{
    "0004/astro-ph0004337_arXiv.txt": {
        "abstract": "Abundances of $\\alpha$-elements such as Ca and Mg in disk and halo stars are usually derived from equivalent widths lines measured on high resolution spectra, and assuming Local Thermodynamic Equilibrium (LTE) . In this paper, we present non-LTE differential abundances derived by computing the statistical equilibrium of CaI and MgI atoms, using high resolution equivalent widths available in the literature for 252 dwarf to subgiant stars. These non-LTE abundances combined with recent determination of non-LTE abundances of iron, seem to remove the dispersion of the [Ca/Fe] and [Mg/Fe] ratios in the galactic halo and disk phases, revealing new and surprising structures. These results have important consequences for chemical evolution models of the Galaxy. In addition, non-LTE abundance ratios for stars belonging to the M92 cluster apparently have the same behavior. More high resolution observations, mainly of globular clusters, are urgently needed to confirm our results. ",
        "introduction": "The determination of abundances of nuclear species at distinct locations in the Galaxy (e.g. halo, disk and bulge) comes mainly from the spectra of late-type star atmospheres. Measured abundances in cool stars at different stages of evolution give not only the understanding of stellar nucleosynthesis, but also provide valuable information about the process of chemical enrichment of the Galaxy.  The archaeological tracers of the chemical evolution of a star system are the elements produced by explosive nucleosynthesis in type II (SNII) and type Ia (SNIa) supernovae events. The interest using such elements as tracers rests on the fact that SNII and SNIa progenitors have different lifetimes; SNII is the final evolution of massive stars and SNIa is a possible final result of evolution of a close binary system of intermediate mass stars. SNII contribute to the enrichment of the interstellar medium (ISM) mainly with elements produced by the capture of $\\alpha$ particles ($\\alpha$ elements) and from the r-process, and SNIa produce elements belonging to the Fe peak. Consequently, the basic tools to constraint the evolution of ISM in the Galaxy are usually the analysis of relations between ratios of heavy elements [element/Fe] and Fe abundance [Fe/H].\\footnote{$\\rm [Fe/H]= log(N_{Fe}/N_H)-log(N_{Fe}/N_H)_{\\odot}$.}  A first glance at the temporal behavior of $\\alpha$ elements shows that the ratio [$\\alpha$/Fe] is approximately constant for halo metal-poor stars ([Fe/H]$\\leq$-1.5) and decreases for metal-rich stars ([Fe/H]$>$-1.5) belonging to the disk. This is reasonably explained by the chemical evolutionary models that assume progressive enrichment of ISM by supernovae: first generation of stars have in their atmospheres the signature of SNII events only (called halo-phase of the Galaxy) and the subsequent generations have a signature of both SNII and SNIa events (disk-phase).  However, a more precise analysis of [$\\alpha$/Fe] vs. [Fe/H] shows a pronounced scatter, mainly in the region of metal-poor stars. This scatter has been interpreted mostly as a consequence of the inhomogeneity of the matter having made stars rather than resulting of poor observational data (Audouze and Silk 1996).  The derivation of abundances based on the analysis of high resolution stellar spectra is usually made under the assumption of Local Thermodynamical Equilibrium (LTE). In the last 15 years, many efforts to estimate errors on abundance determinations caused by LTE assumption have been done. Recently results for Ba II (Gigas 1986, 1988, Mashonkina \\& Bikmaev 1996), Sr II (Belyakova \\& Mashonkina 1997), Na I (Mashonkina et al. 1993), Mg I (Gigas 1986, Mashonkina et al. 1996), Ca I (Drake 1991), B I (Kiselman \\& Carlsson 1996), Al I (Baumuller \\& Gehren 1997) and Fe I and Fe II (Th\\'evenin \\& Idiart 1999, TI99), O I (Mishenina et al. 1999) and Mg I (Zhao, Butler \\& Gehren 1998) demonstrate that most lines can formed far from LTE conditions. So, some important questions arise: what is the influence of non-LTE abundance calculations on the chemical evolution diagrams of the Galaxy? Do these computations add another different constraints to the chemical history of enrichment of the matter in the Galaxy?  In this work, we present non-LTE abundances derived from computation of statistical equilibrium of Ca and Mg atoms, using published equivalent widths (Sect. 3). The atomic data and stellar atmospheric models used are presented in Sect. 2. $\\alpha$ elements like Mg and Ca have well-known enhanced abundances in atmospheres of F-G metal-poor dwarf stars as a result of cumulative stellar generations. Recently, Nissen \\& Schuster(1997) and Jehin et al. (1999) proposed the existence of two sequences of stars having two different [$\\alpha$/Fe] ratios for intermediate stars ([Fe/H] $\\approx$ -1). Their works are based on highly accurate observations of stars having approximately same temperatures and surface gravities. Based on our non-LTE computations, we found different branches or sub-populations of stars not only for intermediate metallicities; consequences for chemical evolution models of the Galaxy are presented in Sect. 4. We drawn our conclusions in Sect. 5. ",
        "conclusions": "We report here non-LTE differential abundances for 252 subdwarf to subgiant stars using published high resolution equivalent widths. [Ca/Fe] and [Mg/Fe] diagrams show remarkable structures, both in the halo and disk phases of the Galaxy, which are not related with observational or atmospheric parameter uncertainties. These results lead us to a possible evolutive galactic scenario of non-homogeneity or incomplete mixing of synthesized SNII yields (Karlsson \\& Gustafsson, 1999, Argast et al. 1999). A surprising result is the behavior of M92 stars, mainly in the Ca diagram, suggesting a common origin for field and cluster stars. Spectroscopic high resolution with good S/N of stars in clusters are needed to confirm or not the sketch of a new chemical evolution model presented in this work. New non-LTE analysis for other $\\alpha$-elements is also necessary to verify if this structural behavior applies to all $\\alpha$-capture SNII products."
    },
    "0004/astro-ph0004047_arXiv.txt": {
        "abstract": "We present spectroscopic observations from the {\\sl Hubble Space Telescope} that reveal for the first time the presence of a broad pedestal of Balmer-line emission in the LINER galaxy NGC 4203.  The emission-line profile is suggestive of a relativistic accretion disk, and is reminiscent of double-peaked transient Balmer emission observed in a handful of other LINERs.  The very broad line emission thus constitutes clear qualitative evidence for a black hole, and spatially resolved narrow-line emission in NGC~4203 can be used to constrain its mass, $M_{BH }\\le 6\\times 10^6$ M$_\\odot$ at 99.7\\% confidence.  This value implies a ratio of black-hole mass to bulge mass $\\lesssim 7\\times 10^{-4}$ in NGC~4203, which is less by a factor of $\\sim 3 - 9$ than the mean ratio obtained for other galaxies.  The availability of an independent constraint on central black-hole mass makes NGC~4203 an important testbed for probing the physics of weak active galactic nuclei (AGNs).  Assuming $M_{BH }$ near the detection limit, the ratio of observed luminosity to the Eddington luminosity is $\\sim 10^{-4}$.  This value is consistent with advection-dominated accretion, and hence with scenarios in which an ion torus irradiates an outer accretion disk that produces the observed double-peaked line emission.  Follow-up observations will make it possible to improve the black-hole mass estimate and study variability in the nuclear emission. ",
        "introduction": "Spectroscopic surveys have revealed that a large fraction of nearby early-type galaxies harbor low-ionization nuclear emission-line regions, or LINERs (Heckman 1980\\markcite{Heckman80}; Ho et al. 1997a\\markcite{Ho97a} and references therein).  The physical understanding of these sources remains rudimentary, but there are strong indications that at least some LINERs are weak versions of the Seyfert or QSO phenomenon, and hence powered by accretion onto a massive black hole.  LINERs are nonetheless distinct from classical AGNs in terms of their characteristic (low) luminosity, emission-line properties, and broad-band spectral energy distribution (Ho 1999\\markcite{Ho99}).  These differences may follow from fundamental disparities in the accretion process operative in luminous and weak AGNs.  The study of emission-line behavior on small spatial scales within galaxy nuclei provides one strategy for probing the energetics, dynamics, and structure of LINERs and related objects.  Here we report on observations acquired with the {\\sl Hubble Space Telescope}\\/ ({\\sl HST}\\/) for the LINER NGC~4203.  The data provide dynamical constraints on a black hole, and reveal line emission that may directly trace an accretion flow in this source.  These observations and future follow-up studies will provide an important framework for testing physical models for the structure of LINERs, and the nature of black holes in galaxy nuclei. ",
        "conclusions": "NGC~4203 is the latest addition to a small set of LINERs that exhibit broad, double-shouldered Balmer lines.  The discovery of two such objects in our {\\sl HST} survey (NGC~4450 being the other; Ho et al. 2000a\\markcite{Ho_et00a}) suggests that such emission may be common in LINERs, but often eludes detection in ground-based apertures for which the contrast with the stellar continuum is weak.  The line profile in NGC~4203 is noteworthy for its resemblance to emission profiles seen in broad-line radio galaxies, which have been interpreted as the signature of a thin outer accretion disk irradiated by an inner ion torus, representing an advection-dominated flow onto a black hole.  The nucleus of NGC~4203 also exhibits spatially resolved emission that can be used to provide kinematic information on the underlying mass distribution.  The existing observations at a single PA can be used to restrict the underlying black-hole mass to $M_{BH} \\le 6 \\times 10^6$ M$_\\odot$.  The availability of a mass estimate for the black hole is of fundamental importance for studying the physics of the accretion process.  Our limiting value for $M_{BH}$ is consistent with a sub-Eddington accretion rate and formation of an ADAF, although a more stringent limit could challenge this picture.  Follow-up observations from space, or with small apertures under good seeing conditions on the ground, will make it possible to improve the estimate of the black-hole mass and study time evolution of the~broad~H$\\alpha$~emission."
    },
    "0004/astro-ph0004271_arXiv.txt": {
        "abstract": "We present the first results of a survey of blue horizontal branch (BHB) stars in the Galactic bulge. In this exploratory study, candidates with $15 \\leq V \\leq 17.5$ covering a wide range in \\bmv\\ color were selected from CTIO Schmidt $UBV$ photometry. Blue spectra were recorded at 2.4\\angs\\ FWHM resolution for 164 stars in a 1.3 sq.\\ dg.\\ field $\\sim$7.5\\dg\\ from the Galactic center. Radial velocities were measured for all stars.  For stars with strong Balmer lines, we devised and applied a spectroscopic technique to determine stellar temperature \\teff, gravity \\logg, and metallicity \\feh\\ independent of reddening. The reddening and distance to each star were then found from $UBV$ photometry.  Reddening proved highly variable, with \\ebmv\\ ranging from 0.0 to 0.55 around a mean of 0.28. The \\bmv\\ colors of cool HB stars of solar metallicity reddenened by \\ebmv\\ $\\geq 0.3$ overlap those of foreground main-sequence stars, but the \\umb\\ vs.\\ \\bmv\\ diagram distinguishes these groups until \\ebmv\\ $>$ 0.5.  Forty-seven BHB candidates were identified with \\teff\\ $\\geq$ 7250\\kel.  Seven have the gravities of \\popi\\ stars, three are ambiguous, and 37 are HB stars, including perhaps a dozen RR Lyraes. The unambiguous BHB stars are all cooler than 9000\\,K. They span a wide metallicity range, from solar to 1/300 solar.  The warmer BHB's are more metal-poor and loosely concentrated towards the Galactic center, while the cooler ones are of somewhat higher metallicity and are situated closer to the center. We detect two cool solar-metallicity HB stars in the bulge of our own Galaxy, the first such stars known. Still elusive are their fainter hot counterparts, the metal-rich sdB/O stars strong in ultraviolet light. ",
        "introduction": "According to classical calculations of single-star evolution \\citep[e.g.,][]{roo73}, the color of a star that has left the giant branch to become a core-helium-burning star on the horizontal branch (HB) depends primarily on its age and metallicity. At lower metallicities, a wide range in color is both found and predicted among BHB stars if modest mass loss is assumed. Among populations of solar metallicity or higher, only HB stars redder than the RR Lyrae instability strip should be produced within a Hubble time.  However, the most metal-rich globular clusters in the Galaxy, such as \\ngc 6388 and \\ngc 6441 near the Galactic center, do show a handful of BHB stars \\citep{ric97}. These are mostly cool or warm BHB's \\citep{moesc99}, with few if any of the hottest, faintest types, the subdwarf sdB stars and their relatively rare successors the sdO stars.  Surprisingly, significant numbers of sdB/O's are found in metal-rich populations in the Milky Way; even more surprisingly, they appear to outnumber cool and warm metal-rich BHB's.  In the field, where it is difficult to distinguish BHB's of intermediate temperature from \\popi\\ A stars, few if any metal-rich cool BHB's are known -- \\citet{gra96}, for example, uncovered only one possible BHB with \\feh\\ $\\ge$ $-0.7$, out of 39 BHB's newly identified. In contrast, there are dozens of field sdB/O stars whose gravities indicate HB status, and whose radial velocities indicate a thick disk rather than halo population \\citep{saf94,saf97}. Among open clusters \\citep[reviewed by][]{fri95}, the most striking occurrence of sdB/O stars is found in \\ngc 6791, with a metallicity 3 -- 4 times solar \\citep{pet98}. In addition to two cooler BHB (or blue straggler) members, the cluster harbors four or five sdB/O stars \\citep{lie94} whose membership is likely given their spatial concentration towards the cluster center \\citep{kal92}.  Hot, metal-rich BHB stars also appear to be present in metal-rich extragalactic systems. As reviewed by \\citet{oco99}, elliptical galaxies and early-type spiral bulges commonly show an upturn in integrated light below 2000\\angs, where old main-sequence turnoff stars have very little flux. Both the smooth spatial distribution and the continuous spectral distribution of the UV upturn indicate that it is not caused by young O and B stars, but rather by hot old stars, sdB/O's and the brighter but rarer blue post-AGB stars \\citep{dor95,bro97}. Among elliptical galaxies, the size of the UV upturn tends to increase with increasing galactic metallicity \\citep{fab83,bur88,lon89}, although the strongest correlation is with line indices based on light elements rather than iron itself, the scatter is large, and there is no continuity between these systems and globular clusters.  The presence of BHB stars in high-metallicity populations and the reversal in their color distribution raise major questions as to their origin. Both observation and theory suggest that BHB stars in metal-rich systems may be produced by channels in addition to those operating in metal-poor systems. Theoretical production of metal-rich sdB's from single-star evolution can be achieved with a rapid increase at  \\feh\\ $> 0$ in either helium abundance \\citep{bre94,yi98} or mass loss \\citep{dor95,dcr96}. It may also be accomplished with deep mixing, and thus accompanied by light-element enhancements \\citep{swe97,kra98}. However, binary mass transfer may play a dominant role at the hot end of the BHB, since a large fraction of field sdB/O stars are found to be binaries \\citep{all94,gre00}. Blue stragglers, thought also to have formed by mass transfer, may therefore be involved in the production of hot HB stars.  Constraints on these mechanisms might be placed by determining the color distribution and light-element ratios of metal-rich versus metal-poor BHB's, and by determining whether blue stragglers of similar metallicity are invariably present. The best place to do this is in the Milky Way bulge itself. The bulge is the only Milky Way population sizable enough to support large numbers of BHB stars. Unlike the centers of external galaxies, it is sufficiently nearby that stars as faint as the sdB's may be resolved individually. It resembles elliptical galaxies and spiral bulges in stellar density and star-formation history \\citep{whi78,fro87,ter90,hou95} as well as in a relatively high abundance of light elements \\citep{mcw94,sad96}. Its metallicity gradient \\citep{ter88,fro90,tys93,tie95} provides a natural testbed of how metallicity drives stellar evolution.  Consequently we have undertaken a survey of the BHB population in four windows of the bulge along and near its minor axis. BHB candidates are selected from $UBV$ photometry. For a representative subset, follow-up moderate-resolution spectroscopy is providing stellar parameters (confirming BHB status and establishing reddening), plus iron and magnesium abundances to as high a temperature as possible. The basic goal is a statistically complete survey of stars on or near the BHB across all temperatures where they occur, from 7250\\kel\\ to 35,000\\kel. We aim for the same degree of completeness for all BHB stars, regardless of their temperature and metallicity. Such a survey would show immediately whether cool BHB's and sdB/O stars exist at all in the bulge, as suggested at the cool end from its RR Lyrae stars \\citep{wt91}. Knowledge of the numbers of hot BHB stars in each field, in conjunction with the metallicity and temperature distribution of cooler BHB stars in the same region, should help greatly in disentangling the relative influence of the various production factors noted above.  We begin with this pilot study of a single region limited in magnitude to the cool end of the BHB, whose results we briefly summarize here. Spectra were obtained during the commissioning phase of the Two-Degree Field (2DF) spectrograph of the \\aat. Forty-seven stars whose Balmer-line profiles indicated temperatures $\\geq$ 7250\\kel\\ were analyzed by comparing their spectra with a grid of theoretical spectra.  Reddening was then found from the model colors as tabulated by Kurucz. We show that reddening varies dramatically from star to star within this field, with $0 \\leq\\ \\ebmv\\ \\leq\\ 0.55$, so that colors alone cannot determine \\teff\\ and \\logg. However, the $UBV$ color-color diagram helps to distinguish BHB stars from main-sequence turnoff interlopers as long as \\ebmv\\ $\\la$ 0.50.  Thirty-seven of the hot stars proved to have temperatures and gravities indicating a position on the horizontal branch. Two more of the hottest stars might be either BHB or main-sequence stars. None of the unambiguous BHB stars in this sample proves to have \\teff\\ $\\geq$ 9000\\,K, which we attribute to the sparseness of our sample at blue colors and faint magnitudes. We estimate roughly a dozen of our coolest hot stars to be RR Lyraes, pulsating variables located just redward of the BHB stars on the horizontal branch. The hotter BHB stars in this sample tend to be more metal-poor and more spatially extended than the cooler ones. We comment on the possible implications for BHB and RR Lyrae production in metal-rich and metal-poor bulge populations. Five HB stars are discerned with \\feh\\ $\\geq$ $-0.5$ in the bulge itself, the first such stars found. Searches to fainter magnitudes are planned to reach the hot sdB stars believed responsible for the UV upturn. ",
        "conclusions": "Of the 47 stars found to have \\teff\\ $\\geq$ 7250\\kel, gravities \\logg\\ indicate that nine are possible \\popi\\ core-hydrogen-burning stars.  (A tenth, star 163 with the unusual gravity \\logg\\ = 3.5 at \\teff\\ = 7750\\kel, is likely to be an RR Lyrae.) Of the nine, three stars have \\teff\\ = 7250\\kel\\ and \\logg\\ = 4.0, and are likely to be foreground young stars given their distances. Four clearly are young \\popi\\ stars, with \\logg\\ = 4.0 and \\teff\\ $<$ 9000\\,K. Two of these show classical peculiarities \\citep{jas90}: 91 is a Sr-Cr-Eu Ap star, with strong \\srj\\ lines and suitable \\bmv\\ colors, and 46 has very weak \\caj\\ for its Balmer-line and metallic-line strengths, as is typical of Am stars. Stars 75 and 91 are broad-lined. If due to rapid rotation, implied \\vsini\\ values are $\\sim$ 110\\kms\\ and $\\sim$ 190\\kms\\ respectively, values in keeping with normal main-sequence A stars (Jaschek \\& Jaschek 1990). Two stars (48 and 182) with \\teff\\ $\\geq$ 9000\\kel\\ and \\logg\\ = 3.5 could be either foreground young stars or \\popj\\ evolved stars. If the latter, neither is in the bulge, but rather in front of it.  Thirty-seven stars are evolved \\popj\\ core-helium-burning stars on the HB: those with 8500\\kel\\ $\\geq$ \\teff\\ $\\geq$ 7250\\kel\\ and 3.0 $\\geq$ \\logg\\ $\\geq$ 2.5. As discussed above, among the cooler stars we must allow for RR Lyraes, which are HB stars but not BHB stars.  We estimate that two stars with \\teff\\ $\\ge$ 8000\\kel, five with 7500 -- 7750\\kel, and seven of the nine with 7250\\kel\\ should be considered variables. This leaves the number of bona-fide BHB stars around 23.  In our sample, no unambiguous BHB stars appear with \\teff\\ $>$ 8500\\kel. The Balmer line strengths reach a maximum near 9500\\kel, that of the hottest star we have discerned, so the possibility exists that we have erroneously assigned hotter stars to the cooler side. We do not think this has happened, because in hot stars both the Balmer line profiles and the strengths of weak atomic lines are well reproduced at the temperatures assigned.  It is premature to conclude that hotter BHB stars are absent from the bulge, however, because of the sparse coverage of the spectroscopic dataset at the colors and magnitudes appropriate for stars hotter than 10,000\\kel\\ reddened by \\ebmv\\ = 0.3. We may judge this from the \\ngc 6791 CMD and the model colors. \\cite{cha99} match isochrones to the \\ngc 6791 CMD, deriving 0.08 $\\leq$ \\ebmv\\ $\\leq$ 0.13 and 13.30 $\\leq$ \\dmv\\ $\\leq$ 13.45. Indeed, the colors \\citep{mon94,kal95} of 2-17, the cool BHB candidate in \\ngc6791, are matched to a few hundredths in $U - B$, % \\bmv, and \\vmi\\ by those of the model with \\teff\\ = 7250\\kel, \\logg\\ = 3.0, and \\feh\\ = +0.3 when \\ebmv\\ = 0.15 is taken. Adopting \\mv\\ = 1.1 for this star, for which $V$ = 15.0, also reproduces \\dmv\\ = 13.45. Stars in the bulge would appear 1 mag fainter at the same reddening, or 1.5 mag fainter in $V$ if \\ebmv\\ = 0.3. Thus a star with \\teff\\ = 11,500\\kel\\ and \\logg\\ = 4.0 would have $V$ = 16.0 in \\ngc 6791, but would have $V$ = 17.5 in the bulge at \\ebmv\\ = 0.3, and \\bmv\\ = 0.20 and \\umb\\ = $-0.06$ based on the models. We have obtained very few spectra of stars this faint and this blue (see Figures 1 and 2), and so should not be surprised to find none this hot.  The HB stars uncovered by our analysis span a wide metallicity range, $-2.5 \\leq \\feh\\ \\leq\\ +0.0$. This is considerably broader than the metallicity distribution for RR Lyrae stars in Baade's Window at $-4$\\dg\\ \\citep{wt91}, which is sharply peaked near $\\langle{\\rm [Fe/H]}\\rangle = -1.0$ and drops quickly for $\\feh > -0.9$. The HB metallicity distribution is also broader than that of K giants, but in the opposite sense: the K-giant mean is $\\langle\\feh\\rangle$ = $-0.11$, and few K giants are found below $-1.0$ \\citep{sad96}.  The true BHB stars with 8500\\kel\\ $\\geq$ \\teff\\ $\\ge$ 8000\\kel\\ span a rather wide range of distances. Excluding as RR Lyrae stars both star 162 and (arbitrarily) star 118 (at 8000\\kel), we find for 14 BHB stars a mean distance of 8.9\\kpc\\ with an individual deviation $\\sigma_i$ = 4.1\\kpc. Their radial velocity dispersion is 100.5\\kms\\ with a mean of 22.0\\kms. They are metal-poor, averaging \\feh\\ = $-1.59 \\pm 0.13$ with $\\sigma_i$ = 0.49\\,dex. The hotter BHB stars in our sample thus are dominated by a metal-poor, extended-bulge population reminiscent of the halo.  The mean metallicity of the cooler HB stars is difficult to judge but is probably higher. The mean and standard deviation for all 19 low-gravity stars with \\teff\\ $\\leq$ 7750\\kel\\ is \\feh\\ = $-1.29 \\pm 0.18$ and $\\sigma_i$ = 0.77. For the ten stars remaining after excluding the most likely RR Lyraes listed above, we find \\feh\\ = $-1.10 \\pm 0.25$ and $\\sigma_i$ = 0.78\\,dex. Because of the large velocity variability of RR Lyraes, exceeding $100$\\kms\\ for $ab$ types, it is premature to examine the cool HB stars' velocity dispersion.  There are five HB stars with \\feh\\ $\\geq$ $-0.5$, including two with solar metallicity. Only one has \\teff\\ $>$ 7500\\kel. The five stars are located an average 8.2\\kpc\\ away, with a 1$\\sigma$ dispersion of 2.9\\kpc. Thus they are indeed bulge stars. We cannot exclude the possibility that both the solar-metallicity stars are RR Lyraes, given their 7500\\kel\\ temperatures. This is not likely, however, in view of the paucity of metal-rich $c$ RR Lyrae variables \\citep{smi95}, as well as the good fits to their spectra and colors. Moreover, the velocities of these two stars are low, $-56$ and $-45$\\kms. Our work has thus identified solar-metallicity HB stars in the bulge for the first time.  Coupled with the dichotomy above between the RR Lyrae and K-giant metallicities in the bulge, this suggests that the warm and cool BHB stars at this angular distance from the Galactic center may be dominated by two different progenitor populations. One is metal-poor and the other metal-rich, and their HB-star output differs considerably. The metal-poor population, less concentrated to the Galactic center, seems to produce many warm and cool BHB stars per K giant. The metal-rich population, more concentrated towards the galactic center, seems to produce far fewer HB stars per K giant, mostly cool ones.  Moreover, the metal-poor population might form RR Lyraes in preference to stable BHB stars of low temperature, while the metal-rich population might do the opposite. This is suggested from the dearth of metal-rich RR Lyraes in both the bulge and the solar neighborhood, and by the low temperature of 7250\\kel\\ \\citep{pet98} found  for the nonvariable metal-rich star 2-17 in \\ngc 6791, if a true BHB. Thus metal-poor and metal-rich RR Lyraes, like metal-poor and metal-rich sdB's, might be formed through a somewhat different assortment of pathways than cool BHB stars. Such a possibility is also suggested by measurements of surface rotation \\vsini. Although none of the 27 field RR Lyrae stars measured by \\citet{pet96} showed \\vsini\\ $\\geq$ 10\\kms, over half of the cool and warm BHB stars do so, in both globular clusters and the field \\citep{pet83,pet85,pet95}. Globular-cluster BHB stars with \\teff\\ $\\geq$ 12,000\\kel\\ also show very low \\vsini\\ values \\citep{beh99,beh00}. However, to clarify the pathways of formation of RR Lyraes and sdB's in metal-rich and metal-poor populations, we must analyze a much larger sample of bulge HB stars in which RR Lyraes are detected photometrically, and which goes blue enough and faint enough to detect BHB stars hotter than 12,000\\kel.  Reaching the sdB's -- the hottest, faintest BHB stars -- requires going to magnitudes well beyond those of our current survey. Given the average magnitude $<V>$ = 18.0 of such stars in \\ngc 6791, the discussion above suggests they are expected near $V$ = 20.5 at the distance of the bulge when \\ebmv\\ = 0.3. At this reddening, the colors of the solar-metallicity Kurucz model with \\teff\\ = 25,000\\kel\\ and \\logg\\ = 5.0 are \\bmv\\ = 0.07 and \\umb\\ = $-0.68$, reaching 0.01 and $-0.86$ at \\teff\\ = 35,000\\kel. To detect and count such stars in the future, we plan to pursue bulge $UBV$ photometry to $V \\sim B \\sim 21$ and $U \\sim 20.5$. We aim to secure spectroscopy for BHB and blue straggler candidates to $V \\sim 18.5$ and perhaps beyond, to characterize their temperatures and metallicities. By so extending this survey, we hope to shed light on the both the mechanism(s) of production of BHB stars and the predisposing factors of the UV upturn phenomenon."
    },
    "0004/hep-th0004134_arXiv.txt": {
        "abstract": "A four-form gauge flux makes a variable contribution to the cosmological constant.  This has often been assumed to take continuous values, but we argue that it has a generalized Dirac quantization condition.  For a single flux the steps are much larger than the observational limit, but we show that with multiple fluxes the allowed values can form a sufficiently dense `discretuum'.  Multiple fluxes generally arise in M theory compactifications on manifolds with non-trivial three-cycles.  In theories with large extra dimensions a few four-forms suffice; otherwise of order 100 are needed.  Starting from generic initial conditions, the repeated nucleation of membranes dynamically generates regions with $\\lambda$ in the observational range.  Entropy and density perturbations can be produced. ",
        "introduction": " ",
        "conclusions": "Compactifications of M-theory generally give rise to multiple four-form field strengths.  We showed that such theories have vacua with discrete but closely spaced values for the cosmological constant. In the Witten GUT scenario, the spectrum will contain values of $\\lambda$ in the observable range if the number of four-forms is of order 100.  (This requires that the cosmological constant to be cancelled is of GUT scale, not weak scale).  In models with large internal dimensions, four or five four-forms suffice, and a weak-scale cosmological constant can be cancelled.  By repeated membrane nucleation, flux configurations with $\\lambda \\approx 0$ arise dynamically from generic initial conditions.  We argued that entropy and density perturbations can be generated in such regions, and showed that the amplitude for the decay of the $\\lambda \\approx 0$ vacuum is negligible.  An attractive feature of this proposal is that it simultaneously addresses two questions that are usually treated as distinct.  The first question is: Why is the cosmological constant not huge?  One would expect a vacuum density $\\lambda$ of order $M_{\\rm Pl}^4$, or at least TeV$^4$ with supersymmetry.  Until recently this was the only cosmological constant problem.  It appeared to require a symmetry ensuring the exact cancellation of all contributions to the cosmological constant.  This is difficult because contributions are expected to come from many different scales.  The second question is: Why is the cosmological constant not zero?  Recent evidence% \\footnote{A review of these observations can be found in Ref.~\\cite{Car00}.} points to a flat universe with $\\Omega_{\\rm m} \\approx 0.3$ and $\\Omega_\\lambda \\approx 0.7$.  The favored value for the vacuum energy is $\\lambda \\approx 10^{-120} M_{\\rm Pl}^4 \\approx (0.003 \\mbox{ eV})^4$.  In particular, a flat universe with vanishing vacuum energy has been ruled out.  But if it is difficult to explain $\\lambda=0$, a small non-zero cosmological constant seems to pose an even greater theoretical challenge.  The mechanism we propose has limited accuracy because of flux quantization, so that a residual cosmological constant is inevitable.  Our proposal has certain features of the Brown-Teitelboim idea, and also certain features of eternal inflation~\\cite{Lin86a}.  Previously, however, both of these ideas have been difficult to realize with a plausible microphysics.  Our proposal allows both to be realized within string theory.  For the Brown-Teitelboim idea, the main problem was the very small energy scale needed in the discretuum; we see that this can be obtained from a normal hierarchy with multiple fluxes. Eternal inflation with generic polynomial potentials requires scalar field expectation values strictly larger than the Planck scale.  In string theory the scale of the field manifold is the string scale, which is no larger than the Planck scale.  The manifold is actually noncompact, but the asymptotic regions generally correspond to decompactification of spacetime, and in this region the effective potential generally ceases to be flat.  We have realized a version of eternal inflation that does not require such a large scalar, and uses elements already present in string theory.% \\footnote{A precursor to the idea of four-form-driven eternal inflation was presented in Ref.~\\cite{BouCha98}.} Moreover, if the membrane charges are large, the high temperature of de~Sitter space before the final membrane nucleation induces Brownian motion of the inflaton field, thus preparing suitable initial conditions for chaotic inflation after the transition.  The main problem with realizing our picture is the stabilization of the compact dimensions, which is of course a ubiquitous problem in string theory.  A positive bulk cosmological constant is a useful ingredient~\\cite{Sun98,ArkDim98b}, but it is not clear that this can be realized in string theory.  It is interesting that the naked singularity proposal~\\cite{Ark00,Kac00} appears to lead in the end to a very similar picture.  The free parameters that correspond to boundary conditions at a naked singularity in a compact space will become, in a four-dimensional effective Lagrangian, variable coupling constants. In the original proposal these were assumed to be continuous and constant in time, but in Ref.~\\cite{PolStr00} it was argued that they are discrete and can change across a domain wall, just as for the fluxes considered here.  In the example~\\cite{PolStr00} there was a potentially large number of states, of order $e^{\\sqrt N}$ where $N$ is at Ramond-Ramond charge of the singularity.  Note, however, that a charge of order $10^5$ is needed to produce a discretuum sufficiently dense to account for the smallness of the cosmological constant.  In Ref.~\\cite{PolStr00} the main focus was on supersymmetric states, which were all degenerate, but with supersymmetry breaking there will again be a spectrum for $\\lambda$.  Again, stabilization will be an issue.  The appearance of the anthropic principle, even in the weak form encountered here, is not entirely pleasant, but we would argue that it is necessary in any approach where the cosmological constant is a dynamical variable.  That is, a small value for the present cosmological constant cannot be obtained by dynamical considerations alone.  The point is that we can follow cosmology at least back to nucleosynthesis, when the present cosmological constant contributed only a fraction $10^{-30}$ to the energy density of the universe, and so was dynamically irrelevant.  At earlier times, including the point where the cosmological constant is to have been determined, the fraction would have been even smaller.\\footnote{One exception is the wormhole idea~\\cite{Col88a}, where the value of the cosmological constant in our universe is determined by the presence of other, empty, universes.  At least one of the authors retains a certain wary fondness for this possibility.}"
    },
    "0004/astro-ph0004029_arXiv.txt": {
        "abstract": "In this work we explore the effects of adopting an initial mass function (IMF) variable in time on the chemical evolution of the Galaxy. In order to do that we adopt a chemical evolution model which assumes two main infall episodes for the formation of the Galaxy. We study the effects on such a model of different IMFs. First, we use a theoretical one based on the statistical description of the density field arising from random motions in the gas. This IMF is a function of time as it depends on physical conditions of the site of star formation. We also investigate the behaviour of the model predictions using other variable IMFs, parameterized as a function of metallicity.  Our results show that the theoretical IMF when applied to our model depends on time but such time variation is important only in the early phases of the Galactic evolution, when the IMF is biased towards massive stars. We also show that the use of an IMF which is a stronger function of time does not lead to a good agreement with the observational constraints suggesting that if the IMF varied this variation should have been small. Our main conclusion is that the G-dwarf metallicity distribution is best explained by infall with a large timescale and a constant IMF, since it is possible to find variable IMFs of the kind studied here, reproducing the G-dwarf metallicity but this worsens the agreement with other observational constraints. ",
        "introduction": "Observational constraints are of fundamental importance to build a realistic chemical evolution model. With respect to these constraints the last years have been of crucial importance (Pagel 1997) and, in the case of the Milky Way, the new observational data required a revision of the previous chemical evolution models (see Pagel and Tautvaisiene 1995 and Chiappini et al. 1997, hereinafter CMG, for a discussion on this point).  In particular, Gratton et al. (1996, 1999) showed that the distribution of the abundances of  $\\alpha$-elements to Fe for a large homogeneous sample of stars in the solar neighbourhood seems to indicate a short timescale for the evolution of the halo and thick disk phases and a sudden decrease in the star formation in the epoch preceding the formation of the thin disk. An analogous result was found by Fuhrmann (1998) for the [Mg/Fe] ratio. Moreover, Beers and Sommer-Larsen (1995) have shown that the thick disk population extends to very low metallicities. Those are very important new information which stimulated us (CMG) to consider a different picture of Galaxy formation than those previously adopted.  Previous models, in fact, (e.g. Matteucci and Fran\\c cois 1989) were based on a pure colapse picture where the disk formed from gas shed from the halo. These models, however, are difficult to conciliate with the new results discussed above indicating a discontinuity between halo and thin disk. We than suggested the so-called {\\it Two-Infall Model}, a model that assumes that the Galactic thin disk was not only formed from gas shed from the halo but was formed mainly from extra-galactic gas. We assume two main infall episodes, the first one is responsible for the formation of the population made of that fraction of the halo and thick disk stars which originated from a fast dissipative collapse, such as suggested by Eggen et al. (1962). The second infall episode forms the thin disk component with a timescale much longer than that of the halo formation.  In this new picture the disk was formed slowly (with a timescale of 7-8 Gyrs at the solar vicinity) and from inside-out (Matteucci and Fran\\c cois 1989). A direct consequence of that is that at high redshift we should expect to see smaller disks in size (Roche et al. 1998). This long timescale for the formation of the thin disk at the solar vicinity, required to produce a good fit of the obseved G-dwarf metallicity distribution (Rocha-Pinto and Maciel 1996) was also confirmed by recent chemical evolution models (eg. Portinari et al. 1998, Prantzos and Silk 1998, Chang et al. 1999) as well as by chemo-dynamical models (eg. Hensler 1998) and is also in agreement with the results showed in this conference by Carraro (2000).  The two-infall model adopted a constant IMF. On the empirical grounds there is at present no clear direct evidence that the IMF in the Galaxy has varied with time. A detailed discussion about possible observed variations in the IMF in different environments is given by Scalo (1998), but such variations are comparable with the uncertainties still involved in the IMF determinations. The present uncertainties in the observational results prevent any conclusion against the idea of an universal IMF.  However, a variable IMF, which formed relatively more massive stars during the earlier phases of the evolution of the Galaxy compared to the one observed today in the solar vicinity, has often been suggested as being one of the possible solutions for the G-dwarf problem (namely the deficiency of metal-poor stars in the solar neighborhood when compared with the number of such stars predicted by the simple model). Such an IMF would also be physically plausible from the theoretical point of view if the IMF depends on a mass scale such as the Jeans mass. Given the uncertainties in both theoretical and observational grounds, the proposed IMFs can in principle be tested by means of a detailed chemical evolution model.  An example of a theoretical approach to the IMF problem is the one proposed by Padoan et al. (1997 - hereinafter PNJ). Since random motions are probably ubiquitous in sites of star formation, PNJ suggested to describe the formation of protostars as the gravitational collapse of Jeans masses in a density distribution shaped by random supersonic motions (but see Scalo et al. 1998). This IMF was already tested in models of elliptical galaxies by Chiosi et al. (1998) and they concluded that a strongly varying IMF could be suitable for such galaxies.  Our goal was to address the question of what time-dependent IMF properties are allowed in order to still match the observational constraints when adopting the two-infall model (CMG). To do this we tested different IMFs in our chemical evolution code, going from those where the variability is contained in the slope of the power-law, assumed to be a function of the metallicity (parametrizations adopted by Scully et al. 1996, Matteucci and Tornamb\\'e 1987), to the one by PNJ which predicts also a change in the stellar mass which contributes most to the IMF as a function of time.  A detailed description of the {\\it Two-Infall} model as well as of the PNJ IMF and the hypothesis needed to introduce it in our chemical evolution code are discussed in detail in Chiappini et al. 2000a (hereinafter CMP; see also CMG and PNJ). ",
        "conclusions": "The main results are summarized below: \\par\\noindent a) We tested the IMF proposed by PNJ in a model for the chemical evolution of the Milky Way (CMG) and we showed that this IMF gives good agreement with the observed properties of the solar vicinity. However, such an agreement is due to the fact that this IMF when applied to the two-infall model shows a time variation that is important only in the early phases of Galactic evolution. This in turn is due to the simplifying assumptions adopted here, like neglecting the dependence of the IMF on the gas temperature which would produce more sharply varying IMF. In these early phases the IMF is biased towards massive stars. \\par\\noindent c) the PNJ IMF combined with the inside-out picture for the thin disk formation predicts a gradient flatter than the one predicted by a model which adopts a constant IMF. This situation cannot be reversed by changing the SFR because in this case the abundance gradient is recovered but the gas density profile is destroyed. \\par\\noindent d) Models which adopt IMFs strongly dependent on metallicity, thus simulating a dependence also on the gas temperature, are not in agreement with the most important observational constraints of the solar vicinity and predict radial gas profiles at variance with observations, therefore they should be rejected. \\par\\noindent e) We conclude that a constant IMF and the assumption of a continuous infall onto the Galactic disk is still the best way to explain the observational constraints in the Milky Way including the G-dwarf metallicity distribution. A probable source of the required infall could be the HVCs (see Burton and Braun 1999 for a discussion on this point).  {\\small"
    },
    "0004/astro-ph0004265_arXiv.txt": {
        "abstract": "We present high-resolution spectra (R $\\sim$ 49,000) of stars that have parallax measurements from the {\\it Hipparcos} satellite and are projected along the line of sight to the two nearest known star forming clouds to the Sun: MBM12 and MBM20.  The spectra were obtained with the FOCES Echelle Spectrograph at the 2.2 meter telescope in Calar Alto, Spain and the wavelength range was chosen to include the interstellar \\ion{Na}{I} D lines at $\\lambda 5889.950$ \\AA\\ and $\\lambda 5895.924$ \\AA. Since the stars are at a range of distances, we use their spectra along with their parallaxes from {\\it Hipparcos} to determine the distance to the molecular gas. The stars in front of the cloud do not show interstellar \\ion{Na}{I} D absorption features while the stars behind the cloud do show interstellar absorption features.  We find that both clouds are somewhat more distant than previously estimated. The revised distance to MBM12 is 58$\\pm$5~pc $< d <$ 90$\\pm$12~pc and the distance to MBM20 is $112\\pm15$ pc $< d <$ $161\\pm21$~pc. ",
        "introduction": "The two nearest known star forming clouds to the Sun are MBM12 and MBM20. The MBM12 complex consists of clouds 11, 12, and 13 from the catalog of Magnani et al.~(1985) and is  located at (l,b) $\\sim$ (159\\fdg4,$-$34\\fdg3) (we will refer to the entire complex as MBM12). It was first identified by Lynds (1962) and appears as objects L1453-L1454, L1457, L1458 in her catalog of dark nebulae. The mass of the entire complex is estimated to be $\\sim$~30--200~M$_{\\odot}$ based on radio maps of the region in $^{12}$CO, $^{13}$CO and C$^{18}$O (Pound et al.~1990; Zimmermann \\& Ungerechts 1990). The cloud MBM20 was first identified as L1642 in the Lynds (1962) catalog of dark nebulae. The mass of MBM20 is $\\sim$~84~M$_{\\odot}$ and it is located at (l,b) $\\sim$ (210\\fdg9,$-$36\\fdg6) southwest of the Orion star forming complex (Magnani et al. 1985).  Both MBM12 and MBM20 are star forming clouds. Several T~Tauri stars have been identified in MBM12 via H$\\alpha$ emission line and X-ray surveys (Herbig \\& Bell 1988; Stephenson 1986; Hearty et al. 2000). Two binary classical T~Tauri stars have been identified in the central region of MBM20 (Sandell et al. 1987). Subsequent high-resolution $^{12}$CO(J=1-0) observations near the center of the cloud have shown that the infrared point source IRAS04325-1419 corresponding to one of the binaries is associated with a red and blue-shifted bipolar outflow (Liljestr\\\"om et al. 1989).   Using {\\it ROSAT} observations, Kuntz et al. (1997) find evidence based on the possible detection of 0.25 keV X-ray shadows of these clouds that both are located either within or at the edge of the Local Bubble. They estimate that distances, $d$, of $60\\leq d \\leq 90$ pc to MBM12 and $d$ $<$~100~pc to MBM20 are consistent with the foreground 0.25 keV emission seen in the direction of each cloud.  The distances are also consistent with previous distance estimates (see Sect.~\\ref{previous}).   Since MBM12 and MBM20 are both star forming clouds which are probably located within or at the edge of the Local Bubble, determining an accurate distance to these clouds is important for studies of stars and of the interstellar medium. In addition, observations of MBM12 are already scheduled for {\\it XMM} (50 ks) and {\\it Chandra} (100 ks) to investigate the diffuse X-ray emission of the Local Bubble. Therefore we present observations to improve the distance estimates to both clouds. In Sect.~\\ref{previous} we review previous distance estimates for these clouds and the {\\it Hipparcos} results that revise them.  In Sect.~\\ref{specsec} we present our spectroscopic observations of {\\it Hipparcos} stars along the line of sight to both of the clouds which we use to improve the previous distance estimates.  In Sect.~\\ref{conclusions} we summarize our investigations and suggest future observations. ",
        "conclusions": "\\label{conclusions} We have investigated the {\\it Hipparcos} parallax measurements of stars to estimate the distance to the two nearest star forming clouds to the Sun: MBM12 and MBM20. The {\\it Hipparcos} observations of previously observed stars suggest that the distances to both of these clouds are not as well constrained as previously thought, although they are consistent with the previous estimates.  Therefore, we obtained high resolution spectra of additional {\\it Hipparcos} stars at intermediate distances and improved the distance estimates to both clouds. The distance to MBM12 is 58$\\pm$5~pc $< d <$ 90$\\pm$12~pc and the distance to MBM20 is $\\sim$ $112\\pm15$ pc $< d <$ $161\\pm21$~pc. Since there are additional stars that were observed with {\\it Hipparcos} that are projected near each cloud (though farther off axis than the stars we observed), future spectroscopic observations like those presented here may further refine our distance estimates."
