , & reindl , b. 2003 , a&a , 404 , 423 templeton , m. 2000 , ph.d . thesis , new mexico state univ . udalski , a. et al . 1999a , acta astronomica , 49 , 223 udalski , a. et al . 1999b , acta astronomica , 49 , 201 wood , p.r . , arnold , a. , & sebo , k.m . 1997 , apj , 485 , l52 notes : + ( 1 ) star i d s and numbers are from kaluzny & rucinski ( 2003 ) for the lw 55 variables , and from kaluzny et al . ( 2006 ) for the lmc disk @xmath0 scuti star . + ( 2 ) periods for the @xmath0 scuti stars in lw 55 were derived in the present paper and have errors of @xmath3 0.02 in @xmath79 , for the disk @xmath0 scuti are from kaluzny et al . + ( 3 ) magnitudes corrected for the tilt of the lmc . + ( 4 ) individual reddenings were estimated from the observed @xmath62 assuming an intrinsic color @xmath102=0.256 ( see section 3 ) . + ( 5 ) tilt - corrected magnitudes , dereddened assuming @xmath78=3.1@xmath4 . + ( 6 ) second overtone pulsator , in brackets the fundamentalized period assuming @xmath930.6223 . + ( 7 ) fundamental mode pulsators . +

the absolute magnitude of the variable was determined from the period - luminosity ( ) relation for galactic scuti stars and from the theoretical modeling of the observed light curves with nonlinear pulsation models . the two methods give distance moduli for the lmc of 18.46 0.19 and 18.48 0.15 , respectively , for a consistent value of the stellar reddening of=0.08.02 . we have also analyzed 24 scuti candidates discovered in the ogle ii survey of the lmc , and 7 variables identified in the open cluster lw 55 and in the galaxy disk by kaluzny et al . we compare the results obtained from the scuti variables with those derived from the lmc rr lyrae stars and cepheids . the corresponding distance moduli are : scuti stars 18.48 0.02 mag ( standard deviation of the weighted average of the three above solutions ) ; rr lyrae stars 18.49.06 mag ; and cepheids 18.53.02 mag . we have assumed an average color excess of = 0.08 mag 0.02 for both scuti stars and cepheids . within the observational uncertainties , the three groups of pulsating stars yield very similar distance moduli . these moduli are all consistent with the long " astronomical distance scale for the large magellanic cloud .

we present results from a well studied scuti star discovered in the large magellanic cloud ( lmc ) . the absolute magnitude of the variable was determined from the period - luminosity ( ) relation for galactic scuti stars and from the theoretical modeling of the observed light curves with nonlinear pulsation models . the two methods give distance moduli for the lmc of 18.46 0.19 and 18.48 0.15 , respectively , for a consistent value of the stellar reddening of=0.08.02 . we have also analyzed 24 scuti candidates discovered in the ogle ii survey of the lmc , and 7 variables identified in the open cluster lw 55 and in the galaxy disk by kaluzny et al . ( 2003 , 2006 ) . we find that the lmc scuti stars define a relation whose slope is very similar to that defined by the galactic scuti variables , and yield a distance modulus for the lmc of 18.50.22 mag . we compare the results obtained from the scuti variables with those derived from the lmc rr lyrae stars and cepheids . the corresponding distance moduli are : scuti stars 18.48 0.02 mag ( standard deviation of the weighted average of the three above solutions ) ; rr lyrae stars 18.49.06 mag ; and cepheids 18.53.02 mag . we have assumed an average color excess of = 0.08 mag 0.02 for both scuti stars and cepheids . within the observational uncertainties , the three groups of pulsating stars yield very similar distance moduli . these moduli are all consistent with the long " astronomical distance scale for the large magellanic cloud .