    },
    "0004/astro-ph0004053_arXiv.txt": {
        "abstract": "We quantitatively cross-associate the 18811 ROSAT Bright Source Catalog (RASS/BSC) X-ray sources with optical sources in the \\usno\\ catalog, calculating the the probability of unique association (\\pid) between each candidate within 75\\arcsec\\ of the X-ray source position, on the basis of optical magnitude and proximity.  We present catalogs of RASS/BSC sources for which \\pid$>$98\\%, \\pid$>$90\\%, and \\pid$>$50\\%, which contain 2705, 5492, and 11301 unique \\usno\\ optical counterparts respectively down to the stated level of significance. Together with identifications of objects not cataloged in \\usno\\ due to their high surface brightness (M31, M32, ...) and optical pairs, we produced a total of 11803 associations to a probability of \\pid$>$50\\%.  We include in this catalog a list of objects in the SIMBAD database within 10\\arcsec\\ of the \\usno\\ position, as an aid to identification and source classification.  This is the first RASS/BSC counterpart catalog which provides a probability of association between each X-ray source and counterpart, quantifying the certainty of each individual association.  The catalog is more useful than previous catalogs which either rely on plausibility arguments for association, or do not aid in selecting a counterpart between multiple off-band sources in the field.  Sources of high probability of association can be separated out, to produce high-quality lists of classes (Seyfert 1/2s, QSOs, RS CVns) desired for targeted study, or for discovering new examples of known classes (or new classes altogether) through the spectroscopic classification of securely identified but unclassified \\usno\\ counterparts.  Low \\pid\\ associations can be used for statistical studies and follow-on investigation -- for example, performing follow-up spectroscopy of the many low-mass stars to search for signatures of coronal emission, or to investigate the relationship between X-ray emission and classes of sources not previously well-studied for their X-ray emissions (such as pulsating variable stars).  We find that a fraction $\\sim$65.8\\% of RASS/BSC sources have an identifiable optical counterpart, down to the magnitude limit of the \\usno\\ catalog which could be identified by their spatial proximity and high optical brightness. ",
        "introduction": "The ROSAT Bright Source Catalog (BSC; \\citenp{voges96,voges99}) contains positions, X-ray countrates, and spectral information of 18811 X-ray sources with countrates $>$0.05 c/s, observed during the ROSAT All-Sky-Survey (RASS).  Efforts to identify the sources of X-ray emission with counterparts in other wavebands are hampered by source confusion. The error-region of the RASS/BSC sources average $\\sim$12\\arcsec (1$\\sigma$), which can contain several candidate objects, any of which may be the source of X-rays.  To date, most efforts to identify the X-ray sources with parent populations -- usually, optical sources -- have been targeted toward sub-groups of known X-ray emitting populations, such as coronal X-ray sources \\cite{berghoefer97,huensch98a,huensch98b,huensch99}; AGN/QSOs \\cite{thomas98,beuermann99}; OB stars \\cite{berghoefer96,berghoefer96er,berghoefer97,motch97a,motch97b}; and high galactic latitude spectrally soft sources \\cite{thomas98}.  As part of a larger effort to identify QSOs, a general spectroscopic survey has also identified stellar type-sources \\cite{bade95,bade98}.  Many of these associations -- though not all -- have been based upon an argument of plausibility.  In the plausibility method, one typically performs imaging photometry and spectroscopy of objects within the X-ray error-box, and finds a plausible counterpart among these; a counterpart is usually considered plausible if the candidate object's class is previously known to emit X-rays, and if the properties (such as magnitude, or implied $L_{\\rm X}/L_{\\rm opt}$) are consistent with those previously observed from other objects within its class.  This method is useful when the parent population is known and rare (much less than one object per average X-ray error-box size). However, this method will not discover X-ray sources independent of object classification. In addition, some of the studies which rely upon plausibility do not measure the level of background contamination, while others do not evaluate the limiting \\pid\\ (the lowest probability of unique identification a prospective counterpart can have, and still be included in the catalog).  None provide a \\pid\\ for each cross-identification, which makes it impossible to quantitatively evaluate the quality of a purported association in future work.  In Table~\\ref{tab:prevwork}, we list previous works which catalog $\\approxgt$ 100 optical counterparts to RASS/BSC objects, or which sought RASS/BSC counterparts for a particular class of sources.  The table includes: (1) the reference; (2) brief description of the cross-identification catalog; (3) the number of cross-identifications found and the estimated number of mis-identified (background) sources in the cross-id list ($N_{\\rm bkg}$); (4) a brief description of the identification algorithm used; (5) and the probability of unique association (\\pid) between the X-ray source and candidate counterpart at the identification limit of the catalog.  For several of these works, no estimation of the probability of cross-identified sources being background sources at the detection limit was given. Of those which do, a probability of confusion with background (that is, unassociated) sources of $\\sim$50\\% is a common limit (below which, the identified counterpart is more likely to be an unrelated background object than actually associated with the X-ray source.  An extensive comparison with several published and unpublished cross-identification catalogs was made in Table~3 of the RASS/BSC work \\cite[\\voges\\ hereafter]{voges99} (for a total of $\\sim$17,000 sources), which we discuss more completely in Sec.~\\ref{sec:v99}.  We include, for comparison, the results of the present work.  We have undertaken a project of off-band identification of ROSAT/BSC X-ray sources -- XID.  The goal of this project is to provide a catalog of cross-identifications, which provides the probability of unique identification (\\pid) between the off-band counterpart for all RASS/BSC sources.  We have additional motivation for performing the present work.  Future databases of sources -- both in the X-ray and in other bands -- will and already do contain $10^{5}-10^{9}$ objects.  This is too great a number of objects on which to perform non-automated methods for identifying counterparts.  We therefore develop and use an automated method for identifying cross-band counterparts, which can be further adapted and used in future studies.  We present the method and results of a statistical cross-identification between RASS/BSC and \\usno\\ catalogs, producing a high, medium, and low-quality cross-identification list.  We summarize these results, and provide a short discussion on the content of the cross-identification catalog.  These catalogs are given in the appendix. ",
        "conclusions": "We have cross-correlated the 18811 RASS/BSC X-ray sources with 321144 candidate \\usno\\ optical counterparts within 75\\arcsec\\ of the RASS/BSC source position, on the basis of $B$ magnitude and source proximity, taking into account the quoted RASS/BSC positional uncertainty.  On this basis, we identify 2705 \\usno\\ objects with \\pid$>$98\\%, with $\\sim$0.66\\% background contamination; 5492 with \\pid $>$90\\%, with $\\sim$2.8\\% background contamination; and 11301 with \\pid$>$50\\%, with $\\sim18\\%$ background contamination.  Thus, we have identified possible optical counterparts to 60\\% of the ROSAT/BSC on the basis of position and photometry alone.  We have also provided -- for the first time -- a probability of unique identification between each of the X-ray sources and their proposed counterpart.  When we include unique ``binary'' identifications, and 30 high-surface brightness objects which were not included in \\usno, we have presented optical associations for a total of 11803 objects, down to a limiting identification probability of 50\\%, which is 62.7\\% of the RASS/BSC catalog objects. More conservatively, we have presented optical associations for 5553 objects, to \\pid$\\ge$90\\%, which is 29.5\\% of the RASS/BSC catalog. The breakdown of these identified sources is listed in Table~\\ref{tab:number}.  The individual identifications are subject to systematic uncertainty of the association between X-ray sources and clustered optical sources (such as clusters of galaxies, open stellar clusters and star formation regions), in which the X-ray source may reside, and the greater than average density of candidate optical counterparts makes the presence of a brighter-than-average source more likely than in a background field.  Thus, the given optical identification should be considered an ``association'' -- and the likelihood that the source of X-ray emission is the identified optical point source directly or a nearby associated object must be evaluated on a case-by-case basis, on the basis of the likelihood of such a secondary association.  For these sources, we have listed the RASS/BSC source-name, and the identified \\usno\\ counterpart.  In addition, we compiled a list of objects in the SIMBAD database within 10\\arcsec\\ of the \\usno\\ counterpart, many of which are likely to be the \\usno\\ counterpart itself.  There are a surprisingly high fraction (25\\% in the \\pid$>$90\\% Catalog) of optical counterparts which are not named in the SIMBAD database.  As these are (photometrically) identical to objects which have been previously classified, the unclassified objects are likely to be the same population.  Thus, a program of classification of these unclassified objects will likely discover new examples of known classes of sources, although they may contain unknown classes as well.  The limit on the fraction of RASS/BSC sources which have counterparts in the \\usno\\ catalog discoverable by this method is $Q\\le$72.2\\%. To improve this identification fraction between X-ray sources and optical data, either additional optical information is required (source classes, spectral colors) which will help distinguish identifiable sources, or improved X-ray localizations (such as from the ROSAT/HRI, or \\chandra), or combining X-ray and optical information to pick out sources of particular classes ($L_x/L_{\\rm opt}$)."
    },
    "0004/astro-ph0004323_arXiv.txt": {
        "abstract": "We present SCUBA observations of the protomultiple system NGC1333/IRAS4 at 450\\micron\\ and 850\\micron.  The 850\\micron\\ map shows significant extended emission which is most probably a remnant of the initial cloud core. At 450\\micron, the component 4A is seen to have an elongated shape suggestive of a disk. Also we confirm that in addition to the 4A and 4B system, there exists another component 4C, which appears to lie out of the plane of the system and of the extended emission.  Deconvolution of the beam reveals a binary companion to IRAS4B.  Simple considerations of binary dynamics suggest that this triple 4A-4BI-4BII system is unstable and will probably not survive in its current form. Thus IRAS4 provides evidence that systems can evolve from higher to lower multiplicity as they move towards the main sequence.  We construct a map of spectral index from the two wavelengths, and comment on the implications of this for dust evolution and temperature differences across the map.  There is evidence that in the region of component 4A the dust has evolved, probably by coagulating into larger or more complex grains. Furthermore, there is evidence from the spectral index maps that dust from this object is being entrained in its associated outflow. ",
        "introduction": "Stars of all ages are commonly found in binary and multiple systems \\cite{duqmay}. In fact, amongst the youngest stars, the frequency of companions appears to be higher than in older systems (eg Ghez, 1995). In order to understand the initial stages of star formation, we need to understand the initial stages of binary star formation. There have been many theories advanced to explain the formation of binary stars (cf Clarke 1995; Bonnell 1999). These generally include either a fragmentation during collapse \\cite{boss,bonnetal}, a fragmentation of a circumstellar disc \\cite{bonnfrag,whitworth} or a post-fragmentation star-disc capture (eg Clarke \\& Pringle 1991). In order to be able to distinguish between these theories, we need to observe the youngest systems. This paper reports on recent observations of a young, protostellar multiple system NGC~1333 IRAS~4 in order to constrain its formation mechanism.  NGC1333/IRAS4 is a well studied protobinary system. It was first identified as a site of star formation by Haschick et al (1986), who observed two variable H$_2$O masers. The distinct core was mapped by Jennings et al (1987).  Sandell et al (1991, hereafter SADRR) mapped the system in the submillimetre with UKT14. They confirmed its multiple nature, finding a 30'' binary system embedded in diffuse emission. They labelled the components 4A and 4B. Component 4A was seen to be elongated, which SADRR interpreted as a massive disk seen at an oblique angle. Dynamical evidence of binarity for IRAS4 is lacking, but the common envelope of material surrounding the two main components suggests that they are the products of a single collapse event, the extended material being identified as the remnants of the precollapse core. This extended envelope then provides the main motivation for considering IRAS 4 to be a multiple system, rather than a chance superposition of cores.  The system is associated with a high velocity outflow, mapped at various molecular transitions by Blake et al (1995, hereafter BSDGMA).  These authors found a high velocity outflow originating from IRAS 4A and aligned with the apparent disk axis.  Minchin et al \\shortcite{minchin} measured polarisation at 800\\micron\\ for both 4A and 4B. They found that the polarisation position angle was similar for both sources and broadly aligned with the elongated circumbinary emission.   Interferometry by Lay et al \\shortcite{lay} revealed that both 4A and 4B are themselves multiple systems. 4A was revealed to be a binary of separation 1.2'' aligned with the direction of elongation of 4A. 4B appeared also to be multiple, but had a more complex nature which could not be determined. ",
        "conclusions": "\\subsection*{Morphology of IRAS4}  The elongation of 4A could be explained as a circular disk inclined at an angle of approximately 53$^{o}$ to the plane of the sky. The elongation of 4B in the raw map appears to be EW. However, upon deconvolving the beam from the map we discovered that this elongation is due to 4B being itself a binary system. The primary component of 4B in the deconvolved 450\\micron\\ map appears to have a slight elongation along the same axis as 4A.  The amount of diffuse material in which the system is embedded is difficult to determine with any accuracy, since the map does not encompass the entire extended ridge, the flux density levels may be affected by the chop throw sampling some of the extended emission, and at 450\\micron\\ the error beam may spread flux density from the compact sources into the surrounding regions.  \\subsection*{Stability of the system}  The separation of the 4B double is approximately 12'' on the sky. If we assume that the triple system is coplanar, so that the 4BI-BII axis has the same inclination to the plane of the sky as the 4A disk, the true BI-BII separation would then be 16'' (5600 AU for a distance of 350pc). The 4A-4B separation is approximately 30''.  The long axis of the 4A disk extends some 8'' from the centre in the deconvolved images (we measure the disk extent as the distance at which the flux density reaches 5\\% of its peak value). Thus the presence of the 4A disk is not incompatible with the interpretation that 4A, 4BI and 4BII have coplanar orbits.     We can also make an assessment of how stable the triple 4A-4BI-4BII system is. We first assume that the system is coplanar and the orbits circular. A criterion for stability in a hierarchical triple was developed by Harrington (1977), \\begin{equation} \\frac{D_{triple}}{D_{binary}} > K \\left\\{ 1 + A \\ln \\left[\\frac{2}{3} \\left( 1 + \\frac{M_3}{M_1 + M_2} \\right) \\right] \\right\\}. \\end{equation} Here, $D_{triple}$ is the periastron distance of the stand-alone star, and $D_{binary}$ the semi-major axis of the binary orbit. $M_3$ is the mass of the single component, $M_1$ and $M_2$ the masses of the binary components. The constants have values $K=3.5$ and $A=0.7$ for a corevolving system, or $K=2.75$, $A=0.64$ for a counterrotating system. For our case, the mass ratio is $1/0.63$ (Section~\\ref{ratios}).  The semimajor axis of the BI-BII system cannot be smaller than the observed separation on the sky of 12''. The true 4A-4BI separation of course depends on the angle of the binary-single star axis.  We have already argued (as have previous authors) that the 4A elongation is an inclined disk.  If we assume that the 4A-4BI orbit is coplanar with the 4A disk (expected if fragmentation formation models apply), then the true 4A-4BI distance would be 30'', which would also be the largest possible periastron distance. Taking these figures we find that the system fails Harrington's stability test easily, even if we consider the more stable counterrotating case.  See Figure~\\ref{stabfig}.   \\begin{figure} \\psfig{{figure=stabgraph.ps,width=3.truein,height=3.truein}} \\caption{\\label{stabfig} Harrington's stability test as applied to IRAS4.  The x-axis is the mass ratio $M_3 / (M_1 + M_2)$. The y-axis is the ratio of periastron distance to binary separation. The two bold lines represent the corotating and counterrotating stability limits. The lower of the two represents the more stable counterrotating case. Systems above these lines will be stable.  The position of the 4A-4BI-4BII system is marked as a point. The error in the mass ratio is calculated from the errors in Table~\\ref{peakmass}. The uncertainty in the ordinate has been estimated as follows. The misalignment between component BI and the axis of elongation of 4A was projected back into the plane of the supposed 4A disk. This then gives the possible discrepency between the 4A-4BI distance measured in the plane of the sky and the inferred 4A-4BI distance in the plane of the disk. The upper error bar shows the possible increase in stability due to this observed misalignment. Since the quantities $D_{triple}$ and $D_{bin}$ have been chosen to maximise stability, the y-position of the IRAS4 system is an upper limit.} \\end{figure}   It is possible to envisage accretion of material stabilising initially-unstable triple systems by modifying the separation ratios of the components (Smith et al, 1997). There certainly seems to be a substantial reservoir of available material in the IRAS4 envelope (see Table~\\ref{extmass}). In the case of a low-mass binary orbiting a higher-mass single star, as here, this mechanism is unlikely to result in stability.  Low specific angular momentum material will tend to be accreted onto IRAS4A, causing the single-binary separation to decrease and destabilizing the system. High specific angular momentum will tend to accrete onto the binary (4BI-4BII), widening it by adding angular momentum and again destabilizing the system. Furthermore, the 4A-4BI-4BII system fails the Harrington stability test by a comfortable margin. Even if the available envelope mass were accreted exclusively onto 4BI and BII, moving the mass ratio towards 1, there is not enough mass to entirely stabilize the system. We conclude that IRAS4 seems to be an unstable multiple which will be disrupted within a few orbits.   \\subsection*{Star forming history of IRAS4}  Could the elongation of the main cloud be explained as a foreshortened disk, in a similar way to the elongation of 4A? The young age (approx 10$^5$ years deduced from the embedded nature), appears to preclude this interpretation. For a circumstellar disk to form requires several of the disks dynamical time and a disk of this size would have a period in excess of 4$\\times$10$^5$ years. This implies that the extended structure must be a remnant of the cloud's initial conditions and not due to the subsequent collapse dynamics.  Based on their polarimetric observations, Minchin et al suggested that the extended emission is an inclined ``pseudo-disk'', of the sort envisaged by for example Galli \\& Shu (1993). In this case, the cloud material is magnetically supported in one direction, but free to collapse in the other. A massive, non-centrifugally supported disk can therefore form in a shorter time than would be expected on dynamical grounds. The ambipolar diffusion timescale would have to be longer than the freefall time of the cloud, and the magnetic field would be an impediment to fragmentation of the central region, because it would tend to enforce solid body rotation. For this reason, we do not favour the ``pseudo-disk'' interpretation.  Simulations indicate that a prolate cloud, with an end-over-end rotation, should undergo fragmentation and form a central binary (Bonnell et al, 1992, Bonnell \\& Bastian 1992). The additional multiplicity of the system (4BI-4BII) is then explainable as being due to an internal disk fragmentation (Bonnell 1994) that occurs in the individual components formed in the prolate cloud fragmentation (Bonnell et al 1992).  This interpretation still leaves us with some unexplained details. Most importantly, the third component to the north east. There are several possible explanations for this component. Firstly, it could be part of the IRAS 4 system, but lie outside of the prolate cloud, in front of the main system and well above the plane of the IRAS 4A disk and inferred 4A-4B orbit.  This poses substantial problems for fragmentation models where collapse occurs preferentially in one direction, reducing the dimension of the cloud and making it unstable to fragmentation (Bonnell~1999). Secondly, it could lie in the same plane as the 4A disk, which would place it well outside the prolate cloud as it is seen in the maps, at a distance of 15,400 AU (44'' at 350pc).  If IRAS 4C is part of the IRAS 4 system, then the most promising explanation is that an independent condensation (see eg, Pringle~1991) was present near the prolate cloud when collapse occured.  There exists the possibility that IRAS 4C is not a member of the IRAS 4 system but just a chance projection. Its separation and the presence of other sources nearby lend credence to this possibility, as does it's spectral index being higher than the surrounding cloud.  A determination of the velocities of the various components could shed further light on the possible relationship of 4C with the central system."
    },
    "0004/astro-ph0004115_arXiv.txt": {
        "abstract": "We performed a series of high-resolution collisionless N-body simulations designed to study the substructure of Milky Way-size galactic halos (host halos) and the density profiles of halos in a warm dark matter (WDM) scenario with a non-vanishing cosmological constant. The virial masses of the host halos range from $3.5 \\times 10^{12}\\ \\msunh$ to $1.7 \\times 10^{12}\\ \\msunh$ and they have more than $10^5$ particles each. A key feature of the WDM power spectrum is the free-streaming length \\fsl\\ which fixes an additional parameter for the model of structure formation. We analyze the substructure of host halos using three \\fsl\\ values:  0.2, 0.1, and 0.05 \\mpc\\ and compare results to the predictions of the cold dark matter (CDM) model. We find that guest halos (satellites) do form in the WDM scenario but are more easily destroyed by dynamical friction and tidal disruption than their counterparts in a CDM model. The small number of guest halos that we find in the WDM models with respect to the CDM one is the result of a lower guest halo accretion and a higher satellite destruction rate. These two phenomena operate almost with the same intensity in delivering a reduced number of guest halos at $z = 0$. For the model with $\\fsl = 0.1\\ \\mpc$ the number of accreted small halos is a factor 2.5 below that of the CDM model while the fraction of destroyed satellites is almost twice larger than that of the CDM model. The larger the \\fsl\\ value the greater the size of these two effects and the smaller the abundance of satellites. Under the assumption that each guest halo hosts a luminous galaxy, we find that the observed circular velocity function of satellites around the Milky Way and Andromeda is well described by the $\\fsl = 0.1\\ \\mpc$ WDM model. In the $\\fsl =0.1- 0.2\\ \\mpc$ models, the surviving guest halos at $z=0$ ---whose masses are in the range  $M_h \\approx 10^9-10^{11}\\ \\msunh$--- have an average concentration parameter $c_{1/5}$ $( =r(M_h)/r(M_h/5) )$ which is approximately twice smaller than that of the corresponding CDM guest halos. This difference, very likely, produces the higher satellite destruction rate found in the WDM models. The density profile of host halos is well described by the NFW fit whereas guest halos show a wide variety of density profiles. A tendency to form shallow cores is not evident; the profiles, however, are limited by a poor mass resolution in the innermost regions were shallow cores could be expected. ",
        "introduction": "Non-baryonic dark matter is an essential ingredient of current inflation-inspired models of cosmic structure formation in the universe. From the point of view of particle physics, there is no obvious preference for any of the predicted dark matter candidates (\\cite{CDW96}), which, according to their rms velocity at the time of their decoupling, can be cold, warm, or hot. From the point of view of structure formation, the most compelling candidate has been the cold dark matter. The CDM scenario for structure formation has successfully accounted for several observational facts, particularly on large scales, without introducing an additional free parameter related to its particle distribution function in phase space. However, on small scales and/or in high-density regions of the universe, the predictions of the CDM models seem to be in conflict with observations.  One of the potential problems of the CDM scenario is that the predicted number of low-mass halos ---where probably dwarf galaxies form--- within a Milky Way-size halo, greatly exceeds the observed abundance of satellite galaxies in the Local Group (Klypin et al. 1999, hereafter \\cite{KKVP99}; \\cite{Moore99a}; see also \\cite{Kauffmann93}). A second problem is that the predicted inner density profiles of CDM halos may disagree with the shallow profiles inferred from the rotation curves of dwarf and low surface brightness galaxies (\\cite{Moore94}; \\cite{FP94}; \\cite{Burkert95}; \\cite{dBM97}; \\cite{HG98}), although the observational data for the latter galaxies are controversial (\\cite{vdB99}; \\cite{SMT2000}; but see \\cite{Firmani2000b}). High-resolution gravitational lensing maps of a cluster of galaxies have also revealed a soft inner mass distribution in the halo of this cluster (\\cite{TKD98}). The rotation curve decompositions of normal galaxies and the Tully-Fisher relation obtained in galaxy formation models as well as the dark mass contained within the solar radius in our Galaxy, also point out to dark halos shallower and/or much less concentrated than those predicted by the CDM model (\\cite{AFH98}; \\cite{Navarro98}; \\cite{NS99}; \\cite{FA2000}; \\cite{MM2000}). If these shortcomings of the CDM scenario are confirmed with more observational and theoretical data, new alternatives (cosmological and/or astrophysical) have to be explored in order to modify the properties of the mass distribution at small scales.  In a recent burst of papers, explored alternatives include modifications to: either the nature of the dark matter candidate (\\eg\\ \\cite{SS99}; Hannestad 1999; \\cite{SomDol00}; \\cite{WC2000}; \\cite{Firmani2000a}; \\cite{HD2000}; \\cite{Moore2000}; \\cite{YSWT2000} ; \\cite{Burkert2000}; \\cite{Peebles2000}; Hannestad \\& Scherrer 2000; Riotto \\& Tkachev 2000), or the generation of the primordial power spectrum (\\eg\\ \\cite{KL99}). More conservative astrophysical mechanisms to overcome the problems mentioned above have also been proposed (\\eg\\ \\cite{NEF96}; \\cite{GS99}; \\cite{BKW2000}; \\cite{BGS2000}). One possible modification is to go from a CDM scenario to a warm dark matter (WDM) one. The WDM particles (warmons) would suppress the power at small scales by free-streaming out of overdense regions limiting the formation of substructure at scales below the free-streaming scale. At large scales, the structure formation would proceed in a very similar way to that of a CDM model. N-body simulations have shown that indeed large-scale structure in WDM models looks similar to that of a CDM model (Colombi et al. 1996). On the other hand, as Hogan \\& Dalcanton (2000) noted, the finite phase density of dark halos inferred from observations could be pointing to a non-negligible DM velocity dispersion at the time of structure formation.  Using the Press-Schechter formalism, Kamionkowski \\& Liddle (1999) have shown that if the CDM power spectrum is filtered at scales corresponding to dwarf galaxies, then the abundance of Milky Way satellites can be reproduced. Recently, \\cite{WC2000} reported results from N-body simulations for WDM models at high redshifts. They found that the abundance of $10^{10} \\msunh$ halos is reduced by a factor of $\\sim 5$ at $z = 3$ with respect to the CDM model when the power spectrum is filtered at $k \\approx 2 \\kmpch$. At the same time they showed that the Ly-$\\alpha$ power spectrum at this redshift is very similar to that of the CDM model, which is in agreement with observations. This apparent contradictory result is explained by the fact that the collapse of large-scale structures, as they go non-linear, regenerates the initially suppressed small-scale modes in the power spectrum (\\cite{WC2000}).  These results encourage us to explore in more detail the predictions of WDM N-body simulations at the present epoch. Does the suppression of power at small scales of a WDM model actually eliminate the excessive degree of substructure predicted by the CDM scenario? Are the WDM halos less concentrated? And if so, do they have a smoother inner mass distribution than their counterpart CDM halos? The main aim of this paper is to give a quantitative answer to the first question. To this end we have carried out high-resolution N-body simulations of Milky Way-size galactic halos in three different WDM models. A host halo of about $2 \\times 10^{12}\\ \\msunh$ has more than $10^5$ particles in the simulations. Since the most successful variant of the CDM models is a flat universe with a non-zero cosmological constant ($\\Omega_{\\Lambda} = 0.7$ and $h = 0.7$), here we also use this cosmological model but instead of CDM we introduce WDM with the extra free parameter \\fsl: these models will be our $\\Lambda$WDM models (for economy we drop off the greek letter $\\Lambda$ hereafter when we refer to either CDM or WDM models). We will also address the questions of concentrations and density profiles of dark halos, although the small number of large high-resolved halos and the small range of masses in the simulations constrain our predictions on this subject.  In Section 2 we discuss the WDM models to be explored in this paper. In Section 3 we briefly describe the numerical technique that we used for the simulations. Section 4 is devoted to the analysis and comparison with observations of the circular velocity function of satellites within host halos of Milky Way-sizes. The concentrations and density profiles of the host and satellite halos are presented in Section 5. In Section 6 we discuss some of the results, and summarize of our main conclusions is given in Section 7. ",
        "conclusions": "Using high-resolution N-body simulations, we have studied the substructure inside Milky Way-size halos in a WDM cosmological scenario. We have also addressed the question whether the density profiles of host and guest halos are different from their corresponding CDM ones. Our main conclusions are:  1. Despite the fact that the power spectrum of fluctuations is suppressed at small scales, a non-negligible number of virialized structures, corresponding to these or smaller scales, form and survive within larger structures. The accretion rate of small halos is found to be less in the WDM scenario than in the CDM one. This is simply explained by the fact that a smaller number of small halos are avalaible for their incorporation into host halos in the WDM models. A higher satellite destruction rate is found in the WDM scenario as compared with the one in the CDM model: it can be accounted for the fact that guest halos are less concentrated by about a factor two in average in the WDM models. The less efficient halo accretion and the higher satellite destruction have almost the same weight as far as the final count of satellites within host halos at $z = 0$ is concerned. The larger the \\fsl\\ value the greater the size of these two effects, and the smaller the abundance of satellites.  2. The predicted maximum circular velocity function of guest halos that seems to best fit the observed one for satellites in the Milky Way and Andromeda is that given by the $\\fsl=0.1\\ \\mpc$ model. This \\fsl\\ value corresponds to a warmon of mass about 1 keV.  3.  For the $\\fsl =0.1$ and 0.2\\ \\mpc\\ models, guest halos ($M_h \\approx 10^9-10^{11}\\ \\msunh$) have a concentration parameter $c_{1/5}$ which is roughly twice smaller than that of the CDM halos. For those guest halos whose density profiles are reasonably well described by a NFW parametric fit ($\\sim 15\\%$), the $c_{\\rm NFW}$ parameter is roughly $1.5-3.0$ times lower than that of the CDM halos. This difference in the concentration parameters, for both $c_{1/5}$ and $c_{\\rm NFW}$, vanishes as we go to more massive halos.  4. The density profile of the host halos ($M_{\\rm vir} \\approx 1-3 \\times 10^{12}\\ \\msunh$) is well described by the NFW profile (in some cases, the inner slope is slightly shallower than $r^{-1}$). The guest halos have a wide variety of density profiles and those whose masses are below the corresponding cut-off mass scale probably present a small shallow core. The poor mass resolution of the simulations at these scales limits our predictions.  In summary, we have shown that in the WDM model with $\\fsl \\approx 0.1 \\mpc$ or $m_W \\approx 1\\ \\kev$ the degree of substructure within a Milky Way-size halo is much lower than in the CDM model and is in agreement with observations, if one assumes that each guest halo hosts a luminous galaxy. The problem of cuspy halos probably still persists in the WDM scenario, although we have found that halos ---in particular the small ones--- are less concentrated than the corresponding CDM halos. If the inclusion of baryonic matter helps to significantly avoid disruption of substructure, then the agreement with observation may continue for less massive WDM candidates for which the rms velocity is larger. In this case, the formation of shallow cores in dwarf galaxy like halos is expected. Nevertheless, as we have shown, halos of larger masses will not have a shallow core."
    },
    "0004/hep-ph0004151_arXiv.txt": {
        "abstract": "\\baselineskip 16pt \\tightenlines  The current Super-Kamiokande data on the D-N asymmetry between the the day event rate and the {\\it Night} ({\\it Mantle}) and {\\it Core} event rates, produced by solar neutrinos which respectively cross the Earth along any trajectory (cross the Earth mantle but do not cross the core), and cross the Earth core before reaching the detector, imply rather stringent constraints on the MSW small mixing angle (SMA) $\\nu_e \\rightarrow \\nu_{\\mu(\\tau)}$ solution of the solar neutrino problem. A simplified analysis shows, in particular, that a substantial subregion of the SMA solution region is disfavored by these data. The {\\it Core} D-N asymmetry data alone allow to rule out at 99.7\\% C.L. a part of this subregion. The constraints on the MSW large mixing angle and LOW $\\nu_e \\rightarrow \\nu_{\\mu(\\tau)}$ solutions as well as on the MSW $\\nu_e \\rightarrow \\nu_{s}$ solution, following from the data on the {\\it Mantle}, {\\it Night} and {\\it Core} D-N asymmetries are also discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004378_arXiv.txt": {
        "abstract": "The Tycho--2 Catalogue, released in February 2000, is based on the ESA Hipparcos space mission data and various ground--based catalogs for proper motions. An external comparison of the Tycho--2 astrometry is presented here using the first U.S.~Naval Observatory CCD Astrograph Catalog (UCAC1). The UCAC1 data were obtained from observations performed at CTIO between February 1998 and November 1999, using the 206 mm aperture 5--element lens astrograph and a 4k x 4k CCD. Only small systematic differences in position between Tycho--2 and UCAC1 up to 15 milliarcseconds (mas) are found, mainly as a function of magnitude. The standard deviations of the distributions of the position differences are in the 35 to 140 mas range, depending on magnitude. The observed scatter in the position differences is about 30\\% larger than expected from the combined formal, internal errors, also depending on magnitude. The Tycho--2 Catalogue has the more precise positions for bright stars (V$\\le 10^{m}$) while the UCAC1 positions are significantly better at the faint end ($11^{m} \\le V \\le 12.5^{m}$) of the magnitude range in common. UCAC1 goes much fainter (to R$\\approx 16^{m}$) than Tycho--2; however complete sky coverage is not expected before mid 2003. ",
        "introduction": "Two new major astrometric catalogs became available in early 2000. The Tycho--2 Catalogue for the brightest 2.5 million stars \\citep{tycho2s}, \\citep{tycho2l} and the first U.S.~Naval Observatory CCD Astrograph Catalog, UCAC1, for 27 million stars on the Southern Hemisphere \\citep{ucac1}. Both catalogs are important steps towards the extension of the optical reference frame \\citep{iau_jd} beyond densities and magnitudes of the Hipparcos Catalogue.  The Tycho--2 Catalogue is a new, extended version of the original Tycho Catalogue \\citep{hip_tycho}, based on a re--reduction of the ESA Hipparcos \\mbox{space} mission Tycho data and over 140 ground--based catalogs for the Tycho--2 proper motions. An external comparison of the Tycho--2 {\\em astrometry} is presented here utilizing the UCAC1 positions. A similar comparison between the Tycho--1 Catalogue and CCD astrograph test data was \\mbox{presented} earlier \\citep{veniceT}.  The U.S.~Naval Observatory CCD Astrograph (UCAC) project was planned and initiated in the mid 1990s \\citep{gaus1}, \\citep{1k_res}, \\citep{veniceU}. Observations started at the South Celestial Pole in 1998 and full sky coverage is expected by mid 2003. For the comparison presented in this paper, these new, high precision observations were available for about 80\\% of the Southern Hemisphere, covering the magnitude range R$\\approx 8-16^{m}$. Thus, particularly, the new, faint extension of the Tycho--2 catalog ($11^{m}$ to $12.5^{m}$) is very well covered by these independent ground--based observations.  Both catalogs are on the Hipparcos system, thus the International Celestial Reference System (ICRS). The epoch difference of about 8 years is bridged by proper motions given in the Tycho--2 catalog at the expense of introducing a third error contribution besides the positional errors of both catalogs. The UCAC is not a photometric catalog, with only approximate magnitudes given in a single bandpass (red).  Therefore {\\em no} external {\\em photometric} comparison of the Tycho--2 data can be presented here. Both Tycho--2 and UCAC1 are of great importance to the general astronomical community and the astrometric comparison presented here is also of benefit for the Tycho--2 and in particular the UCAC projects. Another important catalog comparison between the Tycho--2, the \\mbox{ACT} and the Hipparcos Catalogue is in preparation \\cite{t2acthip}. ",
        "conclusions": "The Tycho--2 and UCAC1 positions overall are in good agreement with systematic differences being very small by current astrometric catalog \\mbox{standards}. This external comparison shows the observed scatter in the position differences to be 10 to 50\\% larger than the formal, internal precisions predict. A significant fraction of the additional errors can be explained by systematic errors present in the current UCAC1, which by all means is a preliminary catalog. The UCAC1 contains the best positions available today (at current epochs) for stars in the 11 to 16 magnitude range. UCAC today covers only 80\\% of the Southern Hemisphere with no data in the North. When completed in 2003 it will form the basis for the FAME \\cite{fame} input catalog. The Tycho--2 Catalogue covers all sky and is the preferred astrometric reference down to about magnitude V=10.5 or wherever no UCAC1 data are available."
    },
    "0004/astro-ph0004072_arXiv.txt": {
        "abstract": "The dependence of star formation rate  on galaxian environment is a key issue in the understanding of galaxy formation and evolution. However, the study of this subject is complex and observationally challenging. This paper reviews some of the current results, drawing mostly from recent large redshift surveys such the LCRS, the MORPH collaboration, and the CNOC1 and CNOC2 redshift surveys. ",
        "introduction": "Tremendous strides have been made in the last few years in understanding the average star formation history of the universe from the present to as far back as 90\\% of the age of the universe (e.g., Steidel et al.~1999, Sawicki et al.~1997). However, the overall picture remains crude, and is still subject to many uncertainties and systematics. Most of the advances have been based on converting the UV luminosity density in different epochs into an average star formation rate (SFR). This requires averaging over volume and extrapolating over luminosity. The universe contains significant large scale structures even at epochs as early as $z\\sim3$ (Giavalisco et al.~1998), as galaxies formed initially in the most massive dark matter halos. Voids, filaments, and rich clusters mark very different environments in which galaxies reside. To come to a full physical understanding of galaxy formation and galaxy evolution, we need to have a more detailed picture, such as  the history of star formation in different environments. This is an observationally challenging task, as currently, we have barely begun to study these issues in the relatively nearby universe.  There are two major observational questions that we ultimately would like to answer. First, how does the SFR depend on the environment?  Second, how does this dependence evolve with time? Currently we have some idea of the answers to the first question, and little or no information on the second. Simple physical arguments provide some expectations on the influence of environments on SFR; however, these are not always clear cut. Processes that lower gas content of a galaxy are expected to decrease SFR; these include, e.g., ram-pressure stripping and evaporation in rich environments; tidal stripping from close encounters of galaxies; and the decrease of the accretion rate of new gas into a galaxy in rich environments. On the other hand, similar processes can also serve to {\\it increase} the SFR: e.g., ram-pressure and tidal shocks, and mergers and harassment of galaxies in close encounters.  To obtain definitive and quantitative conclusions, well-controlled, large samples (in the thousands) of galaxies with redshift, multi-color photometry, and spectroscopic information are required. (Perhaps large, robust photometric redshift samples can provide some needed advancement also.) Examples of completed redshift surveys that meet these goals include the LCRS (e.g., Shectman et al.~1996) for $z\\sim0.1$, the CNOC1 cluster redshift survey (Yee et al.~1996) with 2600 redshifts in fields of galaxy clusters at $0.17<z<0.55$, and the CNOC2 field galaxy redshift survey (Yee et al.~2000) with 6000 redshifts at $0.1<z<0.6$.  In this review, which is not intended to be complete, I will look at three issues. The first is the star formation--environment dependence in the local universe. We will then survey the current knowledge of differential cluster-field evolution, mostly by looking at galaxy populations in rich galaxy clusters at moderate redshifts. Finally some conjectures and remarks are made regarding the evolution of SFR and its environmental dependence. ",
        "conclusions": "The interplay between galaxian environment and star formation is clearly a key issue in the understanding of galaxy formation and evolution. It is also a difficult and complex problem and requires extensive observational resources to make definitive progress. There are clear indications of the effect of the environment when we compare extreme cases of the cores of rich clusters and the field. However, the environmental effect on the rate of galaxy evolution may be subtle, in that over a less extreme range definitive differences in relative evolution have not been observed conclusively. Some of the most intriguing and useful work in this area may come from the direct observations of the effects of a changing environment on star formation and galaxy evolution as found in the infalling galaxies at the outskirts of rich clusters. A thorough investigation of the phenomenon covering  a large span of cluster-centric radius and over a range of redshift may provide a crucial piece of the puzzle. More detailed investigations into the star formation properties of galaxies in clusters, such as those carried out by Moss \\& Whittle (2000, in these proceedings) using H$\\alpha$ narrow-band imaging, has the possibility of providing physical details on how the star formation process is affected when a galaxy falls into a cluster."