astro-ph0401533

c

we present a catalog of ngc 2264 x - ray sources . the observations were taken with the acis - i on board the chandra x - ray observatory . the catalog , consisting of 263 sources , includes x - ray luminosity , optical and infrared photometry and x - ray variability information . we found 41 variable sources , 14 of which have a flare like light curve , and 2 of which have a pattern of a steady increase or decrease . from the optical and infrared counterparts of the x - ray sources , we have learned that most of the x - ray sources have colors consistent with ctts that are younger than 3@xmath67 years . this catalog of x - ray sources will be used to study the relationship between rotational properties and x - ray characteristics of ngc 2264 stars in paper ii @xcite . we plan to discuss correlations of @xmath5/@xmath7 with rotation rate ( period and @xmath75sin@xmath76 ) , disk indicators ( @xmath77 , @xmath3 , @xmath1 , and @xmath0 ) , and mass accretion rate as derived from @xmath1 excess . we will also compare the @xmath5/@xmath7 values found here with those from other young clusters . svr gratefully thanks jeonghee rho , kenji hamaguchi , ettore flaccomio , peter freeman & scott wolk for the useful correspondence about x - ray and chandra data processing and analysis . svr also thanks august muench , paul eskridge , richard pogge , phil appleton , mark lacy , and dario fadda for interesting discussions . we thank the anonymous referee for her / his careful review of the manuscript . financial support for this work was provided by nasa grant go2 - 3011x . this research has made extensive use of nasa s astrophysics data system abstract service , the simbad database , operated at cds , strasbourg , france , and the nasa / ipac infrared science archive , which is operated by the jet propulsion laboratory , california institute of technology , under contract with the national aeronautics and space administration . the research described in this paper was partially carried out at the jet propulsion laboratory , california institute of technology , under a contract with the national aeronautics and space administration . brandt , w. n. , alexander , d. m. , hornschemeier , a. e. , garmire , g. p. , schneider , d. p. , barger , a. j. , bauer , f. e. , broos , p. s. , cowie , l. l. , townsley , l. k. , burrows , d. n. , chartas , g. , feigelson , e. d. , griffiths , r. e. , nousek , j. a. , & sargent , w. l. w. , 2001 , , 122 , 2810 harnden , f. r. , jr . , adams , n. r. , damiani , f. , drake , j. j. , evans , n. r. , favata , f. , flaccomio , e. , freeman , p. , jeffries , r. d. , kashyap , v. , micela , g. , patten , b. m. , pizzolato , n. , schachter , j. f. , sciortino , s. , stauffer , j. , wolk , s. j. , zombeck , m. v. , 2001 , , 547 , l141 reid , i. n. , brewer , c. , brucato , r. j. , mckinley , w. r. , maury , a. , mendenhall , d. , mould , j. r. , mueller , j. , neugebauer , g. , phinney , j. , sargent , w. l. w. , schombert , j. , & thicksten , r. , 1991 , , 661 rosati , p. , tozzi , p. , giacconi , r. , gilli , r. , hasinger , g. , kewley , l. , mainieri , v. , nonino , m. , norman , c. , szokoly , g. , wang , j. x. , zirm , a. , bergeron , j. , borgani , s. , gilmozzi , r. , grogin , n. , koekemoer , a. , schreier , e. , zheng , w. , 2002 , , 566 , 667 rcllcccccccrc 1 & cxorrs j064009.6 + 095338 & 6 40 9.69 & 9 53 38.355 & 9.85 & 14.07 & 0.99 & 40.5 & 0.37 & 0.228e-14 & 29.20 & 100 & + 3 & cxorrs j064012.8 + 094853 & 6 40 12.84 & 9 48 53.952 & 8.78 & 11.07 & 0.97 & 30.2 & 0.16 & 0.986e-15 & 28.83 & 100 & + 4 & cxorrs j064013.5 + 095313 & 6 40 13.57 & 9 53 13.437 & 8.79 & 11.12 & 0.98 & 42.2 & 0.21 & 0.129e-14 & 28.95 & 99 & + 6 & cxorrs j064018.3 + 094416 & 6 40 18.31 & 9 44 16.054 & 9.72 & 13.73 & 0.99 & 31.0 & 0.94 & 0.579e-14 & 29.60 & 100 & + 8 & cxorrs j064018.5 + 095206 & 6 40 18.57 & 9 52 6.259 & 7.34 & 7.72 & 0.98 & 41.9 & 0.52 & 0.320e-14 & 29.35 & 100 & + 9 & cxorrs j064019.3 + 094829 & 6 40 19.34 & 9 48 29.537 & 7.37 & 7.77 & 0.98 & 41.0 & 3.30 & 0.203e-13 & 30.15 & 65 & v , f + 11 & cxorrs j064020.1 + 094511 & 6 40 20.16 & 9 45 11.315 & 8.78 & 11.12 & 0.99 & 39.4 & 4.95 & 0.305e-13 & 30.32 & 98 & + 12 & cxorrs j064020.9 + 094405 & 6 40 20.91 & 9 44 5.391 & 9.36 & 12.69 & 0.99 & 39.3 & 2.01 & 0.124e-13 & 29.93 & 99 & + 13 & cxorrs j064020.9 + 094821 & 6 40 20.93 & 9 48 21.477 & 7.04 & 7.08 & 0.97 & 40.1 & 0.67 & 0.413e-14 & 29.46 & 100 & + 14 & cxorrs j064020.9 + 095519 & 6 40 20.99 & 9 55 19.184 & 8.05 & 9.30 & 0.98 & 40.9 & 0.65 & 0.400e-14 & 29.44 & 100 & + 15 & cxorrs j064021.8 + 095209 & 6 40 21.86 & 9 52 9.178 & 6.54 & 6.15 & 0.97 & 44.0 & 0.34 & 0.209e-14 & 29.16 & 100 & + 16 & cxorrs j064022.2 + 095428 & 6 40 22.28 & 9 54 28.859 & 7.33 & 7.68 & 0.98 & 39.2 & 14.69 & 0.905e-13 & 30.80 & 0 & v , f + 18 & cxorrs j064022.6 + 094945 & 6 40 22.62 & 9 49 45.912 & 6.28 & 5.71 & 0.97 & 38.5 & 0.90 & 0.554e-14 & 29.58 & 99 & + 19 & cxorrs j064022.9 + 095312 & 6 40 22.99 & 9 53 12.159 & 6.62 & 6.30 & 0.97 & 40.9 & 0.41 & 0.253e-14 & 29.24 & 100 & + 20 & cxorrs j064023.3 + 095454 & 6 40 23.39 & 9 54 54.251 & 7.34 & 7.72 & 0.98 & 39.1 & 0.71 & 0.437e-14 & 29.48 & 99 & + 22 & cxorrs j064023.7 + 095523 & 6 40 23.77 & 9 55 23.330 & 7.55 & 8.17 & 0.98 & 37.8 & 6.18 & 0.381e-13 & 30.42 & 96 & + 23 & cxorrs j064024.8 + 095311 & 6 40 24.89 & 9 53 11.135 & 6.17 & 5.51 & 0.97 & 39.8 & 1.32 & 0.813e-14 & 29.75 & 99 & + 24 & cxorrs j064025.5 + 094825 & 6 40 25.52 & 9 48 25.778 & 5.96 & 5.17 & 0.96 & 44.0 & 3.27 & 0.201e-13 & 30.14 & 99 & + 25 & cxorrs j064026.9 + 095138 & 6 40 26.97 & 9 51 38.500 & 5.24 & 4.13 & 0.96 & 40.1 & 0.38 & 0.234e-14 & 29.21 & 100 & + 26 & cxorrs j064027.3 + 095307 & 6 40 27.36 & 9 53 7.553 & 5.59 & 4.62 & 0.96 & 44.8 & 0.25 & 0.154e-14 & 29.03 & 100 & + 27 & cxorrs j064027.6 + 095349 & 6 40 27.66 & 9 53 49.571 & 5.87 & 5.02 & 0.97 & 44.5 & 0.28 & 0.172e-14 & 29.08 & 100 & + 28 & cxorrs j064027.8 + 094119 & 6 40 27.85 & 9 41 19.667 & 10.61 & 16.48 & 1.00 & 38.9 & 0.42 & 0.259e-14 & 29.25 & 99 & + 29 & cxorrs j064028.8 + 094823 & 6 40 28.85 & 9 48 23.232 & 5.22 & 4.08 & 0.96 & 41.7 & 2.89 & 0.178e-13 & 30.09 & 99 & + 30 & cxorrs j064029.0 + 094217 & 6 40 29.01 & 9 42 17.432 & 9.63 & 13.43 & 0.99 & 40.3 & 5.46 & 0.336e-13 & 30.37 & 50 & v + 32 & cxorrs j064029.3 + 094407 & 6 40 29.34 & 9 44 7.148 & 8.01 & 9.20 & 0.98 & 41.8 & 0.23 & 0.142e-14 & 28.99 & 100 & + 33 & cxorrs j064029.4 + 094736 & 6 40 29.45 & 9 47 36.838 & 5.47 & 4.43 & 0.96 & 41.4 & 0.28 & 0.172e-14 & 29.08 & 100 & + 34 & cxorrs j064029.9 + 095010 & 6 40 29.93 & 9 50 10.049 & 4.44 & 3.15 & 0.95 & 45.4 & 1.75 & 0.108e-13 & 29.87 & 99 & + 36 & cxorrs j064030.6 + 095014 & 6 40 30.61 & 9 50 14.085 & 4.28 & 3.00 & 0.95 & 45.5 & 0.65 & 0.400e-14 & 29.44 & 100 & + 37 & cxorrs j064030.6 + 094610 & 6 40 30.68 & 9 46 10.398 & 6.20 & 5.56 & 0.97 & 40.4 & 3.65 & 0.225e-13 & 30.19 & 27 & v , f + 38 & cxorrs j064031.6 + 094449 & 6 40 31.62 & 9 44 49.861 & 7.11 & 7.23 & 0.98 & 39.4 & 0.35 & 0.216e-14 & 29.17 & 100 & + 39 & cxorrs j064031.6 + 094427 & 6 40 31.68 & 9 44 27.278 & 7.44 & 7.92 & 0.98 & 39.5 & 0.89 & 0.548e-14 & 29.58 & 100 & + 40 & cxorrs j064031.6 + 094823 & 6 40 31.70 & 9 48 23.290 & 4.64 & 3.35 & 0.95 & 40.8 & 0.25 & 0.154e-14 & 29.03 & 100 & + cccccc 196 & r 3295 & 155.8 & 0.18 & 0.57 & 14.10 + 113 & r 2942 & 84.5 & 0.30 & 1.36 & 5.25 + 181 & r 3245 & 54.9 & 0.76 & 3.92 & 3.45 + 75 & r 2752 & 42.1 & 0.80 & 5.16 & 2.48 + 94 & r 2840 & 36.4 & 0.40 & 2.21 & 2.34 + 203 & r 3309 & 27.4 & 0.58 & 2.34 & 1.68 + 142 & r 3091 & 24.1 & 1.00 & 4.46 & 1.58 + 251 & r 3470 & 21.3 & 0.34 & 1.25 & 1.26 + 228 & r 3390 & 18.5 & 0.64 & 2.58 & 1.04 + 60 & r 2633 & 18.0 & 0.30 & 1.66 & 0.79 + 208 & r 3323 & 17.6 & 0.25 & 1.11 & 1.03 + 230 & r 3394 & 17.2 & 0.32 & 1.62 & 0.95 + 268 & r 3555 & 15.9 & 0.68 & 3.28 & 0.93 + 16 & r 2173 & 14.7 & 0.41 & 2.23 & 0.83 + 214 & r 3342 & 13.3 & 0.32 & 2.08 & 0.76 + cllrrrrrrrrr r 1817 & 6 40 10.00 & 9 53 41.19 & 1 & 18.47 & 18.36 & 16.93 & 16.00 & 14.94 & 13.53 & 12.94 & 12.73 + r 2065 & 6 40 18.52 & 9 44 18.73 & 6 & 19.65 & 19.09 & 17.50 & 16.51 & 15.48 & 14.06 & 13.40 & 13.15 + r 2066 & 6 40 18.53 & 9 52 05.73 & 8 & & 20.28 & 18.94 & 17.94 & 16.15 & 14.55 & 13.89 & 13.65 + r 2093 & 6 40 19.37 & 9 48 29.83 & 9 & 18.03 & 17.38 & 15.90 & 15.00 & 13.99 & 12.59 & 11.99 & 11.89 + r 2137 & 6 40 20.87 & 9 44 11.29 & 12 & & 21.94 & 20.51 & & 18.06 & 16.32 & 15.84 & 15.22 + r 2163 & 6 40 21.84 & 9 52 09.11 & 15 & & & 19.92 & & 17.00 & 15.41 & 14.76 & 14.59 + r 2173 & 6 40 22.28 & 9 54 28.65 & 16 & 18.20 & 17.16 & 15.62 & 14.61 & 13.56 & 12.29 & 11.56 & 11.31 + r 2182 & 6 40 22.68 & 9 49 45.92 & 18 & 19.41 & 18.57 & 17.23 & 16.20 & 15.12 & 13.88 & 13.12 & 12.87 + r 2190 & 6 40 23.01 & 9 53 12.08 & 19 & & & 20.50 & 18.79 & 17.16 & 15.49 & 14.78 & 14.34 + r 2204 & 6 40 23.48 & 9 54 55.30 & 20 & & 20.18 & 18.54 & 17.29 & 15.44 & 14.01 & 13.35 & 12.97 + r 2217 & 6 40 23.79 & 9 55 23.42 & 22 & 18.59 & 19.03 & 17.55 & 16.32 & 14.92 & 13.12 & 12.26 & 11.81 + r 2251 & 6 40 24.84 & 9 53 11.45 & 23 & 19.89 & 19.23 & 17.49 & & 15.04 & 13.59 & 12.56 & 12.27 + r 2271 & 6 40 25.52 & 9 48 25.88 & 24 & 17.97 & 16.83 & 15.51 & 14.66 & 13.88 & 13.24 & 12.61 & 12.45 + r 2374 & 6 40 28.66 & 9 48 24.31 & 29 & & & 15.76 & 15.02 & 14.09 & 13.00 & 12.32 & 12.04 + r 2383 & 6 40 29.01 & 9 42 17.18 & 30 & 17.80 & 16.79 & 15.47 & 14.66 & 13.92 & 12.94 & 12.24 & 12.05 + r 2391 & 6 40 29.29 & 9 44 07.42 & 32 & 19.17 & 18.85 & 17.68 & 16.62 & 15.40 & 14.06 & 13.34 & 13.06 + r 2401 & 6 40 29.46 & 9 47 36.88 & 33 & 19.79 & 20.54 & 20.15 & 17.74 & 16.77 & 14.54 & 13.89 & 13.50 + r 2419 & 6 40 29.95 & 9 50 10.25 & 34 & 16.51 & 15.62 & 14.44 & 13.78 & 13.16 & 12.22 & 11.68 & 11.56 + r 2442 & 6 40 30.64 & 9 50 14.38 & 36 & & 18.83 & 17.10 & 16.06 & 14.53 & 12.75 & 11.67 & 11.19 + r 2443 & 6 40 30.66 & 9 46 10.59 & 37 & 10.21 & 16.63 & 15.32 & 14.61 & 13.90 & 12.86 & 12.23 & 12.06 + r 2474 & 6 40 31.70 & 9 48 23.26 & 40 & 17.25 & 16.73 & 15.57 & 14.78 & 14.08 & 13.02 & 12.35 & 12.18 + r 2477 & 6 40 31.73 & 9 53 30.09 & 41 & & & & & 17.31 & 14.46 & 13.11 & 12.49 + r 2476 & 6 40 31.73 & 9 49 59.20 & 42 & & & & & 17.39 & 15.57 & 14.90 & 14.61 + r 2507 & 6 40 32.65 & 9 49 32.90 & 44 & & & 20.02 & 18.19 & 16.37 & 14.94 & 14.28 & 14.05 + r 2503 & 6 40 32.42 & 9 41 45.13 & 45 & 20.48 & 20.03 & 18.51 & 17.25 & 15.81 & 14.29 & 13.59 & 13.34 + r 2511 & 6 40 32.85 & 9 51 28.97 & 46 & 18.12 & 17.06 & 15.68 & 14.78 & 13.87 & 12.68 & 12.00 & 11.79 + r 2518 & 6 40 33.18 & 9 49 54.31 & 48 & & & 20.50 & 18.60 & 17.08 & 14.36 & 12.98 & 12.16 + r 2550 & 6 40 33.85 & 9 48 43.64 & 50 & 19.25 & 18.62 & 17.21 & 16.11 & 14.76 & 13.45 & 12.72 & 12.35 + r 2590 & 6 40 34.87 & 9 45 45.11 & 51 & 20.20 & 19.77 & 18.24 & 17.11 & 15.85 & 14.23 & 13.59 & 13.24 + r 2593 & 6 40 34.94 & 9 54 06.87 & 52 & & 20.10 & 18.23 & 17.05 & 15.35 & 13.74 & 13.05 & 12.81 + r 2597 & 6 40 35.23 & 9 51 56.25 & 53 & 19.00 & 17.76 & 16.20 & 15.17 & 14.01 & 12.61 & 11.83 & 11.47 + r 2614 & 6 40 36.06 & 9 47 35.97 & 56 & & & 19.53 & & 16.24 & 14.29 & 13.94 & 13.51 + r 2631 & 6 40 36.58 & 9 50 45.47 & 57 & 17.89 & 18.13 & 16.93 & 15.91 & 14.59 & 12.93 & 12.14 & 11.71 + r 2635 & 6 40 36.69 & 9 48 22.75 & 58 & 16.94 & 16.72 & 15.61 & 14.85 & 14.06 & 13.04 & 12.33 & 12.21 + r 2636 & 6 40 36.71 & 9 52 02.89 & 59 & 17.23 & 17.51 & 16.27 & 15.33 & 14.35 & 12.73 & 11.76 & 11.10 + r 2633 & 6 40 36.67 & 9 47 22.53 & 60 & 13.12 & 13.20 & 12.43 & & 10.91 & 9.65 & 9.00 & 8.59 + rllll 1 & r 1817 & & 2mass j06400993 + 0953415 & + 6 & r 2065 & sung 1487 & 2mass j06401846 + 0944188 & + 8 & r 2066 & & 2mass j06401847 + 0952060 & + 9 & r 2093 & sung 68 & 2mass j06401930 + 0948299 & + 12 & r 2137 & sung 1526 & 2mass j06402081 + 0944114 & + 15 & r 2163 & sung 1539 & 2mass j06402178 + 0952092 & + 16 & r 2173 & sung 78 & 2mass j06402221 + 0954288 & + 18 & r 2182 & sung 79 & 2mass j06402262 + 0949462 & + 19 & r 2190 & sung 1560 & 2mass j06402295 + 0953125 & + 20 & r 2204 & sung 1567 & 2mass j06402342 + 0954555 & + 22 & r 2217 & sung 1578 & 2mass j06402373 + 0955238 & v594 mon , ogura 74 + 23 & r 2251 & sung 1599 & 2mass j06402484 + 0953114 & + 24 & r 2271 & sung 88 & 2mass j06402547 + 0948259 & + 29 & r 2374 & & & ogura 76 + 30 & r 2383 & & & walker 52 + 32 & r 2391 & sung 1705 & 2mass j06402924 + 0944075 & + 33 & r 2401 & sung 1708 & 2mass j06402941 + 0947369 & + 34 & r 2419 & sung 104 & 2mass j06402989 + 0950104 & walker 54,vsb 40 + 36 & r 2442 & sung 1739 & 2mass j06403059 + 0950147 & + 37 & r 2443 & sung 108 & 2mass j06403061 + 0946106 & ogura 81,fx 15,vsb 170 + 40 & r 2474 & sung 113 & 2mass j06403164 + 0948233 & walker 58,v413 mon + 41 & r 2477 & & 2mass j06403168 + 0953304 & + 42 & r 2476 & & 2mass j06403167 + 0949593 & + 44 & r 2507 & sung 1773 & 2mass j06403258 + 0949332 & + 45 & r 2503 & sung 1767 & 2mass j06403237 + 0941449 & + 46 & r 2511 & sung 118 & 2mass j06403280 + 0951293 & fx 17 + 48 & r 2518 & sung 1781 & 2mass j06403311 + 0949547 & + 50 & r 2550 & sung 1798 & 2mass j06403378 + 0948438 & + 51 & r 2590 & sung 1827 & 2mass j06403482 + 0945452 & + 52 & r 2593 & sung 1828 & 2mass j06403489 + 0954071 & + 53 & r 2597 & sung 128 & 2mass j06403518 + 0951567 & ogura 85 + 56 & r 2614 & sung 1841 & & + 57 & r 2631 & sung 129 & 2mass j06403652 + 0950456 & + 58 & r 2635 & sung 130 & 2mass j06403662 + 0948229 & fx 19 + 59 & r 2636 & sung 131 & 2mass j06403665 + 0952032 & + 60 & r 2633 & sung 132 & 2mass j06403667 + 0947225 & walker 66,v780 mon , vsb 46 + llllccccc r 2772 & sung 153 & 2mass 06404102 + 0947577 & & @xmath78 4.11 & @xmath78 0.0955 & @xmath780.0056 & @xmath78 28.61 + r 2898 & sung 173 & 2mass 06404464 + 0948021 & walker 90 , v590 mon , hbc 219 , vsb 62 & @xmath78 3.00 & @xmath78 0.0703 & @xmath780.0041 & @xmath78 28.48 + r 2982 & sung 2066 & 2mass 06404729 + 0947274 & & @xmath7812.63 & @xmath78 0.3940 & @xmath780.0231 & @xmath78 29.22 + r 2491 & sung 1760 & 2mass 06403200 + 094935 & & @xmath78 3.00 & @xmath78 0.1794 & @xmath780.0105 & @xmath78 28.88 + & & & v590 mon , hbc 219 , herbig 25 & @xmath78 4.11 & @xmath78 0.0962 & @xmath780.0056 & @xmath78 28.61 + r 2792 & sung 159 & 2mass 06404156 + 0955174 & vsb 56 & @xmath78 3.00 & @xmath78 0.1984 & @xmath780.0116 & @xmath78 28.93 + & & & ln mon , hbc 215 , herbig 21 & @xmath78 4.82 & @xmath78 0.1193 & @xmath780.0070 & @xmath78 28.71 + r 3167 & sung 2178 & 2mass 06405413 + 0948434 & & @xmath78 4.11 & @xmath78 0.0944 & @xmath780.0055 & @xmath78 28.60 + r 3216 & sung 2206 & 2mass 06405573 + 0946456 & & @xmath78 3.82 & @xmath78 0.0903 & @xmath780.0053 & @xmath78 28.58 + & & & lr mon , hbc 220 , herbig 27 & @xmath78 4.11 & @xmath78 0.1816 & @xmath780.0106 & @xmath78 28.89 + r 2972 & sung 2056 & 2mass 06404694 + 0955036 & & @xmath78 7.24 & @xmath78 0.1697 & @xmath780.0099 & @xmath78 28.86 + r 3022 & sung 195 & mass 06404888 + 0951444 & walker 100 , hd261841 , vsb 72 & @xmath78 7.40 & @xmath78 0.1677 & @xmath780.0098 & @xmath78 28.85 + r 3041 & sung 196 & 2mass 06404953 + 0953230 & walker 104 , vsb 74 & @xmath78 6.78 & @xmath78 0.1561 & @xmath780.0092 & @xmath78 28.82 + r 3073 & sung 2125 & 2mass 06405033 + 0954158 & & @xmath78 4.11 & @xmath78 0.0958 & @xmath780.0056 & @xmath78 28.61 + r 3110 & sung 209 & 2mass 06405155 + 0951494 & walker 109 , hd261878 , vsb 79 & @xmath78 5.82 & @xmath78 0.1328 & @xmath780.0078 & @xmath78 28.75 + & & & fx 67 , vsb 245 & @xmath78 3.94 & @xmath78 0.0926 & @xmath780.0054 & @xmath78 28.60 +