    },
    "0004/astro-ph0004244_arXiv.txt": {
        "abstract": "s{The properties of nearby galaxy clusters limit the range of cosmological parameters consistent with our universe.  We describe the limits which arise from studies of the intracluster medium (ICM) mass fraction \\fICM\\ and consideration of the possible sources of systematic error: $\\Omega_{M}<0.44h_{50}^{-1/2}$ at 95\\% confidence.  We emphasize that independent of Type Ia supernovae (SNe Ia) observations, this cluster study, taken together with published cosmic microwave background (CMB) anisotropy studies, indicates a non-zero quintessence or dark energy component $\\Omega_{Q}>0$.\\\\ We then discuss future galaxy cluster surveys which will probe the abundance of galaxy clusters to intermediate and high redshift. We investigate the sensitivity of these surveys to the cosmological density parameter $\\Omega_{M}$ and the equation of state parameter $w$ of any quintessence component.  In particular, we show that cluster survey constraints from a proposed large solid angle X-ray survey are comparable in precision and complementary in nature to constraints expected from future CMB anisotropy and SNe Ia studies.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004134_arXiv.txt": {
        "abstract": "Increasing evidence suggests that most of the energy density of the universe consists of a dark energy component with negative pressure, a ``cosmological constant\" that causes the cosmic expansion to accelerate. In this paper, we address the puzzle of why this component comes to dominate the universe only recently rather than at some much earlier epoch. We present a class of theories based on an evolving scalar field where the explanation is based  entirely on internal dynamical properties of the solutions. In the theories we consider,  the dynamics causes the scalar field to   lock automatically into a negative pressure state at the onset of matter-domination such that the present epoch is the earliest possible time, consistent with nucleosynthesis restrictions, when it can start to dominate. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004302_arXiv.txt": {
        "abstract": "We present a possible Cepheid-like luminosity estimator for the long gamma-ray bursts based on the variability of their light curves.  To construct the luminosity estimator, we use {\\it CGRO}/BATSE data for 13 bursts, {\\it Wind}/KONUS data for 5 bursts, {\\it Ulysses}/GRB data for 1 burst, and NEAR/XGRS data for 1 burst.  Spectroscopic redshifts, peak fluxes, and high resolution light curves are available for 11 of these bursts; partial information is available for the remaining 9 bursts. We find that the isotropic-equivalent peak luminosities $L$ of these bursts positively correlate with a rigorously-constructed measure $V$ of the variability of their light curves.  We fit a model to these data that accommodates both intrinsic scatter (statistical variance) and extrinsic scatter (sample variance).  We find that $L \\sim V^{3.3^{+1.1}_{-0.9}}$.  If one excludes GRB 980425 from the fit on the grounds that its association with SN 1998bw at a redshift of $z = 0.0085$ is not secure, the luminosity estimator spans $\\approx 2.5$ orders of magnitude in $L$, and the slope of the correlation between $L$ and $V$ is positive with a probability of $1 - 1.4 \\times 10^{-4}$ (3.8 $\\sigma$).  Although GRB 980425 is excluded from this fit, its $L$ and $V$ values are consistent with the fitted model, which suggests that GRB 980425 may well be associated with SN 1998bw, and that GRB 980425 and the cosmological bursts may share a common physical origin.  If one includes GRB 980425 in the fit, the luminosity estimator spans $\\approx 6.3$ orders of magnitude in $L$, and the slope of the correlation is positive with a probability of $1 - 9.3 \\times 10^{-7}$ (4.9 $\\sigma$).  In either case, the luminosity estimator yields best-estimate luminosities that are accurate to a factor of $\\approx 4$, or best-estimate luminosity distances that are accurate to a factor of $\\approx 2$.  Independently of whether or not GRB 980425 should be included in the fit, its light curve is unique in that it is much less variable than the other $\\approx 17$ light curves of bursts in our sample for which the signal-to-noise is reasonably good. ",
        "introduction": "Since gamma-ray bursts (GRBs) were first discovered (Klebesadel, Strong, \\& Olson 1973), thousands of bursts have been detected by a wide variety of instruments, most notably, the Burst and Transient Source Experiment (BATSE) on the {\\it Compton Gamma-Ray Observatory (CGRO)}, which will have detected about 2700 bursts by the end of {\\it CGRO}'s more than 9 year mission in 2000 June (see, e.g., Paciesas et al. 1999).  However, the distance scale of the bursts remained uncertain until 1997, when BeppoSAX began localizing long bursts to a few arcminutes on the sky, and distributing the locations to observers within hours of the bursts.  This led to the discovery of X-ray (Costa et al. 1997), optical (van Paradijs et al. 1997), and radio (Frail et al. 1997) afterglows, as well as host galaxies (Sahu et al. 1997).  Subsequent observations led to the spectroscopic determination of burst redshifts, using absorption lines in the spectra of the afterglows (see, e.g., Metzger et al. 1997), and emission lines in the spectra of the host galaxies (see, e.g., Kulkarni et al. 1998a). To date, redshifts have been measured for 13 bursts.  Recently, Stern, Poutanen, \\& Svensson (1999; see also Stern, Svensson, \\& Poutanen 1997), Norris, Marani \\& Bonnell (2000; see also Norris, Marani \\& Bonnell 1999), and Fenimore \\& Ramirez-Ruiz (2000; see also Ramirez-Ruiz \\& Fenimore 1999) have proposed trends between burst luminosity and quantities that can be measured directly from burst light curves, for the long bursts.  Using 1310 BATSE bursts for which peak fluxes and high resolution light curves were available, Stern, Poutanen, \\& Svensson (1999) have suggested that simple bursts (bursts dominated by a single, smooth pulse) are less luminous than complex bursts (bursts consisting of overlapping pulses); however, see \\S 5.  Using a sample of 7 BATSE bursts for which spectroscopic redshifts, peak fluxes, and high resolution light curves were available, Norris, Marani \\& Bonnell (2000) have suggested that more luminous bursts have shorter spectral lags (the interval of time between the peak of the light curve in different energy bands).  Using the same 7 bursts, Fenimore \\& Ramirez-Ruiz (2000) have suggested that more luminous bursts have more variable light curves.  These trends between luminosity and quantities that can be measured directly from light curves raise the exciting possibility that luminosities, and hence luminosity distances, might be inferred for the long bursts from their light curves alone.  In this paper, we present a possible luminosity estimator for the long bursts, the construction of which was motivated by the work of Ramirez-Ruiz \\& Fenimore (1999) and Fenimore \\& Ramirez-Ruiz (2000). We term the luminosity estimator ``Cepheid-like'' in that it can be used to infer luminosities and luminosity distances for the long bursts from the variabilities of their light curves alone.  We apply this luminosity estimator to every long burst in the current BATSE catalog in an upcoming paper (Reichart et al. 2000).  We rigorously construct our measure $V$ of the variability of a burst light curve in \\S 2.  Qualitatively, $V$ is computed by taking the difference of the light curve and a smoothed version of the light curve, squaring this difference, summing the squared difference over time intervals, and appropriately normalizing the result.  Our variability measure differs from that of Fenimore \\& Ramirez-Ruiz (2000) in three important ways: (1) we define the timescale over which the light curve is smoothed differently; (2) we subtract out the contribution to the variability due to Poisson noise differently; and (3) we combine variability measurements of light curves acquired in different energy bands into a single measurement of a burst's variability differently.  We find that only smoothing timescales that are proportional to burst duration appear to lead to significant correlations between the isotropic-equivalent peak photon luminosity $L$ of a burst and $V$; in particular, smoothing timescales that are a fixed duration in the source frame do not.  We take the smoothing timescale of a burst light curve, acquired in observer-frame energy band $E$, to be the time spanned by the brightest $100f$\\% of the total counts above background, where $f$ is a number between 0 and 1 that we fix using the data in \\S 4.  We schematically illustrate this measure of duration in Figures 1 and 2.  In Figure 1, the 50\\% smoothing timescale of the light curve is given by $T_{f=0.5}^E = T_1 + T_2$, and the 90\\% smoothing timescale of the light curve is given by $T_{f=0.9}^E = T_3$; in Figure 2, the 50\\% smoothing timescale of the light curve is given by $T_{f=0.5}^E = T_4$, and the 90\\% smoothing timescale of the light curve is given by $T_{f=0.9}^E = T_5 + T_6$.  We have chosen this measure of the smoothing timescale over others because it is robust:  small variations in either the light curve or the value of $f$ result in only small variations in the value of the smoothing timescale.  This is not the case with measures like $T_{50}$ and $T_{90}$ (see, e.g., Paciesas et al. 1999).  For example, consider the case of a burst with a precursor.  The value of, say, $T_{90}$ can differ considerable if the precursor's counts above background is $< 5$\\% versus $> 5$\\% of the total counts above background.  Likewise, if the precursor's counts above background is, say, 5\\% of the total counts above background, the duration that one measures can differ considerably if one uses $T_{<90}$ versus $T_{>90}$.  The effect of using a smoothing timescale that is artificially long (short) is measuring a variability that is artificially high (low).  Our measure of the smoothing timescale does not suffer from such robustness problems.  We present our measure of the isotropic-equivalent peak photon luminosity $L$ of a burst in \\S 3.  In \\S 4, we expand the original Ramirez-Ruiz \\& Fenimore (1999) sample of 7 bursts to include a total of 20 bursts, including 13 BATSE bursts, 5 {\\it Wind}/KONUS bursts, 1 {\\it Ulysses}/GRB burst, and 1 NEAR/XGRS burst.  Spectroscopic redshifts, peak fluxes, and high resolution light curves are available for 11 of these bursts; partial information is available for the remaining 9 bursts.  Also in \\S 4, we construct our luminosity estimator, which differs from that of Fenimore \\& Ramirez-Ruiz (2000) in two important ways:  (1) applying the Bayesian inference formalism developed by Reichart (2000), we fit a model to the data that accommodates both intrinsic scatter (statistical variance) in two dimensions and extrinsic scatter (sample variance) in two dimensions; and (2) again applying this Bayesian inference formalism, we determine not only the best estimate for $L$ as a function of $V$, but also the uncertainty in $L$ as a function of $V$, as well as approximate these functions with analytic expressions.  We state our conclusions in \\S 5. ",
        "conclusions": "We have presented a rigorously-constructed measure of the variability of a burst's light curve.  Using this variability measure and a sample of 20 bursts, consisting of every burst for which redshift information is currently available, we have shown that a significant correlation exists between the variability of a burst's light curve, and the burst's isotropic-equivalent peak photon luminosity.  This correlation between variability and luminosity is in agreement with the trends found by Stern, Poutanen, \\& Svensson (1999) and Fenimore \\& Ramirez-Ruiz (1999).  That is, more variable (``complex'' in the terminology of Stern, Poutanen, \\& Svensson 1999) bursts are more luminous, while less variable (``simple'' in the terminology of Stern, Poutanen, \\& Svensson 1999) bursts are less luminous.\\footnote{However, we must draw attention to a potential disagreement between the primary conclusion of Stern, Poutanen, \\& Svensson (1999), and the width of the luminosity distribution of the multiply-peaked bursts with spectroscopically-measured redshifts in our sample.  Stern, Poutanen, \\& Svensson (1999) find that the differential peak count rate distribution of their complex, or multiply-peaked BATSE bursts peaks about a factor of 4 above threshold, while the differential peak count rate distribution of their simple, or singly-peaked BATSE bursts does not have a similar peak.  They interpret this to mean that their complex bursts are more luminous and at higher redshifts than their simple bursts:  they argue that this peak corresponds to the peak in the star-formation history of the universe at $z \\sim 1.5$, and that it is not a threshold effect.  However, as this peak is narrow, spanning less than an order of magnitude, this could only be the case if the luminosity function of the complex bursts is similarly narrow; otherwise, this feature would be washed out. We have visually examined the light curves of the bursts with spectroscopically-measured redshifts in our sample, and find that only GRB 980425 and GRB 970508 appear to be singly-peaked.  The remaining, multiply-peaked bursts span $\\approx 2.5$ orders of magnitude, which appears to contradict their interpretation of this peak that they find in the differential peak count rate distribution of their complex bursts.}  Furthermore, from the correlated variabilities and luminosities of our sample of 20 bursts, we have constructed a possible Cepheid-like luminosity estimator for the long bursts.  If one excludes GRB 980425 from the fits, the luminosity estimator spans $\\approx 2.5$ orders of magnitude in luminosity, and its slope is positive with a probability of $1 - 1.4 \\times 10^{-4}$ (3.8 $\\sigma$).  GRB 980425, however, is consistent with the fitted model.  If one includes this burst in the fits, the luminosity estimator spans $\\approx 6.3$ orders of magnitude in luminosity, and its slope is positive with a probability of $1 - 9.3 \\times 10^{-7}$ (4.9 $\\sigma$).  Future bursts will either increase these probabilities, or possibly disprove the luminosity estimator.  However, since the uncertainty in $L$ as a function of $V_f$ in Figures 10 and 13 is dominated by extrinsic scatter (i.e., $\\sigma_{\\log{V_f}}$), and not by the uncertainties in the fitted values of the model parameters, a larger sample of bursts is unlikely to improve the predictive power of the luminosity estimator.  Currently, the luminosity estimator yields best-estimate luminosities that are accurate to a factor of $\\approx 4$, or best-estimate luminosity distances that are accurate to a factor of $\\approx 2$.  However, independently of whether or not the luminosity estimator is eventually disproved, the light curve of GRB 980425 is unique in that it is much less variable than the other $\\approx 17$ light curves of bursts in our sample for which the signal-to-noise is reasonably good: $\\log{V_{f=0.45}} < -1.5$ with a probability of $1 - 3.4 \\times 10^{-4}$ (3.4 $\\sigma$), and $\\log{V_{f=0.45}} < -1$ with a probability of $1 - 3.5 \\times 10^{-23}$ (9.9 $\\sigma$).  The argument has been made that the association of GRB 980425 with SN 1998bw at the unusually low redshift of $z = 0.0085$ probably is accidental because the light curve of GRB 980425 is no different than the light curves of the cosmological bursts.  On the contrary, we find that the light curve of GRB 980425 is different from the light curves of the cosmological bursts.  Consequently, GRB 980425 may well be associated with SN 1998bw.  If GRB 980425 is associated with SN 1998bw, and if the luminosity estimator is correct, the fact that GRB 980425 is consistent with the fitted model suggests that GRB 980425 and the cosmological bursts may share a common, or at least a related, physical origin, although they cannot share a common redshift distribution and/or luminosity function (Graziani, Lamb, \\& Marion 1999).  This conclusion is made more intriguing by the recent discoveries of supernova-like components to the late afterglows of the cosmological bursts GRB 970228 (Reichart 1999, Galama et al. 2000; Reichart, Castander, \\& Lamb 2000; Reichart, Lamb, \\& Castander 2000) and GRB 980326 (Bloom et al. 1999)."
    },
    "0004/astro-ph0004316_arXiv.txt": {
        "abstract": "We present new X-ray observations of the high mass X-ray binary (HMXRB) pulsar OAO 1657-415, obtained during one orbital period (10.44 days) with the Rossi X-Ray Timing Explorer (RXTE). Using the binary orbital parameters, obtained from Burst and Transient Source Experiment (BATSE) observations, we resolve the fluctuations in the pulse frequency at time scales on the order of one day for the first time. Recent BATSE results showed that OAO 1657-415 has spin-up/down trends in its pulse frequency time series, without any correlation with the X-luminosity at energies $>$20 keV (Baykal 1997). In the present RXTE observations the source is found to be in an extended phase of spin-down. We also find a gradual increase in the X-ray luminosity which is correlated with a marginal spin-up episode. The marginal correlation between the gradual spin-up (or decrease in spin-down rate) and increase in X-ray luminosity suggests that the OAO 1657-415 is observed during a stable accretion episode where the prograde accretion disk is formed. ",
        "introduction": "The high mass X-ray binary source HMXRB source OAO 1657-415 (OAO 1653-40) was first detected by the Copernicus satellite (Polidan et al. 1978) in the 4-9 keV range. The HEAO-1 observations also showed 38.22 sec pulsations in the 1-40 keV and 40-80 keV bands (White $\\& $ Pravdo 1979, Byrne et al. 1981). Observations with Ginga and GRANAT (Kamata et al. 1990, Gilfanov et al. 1991, Mereghetti et al. 1991, Sunyaev et al. 1991) have found pulse period changes. BATSE observations of this source with the Compton Gamma Ray Observatory (CGRO) showed that OAO 1657-415 is in an $10.44^{d}$ binary orbit with an X-ray eclipse by the a stellar companion (Chakrabarty et al. 1993). The observed orbital parameters imply that the companion is a supergiant of spectral class B0-B6. The correlations between X-ray flux and pulse frequency derivatives ($\\dot \\nu $) fluctuations were investigated by using the previously published pulse frequencies and BATSE measurements (Baykal 1997). These correlations suggested that the formation of episodic accretion disks in the case of a stellar wind is the possible accretion mechanism.  In this paper, we present the short term pulse frequency fluctuations and X-ray fluxes of OAO 1657-415 in the light of recent RXTE observations. We have employed background subtraction by using the background models for the RXTE/PCA instrument and galactic ridge emission in the 2-50 keV range. Our X-ray flux and pulse frequency measurements find an increase in the X-ray flux which is correlated with the decrease in the spin-down rate (or marginal spin-up trend). ",
        "conclusions": "OAO 1657-415, has shown strong spin-up/down torques in its time history which can not be explained by wind accretion (Baykal 1997). The formation of temporary accretion disks was therefore considered. If accretion onto the neutron star is from a Keplerian disk (Ghosh \\& Lamb 1979), the torque on the neutron star is given by \\begin{equation} I\\dot \\nu  = n(w_{s})  \\dot M~l_{K}, \\end{equation} where $l_{K} = (GMr_{o})^{1/2}$ is the specific angular momentum added by a Keplerian disk to the neutron star at the inner disk edge $r_{o} \\approx 0.5 r_{A}$ where $r_{A} \\ = (2GM)^{-1/7} \\mu ^{4/7} \\dot M^{-2/7}$ is the Alfven radius, $\\mu$ is the neutron star magnetic moment, $n(w_{s}) \\approx 1.4 (1-w_{s}/w_{c})/(1-w_{s})$ is a dimensionless function that measures the variation of the accretion torque as estimated by the fastness parameter $w_{s}=\\nu  /\\nu _{K}(r_{o}) = 2 \\pi P^{-1} G^{-1/2} M^{-5/7} \\mu ^{6/7} \\dot M^{-3/7}$. Here $w_{c}$ is the critical fastness parameter where the accretion torque is expected to vanish at $w_{c} \\sim 0.35-0.85$ depending on the structure of the disk. According to this model the torque will cause a spin-up if the neutron star is rotating slowly ($w_{s}~<~w_{c}$) in the same sense as the circulation in the disk, or spin-down, if it is rotating in the opposite sense. Even if the neutron star is rotating in the same sense as the disk flow, the torque will spin-down the neutron star if it is rotating too rapidly ($w_{s}~>>~w_{c}$).  In such a model one should see positive correlation between pulse frequency derivative ($\\dot \\nu $) and moderate mass accretion rate ($\\dot M$) if the disk is rotating in the same sense as the neutron star (Baykal 1997).  Recent observations of accreting neutron stars have shown stochastic spin-up/down trends on time scales from days to a few years (Bildsten et al. 1997). Some of the sources switch from spin-up to spin-down states without showing great changes in their mass accretion rates (Bildsten et al. 1997). These unusual behaviors led Nelson et al. (1997) to the possibility of retrograde circulation of accretion disks. GX 4+1 shows correlation between the X-ray flux and the spin-down rate (Chakrabarty et al. 1997) which may suggest a retrograde accretion disk. In our RXTE observations, OAO 1657-415 has shown marginal correlation with accretion rate and pulse frequency change. This positive correlation strongly suggesting that the disk formed in the spin-down episode is in prograde direction. From the BATSE observations, it was concluded that the pulse frequency derivatives and X-ray flux were not well correlated and the 8 days was the minimum for the correlation time scale (Baykal 1997). This RXTE observation implies that the time-scale of correlation is short, only a few days. To see the exact nature of correlations between X-ray flux and pulse frequency derivatives, an even more extensive broad band X-ray observation should be carried out."
    },
    "0004/astro-ph0004120_arXiv.txt": {
        "abstract": "We test accurate models of Comptonization spectra over the high quality data of the \\BS long look at NGC 5548. The data are well represented by a plane parallel corona with an inclination angle of 30$^{\\circ}$, a soft photon temperature of 5 eV and a hot plasma temperature and optical depth of $kT_{\\rm e}\\simeq$ 360 keV and $\\tau\\simeq$ 0.1, respectively. If energy balance applies, such values suggest that a more ``photon-starved'' geometry (e.g. a hemispheric region) is necessary. The spectral softening detected during a flare,  appears to be associated to a decrease of the heating--to--cooling ratio, indicating a geometric and/or energetic modification of the disk plus corona system. The hot plasma temperature derived with the models above is significantly higher than that obtained fitting the same data with a power law plus high energy cut off model for the continuum. This is due to the fact that in anisotropic geometries Comptonization spectra show \"intrinsic\" curvature which moves the fitted  high energy cut-off to higher energies. ",
        "introduction": "The X-ray emission of Seyfert galaxies is commonly believed to be produced by Compton scattering of soft photons on a (thermal or non-thermal) population of hot electrons. The non--detection of Seyferts by Comptel and the high energy cut-offs indicated by OSSE have focused attention on thermal models (e.g. Sunyaev \\& Titarchuk, 1980). In this case the X--ray spectral shape is mainly determined by the temperature $\\Theta=kT_{\\rm e}/m_{\\rm e}c^2$ and the optical depth $\\tau$ of the scattering electrons, while the cut--off energy is related essentially to $\\Theta$. Moreover, if the Comptonizing region and the source of seed photons are {\\it coupled}, one can write an energy balance equation for the hot coronal plasma which determines a roughly constant value of the Compton parameter $y\\simeq 4\\Theta \\tau \\, (1+4\\Theta)(1+\\tau)$ (e.g., Svensson, 1996) and thus a one to one correpondence between $\\theta$ and $\\tau$.  The required value of $y$ depends on the fraction $f$ of the power dissipated in the corona and on geometry. In the following the limiting case $f=1$ is considered.\\\\ \\noindent In the present work (Petrucci et al., 1999, hereafter P99), we test Comptonization models over the high quality data of the (8 days) \\BS long look at the Seyfert I galaxy NGC 5548, deriving constraints on the physical parameters and geometry of the source. These data have already been studied in detail by Nicastro et al. (1999), modelling the continuum with a cut-off power law.  \\begin{figure}[h] \\centerline{\\psfig{file=models-iuesaxosse.ps, width=5cm,angle=-90}} \\caption{{\\bf (a):} Comptonized models for different geometries assuming $\\Theta=0.7$. We have also over-plotted a cut--off power law with $E_{\\rm c}=2kT_{\\rm e}$ in dashed line (cf. text for details).{\\bf (b):} \\BS data set of NGC 5548 with non--simultaneous IUE and OSSE data (from Magdziark et al., 1998) with the corresponding best fits Comptonization model (in slab geometry, solid line) and simple cut--off power law model (dashed line).} \\vspace*{-0.5cm} \\end{figure} ",
        "conclusions": "This \\BS observation of NGC 5548 allowed us to show that i) the temperature $kT_{\\rm e}$ of the Comptonizing plasma can be largely underestimated (up to a factor of 7 here) when derived from simple power law models with high energy cut-off; ii) the data are well fitted by a plane parallel corona model with an inclination angle of $30^{\\circ}$, a soft photon temperature of 5 eV, a hot plasma $kT_{\\rm e}\\simeq$ 360~keV and an optical depth $\\tau\\simeq 0.1$. The latter values suggest however that the hot Comptonizing gas, if in the shape of slab, is not in energy balance. A better agreement is obtained with an hemispherical geometry; iii) the change of state during the central part of the run clearly indicates a variation of the Compton parameter $y$, which could be due, as suggested by the data, to an increase of the cooling."
    },
    "0004/astro-ph0004250_arXiv.txt": {
        "abstract": "We have developed a code to compute multi-colour microlensing lightcurves for extended sources, including the effects of limb darkening and photospheric star spots as a function of spot temperature, position, size and lens trajectory. Our model also includes the effect of structure within the spot and rotation of the stellar source.  Our results indicate that star spots generally give a clear signature only for transit events. Moreover, this signature is strongly suppressed by limb darkening for spots close to the limb -- although such spots can be detected by favourable lens trajectories. ",
        "introduction": "The amplification for a point source microlensing event is dependent only on the projected separation between lens and source. However, if the source size is comparable to the Einstein radius and the projected separation is small, it is necessary to treat the source as finite, with the amplification being calculated as an integral over the source star. Thus, the microlensing lightcurve also contains information about the source surface brightness profile.  Finite source effects were considered by Gould~(1994) for the case in which a lens \\emph{transits} the star, allowing measurement of the lens proper motion and thus distinguishing between MC and Galactic MACHOs. Witt and Mao~(1994) showed that the magnification profile of an extended source with a constant surface brightness profile was significantly different from the point source treatment for a 100 $R_{\\sun}$ source microlensed by a 0.1 $M_{\\sun}$ MACHO. The ability to recover the lens mass, transverse velocity and the source size are considerable assets in extended source microlensing, however rare the events may be.  The opportunity to use such events to provide information about the source star has also been investigated. Simmons \\emph{et al.}~(1995) modelled the polarisation signature of an extended source microlensing event, assuming a pure electron scattering grey model atmosphere. They showed that including polarisation information significantly improves the accuracy of estimates of the source radius, as well as better constraining the radial surface brightness profile of the source. Gould~(1997) suggested the use of extended source events to determine the rotation speed of red giants. The use of extended sources to estimate source limb darkening parameters has also been discussed by several authors (cf. Hendry \\emph{et al.} 1998; Valls-Gabaud 1998). Such an event produces a chromatic signature as the lens effectively \\emph{sees} a star of different radius in different photometric colour bands. This provides an unambiguous signature of microlensing, as opposed to that from a variable star (Valls-Gabaud 1998). ",
        "conclusions": "With the advent of automated `alert' status for candidate events, the prospects for using microlensing as a tool for gravitational imaging and investigating stellar atmospheres are dramatically improved. Of crucial importance in such studies are coordinated global observations to achieve the high level of sampling required to constrain models effectively. Observations of this quality are \\emph{also} precisely what is required to extract useful information from the microlensing of extended sources -- making detailed modelling of such events a timely issue. Moreover, the intensive observation of second caustic crossings in binary events -- currently being pursued with the principal goal of detecting planetary companions -- would also provide powerful constraints on the range of spot-related phenomena discussed in this paper. We are currently, therefore, extending our analyses to consider the photometric, polarimetric and spectroscopic signatures of star spots lensed during line caustic crossing events."
    },
    "0004/astro-ph0004299_arXiv.txt": {
        "abstract": "We analyze the behavior of N/O and C/O abundance ratios as a function of metallicity as gauged by O/H in large, extant Galactic and extragalactic H~II region abundance samples. We compile and compare published yields of C, N, and O for intermediate mass and massive stars and choose appropriate yield sets based upon analytical chemical evolution models fitted to the abundance data. We then use these yields to compute numerical chemical evolution models which satisfactorily reproduce the observed abundance trends and thereby identify the most likely production sites for carbon and nitrogen. Our results suggest that carbon and nitrogen originate from separate production sites and are decoupled from one another. Massive stars (M$>$8~M$_{\\sun}$) dominate the production of carbon, while intermediate-mass stars between 4 and 8~M$_{\\sun}$, with a characteristic lag time of roughly 250~Myr following their formation, dominate nitrogen production.  Carbon production is positively sensitive to metallicity through mass loss processes in massive stars and has a pseudo-secondary character. Nitrogen production in intermediate mass stars is primary at low metallicity, but when 12+log(O/H)$>$8.3, secondary nitrogen becomes prominent, and nitrogen increases at a faster rate than oxygen -- indeed the dependence is steeper than would be formally expected for a secondary element. The observed flat behavior of N/O versus O/H in metal-poor galaxies is explained by invoking low star formation rates which flatten the age-metallicity relation and allow N/O to rise to observed levels at low metallicities. The observed scatter and distribution of data points for N/O challenge the popular idea that observed intermittent polluting by oxygen is occurring from massive stars following star bursts. Rather, we find most points cluster at relatively low N/O values, indicating that scatter is caused by intermittent increases in nitrogen due to local contamination by Wolf-Rayet stars or luminous blue variables. In addition, the effect of inflow of gas into galactic systems on secondary production of nitrogen from carbon may introduce some scatter into N/O ratios at high metallicities. ",
        "introduction": "``It is quite a three pipe problem'' (S. Holmes, quoted in Doyle 1891).  Carbon and nitrogen are among the most abundant of the chemical elements, and of obvious importance for life. Carbon is also a major constituent of interstellar dust. The main nuclear processes which generate these two elements are reasonably well understood -- the carbon must come predominantly from the triple-alpha reaction of helium, and nitrogen by the conversion of carbon and oxygen that occurs during the CNO cycles of hydrogen burning. A lingering problem, though, has been the lack of knowledge of which {\\it sites} are most important for their generation -- in particular do they come mainly from short-lived massive stars, or from longer-lived progenitors of asymptotic giant branch stars? Coupled with this is uncertainty over the form and magnitude of the dependence of their production on the metallicity of the stars in which the reactions take place. Metallicity can essentially be tracked by its principal component -- oxygen -- whose dominant source is the Type~II supernova explosion of massive stars. There is observed variation in the ratios of carbon and nitrogen to oxygen (i.e. C/O, N/O) in both stars (e.g. Gustafsson et al. 1999) and the gas in galactic systems (e.g. Garnett et al. 1999; Henry \\& Worthey 1999), and it is these variations that we use as clues to pin down the synthesis sites.  The threshold temperature of He burning and production of $^{12}$C via the triple alpha process is $\\sim$10$^8$K, a temperature accessible in both massive (M$>$8M$_{\\sun}$) and intermediate mass (1$<$M$<$8M$_{\\sun}$) stars. Thus, these broad stellar groups represent two possible sites for carbon production. Likewise, nitrogen production via the CNO cycles may occur in either of these sites. However, discovering the origin of nitrogen is further complicated by the fact that the seed carbon needed for its production may either have been present when the star was born or is synthesized within the star during its lifetime. We now explore this idea in more detail.  Nitrogen is mainly produced in the six steps of the CN branch of the CNO cycles within H burning stellar zones, where $^{12}$C serves as the reaction catalyst (see a textbook like Clayton 1983 or Cowley 1995 for nucleosynthesis review).  Three reactions occur to transform $^{12}$C to $^{14}$N: $^{12}$C(p,$\\gamma$)$^{13}$N($\\beta$$^{+}$,$ \\nu$)$^{13}$C(p,$\\gamma$)$^{14}$N, while the next step, $^{14}$N(p,$\\gamma$)O$^{15}$, depletes nitrogen and has a relatively low cross-section. The final two reactions in the cycle transform $^{15}$O to $^{12}$C. Since the fourth reaction runs much slower than the others, the cycle achieves equilibrium only when $^{14}$N accumulates to high levels, and so one effect of the CN cycle is to convert $^{12}$C to $^{14}$N. The real issue in nitrogen evolution is to discover the source of the carbon which is converted into nitrogen, and of any oxygen which can contribute through the (slow) side chain $^{16}$O(p,$\\gamma$)$^{17}$F($\\beta$$^{+}$,$\\nu$)$^{17}$O(p,$\\alpha$)$^{14}$N.  The conventional meaning of ``{\\it primary}'' applied to nitrogen is that its production is independent of the initial composition of the star in which it is synthesized. An example is where stars produce their own carbon (and some oxygen) during helium burning, and the carbon (and perhaps oxygen) is subsequently processed into $^{14}$N via the CN(O) cycle. Stars beyond the first generation in a galactic system already contain some carbon and oxygen, inherited from the interstellar medium out of which they formed. The amount of nitrogen formed from CNO cycling of this material will be proportional to its C abundance (and also its O abundance, if the CNO cycling proceeds long enough to deplete the oxygen) and is known as ``{\\it secondary}'' nitrogen. In general, then, primary nitrogen production is independent of metallicity, while secondary production is a linear function of it.  However, these conventional definitions become rather blurred if the evolution of the synthesizing stars, and/or the release of nucleosynthesis products to the interstellar medium, depends on the initial composition of the star - as well it might if stellar wind generation is metallicity-dependent. Thus effective production of carbon could depend on the initial metallicity of the star, although no actual ``seed nucleus'' is involved, and the production of nitrogen might differ from a simple primary or secondary process. Also important is the mass of star in which production takes place, since if {\\it most} production takes place in stars of {\\it low} mass, a significant delay will occur between formation of the source stars and release of the products into the interstellar medium.  Detailed discussion and review of computed stellar yields is left until {\\S\\S}3 and 5 below, but we now refer to current {\\it interpretations} of the observed carbon and nitrogen abundances in stars and galaxies. The literature on the origin of carbon has recently been reviewed by Gustafsson et al. (1999), and Garnett et al. (1999). The former conclude that their own stellar results are ``consistent with carbon enrichment by superwinds of metal-rich massive stars, but inconsistent with a main origin of carbon in low mass stars'', which is pretty much echoed by the latter authors who state that the behavior of C/O ratios as a function of O/H is ``best explained [by] ... effects of stellar mass loss on massive star yields''. They note that theoretical chemical evolution models (Carigi 1994) in which carbon comes from intermediate stars apparently predict too shallow a C/O relation to fit their observed galaxy abundance gradients. Other recent discussions of C/O gradients across our own or other galaxies, or of C/O ratios in low metallicity galaxies, have been given by G\\\"{o}tz and K\\\"{o}ppen (1992), Prantzos, Vangioni-Flam \\& Chauveau (1994), Kunth et al. (1995), Garnett et al. (1995), Carigi et al. (1995), Moll\\'{a}, Ferrini, and D{\\'i}az (1997) and Chiappini, Matteucci \\& Gratton (1997).  The literature on the origin of nitrogen was briefly reviewed by Vila-Costas \\& Edmunds (1993). A major problem with N/O ratios has been to try to explain the {\\it spread} in N/O at a given O/H (see Figure 1B of Section 2), although the reality of a spread at low O/H (12+log(O/H) $\\leq$ 7.6) has been questioned by Thuan et al. (1995) and Izotov \\& Thuan (1999). Two major mechanisms have been proposed for generating a spread - mechanisms that could also apply (with different timescales etc) to C/O ratios. One mechanism invokes a significant time delay between formation of the star which will produce the nitrogen and the delivery of the nitrogen to the interstellar medium. Thus oxygen is expected to be produced predominantly in the SNII explosion of short-lived massive stars, the N/O ratio in the ISM will at first decrease, and then rise again as the nitrogen is released. A delay could affect both primary and secondary nitrogen, and Edmunds \\& Pagel (1978) suggested that N/O ratios might perhaps be an indicator of the age of a galactic system, in the sense that it indicated the time since the bulk of star formation has taken place. This idea continues to find some support (e.g. Kobulnicky \\& Skillman 1996; van Zee, Haynes \\& Salzer 1997). A suspicion of low N/S or N/Si ratios (the S and Si being expected to follow O) in damped Lyman alpha absorption systems (Lu et al. 1996; Pettini, Lipman \\& Hunstead 1995; Pilyugin 1999) has been invoked as evidence of the youth of the systems on the basis of delayed N release. However, if the delay mechanism is to be effective in altering N/O ratios, the delay must be reasonably long - otherwise the probability of catching a system at low N/O, perhaps after a burst of star formation, will be too small. A time scale of {\\it several} 10$^{8}$ or of order 10$^{9}$ years would seem to be necessary. We shall argue later that the dominant source of nitrogen may always be in intermediate mass stars of too high a mass to allow a strong, observable, systematic effect.  The second mechanism for causing N/O variation (or C/O variation, but to avoid repetition we only discuss nitrogen here) is variation in the flow of gas into or out of the galaxy. If the nitrogen is {\\it primary}, then neither inflow of unenriched gas nor outflow of interstellar medium will affect the N/O ratio {\\it except} in the case where the outflow is {\\it different} for nitrogen and oxygen (e.g. Marconi, Matteucci \\& Tosi 1994; Pilyugin 1993). If the nitrogen is {\\it secondary} (or of any composition behavior other than primary) the N/O ratio is still unaffected by any {\\it non-differential} outflow, but can be affected by unenriched inflow (Serano \\& Peimbert 1983; Edmunds 1990). K{\\\"o}ppen \\& Edmunds (1999) were able to place the useful constraint that variation in N/O caused by inflow and secondary nitrogen can be at most a factor of two if the inflow is time-decreasing (as most chemical evolution models tend to assume).  We will find - with no surprise - that nitrogen can be interpreted as having both primary and secondary components. But we will suggest that the nitrogen is (or acts as if it is) secondary on {\\it carbon}, rather than on oxygen. This will allow a rather steep dependence of N/O on O/H at high O/H, possibly aided by inflow effects. That the CNO cycle might not go to completion, but effectively stop at CN equilibrium, was noted in LMC planetary nebulae by Dopita et al. (1996) and (at low metallicity) from observations of old stars by Langer \\& Kraft (1984). We shall not discuss isotope ratios such as $^{15}$N/$^{14}$N here, although they are a subject of some recent interest (Chin et al. 1999; Wielen \\& Wilson 1997).  After reviewing the observational data in Section 2 and stellar yields in Section 3, we give an elementary analytic model which can account well for the general trends of the observational data. The yield parameters of this analytic model can then be compared with published stellar evolution and nucleosynthesis predicted yields with the result that it is possible to choose which are the most realistic yield calculations -- and to identify the stellar sources of carbon and nitrogen. The general form of the analytic models is confirmed by calculations generated by detailed chemical evolution codes in Section 5. After some discussion in Section 6, our identification of the sources and metallicity-dependence of carbon and nitrogen production are summarized in {\\S}7. ",
        "conclusions": "We have satisfactorily and simultaneously reproduced the C/O and N/O behavior in Figs.~3A,B by adjusting the star formation rate to suppress oxygen buildup until nitrogen can be ejected by IMS. Our results clearly suggest that carbon rises relative to oxygen because the former is produced in massive stars in a manner which is directly sensitive to metallicity through the mass loss process. On the other hand, the bimodal behavior of N/O is consistent with primary IMS nitrogen production at metallicities below 12+log(O/H) of 8.3, with metal-sensitive secondary production contributing significantly above this value. Both the C/O and high-metallicity N/O trends are amplified by the decrease in oxygen synthesis by massive stars as metallicity increases. This picture is in complete agreement with the yield predictions for massive and intermediate mass stars discussed in {\\S}3.  The low star formation rate at early times is necessary to allow log(N/O) to rise to about -1.4 at metallicities below those observed, as primary nitrogen production for an IMS population builds up to its maximum level. As metallicities continue to climb, IMS production of nitrogen and massive star production of oxygen remain in equilibrium and the element ratio constant. Prior to 250Myr in our model, the gas surface density is roughly equal to 4~M$_{\\sun}$/pc$^2$, corresponding to a star formation rate of 10$^{-3}$M$_{\\sun}$yr$^{-1}$kpc$^{-2}$.  A reasonable interpretation of this, then, is that the galaxies containing the low metallicity systems studied by Izotov \\& Thuan and Kobulnicky \\& Skillman (1996) and shown in Fig.~1B historically have had relatively low star formation rates due to their low surface densities. Indeed, results in Papaderos et al. (1996) indicate that a typical surface density for a BCG is 2-5~M$_{\\sun}$/pc$^2$, consistent with the surface gas densities in our model at early times. Observations by van~Zee et al. (1997) for low surface brightness galaxies are consistent with this level, as are the total surface densities for the outer disk regions of spirals presented in Vila-Costas \\& Edmunds (1992). Of course the star formation rate in these systems could be reduced because of the lack of environmental effects or spiral structure, or the inability by gas to cool efficiently enough. But whatever the cause, it seems clear that the data are consistent with low star formation rates in these objects.  Thus, our results nicely accomodate the popular picture of blue compact galaxies having an underlying old metal-poor population with star bursts superimposed on the systems. The N/O levels in these systems are set by maximized primary nitrogen production of IMS and oxygen production by massive stars. Our picture contrasts with the explanation by Izotov \\& Thuan which claims that these systems are necessarily very young because of their low metallicities. Then, since their IMS have not had time to begin releasing nitrogen, this forces the conclusion that the N/O ratio at low metallicities must be set by massive star primary nitrogen production. However, this picture neglects the influence of the star formation efficiency in the rate of buildup of metals. According to our analysis, these systems are simply evolving relatively slowly or intermittently due to historically low star formation rates, so that the delay in nitrogen release by IMS becomes insignificant.  More generally, any galaxy or galactic region which evolves slowly will maintain a relatively low metallicity over a significant fraction of a Hubble time, since total metallicity is directly related to the star formation rate integrated over time. This opens up the possibility that objects such as blue compact galaxies and outer regions of spiral disks are in fact several Gyr in age despite their low metallicities. This is consistent with the recent conclusions of Legrand (2000), who employed chemical evolution models to study abundances and continuum colors of I~Zw~18. Legrand found that models characterized by a star formation rate of $10^{-4}$ M$_{\\sun}$ yr$^{-1}$ over 14 Gyr explain the observations. This general explanation may also apply to the outer regions of spiral disks, where CNO abundance ratios are similar to those observed in dwarf galaxies such as blue compact objects (Garnett et al. 1999).  Our numerical models allow us to track the separate contributions of IMS and massive stars to carbon and nitrogen evolution. The left and right sides of Fig.~5 show our predictions for carbon and nitrogen, respectively, for our best numerical model, model~B. For each element the fraction of the total mass of carbon or nitrogen being ejected at a point in time is shown in the upper panels for IMS and massive stars versus time. Meanwhile, the lower panels plot the fraction of accumulated mass of carbon and nitrogen ejected by IMS and massive stars, also versus time.  To gauge the final contributions of IMS and massive stars to carbon and nitrogen synthesis, we can simply read the late-time values off of lower panels in Fig.~5. Clearly, in our models carbon production is dominated by massive stars, with massive stars providing 97\\% (all) of the total carbon. On the other hand, the roles of the stellar types appear to be reversed in the case of nitrogen, with IMS providing roughly 90\\% (all) of this element.  {\\it Our analysis has led to the conclusion that carbon and nitrogen production in the Universe are essentially decoupled from one another, with the former produced in massive stars, while the latter is produced in IMS.} Consensus toward this claim in the case of carbon has been building for several years. For example, work by Prantzos et al. (1994), Garnett et al. (1999), and Gustafsson et al. (1999) show that carbon is largely produced by massive stars, to which we now add our support. In the case of nitrogen, the predicted large IMS yields in conjunction with the large extant database of planetary nebula abundances suggest strongly that the production of this element is dominated by IMS. The most recent problem has been the one presented by the observed constant N/O value at low metallicities, suggesting that massive stars must play at least an important role here, due to supposed significant IMS nitrogen release delays. However, our models have shown that low star formation rates can accomodate IMS nitrogen production with no problem.  Finally, during this investigation we have chosen to ignore until now the large scatter in N/O that is observed at a single O/H value. This matter has been addressed by a number of authors including Garnett (1990), Pilyugin (1993; 1999), and Marconi, Matteucci, \\& Tosi (1994), with the consensus being that at least some of the scatter is real and due to bursts which momentarily lower N/O in the observed H~II regions with sudden injections of fresh oxygen, but as IMS eject nitrogen after the customary lag time, N/O rises again. Therefore, the scatter comes about by observing a large sample of H~II regions in various stages of oxygen and nitrogen enrichment. This picture implies that most points should be concentrated at relatively high N/O values, with fewer points, representing those objects experiencing sudden oxygen enrichment, located below the main concentration, since presumably bursts are followed by relatively long periods of quiescence during which the end point N/O is characteristic.  However, a close look at the vertical distribution of points in N/O in Fig.~1B reveals that most points seem to be clustered along the NO envelope at relatively low values with the concentration falling off as one considers higher N/O values, i.e. exactly the reverse of the standard interpretation of the scatter described above. Indeed, this empirical finding strongly suggests that the ``equilibrium'' or unperturbed locus where most H~II regions reside is the NO envelope, and thus the excursions caused by sudden injections of material are actually {\\it upward} toward the region of fewer points.  Barring an unidentified selection effect, then, the distribution of data points in Fig.~1B would seem to challenge the conventional picture. In fact the data are consistent with the lack of evidence for localized oxygen contamination from massive stars in H~II regions described in Kobulnicky \\& Skillman (1997b). Furthermore, the falloff in points above the NO envelope is more consistent with injections of nitrogen rather than oxygen; in this case the nitrogen source might be Wolf-Rayet stars or luminous blue variable stars, both of which were considered by Kobulnicky \\& Skillman (1997a) in their study of nitrogen-enriched H~II regions in NGC~5253. Their explanation suggests a simultaneous enrichment of helium, and thus H~II regions exhibiting high values of N/O in Fig.~1B should be checked for evidence of helium enrichment. Our general results, though, imply that the contributions of WR and luminous blue variable stars to nitrogen enrichment must be small. We postpone further investigation of this aspect of the N/O distribution, as we plan to take it up in another paper.  Finally, we note the upper limits of log(N/O) for four damped Ly$\\alpha$ systems taken from Lu et al. (1996). While two of the points lie within the NO envelope, the other two points lie about 0.25 dex below it. Since the position of the rising track of a numerical model in Fig.~3B can be forced to the right (higher star formation rates) and left (lower star formation rates), such objects and others at even lower N/O values (if eventually observed) may be explained as representing very early stages of N/O evolution when IMS are just beginning to release nitrogen, i.e. roughly 250 Myr or less after an initial star burst."