the observations were taken with the acis - i camera with an exposure time of 48.1 ks . we present a catalog of 263 sources , which includes x - ray luminosity , optical and infrared photometry and x - ray variability information . we found 41 variable sources , 14 of which have a flare - like light curve , and 2 of which have a pattern of a steady increase or decrease over a 10 hour period . the optical and infrared photometry for the stars identified as x - ray sources are consistent with most of these objects being pre - main sequence stars with ages younger than 3 myr .

we present results of a chandra observation of a field in ngc 2264 . the observations were taken with the acis - i camera with an exposure time of 48.1 ks . we present a catalog of 263 sources , which includes x - ray luminosity , optical and infrared photometry and x - ray variability information . we found 41 variable sources , 14 of which have a flare - like light curve , and 2 of which have a pattern of a steady increase or decrease over a 10 hour period . the optical and infrared photometry for the stars identified as x - ray sources are consistent with most of these objects being pre - main sequence stars with ages younger than 3 myr .

astro-ph0407438

c

we have performed a single dish 21 cm h survey of 20 local lcbgs chosen to be local analogs to the numerous lcbgs studied at intermediate redshifts ( 0.4 @xmath0 z @xmath0 0.7 ) . our findings have verified results from intermediate redshift lcbg studies . we have found that local lcbgs are a morphologically heterogeneous mixture of galaxies . they are typically gas - rich , with median values of m@xmath3 = 5@xmath410@xmath6 m@xmath7 and m@xmath3 l@xmath8 = 0.4 m@xmath7 l@xmath9 . approximately half have mass - to - light ratios approximately ten times smaller than local galaxies of all hubble types at similar luminosities , confirming that these are indeed small galaxies undergoing vigorous bursts of star formation . this proportion is likely an underestimate , as nearly half our sample of galaxies may have dynamical mass overestimates . by comparing line widths and radii with local galaxy populations , we find that local lcbgs are consistent with magellanic spirals , and the more massive irregulars and dwarf ellipticals . measurements of the length of starburst , amount of fading , and ability of these galaxies to retain their interstellar media will help to constrain the evolutionary possibilities of this galaxy class . we begin to address these issues in paper ii @xcite , where we present the results of a molecular gas survey of these same local lcbgs . _ acknowledgments _ we thank the referee for helpful comments which improved the quality of this paper . we thank rick fisher for providing the h spectrum of sdssj0934 + 0014 . we also thank the operators and staff at the gbt for their help with the observing and reduction , and their hospitality . support for this work was provided by the nsf through award gssp02 - 0001 from the nrao . d. j. p. acknowledges generous support from an nsf mps distinguished international postdoctoral research fellowship , nsf grant ast0104439 . r. g. acknowledges funding from nasa grant ltsa nag5 - 11635 . funding for the creation and distribution of the sdss archive has been provided by the alfred p. sloan foundation , the participating institutions , the national aeronautics and space administration , the national science foundation , the u.s . department of energy , the japanese monbukagakusho , and the max planck society . the sdss web site is http://www.sdss.org/. we have made extensive use of hyperleda ( http://www-obs.univ-lyon1.fr/hypercat/ ) and the nasa / ipac extragalactic database ( ned ) which is operated by the jet propulsion laboratory , california institute of technology , under contract with the national aeronautics and space administration ( http://nedwww.ipac.caltech.edu/ ) . the digitized sky surveys were produced at the space telescope science institute under u.s . government grant nag w-2166 . the images of these surveys are based on photographic data obtained using the oschin schmidt telescope on palomar mountain and the uk schmidt telescope . llcccclcc mrk 297 & ngc 6052 , ngc 6064 & 67 & @xmath1321.0 & 0.4 & 20.6 & & sc & y + mrk 314 & ngc 7468 & 30 & @xmath1318.5 & 0.4 & 20.2 & e3 , pec ( polar ring ? ) & e & n + mrk 325 & ngc 7673 , mrk 325 & 49 & @xmath1320.0 & 0.4 & 20.0 & sac ? , pec , h ii starburst & sc & y + mrk 538 & ngc 7714 & 40 & @xmath1320.1 & 0.4 & 20.2 & sb(s)b , pec , h ii liner & sbb & y + sdss j011932.95 + 145219.0 & ngc 469 & 59 & @xmath1318.9 & 0.4 & 20.3 & & & y + sdss [email protected] & & 69 & @xmath1318.8 & 0.5 & 20.2 & & & n + sdss [email protected] & & 67 & @xmath1318.6 & 0.4 & 20.1 & & & n + sdss j072849.75 + 353255.2 & & 56 & @xmath1318.9 & 0.4 & 20.3 & s ? & sbc & n + sdss j083431.70 + 013957.9 & & 59 & @xmath1319.1 & 0.6 & 20.6 & sb(s)b & sbb & y + sdss j090433.53 + 513651.1 & mrk 101 & 68 & @xmath1319.7 & 0.6 & 20.3 & s & sc & n + sdss j091139.74 + 463823.0 & mrk 102 & 61 & @xmath1319.3 & 0.5 & 19.1 & s ? & & n + sdss j093410.52 + 001430.2 & mrk 1233 & 70 & @xmath1319.8 & 0.3 & 19.9 & sb & sbc & y + sdss j093635.36 + 010659.8 & & 71 & @xmath1319.1 & 0.6 & 21.0 & & & y + sdss [email protected] & & 68 & @xmath1319.3 & 0.5 & 20.4 & s0 & s0 & n + sdss j111836.35 + 631650.4 & mrk 165 & 46 & @xmath1318.6 & 0.4 & 19.4 & compact starburst & & n + sdss j123440.89 + 031925.1 & ngc 4538 & 67 & @xmath1319.2 & 0.6 & 21.0 & s pec & sbc & n + sdss j131949.93 + 520341.1 & & 67 & @xmath1318.7 & 0.2 & 20.1 & & & y + sdss j140203.52 + 095545.6 & ngc 5414 , mrk 800 & 61 & @xmath1319.7 & 0.5 & 19.7 & pec & & y + sdss j150748.33 + 551108.6 & & 48 & @xmath1318.9 & 0.4 & 20.8 & s & sbc & n + sdss j231736.39 + 140004.3 & ngc 7580 , mrk 318 & 63 & @xmath1319.3 & 0.6 & 20.4 & s ? & sbc & n + lllll mrk 297 & 4739 @xmath34 17 & 68 @xmath34 0.2 & 362 @xmath34 17 & 6.5 @xmath34 0.08 + mrk 314 & 2081 @xmath34 17 & 30 @xmath34 0.2 & 188 @xmath34 17 & 12 @xmath34 0.1 + mrk 325 & 3427 @xmath34 17 & 49 @xmath34 0.2 & 202 @xmath34 17 & 11 @xmath34 0.2 + mrk 538 & 2798 @xmath34 17 & 40 @xmath34 0.2 & 240 @xmath34 17 & 20 @xmath34 0.2 + sdssj0119 + 1452 & 4098 @xmath34 17 & 59 @xmath34 0.2 & 266 @xmath34 17 & 2.4 @xmath34 0.1 + sdssj0218@xmath130757 & & & & 0.41 @xmath34 0.07 + sdssj0222@xmath130830 & & & & 0.61 @xmath34 0.09 + sdssj0728 + 3532 & 3953 @xmath34 17 & 56 @xmath34 0.2 & 216 @xmath34 17 & 7.9 @xmath34 0.1 + sdssj0834 + 0139 & 4215 @xmath34 160 & 60 @xmath34 2 & 329 @xmath34 160 & 6.9 @xmath34 0.2 + sdssj0904 + 5136 & 4782 @xmath34 103 & 68 @xmath34 2 & 204 @xmath34 103 & 4.5 @xmath34 0.2 + sdssj0911 + 4636 & 4281 @xmath34 17 & 61 @xmath34 0.2 & 151 @xmath34 17 & 2.0 @xmath34 0.1 + sdssj0934 + 0014 & 4860 @xmath34 17 & 69 @xmath34 0.2 & 332 @xmath34 17 & 4.8 @xmath34 0.2 + sdssj0936 + 0106 & 4920 @xmath34 17 & 70 @xmath34 0.2 & 255 @xmath34 17 & 3.3 @xmath34 0.09 + sdssj0943@xmath130215 & 4823 @xmath34 17 & 69 @xmath34 0.2 & 230 @xmath34 17 & 2.6 @xmath34 0.06 + sdssj1118 + 6316 & 3218 @xmath34 17 & 46 @xmath34 0.2 & 152 @xmath34 17 & 2.2 @xmath34 0.07 + sdssj1234 + 0319 & 4685 @xmath34 17 & 67 @xmath34 0.2 & 269 @xmath34 17 & 4.4 @xmath34 0.1 + sdssj1319 + 5203 & 4619 @xmath34 17 & 66 @xmath34 0.2 & 202 @xmath34 17 & 7.7 @xmath34 0.09 + sdssj1402 + 0955 & 4267 @xmath34 17 & 61 @xmath34 0.2 & 343 @xmath34 17 & 6.8 @xmath34 0.2 + sdssj1507 + 5511 & 3373 @xmath34 17 & 48 @xmath34 0.2 & 126 @xmath34 17 & 3.7 @xmath34 0.1 + sdssj2317 + 1400 & 4413 @xmath34 17 & 63 @xmath34 0.2 & 254 @xmath34 17 & 6.2 @xmath34 0.1 + llccccccc mrk 297 & 7.0 @xmath34 0.1 & 42 & 486 & 243 & 8.0 & 11 & 2.0 & 2.7 + mrk 314 & 2.5 @xmath34 0.05 & 65 & 170 & 85 & 3.8 & 0.63 & 0.94 & 0.16 + mrk 325 & 6.3 @xmath34 0.1 & 43 & 242 & 121 & 9.6 & 3.3 & 2.4 & 0.83 + mrk 538 & 7.6 @xmath34 0.1 & 50 & 264 & 132 & 11 & 4.5 & 2.8 & 1.1 + sdssj0119 + 1452 & 2.0 @xmath34 0.1 & 84 & 230 & 115 & 5.6 & 1.7 & 1.6 & 0.48 + sdssj0218@xmath130757 & 0.47 @xmath34 0.08 & 52 & & & 4.9 & & 1.3 & + sdssj0222@xmath130830 & 0.65 @xmath34 0.1 & 48 & & & 4.2 & & 1.2 & + sdssj0728 + 3532 & 6.0 @xmath34 0.1 & 34 & 321 & 160 & 5.4 & 3.3 & 1.2 & 0.74 + sdssj0834 + 0139 & 5.9 @xmath34 0.5 & 83 & 293 & 147 & 6.5 & 3.3 & 1.6 & 0.82 + sdssj0904 + 5136 & 4.9 @xmath34 0.3 & 34 & 301 & 150 & 7.4 & 3.9 & 2.0 & 1.0 + sdssj0911 + 4636 & 1.7 @xmath34 0.1 & 33 & 221 & 111 & 5.6 & 1.6 & 1.2 & 0.35 + sdssj0934 + 0014 & 5.4 @xmath34 0.2 & 51 & 378 & 189 & 6.8 & 5.7 & 1.9 & 1.5 + sdssj0936 + 0106 & 3.8 @xmath34 0.1 & 48 & 293 & 146 & 7.1 & 3.5 & 2.0 & 1.0 + sdssj0943@xmath130215 & 2.9 @xmath34 0.07 & 84 & 194 & 97 & 4.9 & 1.1 & 1.8 & 0.39 + sdssj1118 + 6316 & 1.1 @xmath34 0.04 & 82 & 123 & 61 & 3.1 & 0.27 & 0.91 & 0.08 + sdssj1234 + 0319 & 4.7 @xmath34 0.1 & 52 & 293 & 147 & 6.7 & 3.4 & 2.1 & 1.1 + sdssj1319 + 5203 & 7.9 @xmath34 0.1 & 44 & 239 & 120 & 5.3 & 1.8 & 1.2 & 0.41 + sdssj1402 + 0955 & 6.0 @xmath34 0.2 & 55 & 372 & 186 & 8.7 & 7.0 & 1.6 & 1.3 + sdssj1507 + 5511 & 2.0 @xmath34 0.07 & 44 & 144 & 72 & 7.9 & 0.95 & 1.8 & 0.21 + sdssj2317 + 1400 & 5.8 @xmath34 0.1 & 31 & 420 & 210 & 7.3 & 7.5 & 1.7 & 1.8 + lccccc mrk 297 & 0.03 & 39 & 0.2 & 3 & 0.7 + mrk 314 & 0.2 & 3.9 & 0.7 & 2 & 0.4 + mrk 325 & 0.1 & 16 & 0.4 & 2 & 0.5 + mrk 538 & 0.08 & 17 & 0.5 & 3 & 0.7 + sdssj0119 + 1452 & 0.06 & 5.6 & 0.4 & 3 & 0.9 + sdssj0218 - 0757 & & 5.2 & 0.09 & & + sdssj0222 - 0830 & & 4.3 & 0.2 & & + sdssj0728 + 3532 & 0.09 & 5.6 & 1 & 6 & 1 + sdssj0834 + 0139 & 0.1 & 6.8 & 0.9 & 5 & 1 + sdssj0904 + 5136 & 0.07 & 12 & 0.4 & 3 & 0.9 + sdssj0911 + 4636 & 0.05 & 8.2 & 0.2 & 2 & 0.4 + sdssj0934 + 0014 & 0.05 & 13 & 0.4 & 4 & 1 + sdssj0936 + 0106 & 0.06 & 6.8 & 0.6 & 5 & 2 + sdssj0943 - 0215 & 0.1 & 8.2 & 0.4 & 1 & 0.5 + sdssj1118 + 6316 & 0.2 & 4.3 & 0.3 & 0.6 & 0.2 + sdssj1234 + 0319 & 0.07 & 7.4 & 0.6 & 5 & 1 + sdssj1319 + 5203 & 0.2 & 4.7 & 2 & 4 & 0.9 + sdssj1402 + 0955 & 0.04 & 12 & 0.5 & 6 & 1 + sdssj1507 + 5511 & 0.1 & 5.6 & 0.4 & 2 & 0.4 + sdssj2317 + 1400 & 0.04 & 8.2 & 0.7 & 9 & 2 + . the spectra have been smoothed to a resolution of @xmath112 km s@xmath17 and only the central 800 km s@xmath17 are shown . the dashed lines indicate the 20% crossings used to measure the line width at 20% . the triangles indicate the recessional velocities calculated from sdss redshifts for the sdss galaxies ; velocities from ned redshifts are shown for the non - sdss galaxies . those galaxies with other sources at similar velocities within the gbt beam are indicated by a star in the upper right corner . [ fig4 ] ] ( r @xmath59r@xmath10 ) , for local lcbgs , as measured from h observations , is shown as a gray histogram . for comparison , the range of dynamical masses for local hubble type galaxies are indicated with black boxes . the `` m '' in `` sm '' and `` i m '' indicates magellanic or low - luminosity spirals and irregulars . note that in order to compare lcbg dynamical masses with @xcite s results for local hubble types , the dynamical masses plotted here do not include a line width correction for random motions . [ fig5 ] ] ) to l@xmath21 ratios for local lcbgs is shown as a gray histogram . for comparison , the range of mass - to - light ratios for local hubble type galaxies ( s0a to i m ) @xcite are shown . as in figure 5 , for accurate comparisons , these mass - to - light ratios do not include a line width correction for random motions . many lcbgs tend to have mass - to - light ratios smaller than local normal galaxies , consistent with findings at intermediate redshifts . note that we may have overestimated the dynamical masses for most of the lcbgs with higher mass - to - light ratios . [ fig6 ] ] ) versus absolute blue magnitude ( m@xmath21 ) . local lcbgs are indicated by filled circles . those galaxies with other sources at similar velocities within the gbt beam are circled ; their line widths may be overestimated . the solid line indicates the tully - fisher relationship from @xcite ; the dotted lines indicate their 1 @xmath38 scatter of 0.3 m@xmath21 . the location of our local lcbgs is consistent with the tully - fisher relationship , but with a higher 1 @xmath38 scatter of 0.9 m@xmath21 . [ fig7 ] ] ) versus the the velocity dispersion ( @xmath38 ) is plotted for our local lcbgs ( filled circles ) and intermediate redshift ( 0.4 @xmath0 z @xmath0 1 ) lcbgs ( open circles ) @xcite . the regions occupied by the bulk of other local galaxy types@xmath13ellipticals ( e ) , spirals ( s ) , magellanic spirals ( sm ) , dwarf ellipticals ( de ) , and irregulars ( irr)@xmath13are indicated ( guzmn et al . 1996 , hyperleda ) . we have also indicated the approximate locations for some representative galaxies : draco , ngc 205 and m 31 ( hyperleda ) . local lcbgs are consistent with higher mass irregulars and dwarf ellipticals , and lower mass or magellanic spirals , as well as intermediate redshift lcbgs .

we present single - dish h spectra obtained with the green bank telescope , along with optical photometric properties from the sloan digital sky survey , of 20 nearby ( d 70 mpc ) luminous compact blue galaxies ( lcbgs ) . thesel , blue , high surface brightness , starbursting galaxies were selected with the same criteria used to define lcbgs at higher redshifts . however , at least half have dynamical mass - to - light ratios smaller than nearby galaxies of all hubble types , as found for lcbgs at intermediate redshifts . by comparing line widths and effective radii with local galaxy populations , we find that lcbgs are consistent with the dynamical mass properties of magellanic ( low luminosity ) spirals , and the more massive irregulars and dwarf ellipticals , such as ngc 205 .

we present single - dish h spectra obtained with the green bank telescope , along with optical photometric properties from the sloan digital sky survey , of 20 nearby ( d 70 mpc ) luminous compact blue galaxies ( lcbgs ) . thesel , blue , high surface brightness , starbursting galaxies were selected with the same criteria used to define lcbgs at higher redshifts . we find these galaxies are gas - rich , with m ranging from 5 to 8 m , and m l ranging from 0.2 to 2 m l , consistent with a variety of morphological types of galaxies . we find the dynamical masses ( measured within r ) span a wide range , from 3 to 1 m . however , at least half have dynamical mass - to - light ratios smaller than nearby galaxies of all hubble types , as found for lcbgs at intermediate redshifts . by comparing line widths and effective radii with local galaxy populations , we find that lcbgs are consistent with the dynamical mass properties of magellanic ( low luminosity ) spirals , and the more massive irregulars and dwarf ellipticals , such as ngc 205 .