    },
    "0004/astro-ph0004066_arXiv.txt": {
        "abstract": "Self-similar solutions provide good descriptions for the gravitational collapse of spherical clouds or stars when the gas obeys a polytropic equation of state, $p=K\\rho^\\gamma$ (with $\\gamma\\le 4/3$, and $\\gamma=1$ corresponds to isothermal gas). We study the behaviors of nonradial (nonspherical) perturbations in the similarity solutions of Larson, Penston and Yahil, which describe the evolution of the collapsing cloud prior to core formation. Our global stability analysis reveals the existence of unstable bar-modes ($l=2$) when $\\gamma\\le 1.09$. In particular, for the collapse of isothermal spheres, which applies to the early stages of star formation, the $l=2$ density perturbation relative to the background, $\\delta\\rho({\\bf r},t)/\\rho(r,t)$, increases as $(t_0-t)^{-0.352}\\propto \\rho_c(t)^{0.176}$, where $t_0$ denotes the epoch of core formation, and $\\rho_c(t)$ is the cloud central density. Thus, the isothermal cloud tends to evolve into an ellipsoidal shape (prolate bar or oblate disk, depending on initial conditions) as the collapse proceeds. This shape deformation may facilitate fragmentation of the cloud. In the context of Type II supernovae, core collapse is described by the $\\gamma\\simeq 1.3$ equation of state, and our analysis indicates that there is no growing mode (with density perturbation) in the collapsing core before the proto-neutron star forms, although nonradial perturbations can grow during the subsequent accretion of the outer core and envelope onto the neutron star.  We also carry out a global stability analysis for the self-similar expansion-wave solution found by Shu, which describes the post-collapse accretion (``inside-out'' collapse) of isothermal gas onto a protostar. We show that this solution is unstable to perturbations of all $l$'s, although the growth rates are unknown. ",
        "introduction": "The gravitational collapse of molecular clouds leading to star formation has long been an active area of study. In the early stages of collapse (from $\\rho\\lo 10^{-19}$~g~cm$^{-3}$ to $\\rho\\sim 10^{-12}$~g~cm$^{-3}$) the gas remains approximately isothermal (at temperature $\\sim 10$~K) due to efficient cooling by dust grains (see, e.g., Myhill \\& Boss 1993). The gas dynamics is then specified by two dimensional parameters, the gravitational constant $G$ and the isothermal sound speed $a$, so that the flow is expected to approach a self-similar form in the asymptotic limit, when the memory of initial conditions is ``lost''. Larson (1969) and Penston (1969) found a similarity solution which describes the pre-collapse (i.e., before the central protostar forms) evolution of the cloud, in which the gas collapses from rest, accelerating until it cruises at Mach number of $3.3$ and the density profile reaches a $r^{-2}$ power law. The Larson-Penston solution contains a nonsingular homologous inner core and a supersonic outer envelope. A qualitatively different set of similarity solutions was found by Shu (1977). Of particular interest is Shu's expansion-wave solution which describes the post-collapse accretion of a singular isothermal gas cloud onto a protostar. In this solution, the flow starts from hydrostatic equilibrium (with a $r^{-2}$ density profile) and a rarefaction wave expands from the center and initiates the collapse (the so-called ``inside-out'' collapse); Inside the expansion-wave front, the flow eventually attains the free-fall behavior ($v\\propto r^{-1/2}$) at small radii, with density $\\rho\\propto r^{-3/2}$. The link between the Larson-Penston pre-collapse solution and Shu's expansion-wave solution was elucidated by Hunter (1977), who showed that the Larson-Penston solution can be continued to the post-collapse phase and that there exists an infinite (but discrete) number of pre- and post-collapse solutions of a different type (called ``Type I''; see \\S 2), among which the expansion-wave solution represents a limiting case. Figure 1 illustrates the properties of different self-similar solutions for the collapse and accretion of isothermal spheres.\\footnote{We note that Whitworth \\& Summers (1985) have found a continuum of similarity solutions by relaxing the analyticity condition of the flow at the transonic point; However, these generalized solutions are locally unstable (Hunter 1986; Ori \\& Piran 1988), and therefore may not be realized in astrophysical situations. We also mention that Boily and Lynden-Bell (1995) have constructed similarity solutions for the gravitational collapse of radiatively cooling gas spheres (with emissivity having a power-law dependence on density and temperature).}  With the plethora of possible similarity solutions, it is important to know which, if any, of them are actually realized by collapse of isothermal clouds. One-dimensional hydrodynamical simulations, starting from a regular (Bonner-Ebert) sphere, generally indicate that the collapse resembles the Larson-Penston similarity form in the asymptotic limit (Hunter 1977; Foster \\& Chevalier 1993). This is consistent with the recent finding of Hanawa \\& Nakayama (1997), who showed that the pre-collapse Type I solutions of Hunter's (see Fig.~1) are strongly unstable against global spherical perturbations, and therefore are unlikely to be realized in astrophysical situations or numerical simulations.  Similarity solutions have also been investigated in the context of core-collapse supernovae (Goldreich \\& Weber 1980; Yahil 1983), where the equation of state of the collapsing iron core can be approximated by that of a polytrope, $p=K\\rho^\\gamma$, where $K$ is a constant and $\\gamma\\simeq 4/3$. (In fact, the effective $\\gamma$ is about $1.3$ from the onset of electron capture to the neutrino trapping density, i.e., for $4\\times 10^9$~g~cm$^{-3}\\lo\\rho\\lo 10^{12}$~g~cm$^{-3}$; $\\gamma$ becomes close to $4/3$ when $\\rho\\go 10^{12}$~g~cm$^{-3}$ until nuclear density is reached.) Goldreich \\& Weber (1980) studied the special case of $\\gamma=4/3$, which provides a good description for the inner homologous core; They also performed a global perturbation analysis and showed that the inner core is stable against all radial and nonradial perturbations. Yahil (1983) generalized the Goldreich-Weber solution to general $\\gamma\\le 4/3$; this allows for a proper description of the outer core which collapses supersonically. Since Yahil's solution is the same as the Larson-Penston solution except for different values of $\\gamma$, we shall often refer them as Larson-Penston-Yahil solutions in the remainder of this paper.  The similarity solutions described above (in both star formation and supernova contexts) assume idealized spherical flows. A realistic gas cloud, however, contains nonradial (nonspherical) perturbations,\\footnote{A realistic flow may also contain a non-negligible amount of angular momentum and magnetic fields --- these are neglected in the main text of this paper. In Appendix A we discuss the perturbative effects of rotation on Larson-Penston-Yahil solutions. Terebey, Shu \\& Cassen (1984) have considered how slow rotation affects the expansion-wave solution, and Galli \\& Shu (1993a,b) have studied the perturbative effects of magnetic fields (see also Li \\& Shu 1997).} and it is of interest to understand the behaviors of these perturbations during the collapse/accretion of the cloud. In general, multi-dimensional hydrodynamical simulations are needed to follow the evolution of the perturbed flow, especially when the perturbations become nonlinear. The large dynamical range involved in a collapse makes such simulations particularly challenging (e.g., star formation ultimately involves collapse from $\\rho\\lo 10^{-19}$~g~cm$^{-3}$ to $\\rho\\go 0.1$~g~cm$^{-3}$; even the initial isothermal collapse stage involves seven orders of magnitude increase in densities; see Truelove et al.~1997,1998 and Boss 1998 for a discussion on the numerical subtleties). An alternative, complementary approach is to carry out linear stability analysis to determine whether the flow is unstable to the growth of any nonradial perturbations. Since the unperturbed flow varies in space and time in a self-similar manner, a global analysis is needed to study perturbations which vary on similar scales as the unperturbed flow itself. The stability properties of the flow therefore depend crucially on boundary conditions at different locations of the flow. In this paper we perform global stability analysis for Larson-Penston-Yahil solutions (general $\\gamma$) and for Shu's expansion-wave solution ($\\gamma=1$ only) to determine whether these similarity flows contain growing nonradial modes.  While the stability properties of isothermal similarity collapse solutions (the Larson-Penston solution and the expansion-wave solution) are relevant to the formation of binary (and multiple) stars (see \\S 5), the present study stems from our attempts to understand the origin of asymmetric supernovae and pulsar kicks (Goldreich, Lai \\& Sahrling 1996; see also Lai 1999). Numerical simulations indicate that local hydrodynamical instabilities in the collapsed stellar core (e.g., Burrows et al.~1995; Janka \\& M\\\"uller 1994,~1996; Herant et al.~1994), which can in principle lead to asymmetric matter ejection and/or asymmetric neutrino emission, are not adequate to account for kick velocities $\\go 100$~km~s$^{-1}$ (Burrows \\& Hayes 1996; Janka 1998). Global asymmetric perturbations of presupernova cores may be required to produce the observed kicks. Goldreich et al.~(1996) suggested that overstable g-modes driven by shell nuclear burning may provide seed perturbations which could be amplified during core collapse (see also Lai \\& Goldreich 2000). While the analysis of Goldreich \\& Weber (1980) shows that the inner homologous core is stable against nonradial perturbations, the situation is not so clear for the supersonically collapsing outer core where pressure plays a less important role. It is therefore important to analyse the global stability of Yahil's self-similar solution. Hanawa \\& Matsumoto (1999) have recently found a globally unstable bar mode in the pre-collapse Larson-Penston solution (for isothermal collapse). Our independent calculations confirm their result for $\\gamma=1$. Since the analysis of Hanawa \\& Matsumoto is restricted to perturbations with real eigenvalues and eigenfunctions (see \\S 3), it is not clear whether there exists any other growing modes, nor is it clear whether the growing bar-mode persists for general values of $\\gamma$ (see also Hanawa \\& Matsumoto 2000a).  The remainder of this paper is organized as follows. Section 2 summarizes the basic properties of the (unperturbed) Larson-Penston-Yahil similarity solution. This serves as a preparation for our stability analysis presented in Section 3. (For readers not interested in technical details, the main results are given in \\S 3.4 and Figures 3-5.) In Section 4 we show that Shu's expansion-wave solution for isothermal collapse is unstable to nonradial perturbations of all angular orders. Finally, we discuss the astrophysical implications of our results in \\S 5. Appendix A contains a discussion of the rotational and vortex modes of Larson-Penston-Yahil solutions. ",
        "conclusions": "Early studies by Hunter (1962) and by Lin, Mestel \\& Shu (1965) demonstrated that uniform, pressure-free gas clouds undergoing gravitational collapse are unstable toward fragmentation and shape deformation, with perturbations growing asymptotically as $\\delta\\rho(\\br,t)/\\rho(t)\\propto (t_0-t)^{-1}\\propto\\rho(t)^{1/2}$ in the linear regime, where $t_0$ denotes the epoch of complete collapse, and $\\rho(t)$ is the unperturbed uniform density. However, the presence of even a small initial central concentration and pressure forces significantly alters the evolution of the cloud. If the gas pressure is simply related to the density by a power-law, $p=K\\rho^\\gamma$ (polytropic equation of state), the flow asymptotically approaches the similarity solutions found by Larson (1969), Penston (1969) (for isothermal gas $\\gamma=1$), by Goldreich \\& Weber (1980) (for $\\gamma=4/3$), and by Yahil (1983) (for general $\\gamma$). Since the local Jeans length is of the same order as the length scale at which the flow varies, a global analysis is needed to determine the stability properties of the collapsing cloud. The result (\\S 3) presented in this paper (see also Hanawa \\& Matsumoto 1999) shows that for sufficiently soft equation of state ($\\gamma\\le 1.09$), the Larson-Penston-Yahil similarity flow is unstable against bar-mode perturbations, such that $\\delta\\rho(\\br,t)/\\rho(r,t)\\propto (t_0-t)^sY_{2m}(\\theta,\\phi)$ with $s<0$ ($s=-0.352$ for $\\gamma=1$ and $s$ increases to zero as $\\gamma$ increases to $1.09$, see Fig.~5), where $t_0$ denotes the epoch of core formation. Since the central density increases as $\\rho_c(t)\\propto (t_0-t)^{-2}$, the growth of perturbation, $\\delta\\rho(\\br,t)/\\rho(r,t)\\propto \\rho_c(t)^{-s/2}$, is slow (e.g., for isothermal collapse, $\\delta\\rho/\\rho$ increases by a factor of $1.5$ when $\\rho_c$ increases by a factor of $10$). Such a slow growth (compared with the $\\delta\\rho/\\rho\\propto\\rho^{1/2}$ behavior for the collapse of uniform, pressure-less gas) is a result of the stablizing influence of pressure, despite the large Mach number (about $3$) achieved in the outer region of the cloud.  Our stability analysis applies to the pre-collapse stage (prior to core formation) of the Larson-Penston-Yahil solutions. After the central core forms, the outer core and envelope accrete onto it (see Fig.~1). The gas approaches free-fall as $r\\rightarrow 0$, and the Mach number becomes much greater than unity. In this (accretion) stage, nonradial perturbations (of all scales) grow kinematically as $\\delta\\rho/\\rho\\propto r^{-1/2}\\propto \\rho^{1/3}$, where $r(t)$ is the radius of a fluid element and $\\rho(t)\\propto r^{-3/2}$ its comoving density (Lai \\& Goldreich 2000). Although the fluid element is free-falling, the perturbation grows more slowly compared with the case of uniform pressure-less collapse because the steep velocity gradients provide a stablizing influence on the flow.  The global bar-mode instability for isothermal collapse may have important implications for star formation, particularly in connection with the formation of binary (and multiple) stars (see also Hanawa \\& Matsumoto 1999; Matsumoto \\& Hanawa 1999). Fragmentation is unlikely to occur in a globally spherical collapse because small condensations do not contract fast enough to separate out from the converging bulk flow. Angular momentum (or magnetic field) can obviously make the cloud nonspherical, and thus facilitate fragmentation (e.g., Burkert \\& Bodenheimer 1996; Burkert, Bate \\& Bodenheimer 1997; Truelove et al.~1997,1998; Boss 1998). Observations suggest that many of the molecular cloud cores (with mass of order a few $M_\\odot$ and size $0.1$~pc) have elongated shapes (Myers et al.~1991) and slow rotation rates (with the ratio of rotational to gravitational energies of order $0.02$; Goodman et al.~1993), implying that rotation is probably not a crucial factor in driving fragmentation on scales greater than $200$~AU. Our result on the growth of bar-mode perturbations ($\\delta\\rho\\propto Y_{2m}$) indicates that, even without net angular momentum, the collapsing cloud tends to deform into an ellipsoidal shape (oblate disk or prolate bar, depending on which $m$-mode perturbation is dominant initially). Fragmentation is more likely to occur for such deformed configurations (e.g., Bonnel 1999; Matsumoto \\& Hanawa 1999).  In the context of core-collapse supernovae, our result shows that the homologous inner core and the supersonic outer core are globally stable against nonradial perturbations prior to core bounce at nuclear density and the formation of the proto-neutron star. However, during the subsequent accretion of the outer core (involving $15\\%$ of the core mass) and envelope onto the proto-neutron star, nonspherical perturbations can grow according to $\\delta\\rho/\\rho\\propto r^{-1/2}$ or even $\\delta\\rho/\\rho \\propto r^{-1}$ (Lai \\& Goldreich 2000). The asymmetric density perturbations seeded in the presupernova star, especially those in the outer region of the iron core, are therefore amplified during collapse. The enhanced asymmetric density perturbation may lead to asymmetric shock propagation and breakout, which then give rise to asymmetry in the explosion and a kick velocity to the neutron star (Goldreich et al.~1996; Burrows \\& Hayes 1996).  Our stability analysis (\\S 4) shows that Shu's expansion-wave solution is globally unstable to perturbations of all $l$'s, although the growth rates are unknown at present. The implication of this result is not entirely clear. It is well-known that a static singular isothermal sphere is highly unstable to radial perturbations (A truncated Bonner-Ebert isothermal sphere is unstable when the range of density from the center to the surface is greater than $14.04$; see Bonner 1956, Hunter 1977). Earlier one-dimensional numerical simulations have already shown that a collapsing isothermal cloud does not approach the expansion-wave solution (Hunter 1977; Foster \\& Chevalier 1993). Our stability analysis corroborates this result, and indicates that the expansion-wave solution cannot be realised in a pure hydrodynamical situation.  Magnetic fields play an important role in the current paradigm for forming low-mass stars (e.g., Shu, Adams \\& Lizano 1987; Shu et al.~1999; Mouschovias \\& Ciolek 1999). Ambipolar diffusion of magnetic fields drives the quasi-static contraction of the molecular cloud core with growing central concentration such that the core asymptotically approaches the state of a singular isothermal sphere. When the flux-to-mass ratio drops below certain critical value, a runaway ``inside-out'' collapse ensues, and it is thought that this collapse is well described by the expansion-wave solution (Shu et al.~1999). In reality, there is probably no sharp distinction between the quasi-static contraction and dynamical collapse (e.g., Safier, McKee \\& Stahler 1997; Li 1998), and a real singular isothermal sphere can never be reached. Our global stability analysis of the expansion-wave solution (\\S 4) does not depend on the mathematical singularity of the solution at $r=0$, but depends on the existence of a well-defined rarefaction front and a static isothermal density profile outside the front in the solution. It in not clear whether our idealized hydrodynamical stability analysis can be applied to more realistic situations with (even sub-dominant) magnetic fields (see Galli \\& Shu 1993a,b and Li \\& Shu 1997 for the effects of magnetic field on self-similar ``inside-out'' collapse)."
    },
    "0004/astro-ph0004193_arXiv.txt": {
        "abstract": "The boundaries of the Uranian $\\epsilon$, $\\alpha$, and $\\beta$ rings can be fitted by Keplerian ellipses. The pair of ellipses that outline a given ring share a common line of apsides. Apse alignment is surprising because the quadrupole moment of Uranus induces differential precession. We propose that rigid precession is maintained by a balance of forces due to ring self-gravity, planetary oblateness, and interparticle collisions. Collisional impulses play an especially dramatic role near ring edges. Pressure-induced accelerations are maximal near edges because there (1) velocity dispersions are enhanced by resonant satellite perturbations, and (2) the surface density declines steeply. Remarkably, collisional forces felt by material in the last $\\sim$100 m of a $\\sim$10 km wide ring can increase equilibrium masses up to a factor of $\\sim$100. New ring surface densities are derived which accord with Voyager radio measurements. In contrast to previous models, collisionally modified self-gravity appears to allow for both negative and positive eccentricity gradients; why all narrow planetary rings exhibit positive eccentricity gradients remains an open question. ",
        "introduction": "Each narrow eccentric ring surrounding Uranus is composed of particles moving on nested elliptical orbits. The outer and inner edges of a given ring define ellipses having semi-major axes $a \\pm \\Delta a/2$ and eccentricities $e \\pm \\Delta e/2$, where $\\Delta a \\ll a$, $\\Delta e \\ll e$, and $e \\ll 1$. Observed values of $a$, $e$, $\\Delta a$, and $\\Delta e$ for the Uranian $\\epsilon$, $\\alpha$, and $\\beta$ rings are listed in Table \\ref{param}.  \\placetable{param} \\begin{deluxetable}{ccccc} \\tablewidth{0pc} \\tablecaption{Parameters of Eccentric Uranian Rings\\tablenotemark{a}\\label{param}} \\tablehead{ \\colhead{Ring}  &  \\colhead{$a$(km)}      & \\colhead{$\\Delta a$(km)} & \\colhead{$e\\, (\\times 10^3)$}   &  \\colhead{$\\Delta e\\, (\\times 10^3)$}    }  \\startdata $\\epsilon$ & 51149 & 58.1 & 7.936 & 0.711 \\nl $\\alpha$ & 44718 & 7.15 & 0.761 & 0.076 \\nl $\\beta$ & 45661 & 8.15 & 0.442 & 0.066 \\nl \\tablenotetext{a}{Values taken from Tables I and VII of \\markcite{fetal91}French et al. (1991).}  \\enddata \\end{deluxetable}  Remarkably, the set of ellipses describing an individual ring share a common line of apsides. Apse alignment is surprising because the oblateness of Uranus causes orbits of particles with different semi-major axes to precess differentially. Timescales for differential precession in the absence of other forces are extremely short; in the case of the Uranian $\\epsilon$ ring, the inner edge would precess a full revolution relative to the outer edge in 175 years. Rigid precession of an eccentric planetary ring has remained a problem in ring dynamics for over 20 years.  \\markcite{gt79}Goldreich \\& Tremaine (1979, hereafter GT) proposed that apse alignment is maintained by self-gravity. Their theory predicts that the eccentricity gradient across the ring,  \\begin{equation} q_e \\equiv a \\frac{\\partial e}{\\partial a} \\, , \\end{equation}  \\ni must be positive. A positive eccentricity gradient in an apse-aligned ring implies that the ring is narrowest at periapse and widest at apoapse. Gravitational forces between particles are therefore greatest near periapse. Material in the inner half of the ring pulls radially inward on the outer half at periapse, generating a differential precession which exactly cancels that due to planetary oblateness.  Though the prediction that $q_e > 0$ accords with observations of all known narrow eccentric rings, the standard self-gravity model (hereafter SSG) predicts Uranian ring masses that are too low compared to those inferred from Voyager radio occultations. Ring masses based on observations exceed predictions by factors of at least $\\sim$3 ($\\epsilon$ ring) to $\\sim$50 ($\\alpha$ and $\\beta$ rings) (\\markcite{tetal86}Tyler et al. 1986; \\markcite{g90}Gresh 1990; see also the reviews by \\markcite{espo91}Esposito et al. 1991 and \\markcite{fetal91}French et al. 1991). Low surface densities are particularly problematic for the $\\alpha$ and $\\beta$ rings. With SSG surface densities, torques exerted by inner shepherd satellites would be insufficiently strong to confine the $\\alpha$ and $\\beta$ rings against drag from the distended exosphere of Uranus (\\markcite{gp87}Goldreich \\& Porco 1987, hereafter GP). In addition, as discussed by \\markcite{getal95}Graps et al. (1995), shapes of the $\\epsilon$ ring surface density profiles as derived from occultation light curves do not accord with SSG predictions.  This paper points the way towards resolving these problems. In \\S \\ref{qual}, a theory of collisionally modified self-gravity (hereafter CMSG) is qualitatively described. A simple quantitative model is set forth in \\S \\ref{quant}, in which new surface density profiles are derived for the $\\epsilon$ and $\\alpha$ rings that are in better agreement with observations. In \\S \\ref{impl}, implications of our solutions for torque balance, the role of planetary oblateness, and the value of $q_e$ are discussed. Directions for future research are summarized in \\S \\ref{future}. ",
        "conclusions": "\\label{impl}  \\subsection{Surface Density Profiles and Torque Balance}  Simple CMSG models, while not fully realistic, demonstrate the existence of a new class of self-gravity solution, that obtained by accounting for the modification of ring boundary conditions by interparticle collisions. Remarkably, forces felt by material in the last $\\sim$100 m of a $\\sim$10 km wide ring can increase equilibrium masses by factors up to 100.  Large S-band opacities measured by Voyager, which are incompatible with SSG surface densities (see, e.g., the review by \\markcite{espo91}Esposito et al. 1991), can be reconciled with average CMSG surface densities of $\\sim$75--100$\\gm \\,\\cm^{-2}$ for the $\\epsilon$, $\\alpha$, and $\\beta$ rings. Moreover, CMSG models predict that surface densities near ring edges are higher than those in the interior. This behavior is reminiscent of the ``double-dip'' structure seen in occultation light curves for the $\\epsilon$ and $\\alpha$ rings (see, e.g., the review by \\markcite{fetal91}French et al. 1991).  Greater ring masses as implied by CMSG resolve problems associated with exospheric drag that were pointed out by \\markcite{gp87}GP for rings $\\alpha$ and $\\beta$. For the remainder of this subsection, numerical estimates will be made for the $\\alpha$ ring; similar conclusions hold for the $\\beta$ and $\\epsilon$ rings. Surface densities are scaled to a typical CMSG value in the ring interior of $\\Sigma = 75 \\gm \\,\\cm^{-2}$. An inner shepherd satellite exerts a repulsive, non-linear torque at first-order Lindblad resonances of magnitude  \\begin{equation} T^L_{\\rm nl} \\approx \\frac{10 \\rho_s R_s^3 \\Sigma^2 n^2 a^7}{M_p^2 d} \\approx 6 \\times 10^{17} \\left( \\frac{\\Sigma}{75 \\gm \\,\\cm^{-2}} \\right)^2 \\left( \\frac{R_s}{10 \\km} \\right)^3 \\left( \\frac{\\rho_s}{1.5 \\gm \\,\\cm^{-3}} \\right) \\left( \\frac{500 \\km}{d} \\right) \\erg \\;, \\end{equation}  \\ni where the satellite radius, $R_s$, is scaled to the Voyager upper limit of 10 km (\\markcite{setal86}Smith et al. 1986). The shepherding torque exceeds the magnitude of the drag torque exerted by the Uranian exosphere,  \\begin{equation} \\label{atm} T_{d} \\approx -4\\pi m_H n_H v_T n a^3 \\Delta a \\approx -4 \\times 10^{16} \\left( \\frac{n_H}{10^3 \\,\\cm^{-3}} \\right) \\erg \\, . \\end{equation}  \\ni Here $n_H = 7 \\times 10^{-6} \\, e^{32.4 \\, R_p / a} \\,\\cm^{-3}$ is the number density of hydrogen atoms of mass $m_H$ in the exosphere, and $v_T \\approx 1 \\km \\,\\s^{-1}$ is their thermal speed normal to the orbital plane (\\markcite{betal86}Broadfoot et al. 1986). That $T^L_{\\rm{nl}} > |T_d|$ ensures that the inner shepherd prevents ring particles from spiraling in towards Uranus.  Estimates of viscous torques $T_v$ also require revision. For a ring undergoing Keplerian shear, with minimum kinematic viscosity $n (\\Sigma/\\rho)^2$, the viscous torque is given by  \\begin{equation} T_v \\approx \\frac{3\\pi n^2 \\Sigma^3 a^2}{\\rho^2} \\approx 2.5 \\times 10^{18} \\left(  \\frac{\\Sigma}{75 \\gm \\,\\cm^{-2}} \\right)^3 \\left( \\frac{1.5 \\gm \\,\\cm^{-3}}{\\rho} \\right)^2 \\erg \\end{equation}  \\ni (\\markcite{gp87}GP). That $T_v \\gg |T_d|$ ensures that ring particles on the outer edge press against the inner Lindblad resonance established by the outer shepherd.  Conclusions drawn from comparisons between $T^L_{\\rm{nl}}$ and $T_v$ are on less sure footing. For the choice of scaling parameters, the latter exceeds the former, contrary to the requirement of the standard theory of shepherding that the torques be equal. This might be construed as evidence that the angular momentum luminosity in the ring interior is reduced below $T_v$ by the non-Keplerian shear associated with a non-zero $q_e$ (\\markcite{bgt82}Borderies, Goldreich, \\& Tremaine 1982; \\markcite{gp87}GP). However, the numerical estimates for the two torques differ only by a factor of a few. The shepherding torque should be evaluated using surface densities near the edge, which CMSG predicts are higher than those in the interior; this would increase the estimate of $T^L_{\\rm{nl}}$. Uncertainties in the choice of parameters preclude drawing any conclusion other than that these torques are of the same order of magnitude.  \\subsection{Relative Importance of Planetary Oblateness}  \\subsubsection{$J_2 = 0$ vs. $J_2 \\neq 0$} What does CMSG predict if $J_2 = 0$? Figure 2a displays the answer for the $\\epsilon$ ring, for $c_b = 2$ and 3 $\\cm \\,\\s^{-1}$. In contrast to SSG, a non-vanishing equilibrium surface density does not require a finite planetary oblateness; self-gravity can be balanced entirely by collisional pressure gradients. For the $\\alpha$ and $\\beta$ rings, solutions with and without $J_2$ are practically indistinguishable for $c_b \\geq 0.5 \\,\\cm \\,\\s^{-1}$. The influence of $J_2$ on the equilibrium solution diminishes as $\\Delta a$ decreases or as $c_b$ increases.  \\placefigure{j2} \\begin{figure} \\plotone{f2.eps} \\caption{(a) CMSG $\\epsilon$ ring models for which $J_2$ is reduced from its nominal value of $3.35 \\times 10^{-3}$ {\\protect \\markcite{en84}}(Elliot \\& Nicholson 1984) (solid line) to 0 (dotted line). As $c_b$ is increased from $2 \\, \\cm \\,\\s^{-1}$ (lower two curves) to $3 \\, \\cm \\,\\s^{-1}$ (upper two curves), the influence of $J_2$ diminishes. For the $\\alpha$ and $\\beta$ rings, CMSG models with and without $J_2$ are practically indistinguishable for $c_b \\geq 0.5 \\,\\cm \\,\\s^{-1}$ (data not shown). (b) CMSG $\\epsilon$ ring models for which $q_e = \\pm \\, 0.626$. Contrary to SSG models, a positive $q_e$ is not required to obtain an equilibrium solution.} \\label{j2} \\end{figure}  \\subsubsection{Empirical Scaling Relations for $J_2 = 0$} \\label{scale}  For $J_2 = 0$ and fixed ring geometry, the surface density at quadrature near a given edge scales as  \\begin{equation} \\label{scalesurfb} \\Sigma_b \\,(0 \\leq |x| \\lesssim \\lambda) = \\frac{c_b^2}{G\\lambda} \\;f(|x|/\\lambda) \\, , \\end{equation}  \\ni where $|x|$ measures distance from the edge, $c_b$ and $\\lambda$ are the same free parameters as in Equation (\\ref{prescrip}), and $f$ is a dimensionless function of the similarity variable $|x| / \\lambda$. Well away from ring edges, the surface density at quadrature scales as  \\begin{equation} \\label{scalesurf} \\Sigma_i \\,(|x| \\gg \\lambda) = \\frac{c_b^2}{G\\sqrt{\\lambda \\Delta a}} \\; g(|x|/\\Delta a) \\, , \\end{equation}  \\ni where $g$ is another dimensionless function. The total ring mass scales as  \\begin{equation} \\label{scalemass} M \\sim \\frac{c_b^2 a}{G} \\sqrt{\\Delta a / \\lambda} \\, . \\end{equation}  \\subsection{Value of $q_e$}  \\subsubsection{Sign of $q_e$} Figure 2b displays a CMSG model for the $\\epsilon$ ring obtained by reversing the sign of $q_e$. In contrast to SSG, a positive eccentricity gradient is not necessary in CMSG to obtain an equilibrium solution. This resurrects the problem of why all known eccentric planetary rings, including the Titan and Huygens ringlets around Saturn, are narrowest at periapse and widest at apoapse.  It is possible that equilibria obtained using $q_e < 0$ are unstable. To address this issue, a preliminary investigation of ring stability for an $N=4$ ringlet model has been undertaken. Forces due to pressure gradients are included only for the first and fourth ringlets. Collisional accelerations are treated as if they arise from anti-self-gravity forces (self-gravity with the sign of the acceleration reversed); i.e., collisional shear stresses are ignored. In this crude approximation, equilibria are found to be stable regardless of the sign of $q_e$; small deviations from equilibrium masses result in apsidal librations (\\markcite{bgt83}Borderies, Goldreich \\& Tremaine 1983). It remains to be seen whether collisional shear stresses alter stability properties.  Another possibility is that initial conditions set the sign of $q_e$. If the ring were initially uniform in width as a function of azimuth, then planetary oblateness would determine the initial sense of differential precession within the ring. The resultant narrowing of the ring width near a true anomaly of $f = -\\pi/2$ would cause a positive eccentricity gradient to grow by self-gravity. Under this hypothesis, an $N=2$ ringlet model incorporating forces from self-gravity and planetary oblateness yields the following time evolution for the apse and eccentricity differences between outer and inner ringlets:  \\begin{eqnarray} \\delta \\pomega = - A \\sin \\Omega_{\\rm{lib}} t \\\\ \\delta e = A e ( 1 - \\cos \\Omega_{\\rm{lib}} t ) \\end{eqnarray}  \\ni where $A > 0$ and $\\Omega_{\\rm{lib}}$ are the amplitude and frequency, respectively, of libration (cf. \\markcite{bgt83}Borderies, Goldreich, \\& Tremaine 1983). Note that the time-average of $\\delta e$ is positive. Inelastic collisions would damp librations and the ring would eventually settle into an equilibrium for which $q_e > 0$.  \\subsubsection{Magnitude of $q_e$ Near Ring Boundaries} \\label{magqe}  It has been assumed that the eccentricity gradient, $q_e$, is finite out to the last $\\lambda = c_b/n \\sim 50$ meters of ring material. A finite $q_e$ is necessary to generate a non-zero azimuthal average of the collisional acceleration [see equation (\\ref{azaverage})]. The simple quantitative model of \\S \\ref{quant} employed the observed value of $q_e$ averaged over the entire ring width. The true value over the last few hundred meters of ring material is unknown.  In the case of the best-studied $\\epsilon$ ring, \\markcite{getal95}Graps et al. (1995) combined Voyager photopolarimeter and radio occultation measurements to infer the eccentricity gradient as a function of semi-major axis. They found that $q_e$ decreases over the last $\\sim$5 km from its nearly constant value of $\\sim$0.65 in the interior to $\\sim$0.35 near the edge. The radial resolution of their study was between 1 and 2 km.  A decrease in $q_e$ towards ring boundaries is theoretically plausible. Distortions in a circular ring can be described by the change in separation, $\\delta r$, between neighboring streamlines of the form  \\begin{equation} \\delta r \\propto \\cos m(\\phi - \\Omega_{\\rm pat} t) \\, , \\end{equation}  \\ni where $\\Omega_{\\rm pat}$ is the pattern speed of the distortion and $m$ is an integer. A constant $q_e$ ring that precesses rigidly in the quadrupole field of the central planet is equivalent to a distorted circular ring for which $m = 1$ and $\\Omega_{\\rm pat} = \\langle d\\pomega/dt \\rangle_Q$. Resonant satellite perturbations, which enhance velocity dispersions within a distance $w_r \\sim 1 \\km$ of ring edges, are characterized by much higher values of $m = 2a/3d \\gg 1$ and $\\Omega_{\\rm pat} = \\Omega_s$. Satellite-induced disturbances might therefore reduce the local value of $q_e$. A decrease in $q_e$ over a distance $w_r$ near ring boundaries is roughly equivalent to setting $\\lambda = w_r$ in equation (\\ref{prescrip}). By the scaling relations (\\ref{scalesurf}) and (\\ref{scalemass}), this would reduce surface densities and total ring masses shown in Figure \\ref{surf} by a factor of $\\sqrt{w_rn/c_b} \\sim 4$."
    },
    "0004/astro-ph0004083_arXiv.txt": {
        "abstract": "We present a high-resolution 5-GHz radio continuum map of the starburst galaxy NGC 2146 made with MERLIN and the VLA (A-array), in a search of radio supernovae and supernova remnants expected to be already produced by the most massive stars in the starburst. At 5 GHz, about 20 point sources were detected earlier by Glendenning $\\&$ Kronberg (1986) in the central 800 pc of NGC 2146. Our observations with higher sensitivity and resolution made with MERLIN and the VLA confirms the detection of 18 sources, and resolves 7 of them. Additional 1.6-GHz MERLIN observations disclose 9 sources coincident in position with those detected at 5 GHz, which allows us to derive their spectral indices $\\mathrm{\\alpha_{1.6}^{5}}$, ($\\mathrm{S_{\\nu}\\sim\\nu^{\\alpha}}$). Only 3 sources have indices ($\\alpha$ $<$ 0) consistent with synchrotron emission from supernova remnants or radio supernovae, while the others have very steep inverted spectra ($\\alpha$ $>$ 0). \\\\ We suggest that the sources with positive spectral index are optically thick ultra--compact and/or ultra--dense {H\\,{\\sc ii}} regions with high electron densities and high emission measures (EM $>$ $\\mathrm{10^{7}}$ $\\mathrm{cm^{-6}}$ pc). Minimum energy requirements indicate that these regions may contain up to 1000 equivalent stars of type O6, comparable to the number of stars found in super starclusters. When compared with M\\,82, the galaxy NGC 2146 lacks however a large number of supernova remnants. We suggest that NGC 2146 is experiencing a burst of star formation stronger than that in M\\,82, but being in a younger phase. In this phase, only few stars have already exploded, whereas the others cause strong thermal emission from compact, optically thick ionized gas regions, around the young super starclusters.\\\\ We may, however, not exclude an alternative scenario in which strong free-free absorption at 1.6 GHz in foreground ionized gas with very high emission measures (EM $>$ $\\mathrm{10^{8}}$ $\\mathrm{cm^{-6}}$ pc) hides a certain number of supernova remnants, thus rendering for some sources the observed inverted spectra. ",
        "introduction": "NGC 2146 is a starburst galaxy located at a distance of 14.5 Mpc (1\\arcsec $\\approx$ 70 pc). Its infrared luminosity measured with IRAS at 60 $\\mathrm{\\mu m}$ and 100 $\\mathrm{\\mu m}$ is $6.6\\times10^{10}$ $\\mathrm{L_\\odot}$. This value and the large 25 $\\mathrm{\\mu m}$ to 60 $\\mathrm{\\mu m}$ flux ratio place the object in the lower part of the IR luminosity/spectral index plane populated by superluminous galaxies without active nuclei (Hutchings et al. \\cite{hutch90}). NGC 2146 has a central molecular ring (Jackson \\& Ho \\cite{jack88}, Young et al. \\cite{young88}), and an outflow of hot gas along the minor axis driven by supernova explosions and stellar winds in the starburst region (Armus et al. \\cite{armus95}, Della Ceca et al. \\cite{della99}). These characteristics reveal a strong similarity to the prototype starburst galaxy M\\,82 (at a distance of 3.2 Mpc and with an infrared luminosity of $2.4\\times10^{10}$ $\\mathrm{L_\\odot}$), although NGC 2146 does not have a companion that may have triggered the starburst (Fisher $\\&$ Tully \\cite{fisher76}). It has been suggested that the starburst in NGC 2146 is the result of a far evolved merger (Condon et al. \\cite{condon82}; Young et al. \\cite{young88}, Hutchings et al. \\cite{hutch90}), but a fully convincing kinematic and material trace of the merger has not yet been found.\\\\ The nuclear region is partly obscured by a strongly absorbing dust lane (Benvenuti et al. \\cite{benve75}). The high extinction (A$\\mathrm{_{v}}$ $>$ 5 mag; Young et al. \\cite{young88bis}; Hutchings et al. \\cite{hutch90}; Smith et al. \\cite{smith95}) makes the optical observation of the inner star-forming regions impossible.\\\\ Strong non--thermal radio emission from the centre of NGC 2146 has been detected by Kronberg \\& Biermann (\\cite{kron81}; hereafter KB), by Condon et al. (\\cite{condon82}), and by Lisenfeld et al. (\\cite{lisenfeld96}, \\cite{lisenfeld97}). About 20 point sources were detected at 5 GHz in the central 800 pc of NGC 2146 by Glendenning \\& Kronberg (\\cite{glen86}) using the VLA, and these were interpreted as radio supernovae (RSN) or supernova remnants (SNRs). In this paper we report the detailed spatial distribution of these point sources, together with their number and fluxes derived from radio continuum observations at 5 GHz obtained with MERLIN and the VLA. From additional 1.6 GHz MERLIN observations we obtained the spectral index of 9 of these sources, which allow to identify their nature. Our results open the possibility of a direct comparison with similar observations of M\\,82 (Kronberg et al. \\cite{kron85}; Muxlow et al. \\cite{mux94}) and other strong and nearby starburst galaxies, such as NGC\\,1808, NGC\\,4736 and NGC\\,5253 (Saikia et al. \\cite{saikia90}; Duric \\& Dittmar \\cite{duric88}; Turner et al. \\cite{turner98}). ",
        "conclusions": "The combination of VLA and MERLIN 5 GHz observations of the starburst galaxy NGC 2146 allowed us to resolve seven of the eighteen hitherto unresolved sources detected with the VLA. We have measured the spectral indices between 1.6 and 5 GHz for nine sources. Three sources (37.6+24.2, 38.9+22.5, 41.4+15.0) seem to be good candidates for SNRs or RSN, with total emitted powers at 5 GHz comparable to the strongest SNRs found in M\\,82. Four other sources (36.6+27.5, 39.0+17.9, 40.2+18.1 and 37.7+23.7) have instead been identified with compact or ultra-dense \\ion{H} {ii} regions that have large EMs and a high-frequency turnover. The emitted thermal emission  indicates a large equivalent number of O6 stars, exceeding the number of those found in similar objects in other galaxies. The large number of O6 stars may, however, be consistent with stars in super starclusters. The nature of the last two sources (37.9+23.7 and 39.5+17.7) for which we have spectral index information are still a matter of discussion. It seems realistic that they emit a mixture of thermal and non-thermal radiation, and may manifest the most powerful \\ion{H} {ii} regions ever observed. However, the burst of star formation in NGC 2146 has an extraordinary power, stronger than that in M\\,82. Despite this and the fact that the overall radio spectral index of the galaxy reveals a non-thermal origin, the number of discovered SNRs is much smaller than that in M\\,82. For the explanation of this fact we propose two scenarios: \\begin{itemize} \\item{NGC 2146 is experiencing a burst of star formation stronger than that in M\\,82, but which we are observing in a younger phase. This phase still lacks a large number of SNRs, but instead gives rise to strong thermal emission from optically thick dense ionized gas regions, associated with massive or super starclusters. The origin of the non-thermal extended component must be a trace of past SN explosions, individually too old and faint to be seen, and belonging to a previous burst in the galaxy.} \\item{Strong free-free absorption by very dense \\ion{H} {ii} regions and, for the most compact sources, possibly also\\\\ synchrotron-self absorption, may conceal a certain number of SNRs or RSN, rendering the identification of their nature difficult if only based on measurements of the spectral indices around the turnover frequencies. In this case the global number of SNRs or RSN can be higher, and this would corroborate the similarity between NGC 2146 and M\\,82. This interpretation raises, however, the question why we do not observe similar compact or ultra-dense \\ion{H} {ii} regions in M\\,82 as well.} \\end{itemize} Observations at 15 GHz will be able to better clarify the nature of the sources, and shed light on the unusual phase of the star burst we are observing in NGC 2146."