astro-ph9606111

c

we have found disk and boundary layer solutions with significant departures from keplerian rotation that can reproduce the key spectral features of fu orionis and v1057 cygni . however , observations of these objects were originally modelled with reasonable success using standard keplerian disk theory ( hartmann & kenyon 1985 ; khh ) . here we discuss why it is difficult to compare the relative merits of the two types of solutions on the basis of observations alone . we then outline why we think the boundary layer solutions are nevertheless an important improvement in treating fu orionis accretion physics . both standard thin disk solutions and disk and boundary layer solutions agree quite well with observations . thus , comparisons with observations do not provide a definitive test of the relative merits of the two types of solutions for fu orionis objects . the disk and boundary layer solutions appear to match the differential rotation with wavelength somewhat better than the thin disk solutions of khh ; the thermal pressure support in the inner regions of the disk tends to reduce the ratio of the optical line widths to the near - infrared line widths , as observed . however , as pointed out by khh , there are significant uncertainties in strengths of the infrared co lines in the outer disk due to the possibility of dust formation , and with plausible models of dust contributions a thin disk in keplerian rotation can match the observed differential rotation with wavelength . similarly , differences in how well the two types of solutions match the observed spectral energy distributions , or small differences in the inferred masses , radii , and accretion rates , are not significant in view of the many uncertainties involved . specifically , there are uncertainties in the inclination angles @xmath26 , distances @xmath27 , and extinctions @xmath128 of fu orionis and v1057 cygni . we have used fairly standard inclinations of @xmath136 for fu orionis and @xmath175 for v1057 cygni . these inclinations allow solutions with similar rotational velocities to produce the rather different linewidths seen in these two systems . in fact , the ratio @xmath176 is almost identical to the ratio of the 6170 linewidths ( 57 and @xmath41 ) in these objects . nonetheless , other inclinations , in combination with different choices of the other solution parameters , might produce solutions which match the data . the distances to fu orionis and v1057 cygni are probably only accurate to 1020% , so that the luminosities inferred for these systems could be in error by 2040% . finally , the extinctions of these systems are not well known , and could represent an additional source of error in both the luminosities and the shape of the spectra . all of these considerations limit our ability to distinguish between different disk models . nonetheless , we feel that the disk and boundary layer model presented in this paper offers significant advantages over the standard thin disk model . first , it avoids some uncomfortable assumptions which are implicit in the thin disk model . the two major observational constraints on models of the inner disk in fu orionis systems are the apparent lack of hot boundary layer emission and the lack of slowly - rotating material in the observed line profiles . the absence of these features gives the impression that the boundary layer is not present is fu orionis systems . the standard thin disk model ( hartmann & kenyon 1985 ; khh ) simply did not include any boundary layer region . by omitting the boundary layer , this model made the implicit assumption that the accreting star is rotating at breakup speed . rapid stellar rotation seems unlikely in view of the generally slow rotations of t tauri stars , plus the recognition that fu orionis and v1057 cygni have not accreted enough angular momentum during their outbursts to spin up the entire central star as a solid body ( hartmann & kenyon 1996 ) . this leaves open the possibility that only the outer layers are spun up to keplerian angular velocities . a second implicit assumption of the early model is that the disk can be keplerian all the way in to the stellar surface , which implies that pressure gradients in the inner disk are insignificant . this also seems unlikely ; at the high mass accretion rates found in fu orionis systems , radial pressure gradients will support the accreting material and produce sub - keplerian angular velocities unless @xmath0 is very large . our disk and boundary layer model avoids these assumptions . we have found that by using the slim disk model , we can include the boundary layer region in our calculations of fu orionis disks and still find solutions which agree with the observations . these solutions demonstrate explicitly what the thin disk model implicitly assumed : fu orionis disks do not have a standard accretion disk boundary layer where the angular velocity of the accreting material drops rapidly over a short radial distance . unless @xmath0 is very large , the high accretion rates of fu orionis objects require quite optically thick disks which are very hot in their interiors , resulting in signficant thermal pressure support . thus the narrow , hot boundary layer typically expected in pre - main - sequence stars ( lynden - bell & pringle 1974 ) automatically disappears , for a wide range of stellar rotation rates . instead , the `` boundary layer '' is a broad region where @xmath68 changes gradually ; in our best - fitting solutions , @xmath68 increases gradually and becomes nearly constant as the accreting material approaches the star . as a result , like the early thin disk solutions , our solutions lack the high - temperature , slowly - rotating region expected from a standard boundary layer . but unlike the early solutions , they do not require that the star be rotating at breakup or that @xmath0 be very large . in fact , for @xmath63 , our best solutions have the star rotating at only @xmath177 of breakup speed . 5 shows that as @xmath0 increases , @xmath37 increases , reaching @xmath178 of breakup speed at @xmath179 . this suggests that the keplerian thin disk solutions of hartmann & kenyon ( 1985 ) and khh represent a special , extreme case of the solutions found in this paper , where the values of both @xmath0 and @xmath37 are very large . in addition to eliminating the assumptions of very large values of @xmath0 and @xmath37 which are implicit in the thin disk model , our disk and boundary layer model can offer insight into new areas of fu orionis accretion physics . the model explicitly includes the effects of the rotation rate of the star on the boundary layer region and on the angular momentum accretion rate . since radial advection of energy is included in the slim disk equations , our solutions directly give the rate at which energy is carried into the accreting star . these issues can only be studied when the boundary layer region and the additional physics of the slim disk equations are included in the model . they have important consequences for the evolution of pre - main - sequence stars , which will be discussed in more detail in the following sections . our best - fit solutions for fu orionis and v1057 cygni are equilibrium solutions with @xmath84 , i.e. the central star is neither gaining nor losing angular momentum , and have stellar rotation rates @xmath180 which correspond to rotation periods of 89 days . these periods are comparable to those observed in t tauri stars ( bouvier 1993 , 1995 ) . these low-@xmath78 solutions have important implications for the spin evolution of pre - main - sequence stars , because they support the scenario proposed in paper i , where fu orionis outbursts regulate the rotation rates of t tauri stars . in paper i we showed that the angular momentum accretion rate @xmath78 drops rapidly as the stellar rotation rate @xmath37 increases , and reaches @xmath84 when @xmath181 for fu orionis parameters . for smaller values of @xmath37 , @xmath182 and the star spins up , while for larger values of @xmath37 , @xmath183 and the star spins down . thus , fu orionis outbursts will move @xmath37 toward an equilibrium value where @xmath105 . if fu orionis outbursts dominate mass accretion onto t tauri stars , as seems likely from event statistics , then they should also dominate angular momentum accretion and control the spin evolution of these stars . when the t tauri star is between outbursts , it may spin up or down , but during each outburst it will return to the equilibrium rotation rate . this scenario therefore predicts that fu orionis systems should attain a low-@xmath78 equilibrium state during outbursts , and we find that low-@xmath78 solutions fit the observations . we are unable to make more precise predictions of the spin rates of t tauri stars because of uncertainties in the structure of the central star undergoing fu orionis accretion . as discussed in 5.4 below , it is likely that the advection of large amounts of thermal energy through the inner disk will cause the central star to expand from its equilibrium state . after an outburst ceases , we expect that the star will contract to a smaller radius , and will therefore spin up somewhat in the absence of any angular momentum loss . note that the disk may still remove large amounts of angular momentum from the star and spin the star down substantially during the course of the outburst ; however , at the end of the outburst , we expect that there will always be some spinup due to the contraction of the star . thus the equilibrium rotation rates during the fu orionis phase should correspond to faster t tauri rotation rates . the amount of spinup will depend upon the moment of inertia of the expanded layers of the star , which we are not able to estimate at present . observations of t tauri stars suggest that classical t tauri stars which have accretion disks rotate more slowly than weak - line t tauri stars which lack disks ( bouvier 1993 , 1995 ; edwards 1993 ; eaton , herbst , & hillenbrand 1995 ; choi & herbst 1996 ) . our results , combined with those of paper i , suggest a new way in which the disk can regulate the rotation rate of an t tauri star . previously suggested methods of regulating the stellar rotation rate have relied upon the interaction between the stellar magnetic field and the disk ( knigl 1991 ; cameron & campbell 1993 ; hartmann 1994 ; shu 1994 ) . our method does not depend on the stellar magnetic field , which is presumably not strong enough to disrupt the disk during fu orionis outbursts . however , it does require that outbursts occur , and that they dominate the angular momentum accretion onto the star . we noted in 4.3 that our solutions for fu orionis and v1057 cygni produce luminosities which are substantially smaller than the accretion luminosity . for fu orionis , we found that our solution gave @xmath184 , whereas @xmath185 , so that @xmath186 . for our v1057 cygni solution , @xmath187 and @xmath188 , so we again have @xmath189 . the remaining luminosity goes into heating the disk material . at the stellar surface , the midplane temperature of the disk in our solutions becomes quite large : @xmath190 k in the fu orionis solution , and @xmath191 k in the v1057 cygni solution . this means that the disk material carries energy into the star at the rate @xmath192 for fu orionis and @xmath193 for v1057 cygni . the high temperature of the disk material is due to the large vertical optical depth of the disk and boundary layer in these solutions . if the accreting stars in fu orionis objects are typical pre - main - sequence stars , they should have luminosities @xmath194 , so this advected energy represents a major perturbation to the star . another interesting feature of these solutions is that only a fraction of the luminosity comes from energy liberated by viscous dissipation . popham & narayan ( 1995 ) derived an expression for the total viscous dissipation rate of the disk and boundary layer @xmath195 where we have left out several small terms . the first term in this expression represents the gravitational potential energy and the rotational kinetic energy transferred between the star and the disk . for @xmath84 , the disk is removing a small amount of rotational kinetic energy from the star . both of our solutions have @xmath196 , so the first term is just @xmath197 . in a thin disk , the second term is very small , but in our fu orionis solutions , this term becomes quite large due to the importance of the radial pressure gradient in supporting the accreting material against gravity . in both solutions , we find @xmath198 , so that the viscous dissipation rate is only @xmath199 . nonetheless , we know that both solutions radiate @xmath200 ; the remaining @xmath201 comes from the entropy term in the energy equation @xmath202 the four terms in this equation represent the viscous dissipation , the vertical flux from the disk surface , the advected entropy , and the divergence of the radial flux . if we integrate this equation over the surface of the disk , we have @xmath203 = 0.\ ] ] our boundary conditions assume that the radial fluxes at the inner and outer edges of the disk are small ; at the inner edge we assume that the flux entering the disk is @xmath204 , where we assume @xmath205 k , so @xmath206 , and at the outer edge the radial flux is insignificant . if we neglect these radial flux terms , we can use @xmath207 and the expression for @xmath208 given above to write @xmath209 @xmath210 @xmath211 where the internal energy @xmath212 , @xmath213 , and @xmath214 . this accounts for the difference between the actual luminosities radiated by our solutions and the accretion luminosities for those disk parameters . our best - fit solutions for fu orionis and v1057 cygni have stellar radii of @xmath215 . these are substantially larger than the radii of t tauri stars , which are generally estimated at @xmath216 ( bouvier 1995 ) . since fu orionis outbursts are believed to occur in t tauri star accretion disks , this seems to imply that the star expands rapidly during the outburst , probably as a result of the rapid addition of high - temperature material ( khh ; hartmann , cassen , & kenyon 1996 ) . prialnik & livio ( 1985 ) calculated the effects of accretion onto a @xmath217 , @xmath218 fully - convective main - sequence star . they found that if the accretion carries a sufficient amount of energy into the star , it can cause the star to expand stably or unstably . for accretion rates comparable to those of fu orionis stars , their calculations indicate that the expansion will proceed unstably if the accretion energy is added to the star at a rate faster than about 10% of the accretion luminosity . the rapid expansion is due to the conversion of the normally convective star to a convectively - stable , radiative structure . as the discussion of the luminosities shows , our solutions add energy to the accreting star at an enormous rate - @xmath219 in our solutions for fu orionis and v1057 cygni . this rate of energy transfer will almost certainly produce rapid expansion of the star . the character of this expansion is not well understood . the accreting material comes from a disk , so it carries angular momentum as well as energy , and is added around the star s equator . thus the expansion is likely to be nonspherical . more sophisticated models of this process are needed , but the problem is clearly a difficult one . the large energy input into the star and the resulting stellar expansion should have important implications for the decline from fu orionis outbursts . the star will contract and release the energy it gained during the outburst , which is a substantial fraction of the total accretion energy of the outburst . thus we predict , qualitatively , that at the end of the mass accretion outburst in the disk , there may be a phase in which the central star is overly luminous , and fades over some ( rapid ) time set by the kelvin time of the perturbed portion of the stellar envelope , similar to the original explanation put forth by larson ( 1983 ) for fu orionis outbursts . by including the boundary layer region and the disk structure self - consistently , the solutions presented here represent a substantial advance in our understanding of disks in fu orionis objects and the spectra and line profiles they produce . nonetheless , there are additional explorations which could be made within the context of our assumptions , and improvements which could be made to our model . the solutions we have found do a good job of fitting the observed line profiles and spectra of fu orionis and v1057 cygni . nevertheless , this does not mean that they are the only solutions which would fit the observations of these systems . the number of solution parameters is too great to permit a full exploration of the entire parameter space . accordingly , as described in 4.2 , we have used the standard parameters @xmath59 , @xmath51 , and @xmath52 to guide our search . we have also kept certain parameters constant in order to simplify the fitting procedure . we keep @xmath84 for the reasons discussed in 3 , where we demonstrated that high-@xmath78 solutions fail to produce spectra and line profiles that agree with observations . other values of @xmath78 could produce reasonable solutions ; for instance , in paper i we proposed that fu orionis outbursts may spin down the accreting star , which would require solutions with negative values of @xmath78 . negative-@xmath78 solutions are similar in all respects to @xmath84 solutions ; the boundary layer region has @xmath68 increasing inward throughout and lacks a strong peak in effective temperature . as we demonstrated in paper i , @xmath78 drops rapidly with increasing @xmath37 , so that a @xmath85 solution has @xmath37 only slightly larger than a @xmath84 solution with the same parameters ( see fig . 2 ) . we also use @xmath63 for all of our solutions . as discussed in 3.3 , solutions with smaller values of @xmath0 have thicker disks , so that a larger fraction of the accretion luminosity is advected into the star . thus a smaller fraction is radiated , and the effective temperature is lower . another consequence of changing @xmath0 is that the rotational velocities change ; small values of @xmath0 result in small rotational velocities which would produce centrally peaked line profiles unlike those observed . this suggests that values of @xmath220 may produce acceptable solutions , but those with @xmath0 substantially below @xmath119 may not . this is particularly interesting because bell & lin ( 1994 ) and bell ( 1995 ) found that they needed @xmath221 to obtain satisfactory outbursts using the thermal instability mechanism . our model spectra could also be improved by treating the effects of inclination in more detail . we have implicitly assumed that the disk surface is flat , so that the inclination angle for all parts of the disk surface is the same . in fact , most of our solutions have @xmath222 nearly constant with @xmath223 , with @xmath17 . if we take @xmath224 as representing the surface of the disk , then the disk surface is inclined at @xmath225 to the disk midplane . this means that at inclination angles greater than @xmath226 , a portion of the disk surface will not be visible . even if @xmath227 , the side of the disk closer to the observer will effectively be seen at a larger inclination angle than the opposite side . ( note that when we refer to `` sides '' of the disk here , we are referring not to the top and bottom surfaces of the disk , but to regions of the same surface at different angles around the rotation axis . ) also , one side of the disk will be heated by radiation from the other side . we have also assumed that the disk spectrum does not vary with inclination angle , whereas in fact limb darkening will generally produce a decline in the blue end in the spectrum for systems viewed at large inclination angles . finally , there are improvements which could be made in the physical treatment of the disk and boundary layer . the assumption of a steady disk might be relaxed , although we feel this is unlikely to be a major issue , even though we are comparing our solutions to objects experiencing outbursts , because fu orionis and v1057 cygni have both remained fairly steady over recent years ( khh ; kenyon & hartmann 1991 ) , and we are only addressing the innermost regions of the disk . time - dependent models of boundary layers in disks around pre - main - sequence stars ( godon 1996 ) also seem to agree with our steady - state results ( pnhk ) . calculations of the structure of advection - dominated accretion disks ( narayan & yi 1995 ) have demonstrated that the slim disk equations used in our model provide a fairly good representation of the disk structure even for disks which are quite vertically thick ; still , a two - dimensional model for the boundary layer and disk might provide insights into fu orionis objects that we are unable to make using our current model ( e.g. , kley 1991 ) , especially when considering a more subtle matching of the disk to the inherently two - dimensional star . some consideration should also be given to understanding the effects of rapid accretion on the central star . presumably these include rapid expansion of the star , which could have important effects on the boundary layer region and the angular momentum accretion by the star . ultimately this will have to be a two - dimensional , time - dependent calculation as well , but perhaps some progress can be made with calculations similar to those performed by prialnik & livio ( 1985 ) for main - sequence convective stars .