    },
    "0004/hep-ph0004169_arXiv.txt": {
        "abstract": " ",
        "introduction": "The search for experimental evidence for supersymmetry is currently approaching a transition. For several years now, many of the most incisive experimental searches have been those at LEP~\\cite{lepsusy}, whose constraints on the parameter space of the minimal supersymmetric extension of the Standard Model (MSSM) have grown ever more restrictive, as the centre-of-mass energy of LEP~II has been increased in successive steps. In parallel, improved analyses of data from Run~I of the Tevatron Collider have been providing important complementary constraints~\\cite{sugrarept}. The transition is marked by the termination of the LEP~II experimental programme in late 2000 and the anticipated start of Run~II of the Tevatron Collider in 2001.  The results of experimental searches for different MSSM particles can usefully be compared and combined using the conventional parameterization of the model in terms of supersymmetry-breaking scalar and gaugino masses $m_0, m_{1/2}$, the higgsino mixing parameter $\\mu$, the ratio of Higgs vacuum expectation values (vev's) $\\tan \\beta$ and a universal trilinear supersymmetry-breaking parameter $A$. We work in the framework of gravity-mediated models of supersymmetry breaking, in which it is commonly assumed that the scalar masses $m_0$ and the gaugino masses $m_{1/2}$ are universal at some supersymmetric GUT scale. The assumptions that these supersymmetry-breaking parameters are universal should be questioned, particularly for scalar masses and especially  those of the Higgs supermultiplets, but provide a convenient way of benchmarking comparisons and combinations of different experimental searches. In this paper, we make such comparisons and combinations in variants of the MSSM in which the scalar-mass universality assumption is extended to Higgs fields (UHM, also commonly referred to as mSUGRA or the constrained MSSM or CMSSM), and also without this supplementary assumption (nUHM).  In making such comparisons, we emphasize the importance of including radiative corrections to the relations between these MSSM model parameters $(m_0, m_{1/2}, \\mu, \\tan \\beta, A)$ and the physical masses of MSSM particles. Radiative corrections are well-known to be crucial in the MSSM Higgs sector, but also should not be neglected in the chargino, neutralino, gluino and squark sectors. As we have emphasized previously~\\cite{efgos}, the differences between the domains of MSSM parameter space apparently explored at the tree and one-loop levels are comparable to the differences between the domains explored in successive years of LEP running at higher centre-of-mass energies. In view of the intense experimental effort put into sparticle searches at LEP~II, it is important that the final results of these efforts be treated with the theoretical care they deserve. This issue is also relevant if one wishes to compare the physics reaches for electroweakly-interacting sparticles at LEP and for strongly-interacting sparticles at the Tevatron Collider, in which case one should take into account the important radiative corrections to squark and gluino masses~\\cite{Pierce:1997zz}, as well as to their production cross sections.  In addition to direct searches for the production of MSSM particles, important indirect constraints must also be taken into account. These include other accelerator constraints, such as the measured value of the $b \\rightarrow s \\gamma$ decay rate~\\cite{CLEObsg,ALEPHbsg}, and non-accelerator constraints related to the possible role of the lightest supersymmetric particle (LSP) as cold dark matter (CDM). The lightest supersymmetric particle (LSP) would be stable in any variant of the MSSM which conserves $R$ parity, as we assume here. In gravity-mediated models of supersymmetry breaking, the framework adopted here, the LSP is commonly thought to be the lightest neutralino $\\chi$, and calculations of the cosmological relic density of LSPs, $\\Omega_\\chi$, yield values in the range preferred by cosmology in generic domains of MSSM parameter space~\\cite{EHNOS}. The possibility of supersymmetric CDM provides one of our principal motivations for seeking a deeper understanding of the allowed MSSM parameter space, but is not our only focus in this paper.  The most essential dark-matter constraint is that the relic LSP density not overclose the Universe. The conditions that the universe has an age in excess of 12 billion years and that $\\Omega_{\\rm total}\\le 1$ imply an upper bound on $\\ohsq$ of 0.3.  Further, the convergent indications from astrophysical structure-formation arguments and observations of high-redshift supernovae are that $\\Omega_{CDM} < 0.5$\\cite{Bahcall:1999xn}, whereas the Hubble expansion rate $H_0 = 100 h$~km/s/Mpc: $h = 0.7$ with an error of about 10\\%\\cite{Freedman:1999wh}, so we require $\\Omega_{LSP} h^2 \\le \\Omega_{CDM} h^2 \\le 0.3$. On the other hand, astrophysical structure formation seems to require $\\Omega_{CDM} > 0.2$, so we also require $\\Omega_{LSP} h^2 \\ge 0.1$, while acknowledging that a lower value of $\\Omega_{LSP}$ could be permitted if other CDM particles such as axions and/or superheavy relics are present.  There have recently been some significant developments in the analysis of supersymmetric CDM. One is that the importance of co-annihilation effects involving next-to-lightest supersymmetric particles (NLSPs) such as the $\\tilde \\tau, \\tilde \\mu$ and $\\tilde e$ for calculations of the relic density of a gaugino-like LSP has recently been recognized \\cite{EFOSi,glp}. Another phenomenon whose importance in the CMSSM has recently been underlined is the possible transition of the electroweak vacuum into a charge- and colour-breaking (CCB) minimum \\cite{bbc}. The absence of such an instability is not absolutely necessary, since a transition in the future cannot be excluded. Therefore, we comment on the regions of MSSM parameter space in which the CCB instability is absent, but do not focus exclusively on these regions.  The main purpose of this paper is to prepare for the compilation, comparison and combination of the definitive results from LEP~II and the Tevatron. We illustrate our analysis with the latest available limits from these two experimental programmes~\\cite{lepsusy, sugrarept}, supplemented by educated guesses at their final sensitivities. As we have explained previously, and discuss in more detail below, a key role in constraining the MSSM parameter space is provided by the LEP Higgs search. We express our results as a function of the present LEP lower limit on $m_H$, currently 107.9~GeV~\\cite{Higgs2000}, and the prospective future sensitivity, which may approach 112~GeV. We use our analysis to present lower limits on the LSP mass and on tan$\\beta$. We include a discussion of the implications of relaxing the UHM assumption that the soft supersymmetry-breaking contributions to Higgs masses are also universal. In particular, we investigate whether a light Higgsino LSP is still a viable dark matter candidate, and find that the latest LEP~II data now exclude this possibility. Finally, we discuss the likely future developments in the exploration of the MSSM parameter space in the period before the start-up of the LHC, during which the central role is likely to be played by Run~II of the Tevatron Collider.  The layout of this paper is as follows. In Section 2 we review in more detail the theoretical framework we adopt, discussing the issues of universality and relic coannihilations, and stressing the importance of Higgs mass constraints. In Section 3 we review our implementation of the $b \\rightarrow s \\gamma$ constraint, including, where applicable, the next-to-leading-order (NLO) QCD corrections. We discuss the implications of the latest available constraints from LEP~II in Section 4, combining them in Section 5 with the cosmological and astrophysical constraints $0.1 \\le \\Omega_{CDM} h^2 \\le 0.3$ as well as the $b \\rightarrow s \\gamma$ constraint, and making the UHM assumption. We find \\beq m_\\chi \\ge 51~{\\rm GeV}, \\; \\; \\tan \\beta \\ge 2.2 \\label{summarize} \\eeq and discuss the expanded ranges of $m_\\chi$ and $\\tan \\beta$ that may be explored by the improved Higgs-mass limits that might be obtained from the run of LEP~II in the year 2000. The limits in (\\ref{summarize}) are strengthened when we restrict values of $A_0$ to minimize the parameter space with CCB minima, in which case we find \\beq m_\\chi \\ge 54~{\\rm GeV}, \\; \\; \\tan \\beta \\ge 2.8 \\label{summarize2} \\eeq  \\noindent We further generalize the discussion to non-universal Higgs masses (nUHM) in Section 6, finding that the limits on $m_\\chi$ and $\\tan \\beta$ are relaxed to  \\beq m_\\chi \\ge 46~{\\rm GeV}, \\; \\; \\tan \\beta \\ge 1.9. \\label{summarize3} \\eeq  \\noindent Section 7 is devoted to a discussion of the possibility of Higgsino dark matter in such a nUHM scenario. We find that the LEP~II searches for charginos, neutralinos and Higgs bosons together now  exclude as dark matter an LSP that is more than about 70\\% Higgsino. We turn our attention to the Tevatron Collider in Section 8. We compare the LEP~II and Run~I sensitivities to the MSSM parameters, and discuss and compare the regions of MSSM parameter space to which Tevatron Run~II data should be sensitive~\\cite{sugrarept}. Finally, Section 9 summarizes our conclusions and the prospects for future improvements and extensions of the analysis reported here. ",
        "conclusions": "One of our principal goals in this paper has been to obtain strengthened lower limits on the neutralino mass, combining the latest LEP data with the cosmological dark matter requirement $0.1 < \\Omega_\\chi h^2 < 0.3$. We summarize our limits in Figs.~\\ref{fig:mvtb}, under various different assumptions: universal (UHM) or non-universal (nUHM) soft supersymmetry-breaking scalar masses for Higgs bosons and (in the former case) whether one requires the present vacuum to be stable against transition to a charge- and colour-breaking (CCB) vacuum or not (UHM$_{\\rm min}$). Also, we give limits both for the available 1999 LEP data and with a `realistic' assessment of the likely sensitivity of data to be taken in 2K.  In all cases, for both positive and negative $\\mu$, the lower limits on $m_\\chi$ are relatively insensitive to $\\tan \\beta$ at large $\\tan \\beta$. Here they are determined by the LEP chargino bound, as the LEP Higgs mass bound is weaker than the chargino bound at large $\\tan \\beta$. In fact, in the two UHM cases shown, the points at which the limiting curves bend upward, as one decreases $\\tan \\beta$, are precisely the points at which the Higgs mass bound becomes more stringent than the chargino bound. In the UHM cases, the neutralino mass limits are strong at intermediate values of $\\tan \\beta \\simeq$ 4--7 because, as discussed earlier, the cosmological bound on the relic density prohibits going to large values of $m_0$, and ensures that the Higgs bound places a strong constraint. Below this break point, the lower limit on $m_\\chi$ increases rapidly with decreasing $\\tan \\beta$. Above this break point, the limit on $m_\\chi$ is relatively insensitive to the additional theoretical assumptions made, such as UHM vs. UHM$_{\\rm min}$ or nUHM. However, in the nUHM cases, because one can increase $m_0$ sufficiently to weaken the Higgs mass bound, the break point occurs at a lower value of $\\tan \\beta$. To go to lower values of $\\tan \\beta$ then requires a substantial increase in $m_\\chi$.  \\begin{figure}[htbp] \\begin{center} \\mbox{\\epsfig{file=mvtbn.eps,height=8.0cm}} \\mbox{\\epsfig{file=mvtbp.eps,height=8.0cm}} \\end{center} \\caption[.]{\\label{fig:mvtb}\\it Lower limits on the neutralino mass $m_\\chi$ as functions of $\\tan \\beta$ for (a) $\\mu< 0$ and (b) $\\mu>0$. The curves correspond to the final 1999 LEP results (thin lines) and our `realistic' expectations for the 2K LEP run (thick lines).  We show the UHM case with $A_0 = -m_{1/2}$ to avoid CCB minima (dashed curves): these are the strongest constraints. We also show (dotted lines) the more general UHM$_{\\rm min}$ case where $A_0$ is left free, and we do not require the absence of CCB vacua. We also display additionally the nUHM case, which is the most conservative and allows both $\\mu$ and $m_A$ to be free in addition to $A_0$. } \\end{figure}   In the UHM cases with and without the restriction forbidding CCB vacua, the lower limit on the lightest MSSM Higgs mass, in particular, implies lower limits on $\\tan \\beta$ which are plotted in Fig.~\\ref{fig:tbvmh}. The limits are somewhat stricter for $\\mu < 0$ than for $\\mu > 0$, whether (UHM) or not (UHM$_{\\rm min}$) one requires the absence of CCB vacua. Indeed, Fig.~\\ref{fig:tbvmh} will enable the appropriate conclusion to be drawn from whatever lower limit on the Higgs mass LEP eventually provides. We recall that the existing Higgs mass calculations in the MSSM are believed to be accurate to about 3~GeV.  Therefore in computing the bounds on $\\tb$ for Tables \\ref{tab:tbvmh} and \\ref{tab:tbvmhc}, for example, we have conservatively shifted the exclusion curves of Fig.~\\ref{fig:Higgs} by 3 GeV to the left before reading the values of $\\tb$ off of Fig.~\\ref{fig:tbvmh}. We also show in Fig.~\\ref{fig:tbvmh} the lower bound on $\\tan \\beta$ obtained in the nUHM, which is significantly weaker than in the UHM cases, and essentially independent of the sign of $\\mu$.  \\begin{figure}[htbp] \\begin{center} \\mbox{\\epsfig{file=tbvmh.eps,height=8cm}} \\end{center} \\caption[.]{\\label{fig:tbvmh}\\it Lower limit on $\\tb$ imposed by the experimental and cosmological constraints, as a function of the experimental Higgs mass limit.  The UHM, UHM$_{min}$ and nUHM labels are as in \\protect{Fig.~\\ref{fig:mvtb}}.  The $\\mu>0$ curve in the nUHM case is very similar to the $\\mu<0$ curve.} \\end{figure}  If LEP does achieve the `optimistic' 2K energies and luminosities (\\ref{optimistic}), the above constraints will be somewhat tighter.  The horizontal segments of Fig.~\\ref{fig:mvtb}, corresponding to the chargino limits, increase by a fraction of a GeV;  the vertical branches move to the right, intersecting the horizontal segments at $\\tb=8\\, (7.5)$ for $\\mu<0\\, (\\mu>0)$. And lastly, the lower limits on $\\tb$ improve to   \\begin{table}[htbp] \\begin{center} \\begin{tabular}{|l|c|c|c|} \\hline &UHM&UHM$_{\\rm min}$&nUHM\\\\ \\hline $\\mu<0\\;\\;\\;\\;\\;$&4.7&3.4&2.3\\\\[.2em] $\\mu>0$&4.2&3.0&2.3\\\\ \\hline \\end{tabular} \\caption{\\label{tab:tbvmho} {\\it Limits on $\\tb$, assuming the 'optimistic' 2K energies and luminosities  (\\protect{\\ref{optimistic}})}.} \\end{center} \\end{table}     In many respects, LEP has provided the most stringent constraints on the parameters of the MSSM. This is true, in particular, for its lower limits on the Higgs mass, and the chargino, neutralino and slepton constraints from LEP compare favourably with the Tevatron bounds on squark and gluino masses, once the different mass renormalizations of electroweakly- and strongly-interacting sparticles are taken into account. As we have shown, the present LEP data may be combined to set interesting lower bounds on the lightest neutralino mass, in particular if it is assumed to constitute the dark matter favoured by astrophysics and cosmology, and on $\\tan \\beta$. It may well be that these lower limits will be further strengthened by the LEP run during 2000, as we have discussed in this paper. However, this would be the pessimistic scenario. There is still a chance that sparticles or the Higgs boson may turn up this year, in which case we would be delighted to see our bounds superseded. LEP may not yet have discovered supersymmetry, but it certainly deserves to!  \\vskip 0.5in \\vbox{ \\noindent{ {\\bf Acknowledgments} } \\\\ \\noindent We would like to thank P. Gambino and C. Kao and especially M. Schmitt and M. Spiropulu for useful discussions. This work was supported in part by DOE grant DE--FG02--94ER--40823.  The work of T.F. was supported in part by DOE grant DE--FG02--95ER--40896 and in part by the University of Wisconsin Research Committee with funds granted by the Wisconsin Alumni Research Foundation.}  \\vskip 0.5in"
    },
    "0004/astro-ph0004340_arXiv.txt": {
        "abstract": "BVRI and \\Ha ~imaging and long--slit optical spectroscopic data are presented for four morphologically normal and relatively isolated Sa galaxies, NGC~3626, NGC~3900, NGC~4772 and NGC~5854. VLA \\HI ~synthesis imaging is presented for the first three objects. In all four galaxies, evidence of kinematic decoupling of ionized gas components is found in the long--slit spectroscopic data; the degree and circumstances of the distinct kinematics vary from complete counterrotation of all of the gas from all of the stars (NGC~3626) to nuclear gas disks decoupled from the stars (NGC~5854) to anomalous velocity central gas components (NGC~3900 and NGC~4772). In the three objects mapped in \\HI, the neutral gas extends far beyond the optical radius, $R_{HI}/R_{25}$ $\\ge$ 2. In general, the \\HI ~surface density is very low and the outer \\HI ~is patchy and asymmetric (NGC~3900) or found in a distinct ring, exterior to the optical edge (NGC~3626 and NGC~4772). While the overall \\HI ~velocity fields are dominated by circular motions, strong warps are suggested in the outer regions by bending of the minor axis isovelocity contours (NGC~3900) and/or systematic shifts in position angle between inner and outer rings (NGC~3626 and NGC~4772). In the interior, coincidence is found between the \\Ha ~and \\HI ~in rings, sometimes partial and crisscrossed by dustlanes. Optical imaging is also presented for NGC~4138 previously reported by Jore \\etal ~(1996) to show counterrotating stellar components. The multiwavelength evidence is interpreted in terms of the kinematic ``memory'' of past minor mergers in objects that otherwise exhibit no morphological signs of interaction.  \\vskip 20pt Key words: -- galaxies: evolution -- galaxies: individual (NGC~3626, NGC~3900, NGC~4138, NGC~4772, NGC~5854) -- galaxies: interactions -- galaxies: kinematics and dynamics -- galaxies: spiral ",
        "introduction": "While mergers between comparable mass galaxies may be responsible for some of the most dramatic extragalactic events, minor mergers with small mass satellites may play a subtle but nonetheless critical role in the evolution of garden--variety spirals, including perhaps the Milky Way. It is easy to recognize the fireworks of the former by the tides, starbursts and nuclear activity they trigger. The latter however may lurk hidden to morphological inspection. Minor mergers are suggested to produce the X--shaped structures seen in some peculiar S0's (Mihos \\etal ~1995), the departures from axisymmetry seen in a significant fraction of spiral disks (Zaritsky \\& Rix 1997), the driving mechanism for strong stellar bars (Laine \\& Heller 1999) and the peculiar extended counter--rotating disks seen in some early--type spirals, e.g. NGC~7217 (Merrifield \\& Kuijken 1994; Buta \\etal ~1995), NGC~3593 (Bertola \\etal ~1996; Corsini \\etal ~1998), NGC~3626 (Ciri \\etal ~1995) and NGC~4138 (Jore \\etal ~1996).  Satellites are commonly found in the vicinity of normal galaxies (Zaritsky \\etal ~1997). If the Local Group is at all typical of spiral extragalactic environments, then the characteristics of such satellites are familiar. Interaction of the Milky Way with its low mass satellites is hypothesized to explain a variety of local kinematic peculiarities: the thick disk, the outer flare and warp, the Magellanic Stream and the high velocity \\HI ~clouds.  Several authors have begun to address numerically the complex phenomenon of minor mergers of gas--rich satellites with more massive spiral companions. An immediate worry is how to avoid over--heating the inner disk of the primary or disrupting it entirely. Minor mergers of low--density, small mass, gas--rich companions may avoid vertical heating, thereby preserving the victim disk. Predictions of the observable consequences of the accretion of a small mass companion by a spiral include the warping or actual spreading out of the disk (Quinn \\etal ~1993).  As the satellite material sinks towards the center of the large galaxy, the gas, unlike the stars, may lose angular momentum because of dissipation (Hernquist \\& Mihos 1995). The resultant inflow of gas towards the center may trigger a nuclear starburst (Mihos \\& Hernquist 1994) or fuel nuclear activity (Taniguchi 1999). Because of the predicted heating, even a satellite with one tenth the mass will thicken the disk and possibly build--up the bulge, driving the morphology of the post merger primary towards an earlier spiral type. The discovery of kinematically decoupled disks in early spirals may thus prove the efficacy of minor merger models.  Intrinsic galaxy properties arise from two principal components: absolute ``scale'' (size, luminosity, mass) and ``form'' (morphology, bulge--to--disk ratio, color, gas content, star formation rate). It is well-known (e.g. Roberts \\& Haynes 1994) that the early spiral types show a much greater spread in their characteristic properties related to ``form'' than do their later--type counterparts. As a part of a study to investigate the heterogeneity and dark matter content of the Sa galaxy class, we have investigated the kinematics of gas and stars in a sample of morphologically normal, isolated galaxies classified as Sa in the {\\it Revised Shapley Ames Catalog} (Sandage \\& Tammann 1987; RSA). In the course of that study, we have obtained major and minor axis long slit spectra for both stars and gas for a sample of 20 objects; for nine of them, we have also obtained \\HI ~synthesis maps yielding also two--dimensional \\HI ~velocity fields. A summary of results focussing on the data presentation and mass modelling was presented in Jore (1997). The analysis of the mass modelling and stellar velocities will be presented elsewhere (Jore \\etal, in preparation).  During the course of that work, a number of unusual cases of kinematic decoupling were discovered. In Jore \\etal ~(1996), we discussed the distinctive case of NGC~4138, an isolated Sa with two extended coplanar counter--rotating stellar disks embedded in a huge HI disk that co--rotates with the secondary stellar component. In a separate paper (Jore \\etal ~2000), we discuss five Sa galaxies, NGC~1169, 3623, 4866, 5377 and 5448, which represent a range of environments and morphologies within the Sa class. Of those objects, NGC~1169 and NGC~5448, exhibit no peculiar optical kinematics but their outer \\HI ~disks are strongly warped and kinematically skewed. NGC~3623, whose companions in the Leo Triplet are known to be interacting, exhibits a truncated \\HI ~disk and central gas decoupling. NGC~4866, viewed nearly edge--on, and NGC~5377 show doubling of their central gas components, as well as outer \\HI ~warps. Here, we present four additional examples: NGC 3626, NGC~3900, NGC~4772 and NGC~5854. Along with NGC~4138, each of these morphologically normal Sa galaxies exhibits peculiar kinematics that betray a disturbed past. We propose that each represents a different stage and circumstance of involvement in a relatively minor merger event with a gas--rich companion that is no longer recognizable.  This paper presents evidence of this past interaction through a combination of optical imaging and spectroscopy and \\HI ~synthesis mapping. In Section 2, we discuss the acquisition and reduction of the body of data: broad-- and narrow--band optical images, optical long--slit spectra, and \\HI ~synthesis maps. The kinematic evidence of past interaction is presented for each galaxy in Section 3. Section 4 gives a discussion of the observational results in the context of the minor merger scenario. ",
        "conclusions": "In this paper, we have presented a combined imaging and spectroscopic dataset for four Sa galaxies; images are also included for a fifth, NGC~4138, discussed previously by Jore \\etal ~(1996). All of the galaxies show evidence of departure from kinematic normalcy ranging from large--scale counter--rotation to decoupled central gas and/or stellar components. Despite the difference of detail, several unifying themes seem critical: (1) All of the objects are relatively isolated, morphologically normal, unbarred Sa galaxies. (2) Rings are prominent features both in the optical broadband light and in the distributions of ionized and neutral gas. The apparent rings may also result from tightly wound spiral arms viewed at moderate inclination. (3) The kinematically decoupled gas components appear to be associated with sites of current or at least recent star formation. In all cases where \\Ha ~emission is detectable in the region of interest, it coincides well with the location of kinematically decoupled components.  When there is no \\Ha ~emission in these regions, the decoupled components often show some form of a color change, often towards the red. (4) The gas--rich galaxies (all except NGC~5854) contain moderate \\HI ~masses, but, because the gas is spread over a large area R$_{HI}$/R$_{25} \\ge$ 2, the globally averaged \\HI ~surface densities $<\\sigma_{HI}> = M_{HI}/\\pi R_{HI}^2$ are very low, typically of order 0.5--1.0 M$_\\odot$~pc$^{-2}$. In NGC~3626 and NGC~4772, the \\HI ~is found in two concentric but distinct rings. In NGC~3900, the gas exterior to the optical edge is patchy, but can be traced 50~kpc to the south, giving R$_H$/R$_{25}$ $\\sim$ 4.5. Along the minor axis, the exterior \\HI ~is strongly asymmetric. (5) In all four \\HI ~maps, including that of NGC~4138, the velocity field is dominated by circular rotation, but significant departures from motion in a quiescent disk are evident. In all, the outer \\HI ~velocity contours suggest significant warping of the outer disk, roughly beginning at the optical edge and following Briggs' (1990) rule for warps, as commonly seen in other galaxies. The similiarities found in the five Sa galaxies are consistent with the scenario that the kinematically distinct gas components arise from slow, minor mergers of gas--rich satellites with an already--formed disk galaxy. Differences in their appearance thus arise from differences in the circumstances of the merger events. Retrograde primordial gas infall, as discussed by Thakar \\etal ~(1997) is also offered as a possible solution to produce counter--rotating components, although seemingly more difficult to justify in the growing number of known cases.  Though most numerical work on the merger process has focussed on equal mass progenitors, recent studies have begun to address the minor merger phenomenon (Quinn \\etal ~1993; Hernquist \\& Mihos 1995; Walker \\etal ~1996; Thakar \\etal ~1997; Bekki 1998; Thakar \\& Ryden 1998; Taniguchi ~1999). Although much of this work is preliminary, requiring further refinement of the range of galaxy characteristics and orbital parameters and the influence of gas dissipation, triaxiality and dark matter content, we can nonetheless examine the main characteristics of the five galaxies under the assumption that we may be seeing the results of varying acquisition scenarios.  The extensive \\HI ~distributions seen in four of the five galaxies discussed here may result from the slow accretion of a gas--rich satellite that is tidally stripped before the merger occurs. Accretion of a low density gas-rich satellite would minimize the disk heating problem and might favor the formation of extended \\HI ~disks or rings, but perhaps not inner rings. The formation of rings is seen in a variety of infalling gas models, both due to orbit crowding and to true self--gravitating rings. Also, Quinn \\etal ~(1993) suggest that transport of angular momentum outwards will tend to cause the disk to spread in radius. The newly--acquired gas may reach sufficient densities, if clumpy, to form a new generation of stars. This possibility is suggested by the faint, blue, low-surface brightness outer disk seen in NGC~4772 and a similar excess of light found around the counter-rotating Sab galaxy NGC~7217 (Buta \\etal ~1995).  Clearly the timescales of the merger events responsible for the individual characteristics noted here are likely to vary greatly. Relatively gas--poor and large--bulged, NGC~5854 may represent the result of the accretion of a more massive satellite travelling in a prograde orbit (Walker \\etal ~1996). Merger events might lead to bulge build--up and even induce large-scale star formation, leaving the galaxy in the post-starburst state in which NGC~5854 is found. Furthermore, one of the primary arguments made against the tidal forcing of warps is that the modes would damp out after a few galactic rotations. Perhaps in the cases that suggest strong warping of the \\HI, the minor mergers may be relatively recent events. Over time, the warps may damp, or be reexcited by additional interactions. If there are no impediments to star formation in a counter-rotating gas component, NGC~4138 with its two counter-rotating extended stellar disks may represented the future evolutionary state of NGC~3626.  The similarities among the results for these five galaxies lead to several questions and conundrums. (1) Is it key that the sample is composed of Sa galaxies? The accretion of moderate amounts of gas may also sweep up the gas contained within the primary disk itself, ultimately leading to a smoothing of the spiral structure. Do minor mergers lead to bulge build-up and disk heating, driving the post-merger morphology towards earlier spiral types? (2) Damping the process of disk heating is a challenge for models of mergers, even minor ones, if the accreted satellite is sufficiently massive or dense. The resultant disk heating will be relatively lessened if the satellite is predominantly gaseous. What are the conditions necessary to avoid over-heating the disk and disturbing the overall velocity field and spiral structure? (3) Is the fact that these galaxies are all unbarred important? Loss of angular momentum may drive the material inward toward the central regions where it accumulates in a nuclear disk as is often discussed for the kinematically decoupled cores seen in ellipticals. In the early stages of galaxy evolution, the formation of bars seems inevitable, but if the central concentration of mass is sufficient, the development of an inner Lindblad resonance will inhibit the central flow of gas (Sellwood \\& Moore 1999). Thus, since the progenitor primary was already a well-developed object, infalling material may be stopped at the inner Lindblad resonance, allowing the build--up of gas in a ring with subsequent star formation.  Does this process produce the inner \\Ha ~and \\HI ~rings seen here? Jungwiert \\& Palous (1996) propose that unbarred ringed galaxies must possess weak central oval distortions, while the two--stream instability picture of Lovelace \\etal ~(1997) for galaxies with counter-rotating components suggests the generation of strong m=1 spiral waves. Are the inner kinematic peculiarities evidence of such oval distortions or m=1 streaming? (4) Although these Sa galaxies are relatively isolated, they are all members of loose groups. In fact, Sa galaxies are found typically in denser environments than their later type counterparts. Is the relative morphological segregation seen across the spiral sequence related to the likelihood of minor mergers and their increased probability in environments characterized by moderate density and low velocity dispersion?  While examination of the rate of strongly disturbed galaxies may provide a reliable estimate of the major merger rate (e.g. Keel \\& Wu, 1995), minor mergers may be morphologically hidden, revealed only by the presence of their resultant kinematic peculiarities. Because of the ambiguities involved in disentangling the kinematic clues, it is not possible to trace with any degree of precision the evolutionary history of the multiple stellar and gaseous components evident in these galaxies. Food for thought, however, is the possibility that, over the course of its lifetime, a disk galaxy might undergo multiple minor accretion events with kinematic as well as morphological consequences. In particular, more complicated histories involving multiple minor mergers are likely to produce heterogenous properties, particularly those related to properties of ``form''. We suggest that the heterogeneity of the Sa class in particular is the outcome of such complicated life histories, the memory of which is signaled only by kinematic clues such as those identified here."
    },
    "0004/astro-ph0004389_arXiv.txt": {
        "abstract": "Under the standard model for recombination of the primeval plasma, and the cold dark matter model for structure formation, recent measurements of the first peak in the angular power spectrum of the cosmic microwave background temperature indicate the spatial geometry of the universe is nearly flat. If sources of \\lya resonance radiation, such as stars or active galactic nuclei, were present at $z\\sim 1000$ they would delay recombination, shifting the first peak to larger angular scales, and producing a positive bias in this measure of space curvature. It can be distinguished from space curvature by its suppression of the secondary peaks in the spectrum. ",
        "introduction": "The measurements of the anisotropies of the cosmic microwave background (the CMB) offer extraordinarily powerful tests of the relativistic Friedmann-Lema\\^\\i tre cosmological model and the nature of the early stages of cosmic structure formation (e.g. \\cite{Junetal96} 1996). Indeed, the recent detection of the first peak in the angular power spectrum of the CMB temperature indicates space is close to flat (\\cite{Miletal99} 1999; \\cite{Meletal99} 1999; \\cite{deBer00} 2000). The interpretation is quite indirect, however so we must seek diagnostics for possible complications.  There are relatively few physical effects that can {\\it increase} the angular scale of the first peak.  Because of this fact, the observed large angular scale of the peak is believed to strongly disfavour open universes. Of the fundamental cosmological parameters, only a Hubble constant well in excess of observations can substantially increase the scale of the peak in an open or flat universe (\\cite{HuSug95} 1995).  One possibility is that some process at redshift $z\\sim 1000$ delayed recombination of the primeval plasma. This would increase the sound horizon at last scattering and decrease the angular size distance to last scattering, moving the first peak of the CMB temperature fluctuation spectrum to smaller angular wavenumber (\\cite{HuWhi96} 1996; \\cite{WelBatAlb99} 1999), and biasing the measure of space curvature to a too large (more positive) apparent value. Another consequence of delayed recombination is that it would suppress the secondary peaks due to an increase in the time available for acoustic oscillations to dissipate  (\\cite{Sil68} 1968; \\cite{HuWhi96} 1996).   Because the first peak is not observed to have suffered substantial dissipation, recombination in the delayed model must be rapid compared with the cosmological expansion rate. Thus if recombination were delayed by ionizing radiation from decaying dark matter (e.g. \\cite{SarCoo83} 1983; \\cite{Sciama} 1991; \\cite{Elletal92} 1992) or evaporating primeval black holes (\\cite{pbh} 1987), or by thermal energy input from cosmic string wakes (\\cite{WelBatAlb99} 1999), the source would have to terminate quite abruptly and well before $z\\sim 100$ when the universe starts to become optically thin even if the baryons are fully ionized.  Here we consider a picture that more naturally allows rapid recombination: sources of radiation at $z\\sim 1000$ that produce many more photons in the resonance \\lya line than ionizing photons, in the manner of a quasar. The \\lya photons would increase the population in the principal quantum number $n=2$ levels of atomic hydrogen, increasing the rate of photoionization from $n=2$ by the CMB. Since the rate of thermal photoionization from $n=2$ varies rapidly with redshift at $z \\lsim 1000$, the delayed recombination is rapid, and the residual ionization can be small enough that the optical depth for Thomson scattering after recombination is well below unity. Thus the height of the first peak in the spectrum of CMB temperature fluctuations is little affected by the delayed recombination. The shift in the angular wavenumber at the first peak is modest also, even if early sources produce many \\lya photons per baryon, but the shift can be considerably larger than the projected precision of the measurements. Thus it is fortunate that we seem to have an unambiguous diagnostic for delayed recombination in the suppression of the secondary peaks.  After this work was substantially complete we learned that the recent measurements by the BOOMERanG experiment in fact favor a substantial suppression of the second peak (\\cite{deBer00} 2000). We emphasize that there are many other ways to account for this effect, and that there is no evidence for the early source of \\lya radiation assumed in our picture. Within the conventional adiabatic CDM model the recombination history is well understood (\\cite{SeaSasSco00} 2000 and earlier references therein), so there is good reason to expect the residual fluctuations in the CMB may be related to the cosmology in a simple and computable way. But it is good science to bear in mind the possibility that Nature is more complicated than our ideas, a rule that has particular force in cosmology because of the limited empirical basis.  \\begin{figure} \\plotone{figure1.ps} \\caption{The ionization fraction, $x$, as a function of redshift for $\\epsilon_{i}=0$ and various values of $\\epsilon_{\\alpha}$.} \\end{figure} ",
        "conclusions": "Our model for delayed recombination with $\\epsilon _{\\alpha} = 1$ yields a 5\\% shift of the first peak and a reduction of the secondary peak by 10\\%, an observationally interesting effect. The sources must have $\\epsilon _{\\rm i} \\ll 1$ (see Figure 2). This condition requires that the \\lya sources (perhaps hot stars or quasars) be surrounded by envelopes of neutral primeval material which strongly absorbs the ionizing radiation.  Our picture does not follow from the conventional CDM model for structure formation. Apparent problems with excess small-scale clustering in this model (\\cite{Mooetal99} 1999; \\cite{Klyetal99} 1999) have motivated discussions of modifications (\\cite{SpeSte99} 1999; \\cite{KamLid99} 1999; \\cite{Pee00} 2000; \\cite{HuBarGru00} 2000). These modifications would tend to delay the appearance of the first generation of gravitationally bound systems, however, which is in the opposite direction to what is postulated here. These lines of thought do not rule out early sources of \\lya photons, of course, but they do suggest one might best look for a source outside the adiabatic CDM model. Perhaps cosmic strings produced wakes that were subdominant to the primeval CDM density fluctuations in determining the mass fluctuation power spectrum but did produce occasional non-gaussian density fluctuations (\\cite{ConHinMag99} 1999) large enough to have collapsed to stars or active black holes that could have produced \\lya photons.  To summarize, we have argued that the picture of delayed recombination caused by early sources of \\lya radiation, modeled after quasar spectra, but with suppressed X-ray emission, has the virtue that it naturally preserves the observed first peak in the CMB temperature angular power spectrum while shifting the peak to larger scales. The picture is ad hoc but important as a conceivable complication in the application of an exceedingly powerful cosmological test. The most direct diagnostic for delayed recombination seems to be the suppression of the secondary peaks in the spectrum. The amplitude of the third peak, which may be measured by the BOOMERanG experiment or the MAP satellite\\footnote{\\tt http://map.gsfc.nasa.gov}, is crucial for distinguishing the effect of delayed recombination from that of adjustments of the values of space curvature, the baryon density, the Hubble constant or the shape of the spectrum of the primeval density fluctuations.  If the measured secondary peaks agreed with the standard model for recombination with astronomically acceptable cosmological parameters that fit the primary peak, it would convincingly rule out our model for delayed recombination."
    },
    "0004/astro-ph0004206_arXiv.txt": {
        "abstract": "We present spectra for a sample of radio sources from the FIRST survey, and use them to define the form of the redshift distribution of radio sources at mJy levels. We targeted 365 sources and obtained 46 redshifts (13 per cent of the sample). We find that our sample is complete in redshift measurement to R $\\sim 18.6$, corresponding to $z\\sim 0.2$. Galaxies were assigned spectral types based on emission line strengths. Early-type galaxies represent the largest subset (45 per cent) of the sample and have redshifts $0.15\\la z \\la 0.5$ ; late-type galaxies make up 15 per cent of the sample and have redshifts $0.05\\la z \\la 0.2$; starbursting galaxies are a small fraction ($\\sim 6$ per cent), and are very nearby ($z\\la 0.05$). Some 9 per cent of the population have Seyfert1/quasar-type spectra, all at $z\\ga 0.8$, and there are 4 per cent are Seyfert2 type galaxies at intermediate redshifts ($z\\sim 0.2$).   Using our measurements and data from the Phoenix survey (Hopkins et al., 1998), we obtain an estimate for $N(z)$ at $S_{1.4 \\rm {GHz}}\\ge 1$ mJy and compare this with model predictions. At variance with previous conclusions, we find that the population of starbursting objects makes up $\\la 5 $ per cent of the radio population at S $\\sim 1$~mJy. ",
        "introduction": "An accurate definition of the redshift distribution of radio sources at faint flux densities has become particularly important in the last decade for both radio astronomy and cosmology. It is critical, for example, in testing radio-source unification (e.g. Jackson \\& Wall, 1998), and in large-scale structure studies (e.g. Loan et al. 1997, Magliocchetti et al. 1999) to permit the conversion of angular clustering estimates to the spatial clustering estimates required to evaluate structure formation models.  Classical large-scale-structure studies have used wide-area optical and IR surveys to measure the clustering of galaxies, but these are limited to low redshifts, with a peak redshift selection of $\\la 0.1$ (e.g. APM, Maddox et al., 1990; IRAS, Fisher et al., 1993). Deep small-area surveys such as the Hubble Deep Field North (Williams et al., 1996) or the CFRS survey (Lilly et al., 1995) have been used to measure clustering to much higher redshifts (see e.g. Le Fevre et al., 1995; Magliocchetti \\& Maddox, 1998), but they probe small volumes and can measure clustering only on small scales ($\\la 1$ Mpc). On the other hand radio objects are detected to high redshifts ($z\\sim 4$) and sample a much larger volume for a given number of galaxies, so that they have the potential to provide information on the growth of structure on large physical scales. Recently, clustering statistics in deep radio surveys have been measured (Cress et al., 1997; Loan et al., 1997; Baleisis et al., 1997; Magliocchetti et al., 1998) and these analyses have shown radio sources to be reliable tracers of the mass distribution, even though they may be biased (Magliochetti et al., 1999).  However the relationship between angular measurements and the physically meaningful spatial quantities is highly uncertain without an accurate estimate of the redshift distribution of radio objects $N(z)$ (Magliocchetti et al., 1999). The current estimates of $N(z)$ for radio sources at mJy levels are largely based on predictions from the local radio luminosity functions and evolution (see e.g. Dunlop \\& Peacock, 1990), and are poorly defined.  Indeed the predictions require the knowledge of more than one luminosity function. Deep radio surveys (Condon \\& Mitchell, 1984; Windhorst et al., 1985; Fomalont et al., 1993) have shown a flattening of the differential source count $\\frac{dN}{dS}$ below S $\\sim 10$ mJy. This flattening is generally interpreted as due to the presence of a population of radio sources differing from the radio AGN which dominate at higher flux densities. Condon (1984) suggested a population of strongly-evolving normal spiral galaxies, while others (Windhorst et al., 1985; Danese et al., 1987) claimed the presence of an actively star-forming galaxy population. Observations supported this latter suggestion; the identifications of many of the faint sources are with galaxies with spectroscopic and photometric properties similar to `IRAS galaxies' (Franceschini et al., 1988; Benn et al., 1993). The model predictions of $N(z)$ at faint flux densities from the luminosity functions of Dunlop \\& Peacock (1990) diverge greatly because of inadequate definition of the radio AGN luminosity functions; but in addition the contribution to $N(z)$ of the starburst population needs consideration.  In order to define $N(z)$ and the population mix at mJy levels, we have observed radio sources from the FIRST survey (Becker et al., 1995) using the WYFFOS multi-object spectrograph on the William Herschel Telescope in 8 1-degree diameter fields. We show how these observations constrain both the form of $N(z)$ and the population mix at S $\\sim 1$~mJy for $z\\la 0.3$. Section 2 of the paper introduces this radio sample, while Sections 3 and 4 describe acquisition and reduction of the data. In Sections 5, 6 and 7 we present the radio, spectroscopic and photometric properties of the sample. Section 8 is devoted to the analysis of the observed $N(z)$ and comparison with models; the conclusions are in Section 9. \\begin{figure} \\vspace{19cm}  % \\special{psfile=f02new.ps hoffset=-30 voffset=160 vscale=50 hscale=50} \\special{psfile=f22new.ps hoffset=-30 voffset=0 vscale=50 hscale=50} \\special{psfile=eliasnew.ps hoffset=-30 voffset=-40 vscale=35 hscale=50} \\caption{Area covered by the 8 observed fields. Crosses are for the radio sources while open circles show the positions of the optical fibres. \\label{fig:fields}} \\end{figure} ",
        "conclusions": "We have carried out multi-object spectroscopy of an unbiased selection of FIRST radio sources (S $\\ga 0.8$~mJy) by placing fibres at the positions of 365 sources ($\\sim$ 69~per~cent of the complete radio sample). The spectra obtained have enabled us to measure 46 redshifts, $\\sim 13$~per~cent of the targeted objects.  APM data have provided morphology and photometric data for the corresponding optical identifications. The photometry shows that redshift measurements were obtained only for objects brighter than R $\\simeq 20.5$~mag; from the tight R-$z$ relation observed, the redshift sample is estimated to be $\\sim$100~per~cent complete to R=18.6~mag.  The objects in the spectroscopic sample with R $\\la 20.5$ are a mixture of early-type galaxies at relatively high redshifts, $z\\ga 0.2$ ($\\sim 45$~per~cent of the sample), late-type galaxies at intermediate redshifts, $0.02\\la z\\la 0.2$ ($\\sim15$~per~cent), and very local starburst galaxies with $z\\la 0.05$ ($\\sim 6$~per~cent). We also found a number of Seyfert1/quasar type galaxies, all at $z\\ga 0.8$ ($\\sim 9$~per~cent of the sample), two Seyfert2's (4~per~cent), and 4 stars. The number of objects with featureless spectra are most probably early-type galaxies, given the shape of the continuum, the lack of emission lines and the red colours. Using the R-$z$ relation derived for our sample, we conclude that they are likely to have $z\\ga 0.3$. The redshift incompleteness does not depend on radio flux density; optical apparent magnitude is the only identifiable factor. Using again the R-$z$ relation determined for the sample, we estimate $\\sim$ 100~per~cent completeness for the spectroscopic sample up to $z\\sim0.3\\pm 0.1$.  To define $N(z)$ at S$_{1.4\\,\\rm{GHz}}$ as well as possible, we have combined our sample with the Phoenix spectroscopic sample (Georgakakis et al., 1999), which we estimate to be complete (for non-AGN objects) to $z=0.9$. The combined distribution (Figure~\\ref{fig:N_zgt1}) shows the following:  \\begin{itemize}  \\item The redshift distribution rises up to $z\\simeq 0.05$ and is then approximately leveled to $z=0.3$.  \\item The total number of sources predicted by the luminosity-function models of Dunlop \\& Peacock (1990) agrees with that observed.  \\item None of the models provides a good fit to the shape of $N(z)$. The percentage of objects at $z\\la 0.1$ is seriously overestimated in almost all the models, especially for the pure-luminosity and luminosity/density evolution models (DP 6 and 7) that feature an unrealistic $z\\sim 0$ spike.  \\item The normalization of the models appears to be too high to fit the observed $N(z)$ for $z \\la 0.2$.  This disagreement may imply that the model shape is wrong, and there are more sources at $z\\ga 0.3$ than indicated by the models.  Alternatively, the discrepancy could be due to observing a low density by chance, given a large variance in galaxy density caused by strong clustering.  \\item The $N(z)$ at $S_{1.4\\,{\\rm GHz}} = 1$~mJy is dominated by AGN, and starburst objects constitute less than 5~per~cent of the total. This is a robust conclusion. More starburst galaxies would have substantially raised the proportion of objects with redshift determinations; and a significant intrusion of starburst galaxies at the lowest radio flux densities would have resulted in a higher success rate in redshift determination with decreasing flux density. The great majority of objects in the sample at this level are AGN associated with early-type galaxies whose optical continua and weak-to non-existent emission lines place them at or below the limit of our spectroscopic survey.  \\end{itemize}  The accurate definition of the low$-z$ end of the $N(z)$ relation has impact in four areas: (i) the population mix, which is critical for testing and refining dual-population unified models; (ii) the definition of the local luminosity functions, important for modelling both the form and cosmic evolution of the overall luminosity functions; (iii) the derivation of spatial measurements of the large-scale structure from angular measurements; and (iv) constraints which it enables to be placed on the global star-formation-rate up to $z=0.3$.  \\vspace{1cm} \\noindent {\\bf ACKNOWLEDGEMENTS}\\\\ MM acknowledges support from the Isaac Newton Scholarship. We thank David Helfand for extremely helpful discussions.  The WHT is operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. GC acknowledges a PPARC Postdoctoral Research Fellowship."