we present solutions for the accretion disks and boundary layers in pre - main - sequence stars undergoing fu orionis outbursts . they also have modest stellar rotation rates , comparable to the observed rotation rates of t tauri stars , and angular momentum accretion rates of zero . this supports our earlier suggestion that fu orionis outbursts may regulate the rotation rates of t tauri stars .

we present solutions for the accretion disks and boundary layers in pre - main - sequence stars undergoing fu orionis outbursts . these solutions differ from earlier disk solutions in that they include a self - consistent treatment of the boundary layer region . in a previous paper ( popham 1996 ) , we showed that these stars should stop accreting angular momentum once they spin up to modest rotation rates . here we show that for reasonable values of , these low angular momentum accretion rate solutions fit the spectra and line profiles observed in fu orionis objects better than solutions with high rates of angular momentum accretion . we find solutions which fit the observations of fu orionis and v1057 cygni . these solutions have mass accretion rates of 2 and , stellar masses of 0.7 and , and stellar radii of 5.75 and , respectively . they also have modest stellar rotation rates , comparable to the observed rotation rates of t tauri stars , and angular momentum accretion rates of zero . this supports our earlier suggestion that fu orionis outbursts may regulate the rotation rates of t tauri stars . 2i_k^2(r_in ) 2o_k^2(r_out ) 2s_k^2(r _ _ * )

astro-ph0506434

i

studying the formation of our own solar system and observing the frequency of similar systems associated with other stars are two ways in which we seek to understand our origins . through remote observation and direct exploration , we have developed a much clearer understanding of our solar system . however many questions about the processes involved in the initial formation and subsequent evolution towards the present configuration can not be addressed directly . therefore , we need to study other stars to help place our solar system in context . there are two major zones of debris in the solar system : the asteroid belt at 2 @xmath4 4 au composed of rocky material that is ground up by collisions to produce most of the zodiacal dust cloud and the kuiper belt ( kb ) that consists of small bodies orbiting beyond neptune s orbit at 30 @xmath4 50 au . since their discovery over a decade ago ( jewitt & luu 1993 ) , kuiper belt objects have played an increasingly important role in understanding the formation and evolution of our planetary system ( e.g. malhotra 1993 ; kenyon & bromley 2004 ) . while direct detection of _ in situ _ debris from collisions amongst kuiper belt objects has yet to be confirmed , of order 10% of the solar system s far - ir luminosity could be emitted by kuiper belt dust ( e.g. backman et al . searching for kb - like debris around other sun - like stars as a function of age will determine both the frequency of such systems and provide important insight into the formation of our solar system . most debris disks were found by their mid - to - far infrared emission in excess of the expected photosphere ( e.g. , auman et al . 1984 ; backman & paresce 1993 ; mannings & barlow 1998 ; habing et al . 1999 ; silverstone 2000 ; spangler et al . 2001 , decin et al . 2003 , zuckerman & song 2004 ) . the majority of these debris disks are also associated with hot , luminous stars , since observatories such as iras and iso did not have the necessary sensitivities to detect debris disks around lower - luminosity solar - type stars at distances beyond a few parsecs . the increased sensitivity afforded by the _ spitzer space telescope _ ( _ spitzer _ ; werner et al . 2004 ) has the potential to identify and investigate debris systems that were not detectable with previous observatories . using data from the formation and evolution of planetary systems ( feps ) _ spitzer _ legacy program , we have conducted a search for outer dust disks dominated by temperatures characteristic of the kuiper belt ( @xmath8 k ) . the earliest results from our validation observations are presented in meyer et al . ( 2004 ) , where we identified such debris disks ( exo - kbs ) surrounding the 30 myr - old sun - like star hd 105 and the 0.4 - 1 gyr old star hd 150706 . herein we present five more sun - like stars that exhibit characteristics of exo - kbs , and discuss their properties in the context of the evolution of our own solar system . we also present one star that has a firm detection at 70 @xmath0 m , but the measurement is consistent with the photospheric emission from the star . in this case , we place unprecedentedly low upper limits on the presence of kb - like debris in this system . in 2 , we describe the observations and data reduction . in 3 , we discuss the methodology for identifying the five exo - kuiper belt candidate stars from the feps sample , and we present the resulting spectral energy distributions ( seds ) . we outline our interpretation of the seds in terms of physical models in 4 . we discuss the implications of our results in 5 , and summarize our findings in

we present the discovery of debris systems around three solar mass stars based upon observations performed with the _ spitzer space telescope _ as part of a legacy science program , `` the formation and evolution of planetary systems '' ( feps ) . we also confirm the presence of debris around two other stars . we speculate on the nature of these systems through comparisons to our own kuiper belt , and on the likely planet(s ) responsible for stirring the system and ultimately releasing dust through collisions . the observations place strong upper limits on the presence of any cold dust in this nearby system ( ) .

we present the discovery of debris systems around three solar mass stars based upon observations performed with the _ spitzer space telescope _ as part of a legacy science program , `` the formation and evolution of planetary systems '' ( feps ) . we also confirm the presence of debris around two other stars . all the stars exhibit infrared emission in excess of the expected photospheres in the 70 m band , but are consistent with photospheric emission at m . this restricts the maximum temperature of debris in equilibrium with the stellar radiation to k. we find that these sources are relatively old in the feps sample , in the age range 0.7 3 gyr . based on models of the spectral energy distributions , we suggest that these debris systems represent materials generated by collisions of planetesimal belts . we speculate on the nature of these systems through comparisons to our own kuiper belt , and on the likely planet(s ) responsible for stirring the system and ultimately releasing dust through collisions . we further report observations of a nearby star hd 13974 ( 11 pc ) that is indistinguishable from a bare photosphere at both 24 m and 70 m . the observations place strong upper limits on the presence of any cold dust in this nearby system ( ) .