    },
    "0004/astro-ph0004030_arXiv.txt": {
        "abstract": "In this work we present the predictions of a modified version of the ``two-infall model'' (Chiappini et al. 1997 - CMG) for the evolution of $^3$He, $^4$He and D in the solar vicinity, as well as their distributions along the Galactic disk. In particular, we show that when allowing for extra-mixing process in low mass stars (M$ <$ 2.5 $M_{\\odot}$), as predicted by Charbonnel and do Nascimento (1998), a long standing problem in chemical evolution is solved, namely: the overproduction of $^3$He by the chemical evolution models as compared to the observed values in the sun and in the interstellar medium. Moreover, we show that chemical evolution models can constrain the primordial value of the deuterium abundance and that a value of (D/H)$_p$ $ < $ 3 $\\times$ 10$^{-5}$ is suggested by the present model. Finally, adopting the primordial $^4$He abundance suggested by Viegas et al. (1999), we obtain a value for $\\Delta Y/\\Delta Z$ $\\simeq$ 2 and a better agreement with the solar $^4$He abundance. ",
        "introduction": "As discussed by Tosi (this volume), chemical evolution models are useful both to derive the primordial abundances of D, $^3$He and $^4$He and to give informations on stellar nucleosynthesis. In this work we show the predictions of the two-infall model (CMG) for the chemical evolution of the above elements in the solar vicinity and for their distribution along the galactic disk. We adopt a new version of the two-infall model which includes the contribution by novae enrichment and the new proposed mechanism of extra-mixing in low mass stars (Charbonnel, Sackman, this meeting). The model was calibrated to the solar galactocentric distance of 8 kpc (we were still adopting 10 kpc in CMG to better compare our predictions with the ones of Matteucci and Fran\\c cois 1989). This model assumes two main infall episodes for the formation of the halo (and part of the thick disk) and thin disk, respectively. The timescale for the formation of the thin disk is much longer than that of the halo, implying that the infalling gas forming the thin disk comes not only from the halo but mainly from the intergalactic medium. The timescale for the formation of the thin disk is assumed to be a function of the galactocentric distance, leading to an inside-out picture for the Galaxy disk buildup. The two-infall model differs from other models in the literature mainly in two aspects: i) it considers an almost independent evolution between the halo and thin disk components (see also Pagel \\& Tautvaisiene 1995) and ii) it assumes a threshold in the star formation process (Kennicutt 1989). The last point has important consequences for the predicted abundance gradients (Chiappini, Matteucci \\& Romano 2000 - CMR). ",
        "conclusions": ""
    },
    "0004/astro-ph0004024_arXiv.txt": {
        "abstract": "Far-ultraviolet spectra of Venus and Mars in the range 820~--~1840~\\AA\\ at $\\sim$4 \\AA\\ resolution were obtained on 13 and 12 March 1995, respectively, by the Hopkins Ultraviolet Telescope (HUT), which was part of the Astro-2 observatory on the Space Shuttle {\\it Endeavour}. Longward of 1250 \\AA , the spectra of both planets are dominated by emission of the CO Fourth Positive (\\api\\ -- \\xsig ) band system and strong \\Oone\\ and \\Cone\\ multiplets.  In addition, CO Hopfield-Birge bands, \\bsig\\ -- \\xsig\\ (0,0) at 1151 \\AA\\ and \\csig\\ -- \\xsig\\ (0,0) at 1088 \\AA , are detected for the first time, and there is a weak indication of the \\epi\\ -- \\xsig\\ (0,0) band at 1076 \\AA\\ in the spectrum of Venus.  The $B - X$ band is blended with emission from \\Oone\\ \\lam1152.  Modeling the relative intensities of these bands suggests that resonance fluorescence of CO is the dominant source of the emission, as it is for the Fourth Positive system.  Shortward of Lyman-$\\alpha$, other emission features detected include \\Otwo\\ \\lam 834, \\Oone\\ \\lam 989, \\Hone\\ Lyman-$\\beta$, and \\None\\ \\lam\\lam 1134 and 1200.  For Venus, the derived disk brightnesses of the \\Oone , \\Otwo , and \\Hone\\ features are about one-half of those reported by \\citet{hor91} from {\\it Galileo} EUV measurements made in February 1990.  This result is consistent with the expected variation from solar maximum to solar minimum.  The \\Arone\\ \\lam\\lam 1048,1066 doublet is detected only in the spectrum of Mars and the derived mixing ratio of Ar is of the order of 2\\%, consistent with previous determinations. ",
        "introduction": "Ultraviolet observations of the atmospheres of Venus and Mars, primarily from fly-by or orbiting spacecraft, but also from platforms above the terrestrial atmosphere, have played an important role in the study of the composition and structure of the \\cotwo\\ atmospheres of these planets.  Early spacecraft experiments were broad-band photometers that provided information about the spatial distribution of the emissions in the passband but required remote spectroscopic measurements to assure the identity of the emitting species.  A thorough review of this subject is given by \\citet{pax92}.  Once the emitting species were known, narrowband polychromators were developed and flown on missions such as {\\em Mariner 10} \\citep{bro74} and {\\em Venera 11} and {\\em 12} \\citep{ber81}, although the interpretation of the data from these instruments was far from unambiguous.  With {\\em Mariners 6, 7} and {\\em 9}, flown to Mars in the early 1970s, and {\\em Pioneer Venus Orbiter} launched to Venus in 1978, ultraviolet spectrometers were included in the payloads and provided modest spectral resolution.  Very little additional spectroscopy, particularly from newer generations of Earth-orbiting observatories such as the International Ultraviolet Explorer (IUE) and the Hubble Space Telescope (HST), has been done and is also summarized by \\citeauthor{pax92}.  In February 1990, {\\em Galileo} flew by Venus on its way to Jupiter and the ultraviolet spectrometers on board made observations that were reported by \\citet{hor91}.  The flight of the Hopkins Ultraviolet Telescope (HUT) on the Astro-2 mission on the Space Shuttle {\\it Endeavour} in March 1995 provided an opportunity to measure the ultraviolet disk spectra of both Venus and Mars at a spectral resolution ($\\sim 4$ \\AA) significantly higher than any of the prior spacecraft observations.  Moreover, since the first order spectral range of HUT extended to wavelengths as short as 830 \\AA, the observations of Venus enabled the resolution of the identity of the emissions recorded in the narrow-band photometric channels of the {\\it Venera 11} and {\\it 12} EUV instruments that had been interpreted in terms of analogue terrestrial spectra \\citep{ber81}. This paper presents the HUT disk spectra of Venus and Mars together with the spectral identifications and the disk-averaged brightnesses. CO fluorescence in both the $B - X$ (0,0) and $C - X$ (0,0) Birge-Hopfield bands is identified in both Venus and Mars and shown to be consistent with current models of the CO/\\cotwo\\ mixing ratios in the atmospheres of these planets.  Several other emissions are definitively identified for the first time. ",
        "conclusions": "We have obtained far-ultraviolet spectra of Venus and Mars in the range 820~--~1840~\\AA\\ at $\\sim$4 \\AA\\ resolution with the Hopkins Ultraviolet Telescope (HUT) during the {\\em Astro-2} mission in March 1995. The spectra of both planets are rich in CO band emission, some of the systems being identified for the first time, together with strong \\Oone\\ and \\Cone\\ multiplets.  Resonance fluorescence is identified as the dominant source of the CO emission.  Atomic nitrogen emissions are also identified in the spectrum of Venus, while the \\Arone\\ doublet is seen only in the spectrum of Mars.  These spectra, obtained at higher spectral resolution than was possible from earlier fly-by and orbiting missions to these planets, elucidates and extends the earlier spectroscopic measurements and should provide guidance in the interpretation of the far-ultraviolet spectra obtained during the recent {\\em Cassini} fly-by of Venus \\citep{ste99}."
    },
    "0004/astro-ph0004212_arXiv.txt": {
        "abstract": "Poynting flux driven outflows from magnetized rotators are a plausible explanation for gamma-ray burst  engines.  We suggest a new possibility for how such outflows might transfer energy into radiating particles. We argue that the Poynting flux drives non-linearly  unstable large amplitude electromagnetic waves (LAEMW) which ``break''  at radii  $r_t \\sim 10^{14} $ cm where the MHD approximation becomes inapplicable. In the ``foaming'' (relativisticly reconnecting) regions formed during the wave breaks the random electric fields stochastically  accelerate particles to ultrarelativistic energies which then radiate in  turbulent  electromagnetic fields. The typical energy of the emitted photons is a fraction of the fundamental Compton energy $ \\epsilon \\sim f  \\hbar c/r_e $ with $f \\sim 10^{-3}$ plus additional boosting due to the bulk motion of the medium. The emission properties are similar to synchrotron radiation, with a typical  cooling time $\\sim 10^{-4}$ sec. During the wave break, the plasma is also bulk accelerated in the outward radial direction and at larger radii can produce afterglows due to the interactions with external medium. The near equipartition fields required by afterglow models maybe due to magnetic field regeneration in the outflowing plasma (similarly to the field generation by LAEMW of laser-plasma interactions)  and mixing with the upstream plasma. ",
        "introduction": "A Poynting flux driven outflow from a magnetized rotator is a promising paradigm for gamma-ray burst (GRB) engines and there have been various implementations of this concept (c.f. Usov 1992; Thompson 1994; Blackman et al. 1996; M\\'esz\\'aros \\& Rees 1997; Kluzniak \\& Ruderman 1998). Such models are appealing in light of the fact that neutrino driven GRB emission from compact object mergers (Ruffert \\& Janka 1998; Janka et al. 1999) might fall short in radiative luminosity. Also, Poynting flux models may provide a source of magnetic fields which could help alleviate the  field amplification problem in GRBs themselves and in afterglows which requires fields  of the order of equipartition (\\cite{WG98,FWK99,GPS99}).  Scenarios that could produce  Poynting flux dominated outflows (PFDOs) require a source of magnetic fields  $\\ga 10^{15}$ Gauss, and rotation speeds of order $\\Omega \\sim 10^4$/sec. The total available energy from a  compact ($R_0 \\sim 10^6$ cm, $M \\ge M_{\\odot}$) object is then $ M \\Omega^2 R_0^2/2 \\ge 10^{53}$ ergs and a dipole type luminosity $ 2 B_0^2 \\, R_0 ^6 \\,\\Omega^4/ 3 c^3 \\sim 2 \\times 10^{50}$ ergs/sec. Such combinations could be generated in a number of plausible cases: an accretion torus surrounding a black hole that formed from a neutron star merger (Rees \\&  M\\'esz\\'aros 1997), strongly magnetized neutron stars (``magnetars'' Duncan \\& Thompson 1992) possibly formed from accretion induced collapse of a white dwarf with possible dynamo amplified fields in the hot young neutron star (Usov 1992; Duncan \\& Thompson 1992, Blackman et al. 1996), failed supernova Ib (Woosley 1993),  rapidly rotating magnetized black holes undergoing a Blandford-Znajek type energy extraction from supernovae (Lee et al. 1999,Brown et al., 2000) or hyper-accreting accretion disks (Popham et al. 1999), with an MHD  dynamo (Araya-Gochez 2000).    Without further specifying the nature of the central engine we suppose that GRB are due to such stellar mass objects, so that the  conditions  required for the generation of PFDOs are satisfied. Here we are interested in how Poynting flux models accelerate particles and produce gamma-ray emission.  The suggestion that GRBs are PFDOs whose energy is converted to particles and escaping radiation only at large distances from the central rotator potentially accounts for some important characteristics of GRBs (c.f. Blackman et al. 1996, Rees \\& \t M\\'esz\\'aros 1997) (i) mass loading -  PFDO can  be launched without much matter, and particularly without much baryon contamination (as in pulsars, the wind may carry no baryons at all); (ii) compactness problem - the PFDO converts its energy effectively into particles and gamma-rays at distances many orders of magnitude from the central engine, (iii) collimation -   PFDO have a preferred direction along the rotation axis of the central engine; (iv) jets  observed in pulsars and (probably) AGNs provide examples of operation of PFDO; (v) the energy from Poynting flux can be easily converted into high frequency electromagnetic radiation.  The presence of strong magnetic fields from the rotator provides  electromagnetic energy which might be tapped or converted into the fields which accelerate the radiating particles.  PFDOs may differ from the conventional blast wave models in several important respects. The lepto-baryonic content of the outflow in the inner region is expected to be small by analogy to pulsar winds, and the energy carried by matter in the central region will be less than the energy carried by EM fields. As in the relativistic pulsar wind, the only other well established PFDO,  the baryons may be absent altogether while energy flux in the electron-positron pairs may be as small as one millionth of the Poynting flux.  The conventional estimates of the relativistic expansion Lorentz factors and density of the outflowing plasma used in the framework of blast wave models may not be applicable to the Poynting dominated outflows because of the strong magnetic fields (Paczy\\'{n}ski 1990, Usov 1999). Strong fields dramatically reduce the importance of pair production from photon-photon interactions relative to the photon-magnetic field  interactions. Other exotic QED processes, like photon splitting (e.g., Baring 1991, Adler 1971), will also become important for magnetic fields larger than the critical value of $4.4 \\time 10^{13}$ Gauss. At this moment, we are unaware of the calculations of the Compton scattering cross section in supercritical magnetic fields, but we expect that it will also be strongly suppressed.  The above suggests that  in the central regions of  PFDOs, the magnetic pair production and photon splitting  may be more important than pair production in photon collisions. Note that this situation is realized in pulsars. There acceleration of particles to energies beyond $10^6$ MeV proceeds relatively quietly, without producing a large number of observable photons. Very few rotationally powered  pulsars are observed at high energies. Thus conventional estimates of the Lorentz factors and densities  in GRBs based on two  photon pair production  may be irrelevant.  To estimate a ``revised'' density of the generated pairs in the strong magnetic rotator case applicable to GRB, we use the analogy with pulsars where the number of pairs produced is related to the local Goldreich-Julian density $n_{GJ} = \\Omega B /( 2 \\pi  e c)$ (Goldreich \\& Julian (1969)). The pair number density is usually $\\lambda \\sim 10^3 -10^5$ times larger than $n_{GJ}$ in the inner regions. The factor $\\lambda$ depends sensitively on the magnetic field, its curvature and accelerating potential. A highly curved magnetic field produces denser and less relativistic outflows.   In this  paper we  address the fundamental but  unresolved issue of how Poynting flux is converted into gamma-rays and the subsequent emission characteristics. In section 2 we give the basic picture of the magnetized rotator magnetosphere and the propagation of large amplitude electromagnetic waves. In section 3 we discuss the breaking of these waves, and the predicted emission from random electromagnetic fields is discussed in section 4. There we address the spectrum and time dependence of the emission. The afterglow possibilities from our picture are addressed in section 5. We conclude in section 6. ",
        "conclusions": "The two main points of this work may be summarized as follow: (i) Poynting flux dominated outflows, which so far has been invoked only in the context of central engines, possesses internal instabilities which may also explain radiation generation, (ii) emission from stochastically accelerated particles  in turbulent electromagnetic fields, when electric fields are  as important as magnetic fields in the acceleration and radiation process, is a viable mechanism for the GRB emission.  Our approach may also naturally explain the result that the peak of GRB emission varies only over a small range of value from burst to burst. (Brainerd 1994, 1997).  For our maximally efficient mode of acceleration, the peak, when boosted into observer frame, becomes $ \\epsilon_{max} \\sim  f \\hbar c/r_e  \\Gamma_{\\rm bulk}$. If this maximally efficient regime were operating, the bulk gamma-factor $\\Gamma_{\\rm bulk}$ would not be  much larger than 10. The location of this peak is  a very weak function of the parameters of the underlying rotator, $f \\sim \\Omega ^{1/7} $ and is independent of the progenitor's magnetic field.  This can be compared  to phenomenological internal shock models in which the peak energy is proportional to $\\gamma_{min}^2 \\om_B \\Gamma_{bulk}$, where $\\gamma_{min}$ is the low energy cutoff to the electron power law spectrum. All three quantities here are usually taken to be free parameters. It seems unlikely for their combination to be almost constant from burst to  burst, and within each pulse of a given burst unless some physics dictates this to be the case. We have tried to add some of this physics. In our approach the remaining free  parameter is the Lorentz factor of the bulk flow.  PFDOs may also resolve the problem of the magnetic field generation since in the  Poynting flux driven  outflows the large magnetic fields are supplied by the source. These fields do not necessarily have to be generated in the external shock for the afterglow since the shock is a current sheet through which outflow and ambient particles mix (Smolsky \\& Usov 2000).  At the same time, however there exists an unexplored analogy between field generation mechanisms in laser driven plasmas and GRB outflows that will have to be understood to determine the scale and structure of the outlfow  magnetic field.   Other observational properties of the GRBs that may be explained in our framework: (i) GRBs show no correlation between the spectra and other micropulse characteristics (e.g., Lee \\& Bloom 2000) - this is a direct consequence of the turbulent EM acceleration which produces photons with the frequency $\\sim c/r_e$; (ii) the fact that pulses peak earlier at high frequencies and that bursts have shorter duration  at higher energies  may be due to the initial development of the turbulence: if the turbulence develops from large scales to small scales then initially the acceleration may be  more effective since it is due to large scale electric fields (with coherence length larger than the radiative length). In this case  initially particles are accelerated to the limiting energies $\\gamma \\sim \\sqrt{r_L /r_e}$ while later, when the  coherence length of the turbulence becomes smaller than the  radiative length, the particle distributions soften  and radiation spectra emanates at lower frequencies. The typical time for such a cascade should be of the order of the \"vortex overturn time\" on the largest scale of the turbulence, which may be on the order of seconds. (iii) the absence of correlation between the pulses and overall burst characteristics, interpreted as arising from random and independent  emission episodes (Lee \\& Bloom 2000), is natural in our model since each wave overturn happens independently. (iv) The temporal characteristics of the microbursts, FREDs,  and the hard-to-soft spectral evolution is a consequence of synchrotron cooling of the  reservoir of energy released during wave overturn. (v) composite structure of GRBs (a burst being a sum of many independent  emission events) (Stern \\& Svennson 1996) naturally follows from the model - each overturning region produces an independent microburst. (vi)  the average power density spectrum of GRBs is well described as being due to selfsimilar turbulent-type process near marginal stability (Stern 1999). This is reminicent of the cellular automata model of solar reconnecting regions (Lu \\& Hamilton 1991) and may be related to the reconnecting \"foaming\" regions in our model.   We also would like to point out that the particular mechanism of the wave instability, the wave overturn during MHD-wave transition, may not necessarily be the only one. Other plasma instabilities may contribute to the generation of EM turbulence."
    },
    "0004/astro-ph0004162_arXiv.txt": {
        "abstract": "We present a new sample of 415 bright QSOs and Seyfert~1 nuclei drawn from the Hamburg/ESO survey (HES). The sample is spectroscopically 99\\,\\% complete and well-defined in terms of flux and redshift limits. Optical magnitudes are in the interval $13\\la B_J\\la 17.5$, redshifts range within $0< z < 3.2$. More than 50\\,\\% of the objects in the sample are new discoveries. We describe the selection techniques and discuss sample completeness and potential selection effects. There is no evidence for redshift-dependent variations of completeness; in particular, low-redshift QSOs -- notoriously missed by other optical surveys -- are abundant in this sample, since no discrimination against extended sources is imposed. For the same reason, the HES is not biased against QSOs multiply imaged due to gravitational lensing. The sample forms the largest homogeneous set of bright QSOs currently in existence, useful for a variety of statistical studies. We have redetermined the bright part of the optical quasar number-magnitude relation. We confirm that the Palomar-Green survey is significantly incomplete, but that its degree of incompleteness has recently been overestimated. ",
        "introduction": "While the total number of known quasars is growing rapidly, most of these are too faint for detailed investigations. Existing catalogues are mostly inhomogeneous in composition and not representative, featuring disproportionately high fractions of radio-loud, X-ray- and infrared-selected QSOs especially among the brightest known QSOs. To homogeneously sample the bright part of the quasar population in the optical regime, substantial fractions of the extragalactic sky have to be covered with efficient surveying techniques. The only presently available such catalogue, the Palomar-Green Bright Quasar Sample (BQS; Schmidt \\& Green \\cite{schm+gree:83:QEBQS}), is now known to be substantially incomplete (Goldschmidt et al.\\ \\cite{goldschmidt*:92:SDBQ}; K\\\"ohler et al.\\ \\cite{koehler*:97:LQLF}). A similar project in the southern hemisphere, the Edinburgh/Cape Survey (ECS; Stobie et al.\\ \\cite{stobie*:97:EC1}), has just produced first output, and the completeness of this survey remains to be assessed. Moreover, both BQS and ECS discriminate against low-$z$ QSOs with extended host galaxies, and the photometric UV excess selection technique confines both surveys to $z<2.2$ QSOs.  In 1990 we have started the `Hamburg/ESO survey' (HES), a new wide-angle survey for bright QSOs and Seyferts, as an ESO key programme. The survey is based on digitised objective-prism photographs taken with the ESO Schmidt telescope, covering essentially the entire southern extragalactic sky. QSO candidates are selected with largely automated procedures, minimising human interaction and establishing well-defined input samples for follow-up slit spectroscopy. A description of initial design and first results of the HES was given by Wisotzki et al.\\ (\\cite{wisotzki*:96:HES1}; hereafter Paper~1), where also the driving science objectives behind the survey have been listed. A list of 160 newly discovered QSOs and Seyfert~1 galaxies was published by Reimers et al.\\ (\\cite{reimers*:96:HES2}; hereafter Paper~2). The first flux-limited sample of 55 QSOs, distributed over an effective area of 611\\,deg$^2$, was presented and analysed by K\\\"ohler et al.\\ (\\cite{koehler*:97:LQLF}), to measure the surface density of bright quasars and to estimate the combined local luminosity function of quasars and Seyfert~1 nuclei.  In this paper we present a major expansion of our earlier work. The survey area has been sixfolded, and the QSO selection techniques have evolved significantly. We have constructed a new flux-limited sample of 415 optically bright QSOs that will be useful for a wide variety of statistical investigations. We describe the selection techniques used to build the sample and discuss its completeness, in comparison with other surveys as well as in terms of expected redshift-dependent selection effects. In providing these details, this paper serves also as a reference for future publications of HES-selected QSO samples. The paper closes with a reassessment of the surface density of bright QSOs. In a companion paper (Wisotzki \\cite{wisotzki:99:BQLF}) we study the impact of this new sample on the bright end of the QSO luminosity function and its evolution. Further work is in progress to follow up on several issues like host galaxy characteristics, radio properties, or the incidence of gravitational lensing events.     \\begin{figure}[tb] \\epsfxsize=\\hsize \\epsfclipon \\epsfbox[90 190 525 602]{h1831f1.ps} \\caption[]{Distribution of HES survey fields in the southern hemisphere, in area-conserving azimuthal projection. The South pole is in the centre, and meridians of constant right ascension are marked. The black areas correspond to the 207 fields forming the basis for the flux-limited sample, while greyshaded zones indicate further HES fields where survey work is not yet completed. Fields left white are in or too close to the Milky Way.} \\label{fig:area} \\end{figure} ",
        "conclusions": "The Hamburg/ESO survey is an ongoing project that provides substantial improvements over previous searches for optically bright QSOs. We have constructed a new flux-limited quasar sample useful for a wide variety of statistical studies, consisting of 415 objects distributed over 3700~deg$^2$ in the sky. Greatest care was taken to minimise selection biases that could lead to redshift-dependent incompleteness. Key properties of the HES in this respect are (i) very relaxed UV excess selection criteria, equivalent to $(U-B)<-0.18$, (ii) the small errors of determining the UV excess in individual objects, and (iii) the high spectral resolution of the ESO objective prism spectra which ensures low stellar contamination without sacrifices in completeness.  Comparing the sample with the results of other surveys, in terms of general number counts as well as by detailed comparison of common survey regions, we find no evidence that the HES might miss a sizeable fraction of the known QSO population, within its redshift limits. The surface density of detected bright quasars, for given magnitudes, is consistent with most recently published work, but the HES sample greatly improves the statistics at very bright magnitudes. We confirm that the Palomar-Green BQS is incomplete (although the degree of incompleteness has been overestimated in the past), but alert the reader that there is also evidence for significant overcompleteness.  Several precautions were taken to reduce potential selection biases associated with low-redshift QSOs. Unlike most other optical QSO surveys, there is no discrimination of objects with extended morphological structure; these objects are subjected to the same selection criteria as the point sources. In the extraction of spectra, nuclear properties are enhanced relative to possibly `red' galaxy contributions. This also affects the derived magnitudes which are dominated by the nuclear emission. Furthermore, a specific selection criterion has been implemented that is sensitive to `Seyfert-like' spectral shapes. All these factors enable the HES to produce the first well-defined optically selected samples of low-$z$ QSOs and bright Seyfert~1 nuclei. At the same time, selection biases against gravitationally lensed QSOs are avoided.  The HES substantially improves the sampling in the high-luminosity, high-redshift domain unaccessible to traditional UV excess surveys. While the number of $z>2.5$ sources in the present flux-limited sample is still small, there is a large reservoir of unambiguously detected high-$z$ quasars. By using dedicated flux limits optimised for these objects, combined with a detailed determination of the survey selection function, the total number of bright QSOs in the redshift range $2.5\\la z\\la 3.3$ that are part of a well-defined flux-limited sample will be increased by at least an order of magnitude in the near future."
    },
    "0004/astro-ph0004118_arXiv.txt": {
        "abstract": "We test accurate models of Comptonization spectra over the high quality data of the \\BS long look at NGC 5548, allowing for different geometries of the scattering region, different temperatures of the input soft photon field and different viewing angles.  We find that the \\BS data are well represented by a plane parallel or hemispherical corona viewed at an inclination angle of 30$^{\\circ}$. For both geometries the best fit temperature of the soft photons is close to 15$^{+3}_{-9}$ eV. The corresponding best fit values of the hot plasma temperature and optical depth are $kT_{\\rm e}\\simeq$ 250--260 keV and $\\tau\\simeq$ 0.16--0.37 for the slab and hemisphere respectively.  These values are substantially different from those derived fitting the data with a power-law + cut off approximation to the Comptonization component ($kT_{\\rm e}\\lta$ 60 keV, $\\tau\\simeq$ 2.4).  In particular the temperature of the hot electrons estimated from Comptonization models is much larger.  This is due to the fact that accurate Comptonization spectra in anisotropic geometries show \"intrinsic\" curvature which reduces the necessity of a high energy cut-off.  The Comptonization parameter derived for the slab model { is} larger than predicted for a two phase plane parallel corona in energy balance, suggesting that a more ``photon-starved'' geometry is necessary. The case of a hemispheric corona is consistent with energy balance but requires a large reflection component.  The spectral softening detected during a flare which occurred in the central part of the observation corresponds to a decrease of the Comptonization parameter, probably associated with an increase of the soft photon luminosity, the { hard} photon luminosity remaining constant. The increased cooling fits in naturally with the derived decrease of the coronal temperature $kT_{\\rm e}$ in the high state. ",
        "introduction": "The X-ray emission of Seyfert galaxies is commonly believed to be produced by Compton scattering of soft photons on a population of hot electrons { (Pozdniakov et al., \\cite{poz76}; Shapiro et al., \\cite{sha76}; Liang \\& Price, \\cite{lia77}; Sunyaev \\& Titarchuk, \\cite{sun80})}. Although high energy observations of Seyfert galaxies may be consistent with both thermal and non--thermal models (Zdziarski et al. \\cite{zdz94}; Malzac et al. \\cite{mal98}) the non--detection of Seyferts by Comptel and the high energy cut-offs indicated by OSSE (Jourdain et al. \\cite{jou92}; Maisack et al., \\cite{mai93}) and {\\it Beppo}SAX (Matt \\cite{mat99}) have focused attention on thermal models. In fact it was shown early on that assuming thermal equilibrium and a plane parallel geometry for the Comptonizing plasma above an accretion disk, and neglecting direct heating of the disk, the average properties of the X-ray emission of Seyfert galaxies could be naturally accounted for (Haardt \\& Maraschi \\cite{haa91}).  The spectral slope is mainly determined by two parameters, the temperature $kT_{\\rm e}$ and optical depth $\\tau$ of the scattering electrons, while the cut--off energy is related essentially to $kT_e$. Thus, simultaneous measurements of the slope of the X-ray continuum {\\it and} of the cut--off energy are necessary to determine the physical parameters of the Comptonizing region.  Moreover, in a disk plus corona system, the Comptonizing region and the source of soft photons are {\\it coupled}, as the optically thick disk necessarily reprocesses and reemits part of the Comptonized flux as soft photons which are the seeds for Comptonization.  The system must then satisfy equilibrium energy balance equations, which depend on {\\it geometry} and on the ratio of direct heating of the disk to that of the corona. In the limiting case of a \"passive\" disk, the amplification of the Comptonization process, determined by the Compton parameter $\\displaystyle y\\simeq 4\\left(\\frac{kT_e}{m_ec^2}\\right)\\, \\left[1+4\\left(\\frac{kT_e}{m_ec^2}\\right)\\right] \\tau (1+\\tau)$, is fixed by geometry only. Therefore, if the corona is in energy balance, the temperature and optical depth must satisfy a relation which can be computed for different geometries of the disk+corona configuration (e.g., the review of Svensson \\cite{sve96}). It is then theoretically possible to constrain the geometry of the system and verify the selfconsistency of the model, provided that $kT_e$ and $\\tau$ are known with sufficient precision.  The extent of spectral variability during luminosity variations is an additional, in principle powerful, diagnostic tool that can be used to test existing models, as it provides direct insight into the way the emitting particles are heated and cooled. For example, if the plasma is pair dominated, the value of $\\tau$ is fixed by the compactness parameter (i.e. by the luminosity for fixed geometry). This yields a definite relation between spectral variations and intensity, predicting modest spectral changes for large (factor 10 at least) variations in the intensity. On the contrary, for low pair density plasmas, significant spectral variations are possible even in the absence of luminosity variations.  { In the latter case, for constant geometry (i.e. at a constant Compton parameter $y$), the spectral index is expected to increase when the electron temperature decreases (Haardt, Maraschi \\& Ghisellini \\cite{haa97})}.  In the present paper, our aim is to test Comptonization models over high quality data, deriving { further} constraints on the physical parameters and geometry of the source.  To achieve such goal, the long look at the Seyfert I galaxy NGC 5548 performed by \\BS provides an ideal dataset. Previous studies of the source conducted over several years using EXOSAT, GINGA, ASCA and RXTE and, in some cases, coordinated observations first with IUE and later with EUVE (Walter \\& Courvoisier, \\cite{wal90}; Nandra et al., \\cite{nan91}; Chiang et al., \\cite{chi00}) support the general Comptonization picture, revealing the presence of several components in the X-ray spectra (neutral Iron line, reflection hump, soft excess...). No information on the high energy and of the Comptonization component could however be obtained from the above studies. Magdziarz et al. (\\cite{mag98}), using average OSSE data and non simultaneaous GINGA observations, suggested a temperature of 50~keV for the hot corona.  NGC 5548 was observed by \\BS in a single long (8 days) observation, with a net exposure of 314 ks on source.  The high quality of the \\BS data allows a detailed study of the spectrum over a very wide energy range, from 0.2 to 200 keV, and offers the opportunity to study spectral variability, since a conspicuous flare occurred in the middle of the observation. A detailed analysis of these data has been presented by Nicastro et al. (\\cite{nic99}, hereafter N99). The differences between the latter analysis and our results concern only the modelling of the continuum and will be discussed in the course of the paper.  Our main progress here is to adopt and fit directly to the data a detailed model of the Comptonized spectrum, for which the commonly adopted representation of a simple power law with a high energy cut--off turns out to be a rather poor approximation. Effectively, Comptonized spectra intrinsically show additional features, such as bumps due to different scattering orders, and an anisotropy \"break\" due to the (plausible) anisotropic nature of the soft photon input.  The paper is organized as follows. In section \\ref{comptonmodel} we briefly summarize the main characteristics of Comptonization models, and compare results of different approaches and geometries. The analysis of the \\BS data is presented in section \\ref{dataanal}.  We will not be concerned here in detail with the warm absorber features already discussed in N99. In section \\ref{iueosse}, we compare the \\BS data with non--simultaneous IUE and OSSE data. We discuss our results and their physical interpretations in section \\ref{discussion}. We then conclude in the last section. ",
        "conclusions": "The aim of this paper was to test realistic Comptonization models over the high signal to noise \\BS observations of NGC 5548 (which has been detailed by N99). Our main effort was to adopt detailed Comptonization codes that treat carefully the Compton processes in different geometries, especially for what concerns the possible anisotropy of the soft photon field. This last point is crucial, since, in geometries like slabs or hemispheres, the observed first order scattering humps is highly reduced in comparison to the others, producing an anisotropy break just above the averaged energy of twice scattered photons. Specifically, the \\BS observation of NGC 5548 has allowed us to show that: \\begin{itemize} \\item the data are well fitted by a plane parallel corona model with an inclination angle of $30^{\\circ}$ and a soft photon temperature between 5 and 15 eV. The corresponding best fit values of the hot plasma temperature and optical depth are $kT_{\\rm e}\\simeq$ 250 keV and $\\tau\\simeq$ 0.1 respectively. These values of $kT_{\\rm e}$ and $\\tau$ are however not in agreement with a radiatively balanced two phases disk/corona system in plane parallel geometry, even under the assumption that all the gravitationnal power is dissipated in the hot corona. Instead the data suggest that the hot Comptonizing gas is photon-starved. A better agreement is effectively obtained with a hemispherical geometry, which is naturally undercooled with respect to the slab model. However for the hemisphere the anisotropy break is larger, yielding a large reflection component (R=1.6$\\pm$0.3). This suggests that the real configuration is probably more complex than these two ideal cases.  \\item the flare occurring during the central 70 ks part of the run enabled us to interpret the spectral change of the source between different states in terms of changes in the physical parameters of the system.  Within the framework of anisotropic thermal Comptonisation the change of state suggests a variation of the Compton parameter $y$, and thus a change of the central configuration.  The spectral variability can be most naturally explained by an increase of the soft photon flux causing a reduction of the coronal temperature. % be the first  \\item the zero order cut--off power law approximation commonly used to Comptonization spectra leads to quite different results. The corresponding best fit values of the temperature and optical depth of the corona are $kT_e\\simeq$ 60 keV and $\\tau\\simeq$ 2.4. The bright state is consistent with the Compton parameter remaining constant during the spectral transition, but suggets an {\\it increase of the coronal temperature} and a {\\it decrease of the optical depth}. This model is statistically acceptable.  \\item We expect that the substantial improvements in the spectral data at low and medium X-ray energies achievable with CHANDRA and XMM-Newton could discriminate between these alternative models if complemented by simultaneous measurements at higher energies possibly provided by \\BS.  A decisive progress could come from improved soft $\\gamma$--ray observations like those expected in the near future by INTEGRAL,  \\end{itemize}  \\noindent {\\sl Acknowlegements:} We gratefully acknowledge J. Poutanen for providing us his code and the anonymous referee for his helpful comments. POP, FH and LM were supported by the European Commission under contract number ERBFMRX-CT98-0195 ( TMR network \"Accretion onto black holes, compact stars and protostars\") and by the Italian Ministry for University and Research (MURST) under grant COFIN98-02-154100. GM and GCP were supported by the Ministry for University and Research (MURST) under grant COFIN98-02-32."
    },
    "0004/astro-ph0004097_arXiv.txt": {
        "abstract": "Motivated by the recently improved knowledge on the kinematic and chemical properties of the Galactic metal-poor stars, we present the numerical simulation for the formation of the Galactic stellar halo to interpret the observational results. As a model for the Galaxy contraction, we adopt the currently standard theory of galaxy formation based on the hierarchical assembly of the cold dark matter fluctuations. We find, for the simulated stars with [Fe/H]$\\le-1.0$, that there is no strong correlation between metal abundances and orbital eccentricities, in good agreement with the observations. Moreover, the observed fraction of the low eccentricity stars is reproduced correctly for [Fe/H]$\\le-1.6$ and approximately for the intermediate abundance range of $-1.6<$[Fe/H]$\\le-1.0$. We show that this successful reproduction of the kinematics of the Galactic halo is a natural consequence of the hierarchical evolution of the subgalactic clumps seeded from the cold dark matter density fluctuations. ",
        "introduction": "Structure and dynamics of the metal-deficient halo component in the Galaxy provide valuable information on the early evolution of the Galaxy (e.g., Freeman 1987; Majewski 1993). Accordingly, the origin of the Galactic stellar halo has been extensively discussed by many authors since Eggen, Lynden-Bell, \\& Sandage (1962, hereafter referred to as ELS) reported a strong correlation between metal abundances and space motions of the high-velocity stars in the solar neighborhood. ELS argued that the contraction of the Galaxy must have been monolithic and rapid within a free-fall time ($\\sim$ $10^{8}$ yr). Several authors however pointed out that the collapse time scale estimated by ELS is greatly affected by their selection criterion against the halo stars having high angular momentum (e.g., Yoshii \\& Saio 1979; Norris, Bessell, \\& Pickles 1985; Chiba \\& Yoshii 1998, hereafter CY). Alternatively, Searle \\& Zinn (1978, hereafter SZ) proposed that the Galactic halo was formed slowly ($\\sim$ $10^{9}$ yr) by chaotic merging/accretion of several subgalactic fragments. It is yet unsettled whether either ELS's monolithic or SZ's merging picture (or both, e.g., Norris 1994; Freeman 1996; Carney et al. 1996) is more plausible and realistic for describing the early evolution the Galaxy.  Kinematic aspects of the Galactic metal-poor stars have been greatly improved by the recently completed $Hipparcos$ Catalogue (ESA 1997) and various ground-based catalogs (e.g., Platais et al. 1998; Urban et al. 1998) which provide unprecedentedly accurate proper motion data for a wealth of metal-poor stars (Beers et al. 2000). Using the non-kinematically selected sample of stars having available proper motions, CY and Chiba \\& Beers (2000, hereafter CB) revisited the relation between metal abundances and orbital eccentricities of the halo stars and found no evidence for a strong correlation between these quantities, in contrast to ELS's finding. CB also showed a clear evidence for internal structure of the halo: the outer halo shows no systematic rotation and nearly spherical density distribution, whereas the inner halo has a prograde rotation and a highly flattened density distribution. Also, Helmi et al. (1999) discovered a statistically significant clumpiness of the nearby halo stars in the angular momentum space, and argued that about 10 \\% of the halo come from a single small galaxy that was disrupted during or soon after the Galaxy formation. These new findings not only provide constraints on the formation of the Galactic stellar halo but also improve our understanding of how disk galaxies like our own form.  In this Letter, we address the question whether the above kinematic and chemical properties of the Galactic halo are understandable in the context of the currently favored theory of galaxy formation based on the hierarchical assembly of cold dark matter (CDM) halos (White \\& Rees 1978). Most of the CDM-based numerical models on disk galaxy formation have focused on only the fundamental properties of a disk, such as an exponential density profile (Katz 1992) and Tully-Fisher relation (Steinmetz \\& Navarro 1999). The spatial structure of the stellar halo has been examined by Steinmetz \\& M\\\"uller (1995), but the detailed internal kinematics of the halo stars in the simulated model remained unknown. Here we explore a numerical simulation for the formation of the Galactic halo, to investigate whether the CDM model can successfully explain the kinematic and chemical properties of the Galactic halo. We particularly focus on the evolution of hierarchically clustered subgalactic clumps seeded from the CDM density fluctuations and investigate their evolutionary effects on the orbital distribution of the stars in conjunction with the metal enrichment. More extensive analyses and results of the numerical simulations will be presented elsewhere (Bekki \\& Chiba 2000). ",
        "conclusions": "Although both the ELS monolithic and SZ merger scenarios have offered the basic ingredients for describing the early evolution of the Galaxy, either model alone does not comprehensively explain the currently improved knowledge on the fundamental properties of the halo (e.g., Freeman 1996; CB). For example, the lack of the abundance gradient in the halo stars (Carney et al. 1990; CY) and no significant correlation between [Fe/H] and $e$ (CY; CB) are difficult to interpret in the context of the ELS scenario. The SZ scenario seems unlikely to explain a large vertical gradient of the mean rotational velocity $<V_{\\phi}>$ in the halo component (CB). It is also unclear how the rapidly rotating disk component subsequently formed after the totally chaotic merging of ``Searle \\& Zinn'' fragments. In contrast, to explain the dual nature of the observed halo in its density, kinematics, and age (Norris 1994; Carney et al. 1996; CB), one requires the sort of hybrid picture, combining aspects of both the ELS and SZ scenarios.  As a possible candidate model to achieve the above hybrid picture, we have considered the CDM model, which invokes both the hierarchical assembly of subgalactic clumps and the dissipative process of gas inside the clumpy protogalactic system. As a first step toward understanding the formation of the Galactic halo in this context, we have investigated the orbital properties of the stars in the simulated Galactic halo, and have shown that the hierarchical merging of CDM clumps in the course of the expansion and contraction of the overall protogalactic sphere plays a vital role in determining the observed [Fe/H]-$e$ relation of the metal-poor stars. It is also found that the dissipative merging of the clumps is important for the development of the characteristic structure of the halo and also for the subsequently formed disk component (Bekki \\& Chiba 2000).  While we have reproduced the most basic relation between metal abundances and orbital eccentricities of the halo stars based on the currently favored picture of galaxy formation, there are still a couple of points to be clarified for the comprehensive understanding of the halo formation. For example, Sommer-Larsen et al. (1997) reported that the velocity ellipsoid of the metal-poor stars changes from radial anisotropy near the Sun to tangential anisotropy in the outer part of the Galactic halo. This may be explained via the anisotropic, {\\it dissipative} merging between protogalactic gas clouds in a collapsing galaxy (Theis 1997), but it is yet unsettled as to whether the similar process is equally applied in the case of the CDM clumps consisting of both dissipationless particles (dark matter and stars) and dissipative gas. Also, the non-kinematically selected sample of the nearby stars shows a remarkable discontinuity of the mean rotational velocity, $V_{rot}$, at [Fe/H]$\\sim-1.7$: the stars at [Fe/H]$<-1.7$ shows an approximately constant rotation, whereas those at [Fe/H]$>-1.7$ show the linear increase of $V_{rot}$ with increasing [Fe/H] (e.g., CB). We will further discuss in the forthcoming paper  whether these other kinematic properties of the Galactic halo are also explained by the dynamical evolution of the system of subgalactic clumps seeded from the CDM density fluctuations."
    },
    "0004/astro-ph0004233_arXiv.txt": {
        "abstract": "\\noindent Measurements of electron, muon, and hadron lateral distributions of extensive air showers as recorded by the KASCADE experiment are presented.  The data cover the energy range from $5\\times10^{14}$ eV up to almost $10^{17}$ eV and extend from the inner core region to distances of 200~m.  The electron and muon distributions are corrected for mutual contaminations by taking into account the detector properties in the experiment.  All distributions are well described by NKG-functions.  The scale radii describing the electron and hadron data best are $\\simeq 30$~m and $\\simeq 10$~m, respectively.  We discuss the correlation between scale radii and `age' parameter as well as their dependence on shower size, zenith angle, and particle energy threshold. ",
        "introduction": "The KASCADE ({\\em KArlsruhe Shower Core and Array DEtector\\/}) experiment is located at Forschungszentrum Karlsruhe, Germany, at an altitude of 110~m a.s.l.\\ and has been described in detail in \\cite{Doll-1990b,AAA067.022.145}.  The experiment has three major components: an array of electron and muon detectors, a central detector mainly for hadron measurements but with substantial muon detection areas, and a tunnel with streamer tube muon telescopes. Since the latter have only been completed at the time of writing this article, no data from the muon tunnel are included in the present analysis.  The KASCADE array covers an area of about 200$\\times$200~m$^2$ and consists of 252 detector stations located on a square grid of 13~m separation.  These are organized in 16 clusters of 16 stations each, except for the inner four clusters where the location of one station is blocked by the central detector.  The stations contain two types of detectors, liquid scintillation counters ($e/\\gamma$ detectors) of 0.79~m$^2$ area each and 5~cm thickness with little shielding above and plastic scintillators of 0.81~m$^2$ area each and 3~cm thickness (muon detectors) below a shielding of 10~cm lead and 4~cm steel.  The inner four clusters are instrumented with four $e/\\gamma$ detectors per station but without muon detectors while the outer 12 clusters house two $e/\\gamma$ and four muon detectors per station.  For both detector types the sum of photomultiplier signals, the earliest time, and the hit pattern are recorded for all stations fired.  The central hadron calorimeter is of the sampling type and has a fiducial area of 16$\\times$19~m$^2$.  A detailed description can be found in \\cite{Engler-1999}.  The energy is absorbed in an iron stack and sampled in eight layers of liquid ionization chambers with anode segments of 0.25$\\times$0.25~m$^2$ (appr.\\ 38\\,500 channels).  The thickness of the iron slabs increases from 12 to 36~cm towards the deeper parts of the calorimeter, amounting to 154~cm in total.  The 8$^{\\rm th}$ layer is located below an additional concrete ceiling of 77~cm thickness.  On top, a 5~cm lead layer filters off the soft electromagnetic component.  The ionization chambers are read out by logarithmic amplifiers and 13 bit ADCs, achieving a dynamic range of $6\\mal10^4$.  Signals starting from single minimum ionizing muons up to energy deposits of 10~GeV in a chamber are read out without saturation.  The response curve of each channel is calibrated with a reference capacitor coupled to the preamplifier, injecting known charges into the electronics chain.  Below the 8$^{\\rm th}$ calorimeter layer, two layers of multiwire proportional chambers (MWPCs), vertically separated by 38 cm, are used as muon detectors.  In total, 32 chambers are operated with 129~m$^2$ total area per layer.  Hits are registered on anode wires and two layers of cathode strips at angles of $\\pm34^\\circ$ with respect to the wires.  A total of 456 plastic scintillation counters of $0.48 \\times 0.95$~m$^2$ area each and 3~cm thickness are used within the calorimeter (below 5~cm lead and 36~cm steel) to trigger the calorimeter and the MWPCs.  They also serve as muon counters.  On top of the calorimeter, an additional 50 such counters fill the central gap of $e/\\gamma$ detectors but their data are not included in the present analysis.  A summary of the detector components used in this article together with their most relevant parameters is given in Table~\\ref{tab:detectors}.  \\begin{table}[t] \\caption{KASCADE detector components used in this analysis. Detection thresholds refer to particle energies above the absorber material of the detectors.} \\label{tab:detectors} \\def\\secz{~$\\times\\sec\\theta$} \\def\\pscz{\\phantom{\\secz}} \\vspace*{3mm} \\begin{tabular}{lcccrc} \\hline Detector & channels & separation & total area & threshold $E_{\\rm kin}$~~ & for \\\\ \\hline array $e/\\gamma$ & 252     & 13~m       & 490~m$^2$  &    5 MeV\\pscz  & $e$ \\\\ array $\\mu$      & 192     & 13~m       & 622~m$^2$  &  230 MeV\\secz  & $\\mu$ \\\\ trigger          & 456     &   ---      & 208~m$^2$  &  490 MeV\\secz  & $\\mu$ \\\\ MWPCs            & 26\\,080 &   ---      & 129~m$^2$  &  2.4 GeV\\secz  & $\\mu$ \\\\ calorimeter      & 38\\,368 &   ---      & 304~m$^2$  &   50 GeV\\pscz  & hadrons \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "Measurements of electron, muon, and hadron lateral distributions as recorded by the KASCADE experiment have been presented for radial distances of up to 200~m and for the energy range $5\\times10^{14} {\\rm eV} < E < 10^{17} {\\rm eV}$.  Detector simulations were performed to account for effects like muon and hadron contamination in signals of the $e/\\gamma$-scintillators and for punch-through of electrons and hadrons into the muon detectors.  All types of lateral distributions are well described by NKG-functions using different scale radii $r_{i}$ (with $i=e,\\mu,h$) for the different air shower components. \\begin{itemize}  \\item A study of the electron LDFs shows that optimum fits are not obtained for the canonical value of $\\Rmol = 79$~m, but for $r_{e} \\simeq 20-30$~m, i.e.\\ we observe a stronger curvature in the experimental data than in the conventional NKG function.  This imposes a systematic effect of up to 5\\,\\% in the integrated number of shower electrons.  Due to the strong correlation of $r_{e}$ and $s$, the preferred lower scale radius is accompanied by a larger age parameter of $s \\simeq 1.65$.  The optimum set of parameters depends on shower size and zenith angle and may be used to infer the mass of the primary particle. For practical reasons and because of limited statistics within single events, information about the shape of the electron LDF is usually extracted by fixing $r_{e}$ and fitting only the age parameter $s$.  A problem specific to NKG functions in this approach and with small scale radii $r_{e}$ is, that upwards fluctuations of $s$ easily lead to ill defined shower sizes.  As an alternative, we have demonstrated that a fixed age parameter but variable scale radius provides an equally good fit to the data.  The parameter $r_{e}$ then changes in a characteristic way and also exhibits a distinct structure at shower sizes corresponding to the knee position.  \\item Within the fiducial area of KASCADE, muon LDFs are well described by a NKG function, but with a scale radius of $r_{\\mu} = 420$~m.  Because of the limitation to 40~m $< r <$ 200~m, the experiment is not very sensitive to the actual value chosen and the data are also equally well described by a Greisen parametrization.  Significant differences would only occur at radial distances outside the acceptance of the experiment.  The unknown flat shape of the muon LDF at large distances imposes serious problems (even for much larger surface detector arrays) when calculating the total number of muons within an air shower. Most importantly, $N_{\\mu}$ is subject to systematic shifts and increased fluctuations, thereby deteriorating the shower size and primary energy resolution.  Thus, for classifying events, we have introduced the truncated muon number, $N_{\\mu}^{\\rm tr}$, obtained from integrating the LDF only within the experimental acceptance of 40-200~m.  A rough scan of the low energy muon spectrum has been performed by analysing LDFs at $E_{\\mu} \\ge 230$, 490, and 2400~MeV. Similarly to electrons, a steepening of the muon LDF is observed with increasing shower size and decreasing zenith angle, as is expected for observations being increasingly closer to the shower maximum.  \\item Quite interestingly, also hadronic lateral energy density and particle number distributions are well approximated by the NKG form up to distances of at least 90~m.  The scale radius for $E_{\\rm th} \\ge 50$~GeV is $r_{h} \\simeq 10$~m and scales roughly proportional to $E_{\\rm th}^{-1}$, as expected by a simple dimensional comparison of electromagnetic multiple scattering and hadronic interactions.  \\item The interrelation between the electromagnetic and hadronic EAS component may explain the `unconventional' small preferred scale radius of the electron LDF of 20-30~m as compared to the classical value of $\\Rmol \\simeq 80$~m.  It should be kept in mind that the classical Moli\\`ere radius has been derived for pure electromagnetic showers and for zero energy threshold only. However, extensive air showers are mostly initiated by primary hadrons.  Therefore, the shower evolution is mostly driven by the substantially narrower hadronic component, and the effective lateral scale radius of observed electrons is expected to be smaller than for $E_{\\rm kin} \\ge 0$ electrons in pure $\\gamma$-initiated showers. \\end{itemize}  The present paper is not focussed to detailed analyses in terms of predictions of the EAS developments from Monte Carlo simulations and to a comparison of different theoretical high-energy interaction approaches like VENUS \\cite{venus}, QGSJET \\cite{qgsjet} and SIBYLL \\cite{sibyll}.  These models, continuously in the process of refinement, are generators implemented into the Karlsruhe EAS Monte Carlo code CORSIKA \\cite{CORSIKA}.  However, the presented results provide a coherent experimental basis for serious tests considering simultaneously the three main EAS components, not only concerning the interaction but also the particle propagation procedures.  It may be noted that the muon lateral distributions are experimentally given for three different energy detection thresholds of the registered muons, thus implying also some sensitivity to the low energy spectrum.  Most valuable for such tests are observations based on the hadronic component.  An example of first analyses in this scope were presented in \\cite{kascade-ww} and a remarkable agreement of lateral distributions of hadrons for primary protons and Fe nuclei was observed.  In particular, the absence of peculiar features, in contrast to earlier observations by Danilova et al.\\ \\cite{Danilova-1985} and Arvela and Elo \\cite{Arvela-1995} can be stated, even at energies as high as 10 PeV. Such results support the trust in a correct handling of the particle propagation and of the development of the hadronic component at least for hadron energies above 50 GeV. More detailed comparisons of lateral distributions with CORSIKA simulations are under study and will be subject of a forthcoming publication."
    },
    "0004/astro-ph0004005_arXiv.txt": {
        "abstract": "Low-radio-frequency observations played a remarkable role in the early days of radio astronomy; however, in the subsequent three or four decades their usefulness has largely been in terms of the {\\sl finding-frequency} of surveys.  Recent technical innovation at the VLA has meant that {\\sl spatially well-resolved} imaging at low frequencies is now possible.  Such imaging is essential to understanding the relationship between the hotspot and lobe emission in classical double radio sources, for example.  We here present new images of 3C radio sources at 74\\,MHz and 330\\,MHz and discuss their implications. ",
        "introduction": "Low-frequency radio surveys play a key role in selecting samples of radio sources, dominated by optically thin synchrotron emission, which are free of orientation biases.  Examples of such samples, which have been pivotal in advancing our understanding of the nature of radio sources, are the celebrated 3C sample (revised by Laing, Riley and Longair 1983) and very recently the much fainter 7C sample (Rawlings et al.\\ 1998).  Following recent technical innovation at the VLA, it is now possible to make spatially well-resolved images of radio sources at low frequencies: at 74\\,MHz, images can now routinely be made with an angular resolution of 25$^{\\prime\\prime}$.  It is in this low-frequency regime that models of the energy supply to the lobes from the hotspots in classical doubles may be tested, and where their energy budgets may be investigated, for example by the detection or otherwise of steep-spectrum halos surrounding these objects. ",
        "conclusions": ""
    },
    "0004/astro-ph0004413_arXiv.txt": {
        "abstract": "The first sources of ionizing radiation to condense out of the dark and neutral IGM sent ionization fronts sweeping outward through their surroundings, overtaking other condensed objects and photoevaporating them. This feedback of universal reionization on cosmic structure formation is demonstrated here by gas dynamical simulations, including radiative transfer, for the case of a cosmological minihalo of dark matter and baryons exposed to an external source of ionizing starlight, just after the passage of the global ionization front created by the source. ",
        "introduction": "\\label{sec:constraints}  The neutral, opaque IGM out of which the first bound objects condensed was dramatically reheated and reionized at some time between a redshift $z\\approx50$ and $z\\approx5$ by the radiation released by some of these objects. When the first sources turned on, they ionized their surroundings by propagating weak, R-type ionization fronts which moved outward supersonically with respect to both the neutral gas ahead of and the ionized gas behind the front, racing ahead of the hydrodynamical response of the IGM, as first described by Shapiro (1986) and Shapiro \\& Giroux (1987). These authors solved the problem of the time-varying radius of a spherical I-front which surrounds isolated sources in a cosmologically-expanding IGM analytically, taking proper account of the I-front jump condition generalized to cosmological conditions. They applied these solutions to determine when the I-fronts surrounding isolated sources would grow to overlap and, thereby, complete the reionization of the universe (Donahue \\& Shull 1987 and Meiksen \\& Madau 1993 subsequently adopted a similar approach to answer that question). The effect of density inhomogeneity on the rate of I-front propagation was described by a mean ``clumping factor'' $c_l>1$, which slowed the I-fronts by increasing the average recombination rate per H atom inside clumps. This suffices to describe the rate of I-front propagation as long as the clumps are either not self-shielding or, if so, only absorb a fraction of the ionizing photons emitted by the central source. Numerical radiative transfer methods are currently under development to solve this problem in 3D for the inhomogeneous density distribution which arises as cosmic structure forms, so far limited to a fixed density field without gas dynamics (e.g.\\ Abel, Norman, \\& Madau 1999; Razoumov \\& Scott 1999). The question of what dynamical effect the I-front had on the density inhomogeneity it encountered, however, requires further analysis. ",
        "conclusions": ""
    },
    "0004/astro-ph0004139_arXiv.txt": {
        "abstract": "\\noindent Mapping the central 5 pc region of the nearby radio galaxy NGC 6251 with a 0.2 pc resolution using VLBI at two radio frequencies, 5 GHz and 15 GHz, we have found the sub-pc-scale counter jet for the first time in this radio galaxy. This discovery allows us to investigate the jet acceleration based on the relativistic beaming model (Ghisellini et al. 1993). ",
        "introduction": "The genesis of powerful radio jets from active galactic nuclei is one of the long standing problems in astrophysics (Bridle \\& Perley 1984). Although global morphological properties give us very important information, very inner regions in the radio jets also provide hints to understanding the genesis of radio jets. In order to investigate radio galaxies at the sub-milli arcsecond angular resolution, we have performed new high-resolution VLBI observations of NGC 6251 using HALCA (Hirabayashi et al. 1998). We use a distance to NGC 6251, 94.4 Mpc (for a Hubble constant  $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$). Note that 1 mas (milli arcsecond) corresponds to 0.48 pc at this distance. ",
        "conclusions": "We assume for simplicity that the core is surrounded by a plasma sphere with a radius of $a$ and the radio emission from the inner part of the jets (i.e., both the jet and the counter jet) suffer from the free-free absorption. The approaching jet escapes from the plasma sphere at a projected distance of $x = a \\times \\sin \\theta_{\\rm jet} $ and thus  the spectral index becomes to be intrinsic here. The counter jet suffers the effect of free-free absorption until $x = -a $. It is also noted that the path length is longest at $x = - a \\times \\sin \\theta_{\\rm jet} $. As shown in Figure 2a, we define the  following projected distances; $X_{\\rm jet} = X_{\\rm peak} =  a \\times \\sin \\theta_{\\rm jet}$ and $X_{\\rm cjet} = a$. Then we are able to estimate $ \\theta_{\\rm jet} = \\sin^{-1} \\left({\\frac {X_{\\rm jet}}{X_{\\rm cjet}}} \\right) $. We adopt the 15 GHz brightest peak as the core, because the optical depth at 15 GHz is enough to be small in the case of free-free absorption. In Figure 2b, we show the observed spectral index variation along the radio jet. We estimate $X_{\\rm peak} \\approx$ 0.24 mas, $X_{\\rm jet} \\approx$ 0.41 mas, and $X_{\\rm cjet} \\approx$ 0.80 mas. Then we obtain $\\theta_{\\rm jet} \\simeq 31^\\circ$.  \\begin{figure} \\begin{center} \\epsfile{file=sudou_fig2.eps,scale=0.52} \\caption{A schematic illustration of the spherical-absorption model (a), and spectral index variation along the radio jet of NGC 6251 (b).} \\label{Fig2} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\epsfile{file=sudou_fig3.eps,scale=0.46} \\caption{The jet/counter jet flux density ratio $R$ as a function of the projected distance from the brightest peak at 15 GHz for the 5GHz data ({\\it filled triangle}) and for the 15 GHz data ({\\it filled circles}) (a), and jet velocity \\bj ~ as a function of projected distance from the brightest peak at 15 GHz (b).} \\label{Fig2} \\end{center} \\end{figure}  Adopting the so-called Doppler beaming model, the jet to the counter jet intensity ratio $R$ can be written as a function of projected distance $x$, $R(x) = \\left[\\frac {1 + \\beta_{\\rm{jet}}(x) \\cos \\theta_{\\rm{jet}}(x)} {1 - \\beta_{\\rm{jet}}(x) \\cos \\theta_{\\rm{jet}}(x)}\\right]^{2 - \\alpha(x)}$ , where \\bj($x$) = $v_{\\rm jet}(x) / c$. In Figure 3a, we measured $R$ as a function of $x$. It is shown that $R$ is estimated to be systematically larger at 5 GHz than those at 15 GHz. This is probably due to stronger absorption at 5 GHz than at 15 GHz. It is shown that $R$ at 15 GHz increases from 1.7 at 0.85 mas  (0.4 pc) to 10 at 1.25 mas (0.6 pc) with projected distance from the core. If this increase is just caused by the Doppler beaming, it is suggested that \\bj ~ increase with  distance, because it cannot be seen that \\tj ~ varies significantly. Figure 3b shows that the jet is accelerated from \\bj ~ $\\approx 0.1$ at 0.4 pc to \\bj ~ $\\approx 0.5$ at 0.6 pc. This provides the first direct evidence for the acceleration at the sub-pc-scale radio jet."
    },
    "0004/astro-ph0004280_arXiv.txt": {
        "abstract": "The atmospheric \\C imaging technique has been used to search for point-like and diffuse TeV gamma-ray emission from the southern supernova remnant, W28, and surrounding region. The search, made with the CANGAROO 3.8\\,m telescope, encompasses a number of interesting features, the supernova remnant itself, the EGRET source 3EG J1800$-$2338, the pulsar PSR~J1801$-$23, strong 1720 MHz OH masers and molecular clouds on the north and east boundaries of the remnant. An analysis tailored to extended and off-axis point sources was used, and no evidence for TeV gamma-ray emission from any of the features described above was found in data taken over the 1994 and 1995 seasons. Our upper limit ($E>1.5$ TeV) for a diffuse source of radius 0.25$^\\circ$ encompassing both molecular clouds was calculated at 6.64$\\times10^{-12}$ cm$^{-2}$s$^{-1}$ (from 1994 data), and interpreted within the framework of a model predicting TeV gamma-rays from shocked-accelerated hadrons. Our upper limit suggests the need for some cutoff in the parent spectrum of accelerated hadrons and/or slightly steeper parent spectra than that used here ($-$2.1). As to the nature of 3EG J1800$-$2338,  it possibly does not result entirely from $\\pi^\\circ$ decay, a conclusion also consistent with its location in relation to W28. ",
        "introduction": "\\label{sec:intro}  Supernova remnants (SNRs) have long been thought to be the dominant source of cosmic rays (CR) at energies below 100 TeV (for a review see e.g. Blandford \\& Eichler \\cite{Blandford:1}).  SNR, via the diffusive shock process, are able to accelerate electrons and hadrons and meet the energetics of the observed cosmic rays. The TeV gamma-ray flux predicted from SNR is the most accessible tracer of CR acceleration and its detection would be convincing evidence for the SNR origin of galactic CR. Models of the TeV gamma-ray emission from SNR predict distinct spectral features, according to the hadronic and/or electronic nature of the parent CR accounting for the gamma-ray flux (see Drury \\etal \\cite{Drury:1}, Naito and Takahara \\cite{Naito:1}, Baring \\etal \\cite{Baring:1}, and references therein for a summary).  Ground-based surveys of SNR at gamma-ray energies (TeV to PeV) have been carried out on several promising northern hemisphere candidates (e.g. IC443, Tycho's SNR, W51, W44, G78.2$+$2.1). The Whipple (Buckley \\etal \\cite{Buckley:1}, Lessard \\etal \\cite{Lessard:1}), HEGRA (Hess \\etal \\cite{Hess:1} at TeV energies, and Prosch \\etal \\cite{Prosch:1} at multi-TeV energies), CAT (Goret \\etal \\cite{Goret:1}) and CYGNUS (Allen \\etal \\cite{Allen:1}) groups have reported upper limits. Recently however, the HEGRA has seen marginal evidence for TeV gamma-rays from the young SNR Cas-A, after deep observation (P\\\"{u}hlhofer \\etal \\cite{Puhlhofer:1}). In the southern hemisphere, the CANGAROO has reported the detection of TeV gamma-rays from SNR SN1006 (Tanimori \\etal \\cite{Tanimori:1}) and SNR RX J1713.7$-$3946 (Muraishi \\etal \\cite{Muraishi:1}), and if confirmed, will be strong evidence in favour of the production of cosmic rays electrons in SNR.  W28 (also SNR G6.4$-$0.1 from Green \\cite{Green:1}) is considered an archetypal composite (mixed or M-type) supernova remnant, characterised by a centrally filled X-ray and shell-like radio morphology (Rho \\& Petre \\cite{Rho:1}, Long \\etal \\cite{Long:1}).  The ROSAT X-ray emission appears best explained by a thermal model (Rho \\etal \\cite{Rho:2}) although Tomida \\etal (\\cite{Tomida:1}) from the analysis of ASCA data, has suggested the presence of a weak a non-thermal component in the south west region. The limb-brightened radio emission (20, 6 \\& 2cm) shows a synchrotron spectrum of varying spectral index (Andrews \\etal \\cite{Andrews:1}). A radio point source at $l=6.6^\\circ, b=-0.16^\\circ$ (G6.6$-$0.1) is defined (Altenhoff \\etal \\cite {Altenhoff:1}, Andrews \\etal \\cite{Andrews:1}), hereafter referred as A83 in this paper. A glitching radio pulsar, PSR~J1801$-$23 (PSR~B1758$-23$, $P=$416ms, $\\dot{P}=113\\times 10^{-15}$ ss$^{-1}$), lies at the northern radio edge (Kaspi \\etal \\cite{Kaspi:1}). An upper limit to this pulsar's characteristic age is estimated at 58\\,000 years, and it's spin-down luminosity ($\\dot{E} \\sim 6.2\\times 10^{34}$ erg s$^{-1}$) is at the lower edge of luminosity values when compared to the known gamma-ray (EGRET \\& COMPTEL) pulsars.  The age of W28 is estimated (Kaspi \\etal \\cite{Kaspi:1}) in the range 35\\,000 to 150\\,000 years, with upper and lower limits taken from the assumptions that W28 is currently in either the radiative or Sedov phases of expansion. According to Kaspi \\etal \\cite{Kaspi:1}, the distance of PSR~J1801$-$23 (9 to 16.5 kpc) derived from it's dispersion measure (DM) appears inconsistent with that derived for the remnant. Estimates for the remnant distance are set at 1.8 kpc (Goudis \\cite{Goudis:1} $\\Sigma$-D relation) and 3.3 kpc respectively (Lozinskaya \\cite{Lozinskaya:1}, from mean optical velocities), indicating that the pulsar/W28 association is possibly a line-of-sight coincidence. However, Frail \\etal (\\cite{Frail:1}) have noted the large uncertainty in using the DM as a distance estimate for this pulsar due to the high concentration of ionised material in the line of sight, and conclude there is sufficient evidence for the pulsar/remnant association. The unidentified EGRET source 3EG J1800$-$2338 (95\\% error circle 0.32$^\\circ$ radius) (Hartman \\etal \\cite{Hartman:1}), listed as 2EG J1801$-$2312 in the second EGRET catalogue (Thompson \\etal \\cite{Thompson:1}), lies on the edge of the radio shell and was thought to be associated with the remnant (Esposito \\etal \\cite{Esposito:1}, Zhang \\& Cheng \\cite{Zhang:1}). 3EG J1800$-$2338 has a relatively hard spectral index (Hartman \\etal \\cite{Hartman:1}) with no apparent sign of a turnover at 1\\,GeV (Merck \\etal \\cite{Merck:1}). Lamb \\& Macomb (\\cite{Lamb:1}) also point out that 3EG J1800$-$2338 is visible above 1 GeV at 5.4$\\sigma$ significance, and is centred very close to the A83 radio position. The 3EG position of the EGRET source is displaced by about 0.5$^\\circ$ relative to the 2EG position, yet still lies comfortably within the SNR radio shell, and remains a strong example of an EGRET source/SNR association (Romero \\etal \\cite{Romero:1}). The 3EG error circle however, now excludes PSR~J1801$-$23 and the molecular clouds.  W28 lies in a complex region of the galactic plane with many HII regions and dense molecular clouds (Wootten \\cite{Wootten:1}) contributing to the ISM surrounding the SNR. Over forty OH (1720\\,MHz) maser emission sites are concentrated at the eastern and northern edges of the SNR (Claussen \\etal \\cite{Claussen:1}), along the SNR and molecular cloud interface. The distribution of shocked and unshocked gas in this region is also consistent with the idea of the SNR shock passing through the cloud (Arikawa \\etal \\cite{Arikawa:1}). OH maser emission (1720 MHz) is considered a strong indicator of collisional pumping with matter densities $\\sim10^5$ cm$^{-3}$ (Claussen \\etal \\cite{Claussen:1,Claussen:2}). Enhanced levels of TeV $\\gamma$-ray emission via the decay of neutral pions may be expected from such areas associated with the masers and molecular cloud (Aharonian \\etal \\cite{Aharonian:1}). Fig. ~\\ref{fig:w28_claussen} indicates the sites of interest in relation to the radio continuum emission (327 MHz). The presence of these interesting objects make W28 a prime southern hemisphere candidate for study at TeV gamma-ray energies.  \\begin{figure} \\vspace{8cm} \\special{psfile=w28_claussen.ps hoffset=-120 voffset=-240 hscale=80 vscale=80} \\caption{Radio continuum (327 MHz) for W28, adapted from Claussen \\etal (\\protect\\cite{Claussen:1}). Included are the positions (and error circles) of the EGRET source 3EG\\,J1800$-$2238 (2EG\\,J1801$-$2312 in the 2nd EGRET catalogue), OH maser sites A to F, PSR\\,J1801$-$23 (also the tracking position for 1994 data) and the radio point source, labelled A83 (Andrews \\etal \\cite{Andrews:1}), is the tracking position for 1995 data. See Sect.~\\ref{sec:res} for a discussion.} \\label{fig:w28_claussen} \\end{figure}  We report here on the comprehensive analysis of data taken in 1994 and 1995 with the CANGAROO 3.8 metre telescope.  This work follows analysis of data taken in 1992 (Kifune \\etal \\cite{Kifune:1}) in which weak evidence for a gamma-ray signal was reported. At that time, only ON source data were collected, making an estimation of the background rate difficult.  Mori \\etal (\\cite{Mori:1}) reported briefly on an analysis of 1994 data centred on PSR~J1801$-23$, in which a $\\pm0.7^\\circ$ field of view was searched. Both ON and OFF source data were collected and no evidence for TeV $\\gamma$-ray emission was seen from various point-like sources including the pulsar and both radio and X-ray maxima. The 1995 data were centred on the radio position A83, located $\\sim0.3^\\circ$ away from PSR~J1801$-23$. A search for point-like and diffuse sources of TeV emission was carried out on the 1994 and 1995 datasets out to $\\pm 1^\\circ$ from the tracking centre of each dataset, using an extended source analysis. We have used an improved set of cuts to those used in the analysis of data taken on the Vela Pulsar/Nebula (Yoshikoshi \\etal \\cite{Yoshikoshi:1}). These cuts were designed to minimise the loss of gamma-ray sensitivity for off-axis sources and in particular maintain reliable statistics over the search region. ",
        "conclusions": "A search for TeV gamma-ray emission from the southern SNR W28 was carried out by the CANGAROO over two observation seasons (1994 and 1995) using the atmospheric \\C imaging technique. An analysis providing a consistent gamma-ray acceptance and quality factor for extended sources was used. A number of sites within a search region of $\\pm1^\\circ$ were considered as potential point-like and diffuse sources of TeV gamma-ray emission. No evidence was found for the emission of TeV gamma-rays at any of these sites which include those from the strongest two masers, an EGRET source, a radio pulsar (all as point sources) and a diffuse region containing the molecular clouds that appear to be interacting with the remnant. Our 3$\\sigma$ upper limit from this diffuse region at 6.64$\\times10^{-12}$ cm$^{-2}$s$^{-1}$ $> 1.5$ TeV, and the flux of the EGRET source 3EG J1800$-$2338 were compared with gamma-ray flux predictions from the model of Naito and Takahara (\\cite{Naito:1}). Under this framework, our upper limit rules out a straight extrapolation of the EGRET flux to TeV energies. It also constrains somewhat the flux expected from the shocked region of gas in the molecular cloud, placing limits on the parent spectra for hadrons and/or cutoff energy. Our results suggest that the EGRET source probably does not result entirely of $\\pi^\\circ$ gamma-rays. This fact is supported by it's location in relation to the molecular clouds. In a later paper, we will consider electronic bremsstrahlung and inverse Compton scattering and discuss the broader implications of our results in relation to the origin of galactic cosmic rays.  Further data on W28 will no doubt be taken with the recently completed CANGAROO II telescope (Yoshikoshi \\etal \\cite{Yoshikoshi:3}), which, at the very least will provide tighter constraints on models of gamma-ray production for this interesting source."
    },
    "0004/astro-ph0004249_arXiv.txt": {
        "abstract": "N-body codes for performing simulations of the origin and evolution of the large scale structure of the universe have improved significantly over the past decade in terms of both  the resolution achieved and the reduction of the CPU time. However, state-of-the-art  N-body codes hardly allow one to deal with particle numbers larger than a few  $10^7$, even on the largest parallel systems. In order to allow simulations with larger resolution, we have first reconsidered the grouping strategy as described in J. Barnes (1990, {\\it J. Comput. Phys.} {\\bf 87}, 161) (hereafter B90) and applied it with some modifications to our WDSH-PT (Work and Data SHaring - Parallel Tree) code (U. Becciani et al., 1996, {\\it Comput. Phys. Comm.} {\\bf 99},1). In the first part of this paper we will give a short description of the code adopting the algorithm of J. E. Barnes and P. Hut (1986, {\\it Nature}, {\\bf 324}, 446)  and in particular  the memory and work distribution strategy applied to describe  the {\\it data distribution} on a CC-NUMA machine like the CRAY-T3E system. In very large simulations (typically $ N \\geq 10^7$), due to network contention and the formation of clusters of galaxies, an uneven load  easily verifies. To remedy this, we have devised an automatic work redistribution mechanism which provided a good dynamic load balance without adding significant overhead. In the second part of the paper we describe the modification to the Barnes  grouping strategy we have devised to improve the performance of the WDSH-PT code. We will use the property that nearby particles  have similar interaction lists. This idea has been checked in B90, where an interaction list is built which applies everywhere within a cell $C_{group}$  containing a small number of particles $ N_{crit}$. B90 reuses this interaction list for each particle $ p \\in C_{group}$ in the cell in turn. We will assume each particle $ p$ to have the same interaction list. We consider that the agent force ${\\bf F}_p$ on a particle $ p$ can be decomposed into two terms ${\\bf F}_p = {\\bf F}_{far} + {\\bf F}_{near}$. The first term ${\\bf F}_{far}$ is the same for each particle in the cell and  is generated by the interaction between a hypothetical particle placed in the center of mass of the $C_{group}$ and the farther cells contained in the interaction list. ${\\bf F}_{near}$ is different for each particle $p$ and  is generated by the interaction between $ p$ and the elements near $C_{group}$. Thus it has been possible to reduce the CPU time and increase the code performance. This enable us to run simulations with a large number of particles ($ N \\sim 10^7 \\div 10^9$) in nonprohibitive CPU times. ",
        "introduction": "N-body codes are one of the most important tools of theoretical cosmology \\cite{ber91} because they offer the possibility of simulating most of the gravitational processes driving the formation of the  large scale structure of the universe (hereafter LSS) \\cite{bar86}\\cite{her87}\\cite{dub88}. These simulations are often used to check cosmological models and to constrain the free parameters of these models which cannot be fixed either theoretically or observationally.\\\\ The typical mass scale for gravitational instability, the Jeans mass, has a value of $\\approx 10^{6.5}$ solar masses (1 solar mass $\\approx 1.9\\times 10^{33} g$) at the recombination epoch, and it gives the approximate size of the first objects forming by gravitational collapse at that epoch. On the other hand, the largest structure we see in our Universe today, the ``Supercluster'' of galaxies, has a mass in excess of $\\approx 10^{18}$ solar masses. Moreover, the gravitational force has a truly long-range character, which makes it impossible to introduce reasonable upper cutoffs in the mass range. For all these reasons,  one would like to be able to perform simulations spanning more than 12 orders of magnitude in mass, but present-day state-of-the-art software and hardware technology does not allow simulations with more than $\\approx 10^9$ bodies. For these reasons, the quest for increasingly efficient algorithms is still in progress. However, the importance of making N-body simulations is clear to several authors \\cite{dub96}\\cite{kra97}\\cite{kuh96}. During the past years N-body codes have been much improved and applied successfully to various problems in galaxy dynamics, galaxy formation, and cosmological large structure formation. Nevertheless, the computational expense has remained prohibitive for $ N > 10^9$, even using tree-based algorithms on the most powerful computers.\\\\ The situation is even worse for other N-body algorithms. The N-body direct evolution method scales as $O(N^2)$, which  makes it impossible to run simulations with more than $10^4$ particles. To overcome this difficulty, and when high accuracy is required, alternative numerical methods based on hierarchical force-computation algorithms are widely used. The recent effort has  addressed  the production of new software and algorithms for the new  generation of high-performance computer systems. The ultimate target is  an implementation of the tree N-body algorithm to run simulations with higher  accuracy and particle number, decreasing the cost of the simulation in terms of CPU time and increasing  performance in terms of number of particles/second elaborated when running on MPP systems. \\\\ Among the tree algorithms designed to compute the gravitational force in N-body systems, one of the most used and powerful in modern cosmology is that by Barnes and Hut (BH) \\cite{barh86}. The BH octal-tree recursive method is inherently adaptive and allows one to achieve a higher mass resolution even if parallel implementation of this algorithm \\cite{bec96}\\cite{sal90}  suffers from a serious drawback: it can easily run into imbalance as soon as the configuration evolves, causing  performance degradation. In this paper we present a modified version of the BH algorithm in which we have introduced an enhanced grouping strategy. We will show how this feature allows an increase in performance when we consider N-body  simulation with a large number of particles ($ N \\geq 10^6$). The code we present incorporates  fully periodic boundary conditions using the Ewald method, without the use of fast Fourier transform techniques \\cite{her91}.\\\\ In Section 2  we give a brief description of our N-body parallel code, based  on the BH tree algorithm, and the dynamic load balance (DLB) policy adopted. In Section 3 we describe our enhanced grouping strategy.  In Section 4 we show the results of our tests and in Section 5 we report our conclusions. ",
        "conclusions": "The code  WD99  is mainly used for LSS studies, but it could be tested and used for other applications where  accuracy not higher  than 1\\% is necessary. Considering the high performances we obtained, the WD99 method may be  very successfully  applied  when clustered configurations such as galaxies or clusters of galaxies have to be studied. The new approach  could be applied also to other fields of physics where collisionless systems are to be simulated, as in plasma and hydrodynamic studies.\\\\ The code is written in Fortran 90 with PGHPF/CRAFT, but the latest version (written in F90 and C languages) uses the one-side communication library SHMEM, allowing it to run on the ORIGIN 2000 systems. A new version will be implemented using dynamical array allocation, and we are studying the implementation of the parallel out-of-core \\cite{sal97}, moving data in the disk. This version will be developed for a CC-NUMA machine with MPI-2. We plan to have a freely available version of WD99 in October 2000."
    },
    "0004/astro-ph0004294_arXiv.txt": {
        "abstract": "Radio recombination lines (RRL) at 8 GHz and 15 GHz detected from four starburst galaxies are shown to arise in compact high density HII regions, which are undetectable below $\\sim$4 GHz. Detection of an RRL at 1.4 GHz towards one galaxy and upper limits in the other three are consistent with the presence of an equal amount of low density diffuse gas. Continuum flux density measurements using the GMRT will be important in constraining the properties of the diffuse gas. ",
        "introduction": "\\vspace{-0.3cm} RRL and radio continuum studies of nuclear starbursts in galaxies are proving to be useful not only because of the absence of extinction but also because different density components of the ionized gas can be accessed through observations of RRLs at different frequencies (Zhao et al. 1996). We report a multi-frequency RRL and continuum study of four starburst galaxies. \\vspace{-0.5cm} ",
        "conclusions": "\\vspace{-0.3cm} The models show that the 8.3 GHz and 15 GHz lines originate in a population of compact (0.1$-$5 pc) high density (5000$-$50000 cm$^{-3}$) HII regions with low total volume filling factor ($<$~10$^{-4}$). These lines arise from internal emission in the HII regions. Since the NIR and optical data imply much lower densities, this component is probably not detected in these bands due to high extinction. In all four galaxies, the photon flux neccesary for ionizing this gas is equal to or greater than that derived using conventional means. This result could lead to an upward revision of their star formation rates.  The above gas is practically undetectable at frequencies \\raisebox{-0.8ex}{$\\stackrel{<}{\\sim}$}4 GHz. Modeling the detected H166$\\alpha$~line at 1.4 GHz from NGC 253 indicates that this line arises from low density (10$-$100 cm$^{-3}$) diffuse (5$-$100 pc) HII regions with an area filling factor $>$~0.1. The upper limits to the 1.4 GHz line for the other three galaxies give very similar results. Although the detection of lines from this component in the other three galaxies might be difficult with current instruments, measuring the continuum flux densities at $\\nu<$1 GHz from the nuclear region using the GMRT will enable us to strongly constrain the properties of this gas.  These observations are a first step towards deriving properties of the ionized gas at different densities in starburst regions and hence towards studying the star formation at different time scales since the lifetime of an HII region depends on its density. \\vspace{-0.6cm}"
    },
    "0004/astro-ph0004407_arXiv.txt": {
        "abstract": "We model the formation of magnetospheric components of millisecond pulsar light curves, deriving an approximate model for the curved space, `swept-back' dipole field and following photon emission, propagation, and scattering.  Magnetospheric pulse components are strongly affected by rapid rotation and Schwarzschild effects. ",
        "introduction": "With the launch of {\\it CXO} and {\\it XMM-Newton}, there has been a dramatic increase in sensitivity for high resolution X-ray studies of compact objects.  Rapidly rotating neutron stars are targets of particular interest, since measurements of the stellar surface can probe the neutron star evolution and the equation of state at high density.  Several relativistic effects may also be visible, in some cases.  A number of millisecond pulsars (MSP) have been observed in the X-rays, both as isolated rotation powered objects \\citep{bec99} and as accretion powered LMXBs \\citep[e.g.][]{wij98}.  In \\citet*{bra00}, we explored the distortions of the surface emission imparted by rapid rotation, including the important effects of Doppler boosting, aberration, and gravitational focusing and time delays. More subtle effects induced by frame dragging were also considered.  For rotation-powered pulsars, the surface radiation must traverse the magnetosphere, where resonant scattering can introduce significant distortions in the pulse profile \\citep{raj97}.  Further, several X-ray MSP have narrow, non-thermal pulse components, suggesting a direct origin in the magnetosphere.  In this paper, we extend the treatment of rapid rotation effects to include scattering and emission in a surrounding dipole magnetosphere.  Our model assumes a spherical star, which we take to be sufficiently centrally condensed that a Schwarzschild (or Kerr) metric is an adequate model of the external spacetime (the `Roche approximation'). Because we wish to extend the modeling to small spin periods $P_*$, the effects of `sweep-back' on the magnetic field, even for $r$ a few times the stellar mass $M$, can be substantial.  In addition, the rapid rotation develops large electro-motive forces.  We follow the common assumption that this EMF causes pair production such that the closed zone of the magnetosphere is filled with a charge-separated pair plasma, whose charge distribution cancels the rotational EMF.  For the magnetic field, we assume a point dipole located at the stellar center and generalize the computation of the field structure for the non-rotating dipole in the Schwarzschild spacetime \\citep[e.g.,][]{bar73,pra97,mus97}.  Our result passes smoothly to the Schwarzschild magnetic field {\\boldmath $B$} as the stellar angular velocity $\\Omega_*$ vanishes and to the flat space rotating dipole solution as $r \\rightarrow \\infty$.  We assume that the plasma co-rotates (is stationary) in the closed zone, which is defined by the `last closed' field lines traced from tangent approach to  the light cylinder at $r_{\\rm LC}=c\\,P_*/2\\pi$ to the stellar surface.  Our light curve modeling follows the procedure described in \\citet{bra00}.  This Monte Carlo code starts with photons randomly drawn from the surface emission zone, with initial directions drawn from a model limb-darkened distribution.  We aberrate these surface photons and propagate to infinity through curved space, using the Schwarzschild or Kerr metric, as appropriate.  As the photon passes through the magnetosphere, we monitor for local cyclotron resonance. At the resonance position photons are re-emitted with the proper boosted scattering angular distribution.  We also follow photons arising directly from the magnetosphere --- from acceleration gaps or other non-thermal sources.  After including the gravitational and time-of-flight delays, the photons are assigned to energy and rotational phase bins to produce maps of the radiation on the sky. Slices through these maps at the Earth's viewing angle provide pulsar light curves and phase-resolved spectra. ",
        "conclusions": "We have improved the modeling of light curves of rapidly rotating neutron stars (MSP) by extending the description of the surrounding vacuum dipole magnetosphere to include curved space and rapid rotation effects up to corrections of order $\\mathcal{O}[(\\Omega_*r/c)^2(GM/rc^2)]$.  These effects produce substantial changes in the pulse components introduced by magnetospheric resonant scattering (for thermal surface emission) or direct emission from the rotating magnetosphere (high altitude gap radiation).  We have developed a Monte Carlo code to compute light curves illustrating these effects including the energy dependence of the scattered photon pulse shapes.  Higher order effects from frame dragging have also been computed, but these are likely too subtle to be discernible in magnetospheric pulse components unless the S/N is very high.  Scattering perturbations, though small, will be greatly enhanced if the magnetosphere can support a plasma with $\\tau > \\tau_{\\rm GJ}$.  In this paper, we have only considered the minimum possible scattering perturbation.  With new detector technologies \\citep[e.g.][]{rom99}, we should see a marked increase in the quality of optical light curves.  With enhanced sensitivity and spectral resolution (and possibly higher scattering $\\tau$), light curve perturbations due to rapid rotation effects may be promoted to important probes of pulsar physics."
    },
    "0004/astro-ph0004361_arXiv.txt": {
        "abstract": "\\footnotesize We have combined 2MASS and POSS II data in a search for nearby ultracool (later than M6.5) dwarfs with $K_s<12$.  Spectroscopic follow-up observations identify 53 M7 to M9.5 dwarfs and seven L dwarfs.  The observed space density is $0.0045 \\pm 0.0008$ M8-M9.5 dwarfs per cubic parsec, without accounting for biases, consistent with a mass function that is smooth across the stellar/substellar limit. We show the observed frequency of H$\\alpha$ emission peaks at $\\sim 100\\%$ for M7 dwarfs and then decreases for cooler dwarfs. In absolute terms, however, as measured by the ratio of H$\\alpha$ to bolometric luminosity, none of the ultracool M dwarfs can be considered very active compared to earlier M dwarfs, and we show that the decrease that begins at spectral type M6 continues to the latest L dwarfs. We find that flaring is common among the coolest M dwarfs and estimate the frequency of flares at 7\\% or higher. We show that the kinematics of relatively active ($EW_{H\\alpha}>6$ \\AA) ultracool M dwarfs are consistent with an ordinary old disk stellar population, while the kinematics of inactive ultracool M dwarfs are more typical of a 0.5 Gyr old population. The early L dwarfs in the sample have kinematics consistent with old ages, suggesting that the hydrogen burning limit is near spectral types L2-L4. We use the available data on M and L dwarfs to show that chromospheric activity drops with decreasing mass and temperature, and that at a given (M8 or later) spectral type, the younger field (brown) dwarfs are less active than many of the older, more massive field stellar dwarfs. Thus, contrary to the well-known stellar age-activity relationship, low activity in field ultracool dwarfs can be an indication of comparative youth and substellar mass. ",
        "introduction": "}  Catalogs of nearby stars \\citep{gj91,khs95,rhg95} and high proper motion stars \\citep{l79} are grossly deficient in very low mass (VLM) dwarfs. With spectral types of M7 and later, these objects, sometimes called ``ultracool M dwarfs,''\\footnote{ All spectral types in this paper are on the \\citet{khm91} M dwarf and \\citet{k99} L dwarf systems.  L dwarfs are cooler than ``ultracool'' M dwarfs.} are so optically faint that even nearby ones eluded searches based on the older (pre-1980s) sky surveys. These dwarfs have particular importance because they lie at or below the hydrogen burning limit -- and have proven not only to be estimators of the numbers of dark brown dwarfs, but also present interesting astrophysical challenges in their own right.  The new generation of sky surveys allows this deficiency to be addressed and large samples of nearby VLM dwarfs to be identified. The Two Micron All-Sky Survey\\footnote{2MASS data and documentation are available at \\url{http://www.ipac.caltech.edu/2mass}} (Skrutskie et al., in prep.; hereafter 2MASS) provides reliable photometry in the JHK$_s$ passbands, close to the peak of emission for these cool dwarfs.  Furthermore, the Second Palomar Sky Survey \\citep{poss2}, hereafter POSS II) provides B$_J$,R$_{F}$, and I$_N$ photographic photometry in the northern hemisphere.  In the southern hemisphere, the UK Schmidt and ESO sky survey plates provide $B_J$ and $R_F$ magnitudes. In sum, it is becoming possible to identify both the least luminous stars and young massive brown dwarfs by their optical and near-infrared colors alone over most of the sky.  We present first results of a search using near-infrared and optical sky survey data aimed at completing the nearby star catalog for the ultracool M dwarfs. We discuss the sample selection and spectroscopic followup in Section~\\ref{data}.   Although the sample discussed in this paper includes only a small fraction of the total population of nearby ultracool dwarfs,  it represents a fourfold increase in the number of such sources known. We discuss some preliminary results concerning the statistical properties of these sources in the latter sections of this paper.  We discuss some stars of special interest in Section~\\ref{special} and the 2MASS colors of ultracool M dwarfs in Section~\\ref{colors}. The local space density of VLM dwarfs is discussed in Section~\\ref{lf}, their activity and kinematics are discussed in Section~\\ref{activity}, and finally our conclusions and future prospects are discussed in ~\\ref{summary}. ",
        "conclusions": "}  We show that a sample of bright, nearby ultracool M and L dwarfs can be selected without proper motion bias using 2MASS and PMM scans. Our initial samples include high proper motion objects, visible on the POSS plates, that should be added to an updated version of the  LHS Catalogue, and one M8.0 dwarf with a photometric parallax that places it within 10 parsecs.  We intend to continue this study in order to complete the nearby star catalog for the lowest mass stars.  Using our initial sample, we estimate the space density of dwarfs near the hydrogen-burning limit.  We show that the dropoff near the hydrogen burning limit in the five and eight parsec nearby star samples is likely to be due to incompleteness. This is more consistent with a smooth relation across the hydrogen burning limit.  Trigonometric parallaxes and searches for companions will help improve the space density estimate.  Most importantly, we use our spectroscopic observations of our well-defined sample to explore the relationships between age, kinematics, and chromospheric activity for the ultracool M and L dwarfs.  We show that the observations can be understood if activity is primarily related to temperature and secondarily mass and age, and that lower mass (substellar) objects have weaker chromosperes. Thus, the classical relation that strong H$\\alpha$ emission implies youth is not valid for these dwarfs.  Instead, strong H$\\alpha$ emitters in the field are likely to be old ($\\gtrsim 1$  Gyr) stars, while weaker emitters are often young ($<1$ Gyr), lower-mass brown dwarfs. This does not exclude the idea that for a given dwarf, H$\\alpha$ activity declines with age -- but spectral type (temperature) is the observable in the field. The local population of ultracool M dwarfs apparently consists both of the most massive (lithium burning) brown dwarfs and the lowest mass (hydrogen burning) stars, with the substellar objects making up a significant fraction of the sample. The early L (L0-L4) dwarfs are consistent with an old, at least partially stellar population.  The evidence thus suggests, as do some models, that early L dwarfs can be stable hydrogen-burning stars. Expansion of the sample with follow-up observations should clarify the relative contribution of stars and brown dwarfs to these temperature ranges."
    },
    "0004/astro-ph0004011_arXiv.txt": {
        "abstract": "An all high-latitude sky survey for  cool carbon giant stars in the Galactic halo  has   revealed  75  such  stars,   of  which  the   majority  are  new detections. Of these, more than half  are clustered on a Great Circle on the sky which intersects the center  of Sagittarius dwarf galaxy and is parallel to  its proper  motion  vector, while  many  of the  remainder are  outlying Magellanic  Cloud  carbon  stars.   Previous numerical  experiments  of  the disruption  of the  Sagittarius dwarf  galaxy (the  closest of  the Galactic satellite galaxies)  predicted that the  effect of the strong  tides, during its repeated close encounters with the Milky Way, would be to slowly disrupt that  galaxy.   Due  to  the  small velocity  dispersion  of  the  disrupted particles, these  disperse slowly along (approximately) the  orbital path of the progenitor, eventually giving rise to a very long stream of tidal debris surrounding  our Galaxy.   The  more recently  disrupted  fragments of  this stream should contain a similar mix  of stellar populations to that found in the progenitor, which includes giant carbon stars.  Given the measured position and  velocity of the Sagittarius dwarf, we first integrate its  orbit assuming  a standard spherical  model for  the Galactic potential, and find both that the  path of the orbit intersects the position of the stream, and that the radial velocity of the orbit, as viewed from the Solar position, agrees very well  with the observed radial velocities of the carbon  stars. We  also present  a pole-count  analysis of  the  carbon star distribution, which clearly  indicates that the Great Circle  stream we have isolated  is statistically  significant, being  a  5-6$\\sigma$ over-density. These two arguments strongly support our conclusion that a large fraction of the Halo carbon stars originated in the Sagittarius dwarf galaxy.  The stream orbits the Galaxy between the present location of the Sagittarius dwarf, $16\\kpc$ from the Galactic center, and the most distant stream carbon star, at  $\\sim 60\\kpc$.  It  follows neither a  polar nor a  Galactic plane orbit, so  that a  large range in  both Galactic  $R$ and $z$  distances are probed.  That  the stream is observed  as a Great Circle  indicates that the Galaxy does not exert a significant  torque upon the stream, so the Galactic potential  must   be  nearly  spherical   in  the  regions  probed   by  the stream. Furthermore, the  radial mass distribution of the  Halo must allow a particle  at the position  and with  the velocity  of the  Sagittarius dwarf galaxy to reach the distance of  the furthest stream carbon stars. Thus, the Sagittarius  dwarf  galaxy tidal  stream  gives  a  very powerful  means  to constrain the mass distribution it resides in, that is, the dark halo.  We  present  N-body experiments  simulating  this  disruption  process as  a function of  the distribution  of mass in  the Galactic halo.   A likelihood analysis  shows that, in  the Galactocentric  distance range  $16\\kpc <  R < 60\\kpc$, the dark halo is most likely almost spherical. We rule out, at high confidence levels,  the possibility that  the Halo is  significantly oblate, with  isodensity  contours of  aspect  $q_m <  0.7$.  This  result is  quite unexpected and contests currently popular galaxy formation models. ",
        "introduction": "One  of the  most intriguing  puzzles of  modern astronomy  is  the apparent existence of unseen  mass on galactic and larger scales.  Even for the Milky Way, our  knowledge of the dark  component is poor, the  little known having been gleamed from indirect, mostly kinematic, probes.  The  best constraints  on  the mass  distribution  in the  Galaxy have  been derived from the  kinematics of tracers of the  Galactic disk, in particular the \\ion{H}{1} gas,  which samples the Galaxy from almost  its center to the outermost  parts of  its disk.   Indeed, the  success of  this  technique in fitting the  \\ion{H}{1} velocities at Galactic longitudes  $\\ell > 30^\\circ$ strongly  suggests  \\citep{dehnen}  that   the  Galaxy  is  axisymmetric  at Galactocentric distances  greater than $\\sim 5  \\kpc$ --- where  most of the mass resides.  Information  on the vertical mass gradient  has been obtained from the  study of tracers  of the Galactic spheroid  \\citep{wyse, binney87, morrison90,  sommer-larsen, vandermarel,  amendt, morrison00}.   Yet despite considerable observational and theoretical effort, the vertical structure of the  spheroid  component  is  still  poorly constrained;  dependent  on  the technique employed, the  spheroid's oblateness is found to  lie in the range $0.5 < q  < 1$.  (It is  usually assumed that the density  of the spheroidal components   of    the   Galaxy   is   of   the    form   $\\rho(s)$,   where $s=(R^2+(z/q)^2)^{1/2}$, and $R$, $z$ are Galactic cylindrical coordinates). This  paucity of information  on the  vertical structure  of the  Galaxy, it transpires, is the primary cause  of the large uncertainties on the Galactic mass models \\citep{dehnen}.  The above studies of the stellar spheroid do not, however, necessarily place strong constraints  on the massive dark  halo.  This is  because the stellar spheroid  is  a partially  self-gravitating  system,  with  its own  density profile and kinematics.  Also, at large  radii, the stellar halo will not be dynamically relaxed, so  we should expect to see  a light distribution which reflects the formation  process of this component (which  may well have been quite different to  that of the dark halo).  Strong  constraints on the dark material   must  therefore   come   from  dynamical   analyses.   At   large Galactocentric radii,  satellites provide the  only means to probe  the dark halo \\citep{zaritsky89, lin, zaritsky94}. Within  the end of a galaxy's disk other  analyses are  possible,  probing the  self-gravity  of stellar  disks \\citep{ueda}, or  of HI  disks \\citep{becquaert}.  For  the case of  our own Galaxy, it has  already been possible to place constraints  on the dark halo distribution through  an analysis of  the HI disk  thickness \\citep{olling}, which suggests  that the mass  distribution is not  substantially flattened, with $q \\sim 0.8$.  In this contribution,  we analyze kinematic and distance data  of a group of newly-discovered stars that  are part of a stream of  material that has been tidally stripped  from the  Sagittarius dwarf galaxy.   A stream  gives much stronger  constraints  on  the  underlying  mass distribution  than  does  a uniformly distributed spheroid star population, as it implies a low velocity dispersion  and  a  continuity  between  the  stream  members.  It  provides essentially the trace of an orbit through the Halo.  The stream studied here is  now distributed  through the  Halo between  Galactocentric  distances of $16\\kpc$ to $\\sim 60\\kpc$, probing a  region between the end of the Galactic disk, and the Magellanic clouds.  In section~2,  we discuss  the all-sky survey  that reveals the  Halo carbon star population; section~3 reviews previous analyses of the evolution of the Sagittarius  dwarf galaxy and  their predictions  about its  stellar stream; section~4 makes a  simple first-pass comparison of the  expected stream with the  carbon  star  dataset  and   discusses  the  probability  of  a  chance alignment. Section~5  presents a simple  analytic argument to  constrain the Halo shape  and mass; section~6 provides  a fuller analysis with  the aid of N-body  simulations  covering  a  range  of  static  Halo  models.  Finally, section~7 presents the conclusions of this work. ",
        "conclusions": "The results  presented above  provide strong evidence  that the  dark matter halo   surrounding  our   Galaxy   is  not   significantly  oblate   between Galactocentric radii of $\\sim 16\\kpc$  to $\\sim 60\\kpc$.  Flat halos $q<0.7$ are  ruled  out at  very  high  confidence  levels. Therefore,  dark  matter candidates  such as  cold molecular  gas  \\citep{pfenniger94a, pfenniger94b, combes} or  massive decaying neutrinos \\citep{sciama}, that  would give rise to  a highly  flattened component,  cannot contribute  significantly  to the galactic mass budget.  Our result is at odds with  some earlier studies that used spheroid  tracers to  probe the  dark halo; the  large spread  of $q_m$ between different  teams adopting  that approach is  likely a result  of the different  assumptions  employed to  reduce  the  complexity  of the  Jeans' equations, and also  of the technical difficulty of  measuring the necessary input quantities to those equations.  In contrast, our measurement hinges on the very simple physical principle  of conservation of angular momentum in a spherical potential, which, we believe, provides a much cleaner test.  \\citet{olling}  have recently  summarized  extant measurements  of the  dark matter shapes of galaxies. The methods used to date are confined to analyses of the  flaring of the galactic gas  layer, warping of the  gas layer, X-ray isophotes, polar  ring analysis, and  precession of dusty disks.   They note that the derived answer correlates strongly with the technique employed (see their Figure~1). It is interesting  that these measurements are derived from data at  smaller galactocentric distance  than the stellar  stream presented here. It  is possible that this reflects  a radial gradient in  the shape of dark  matter  halos,  the   inner  regions  being  more  flattened.   Better statistics,   perhaps   from    gravitational   weak   lensing   experiments \\citep{fischer}, may help solve this issue.  Is our result unexpected on theoretical grounds?  Numerical simulations with purely  collisionless matter  \\citep{katz91a,  dubinski91, warren,  katz91b, summers, dubinski94}  produce flattened triaxial halos,  with a distribution of flattening ratios that peaks near  $q_m = 0.7$. However, it is found that the presence of  a realistic fraction of dissipational  gas particles in the simulations  \\citep{katz91b, summers,  steinmetz} alters  the orbits  of the dark matter  particles, giving  rise to oblate  and slightly  flatter halos. The conclusion  of this study, albeit  a single measurement,  is somewhat at odds with these widely popular theories;  we find in all possible cases that the Galactic  dark matter  halo is closer  to being spherical.   The initial conditions,  and  the  physics  and  numerical techniques  employed  in  the subsequent  evolution  of  those  galaxy  formation  simulations  should  be re-addressed in the  light of these observational results  if other galaxies are also deduced to have nearly spherical mass distributions at large radii.  Further  surveys should attempt  to discover  the other  stellar populations associated with the Great Circle  streams identified in this study. Stars of interest will  include RRLyrae variables, which are  better standard candles than  C-stars, and  being  much older,  trace  the more  ancient history  of disruption events in our Galactic halo. These stars could be identified from the Sloan  Digital Sky  Survey (SDSS) for  instance.  Furthermore,  SDSS (or other surveys) could help improve  the stream statistics by finding the much more numerous fainter stars  spatially and kinematically associated with the C-stars presented above. \\footnote{After   the  submission   of  this   article,   \\citet{yanny}  and \\citet{ivezic} discovered  the presence of  a large population  of A-colored Halo stars in the SDSS dataset,  distributed apparently in a ring around the Milky Way. In a companion  paper \\citep{ibata00b}, we argue that those stars are most likely the RR-Lyrae  members of the Sagittarius stream discussed in the  present contribution.  Analysis of  the  (as yet  very limited)  region covered by the  SDSS survey shows a distribution in  good agreement with the simulation of Figure~8 (where $q_m=0.9$).}  On a  large telescope equipped  with a wide  field camera (such  as Subaru's Suprime-Cam), it  will also  be possible to  conduct a  complementary C-star survey around our  neighboring galaxy M31.  It will  be very interesting to see  whether  similar  Great Circle  streams  are detected  in  a  different environment.  With the advent  of the next generation of  space astrometric missions, such as the National Aeronautics and  Space Administration's SIM and the European Space Agency's GAIA,  it will be possible to  obtain accurate proper motions and  even distances  for  stars from  the  Halo streams  identified in  this contribution.    With   the  resulting   full   6-dimensional  phase   space information,  we can  expect to  be able  to map  the distant  Galactic mass distribution  in exquisite detail,  and clarify  how the  Milky Way  and its satellite galaxies  formed, and  how the  latter came to  be torn  apart and their contents flung across the sky."
    },
    "0004/astro-ph0004227_arXiv.txt": {
        "abstract": "{We show that our Universe may be inhomogeneous on large sub-horizon scales  without us being able to realise it. We assume that a network of domain walls permeates the  universe dividing it in domains with slightly different vacuum energy densities.   We require that the energy scale of the phase transition which produced the domain walls is sufficiently low so that the walls  have a negligible effect on structure formation. Nevertheless,  the different vacuum densities of  different domains will lead to different values of the cosmological  parameters $\\Omega_\\Lambda^0$, $\\Omega_m^0$ and $h$, in each patch  thus affecting the growth of cosmological perturbations at recent times. Hence, if our local patch of the universe (with uniform vacuum density) is big enough -- which is likely to happen given that we should have on average about one domain per horizon volume -- we might not notice these large-scale inhomogeneities. This happens because in order to see a patch with a different vacuum density one may have to look back at a time when the universe was still very homogeneous. } ",
        "introduction": "\\label{secintro}   The last year or so has seen the first serious attempts to provide some direct connections between  ``fundamental'' high-energy physics \\cite{Polchinski} and ``low-energy'' standard cosmology \\cite{Kolb}. Although this ``top-down'' approach is still at a very early stage, a number of crucial general trends already became apparent. For example, since high-energy theories are formulated in higher dimensions, any low energy limit will necessarily involve dimensional reduction, and possibly also compactification \\cite{Banks}. This turns out to be crucial because as a result of this process the low energy, four-dimensional coupling constants become functions of the radii of extra dimensions, which are often variables. One can therefore end up with low-energy effective models in which some of the ``fundamental'' constants of nature are time and/or space-varying quantities. There are a number of known examples of such models \\cite{Chodos,Marciano,WuWang,Kiritsis,Alexander}. On the other hand, there are recent tentative suggestions of a time-variation of the fine structure constant \\cite{Webbetal}, but these  require further confirmation. It is therefore interesting to study the possible observational signatures of such variations, and in particular to find out how such observational signatures constrain the possible models. It turns out to be easier to study this issue by constructing simple ``toy models''. This provides a ``bottom-up'' approach, in which one gives up the possibility of testing particular assumptions from first principles, but instead has the possibility of exploring a larger patch of parameter space. This idea goes back at least to Dirac, and had its first detailed realization with the Brans-Dicke model \\cite{BransDicke}, which has a varying $G$. A number of toy models have recently been constructed to analyse possible variations of the fine structure constant \\cite{Hannestad,Kaplinghat,Bergstrom}, the speed of light \\cite{Mof92,Mof98,BarMag98,AlbreMag99,Bar99,AveMar99,BarMag99,AveMarRoc,AveMar00a} and electric charge \\cite{Bekenstein}.  A somewhat related approach is that of ``quintessence'' (see for example \\cite{PeeblesRatra,Wetterich,Caldwell,Hueyetal}). These are essentially models with a time-varying cosmological constant.  Here we consider the possible observational effects of having a universe made up of different domains, each with a different value of the cosmological constant. Such a structure would dramatically influence the future evolution of the universe \\cite{Starkman,AvedeCarMar}. Recent observations of Type Ia Supernovae up to redshifts of about $z \\sim 1$ \\cite{Perlmutter,Riess,Garnavich}, when combined with CMBR anisotropy data, seem to indicate that our patch of the universe is currently characterised by the parameters $\\Omega_\\Lambda^0 \\simeq 0.7$ and $\\Omega_m^0 \\simeq 0.3$, implying that the cosmological constant become important only very recently. As has been pointed out before, the Supernovae measurements are local, and so they can not be extrapolated all the way to the horizon. For example, we could be living in a small, sub-critical bubble, and our local neighbourhood could have a value of $\\Omega_\\Lambda^0\\simeq0.7$ that is uncharacteristic. Here we discuss some basic consequences of such a scenario. We shall assume that different regions of space have different values of the vacuum energy density, separated by domain walls. This can be achieved if there is a scalar field, say $\\phi$, which within each region sits in one of a number of possible minima of a time-independent potential. The above simplifying assumptions could be relaxed; for example one could instead consider quintessence-type fields. This would introduce quantitative differences, but would not change the basic qualitative results we are discussing.  In the following section we describe our numerical simulations of the evolution of the domain wall network. We then proceed to discuss the basic features of the structure formation mechanism for this scenario in section~\\ref{growth}. Finally we present our results in section~\\ref{results} and discuss our conclusions in section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this paper we have provided a simple example of a cosmological scenario where the universe becomes inhomogeneous at a very recent epoch, in a way which is perfectly consistent with current observations assuming that we are not very close to one of the boundaries. The inhomogeneity arises due to the onset of vacuum energy domination, if the value of the `cosmological' constant is different in different domains. This in turn implies that the subsequent dynamics of each patch of the universe will be different, leading to different values of other cosmological parameters in each patch, such as the matter density and the Hubble constant. The size of each patch is determined by the dynamics of the domain walls but is expected to be of the order of the horizon (although causality prevents it from exceeding it). This fact makes the detection of the cosmological signatures of this kind of model more difficult given that the contribution of the vacuum energy density rapidly becomes negligible with increasing  red-shift."
    },
    "0004/astro-ph0004132_arXiv.txt": {
        "abstract": "According to the hierarchical scenario, galaxies form via merging and accretion of small objects. Using $N$-body simulations, we study the frequency of merging events in the history of the halos.  We find that at $z \\la 2 $ the merging rate of the overall halo population can be described by a simple power law $(1+z)^{3}$.  The main emphasis of the paper is on the effects of environment of halos at the present epoch ($z=0$). We find that the halos located inside clusters have formed earlier ($\\Delta z \\approx 1$) than isolated halos of the same mass. At low redshifts ($z<1$), the merger rate of cluster halos is 3 times lower than that of isolated halos and 2 times lower than merger rate of halos that end up in groups by $z=0$.  At higher redshifts ($z\\sim 1-4$), progenitors of cluster and group halos have 3--5 times higher merger rates than isolated halos. We briefly discuss implications of our results for galaxy evolution in different environments. ",
        "introduction": "\\label{intro}  \\footnotetext{Hubble Fellow} A significant fraction of mass in the universe is believed to be in the form of dark matter (DM).  According to the standard theoretical paradigm of structure formation, small-mass DM perturbations collapse first and the resulting objects then merge to form increasingly larger DM halos.  Baryonic matter (gas) is assumed to follow the gravitationally dominant dark matter.  Galaxies, thus, could have been formed within dense DM halos when the infalling gas reaches sufficiently high overdensities to cool, condense, and form stars.  The most convincing observational evidence for substantial amounts of dark matter even in the very inner regions of galaxies comes from HI studies of dwarf and low surface brightness galaxies. The gravitational domination of DM on the scale of galaxy virial radius implies that collisionless simulations can be used to study the formation of the DM component of galaxies.  Interactions  between halos, such   as mergers, collisions, and  tidal stripping,  are thought to   play a crucial role  in  the evolution of galaxies.   In   particular, there  is    a substantial evidence  that elliptical galaxies may have  formed   by mergers of disk  systems  (e.g., Barnes   1999). Observations of  faint   distant systems indicate that interaction rate rapidly increases with redshift (e.g., Abraham 1999). Intuitively, one could expect that the merging rate of galaxies should depend  on environment (in    particular,  on the  local  density  and velocity dispersion).  For example, Makino  \\&  Hut (1997) found under some simplifying  assumptions that  the  merging rate in  clusters  is proportional  to $n^2\\sigma^{-3}$ where $n$   is the number density of galaxies in the cluster and $\\sigma$ is their one dimensional velocity dispersion of galactic velocities.  Since the environment changes with time one could also  expect dramatic changes  in the evolution of  the merging rate.  In order to study the evolution of the merging rate and its dependence on   environment one must   follow   the evolution    of  halos in   a representative cosmological   volume.  Moreover,  the simulation  must have  sufficiently   high  mass and   force  resolutions. Insufficient resolutions leads to structureless virialized halos instead of systems similar  to observed groups  and  clusters of  galaxies with wealth of substructure. This effect is well   known as the overmerging   problem (e.g., Moore \\etal 1996, Frenk \\etal 1996, Klypin \\etal 1999).  Cosmological scenarios with cold dark matter (CDM) alone cannot explain the structure formation both on small and very large scales.  Variants of the CDM model with a non-zero cosmological constant, $\\Lambda$, have proven to be very successful in describing most of the observational data at both low and high redshifts. Moreover, from a recent analysis of 42 high-redshift supernovae Perlmutter \\etal (1999) found direct evidence for $\\Omega_{\\Lambda}= 0.72$, if a flat cosmology is assumed. Also, from the recent BOOMERANG data Melchiorri \\etal (1999) found strong evidence against an open universe with $\\Lambda = 0$. For our study we have chosen a spatially flat cosmological model with a cosmological constant $\\Omega_{\\Lambda}= 0.7$ and the present-day Hubble constant of $H_0=70$~km/s/Mpc.  The goal of this study is to determine  the distribution with redshift of merging  events for  halos   which exist  at $z=0$. We   study this merging rate and its dependence on environment of  halos at $z=0$. The paper is  organized  as follows. In the   next  section we  define the merging events studied in this paper. In \\S~3 we describe the cosmological model and  the numerical simulation.  We briefly describe our  halo  finding algorithm,  the  definition of  environment and the detection of progenitors  of   halos. The technical details   of these procedures are   presented   in the  Appendix.  We  use   the extended Press-Schechter formalism to  test our  procedure. In \\S~4 we discuss  the merging of halos  found in  the simulation and compare our results  with  observations. In   \\S~5 we summarize  our results and briefly discuss their implications. ",
        "conclusions": "\\label{concl}  We have analyzed a high-resolution collisionless simulation of the evolution of structure in a $\\Lambda$CDM model. We have followed the formation and evolution of DM halos in different cosmological environments and estimated the evolution of the major merger rate of dark matter halos.  We have found that regardless of their present-day mass, halos that end in clusters form earlier than isolated halos of the same mass (Figs.~3 and 4). We find that at redshifts $z \\la 2$ major merger rate evolves as $(1+z)^{\\sim 3.0}$, in good agreement with observations.  Finally, we have calculated the merging rate evolution as a function of halo environment at $z=0$ (Fig.~9).  The merger rate of halos located in clusters or groups at present increases faster back in time than that of isolated halos. The cluster and group halos are therefore predicted to have a higher rate of major merger events in the past. At $z\\la 1$, the merger rate of cluster and group halos drops very quickly, while numerous major merger events for isolated halos have been detected down to $z=0$. This implies possible systematic differences between cluster and field ellipticals. Evidence for such differences was found by de Carvalho \\& Djorgovski (1992), while Bernardi \\etal 1998 detected close similarity between cluster and field early type galaxies.  The agreement between theoretical predictions and observations are encouraging and supports the validity of the hierarchical structure formation scenario. Future, higher resolution simulations should extend the predictions presented here to higher redshift and, in case of gasdynamics simulations, provide a more straightforward connection to observations. On the observational side, the ever increasing size of the high-redshift galaxy samples should also allow estimates of the merger rate at high redshifts in the near future."
    },
    "0004/astro-ph0004304_arXiv.txt": {
        "abstract": "Clusters of galaxies are the most massive objects in the Universe and mapping their location is an important astronomical problem. This paper describes an algorithm (based on statistical signal processing methods), a software architecture (based on a hybrid layered approach) and a parallelization scheme (based on a client/server model) for finding clusters of galaxies in large astronomical databases. The Adaptive Matched Filter (AMF) algorithm presented here identifies clusters by finding the peaks in a cluster likelihood map generated by convolving a galaxy survey with a filter based on a cluster model and a background model. The method has proved successful in identifying clusters in real and simulated data.  The implementation is flexible and readily executed in parallel on a network of workstations. ",
        "introduction": "Clusters of galaxies are the largest objects known to humans (see Figure \\ref{fig:colley_cluster}).  They are the ``mountains'' of the Cosmos, and like terrestial mountains they lie in great ranges that define the cosmic ``continents'' and ``oceans'' (see Figure \\ref{fig:ncsa_sheets}). Mapping clusters of galaxies is very much akin to surveying our own world and allows us to understand the creation, evolution and eventual fate of our Universe \\cite{Bahcall88}.  The process by which astronomers detect clusters of galaxies begins with assembling large images of the sky, which are the result of hundreds of nights of observing through a telescope. These pictures are analyzed to produce a database of galaxies $X = \\{ \\xb_i : i = 1, \\ldots, N_X \\}$, where $N_X \\sim 10^8$. Each record $\\xb_i \\in X$ consists of a position on the sky, brightness measurements in one or more bands and possibly hundreds of additional measurements describing the shape and composition of the galaxy.  Clusters are local density peaks in the three dimensional distribution of galaxies across the Universe.  In 3D data, clusters are easy to detect.  Unfortunately, the majority of distances to individual galaxies are not known and can only be inferred statistically from empirical models of their brightness.  Thus, it is difficult to differentiate small nearby clusters from large far away clusters. The goal of cluster detection and estimation is to create a catalog, $\\Theta$, consisting of thousands of clusters $\\Theta = \\{ \\thetab_i : i = 1, \\ldots, N_\\Theta \\}$, where $N_\\Theta \\sim 10^4$.  Each cluster in this list, $\\thetab_i$, consists of a position on the sky, a distance estimate, a size estimate and perhaps additional estimated properties of the cluster.   The first catalog of galaxy clusters was compiled by Abell \\cite{Abell58}, and has proved extremely useful to astronomers over the past four decades.  Abell's catalog was created by visually inspecting hundreds of photographic plates taken from the first Palomar Observatory Sky Survey (POSS). Modern galaxy databases are too large for such methods to be used today.  Subsequent efforts to detect clusters have relied on Matched Filter techniques taken from statistical signal processing (e.g., \\cite{Lumsden92}, \\cite{Dalton94}, \\cite{Postman96}, \\cite{Kawasaki97} and \\cite{Bramel00}.  These methods have a strong mathematical foundation, but require extensive prior information and are often computationally prohibitive as they test every possible location in the domain of the cluster space $\\Omega_\\Theta$ for the presence of a cluster. The Adaptive Matched Filter \\cite{Kepner99} that is described later in this paper is a variation on the Matched Filter that uses a hierarchical set of filters, as well as software coding and parallel computing techniques that address some of the Matched Filter's drawbacks. The AMF is adaptive in two ways. First, the AMF uses a two step approach that first applies a coarse filter to find the clusters and then a fine filter to provide more precise estimates of the distance and size of each cluster. Second, the AMF uses the location of the data points as a ``naturally'' adaptive grid to ensure sufficient spatial resolution.  A variety of other techniques have also been applied to the cluster finding problem.  The compact nature of clusters make Wavelet based signal processing approaches an appealing alternative \\cite{Fang97} and \\cite{Fadda97}. Geometric approaches such as Voronoi tessellation \\cite{Ramella98} have also been used.  In this method each $\\xb_i$ is the seed for the tessellation.  Clusters are then found by computing the volume of each tessel and selecting the points with the smallest volume, which presumably have the highest density.  Such geometric methods have the advantage that they require very little prior information.  The Matched Filter, Adaptive Matched Filter, Wavelet and Voronoi Tessellation approaches all use the three to five high affinity dimensions of $X$ (i.e., angular position and brightness measurements). These dimensions are continuous real variables that lend themselves to Euclidean distance metrics.  Working in these lower dimensions allows more compute intensive techniques which are necessary to de-project clusters from the observed data domain $\\Omega_X$ to the desired underlying domain in which clusters exist $\\Omega_\\Theta$, i.e., angular position, {\\it distance} and size. More recently, there has been interest in exploiting the low affinity dimensions that are also available in galaxy databases \\cite{Djorgovski97,Gal99} to enhance detection.  As the understanding of these methods increases, the exploitation of many dimensions should be possible using advanced datamining techniques (see e.g. \\cite{Fayyad96} and \\cite{Dasarathy99} and references therein). These methods have enormous potential for detecting new clusters and possibly separating them into distinct groups thus revealing new classes of galaxy clusters.   The rest of this paper presents in greater detail the AMF algorithm, its implementation and results.  In section two a detailed derivation of the AMF is given.  The derivation is meant to be sufficiently general that it can lend itself to other types of databases. Section three presents the implementation of the AMF using a layered software architecture and a client/server parallelization model.  Again, these methods are not limited to the specific problem presented here and are applicable to a variety areas. In section four the results of applying the AMF on simulated and real data are discussed.  Finally, section five gives the summary and conclusions. ",
        "conclusions": "We have presented the Adaptive Matched Filter method for the automatic selection of clusters of galaxies from a galaxy database. The AMF is adaptive in two ways. First, the AMF uses a two step approach that first applies a coarse filter to find the clusters and then a fine filter to provide more precise estimates of the distance and size of each cluster. Second, the AMF uses the location of the data points as a ``naturally'' adaptive grid to ensure sufficient spatial resolution.  Matched Filter techniques have a firm mathematical basis in statistical signal processing.  The AMF uses a hierarchy of two filters (each mathematically correct under its assumptions).  Combining these filters allow the AMF to maximize computational performance and accuracy. The AMF also provides two estimates for each cluster which can be compared as an additional check.  This is particular effective for these filters because they react differently when given insufficient data.  The AMF relies heavily on models for both the cluster and the background field.  This prior information is quite extensive and makes the AMF complex to implement and difficult to adapt to new data sets. To alleviate this coding challenge a hybrid coding approach was used to leverage the ease of use of interpreted languages along with the compute performance  of compiled languages.  In this way the complex task of testing model inputs and observing their effect through the data processing pipeline can be done quickly without sacrificing the compute efficiency necessary to complete the application in a timely manner.  A further benefit of the hybrid approach is that it makes available to the compiled code a wide variety of parallel software libraries and tools.  A parallel implementation is critical to the application because matched filter techniques work by testing every possible location in the cluster space for the presence of a cluster.  This is a compute intensive operation, but also provides a high degree of parallelism. The parallelization scheme used for the AMF application is a client/server approach which is a very effective on Network-Of-Workstations.  The TNT client/server software used is lightweight and efficient, and provides a naturally load balancing and fault tolerant framework.  The AMF has been extensively tested on simulated data.  These results indicate that it robustly and accurately detects clusters and estimates their positions while having few false positives. The AMF is now being applied to the first results of the Sloan Digital Sky Survey \\cite{Kim00}.  These tests have shown that the AMF detects all previously known clusters in this data and performs at or above other cluster finding methods.  The AMF hybrid application architecture has proven effective in supporting the implementation of new datasets.  The TNT based client/server parallelization scheme has also demonstrated significant speedups which will make it feasible for this application to address to the entire SDSS when it becomes available."
    },
    "0004/astro-ph0004074_arXiv.txt": {
        "abstract": "The KPNO International Spectroscopic Survey (KISS) is a new objective-prism survey for extragalactic emission-line objects.  It combines many of the features of previous slitless spectroscopic surveys that were carried out with Schmidt telescopes using photographic plates with the advantages of modern CCD detectors.  It is the first purely digital objective-prism survey, and extends previous photographic surveys to substantially fainter flux limits.  In this, the first paper in the series, we give an overview of the survey technique, describe our data processing procedures, and present examples of the types of objects found by KISS.  Our first H$\\alpha$-selected survey list detects objects at the rate of 18.1 per square degree, which is 181 times higher than the surface density of the Markarian survey. Since the sample is line-selected, there is an imposed redshift limit of z$\\simlt$0.095 due to the filter employed for the objective-prism observations.  We evaluate the quality of the observed parameters derived from the survey data, which include accurate astrometry, photometry, redshifts, and line fluxes.  Finally, we describe some of the many applications the KISS database will have for addressing specific questions in extragalactic astronomy.  Subsequent papers in this series will present our survey lists of emission-line galaxy candidates. ",
        "introduction": "\\par Active galactic nuclei (AGN) and galaxian starbursts are among the most energetic phenomena known in the universe.  From Blue Compact Dwarf galaxies to QSOs, galaxian activity manifests itself on all scales and over the entire range of the electromagnetic spectrum.  Activity in galaxies appears to be common, with between 5 and 10\\% of all galaxies showing evidence for it via the presence of unusually strong emission lines in their spectra (e.g., Salzer 1989, Gregory \\etal 2000).  The study of these objects has been a major research topic for many years, and a great deal has been learned about the physical processes occurring in the active regions of both AGN and starburst galaxies.  Much still remains to be understood.  How does the activity, be it central or global, shape the overall evolution of the host galaxy? Do all galaxies cycle through periods of activity?  What role does environment play in whether or not a galaxy becomes active?  Given that many of the galaxies we observe at high redshift are active, can we understand their evolutionary status in the context of what we observe locally?  How has the global star formation density evolved as a function of cosmic epoch?  \\par Central to the study of AGN and starburst activity have been the surveys which have cataloged large samples for subsequent investigation.  Few types of extragalactic surveys have been as scientifically fruitful as the various objective-prism surveys for UV-excess and emission-line galaxies (ELGs) carried out with Schmidt telescopes.  Much of what we know about Seyfert galaxies, starburst galaxies, and even QSOs has been learned by studying objects originally discovered in surveys with such familiar names as the Markarian, Tololo, Wasilewski, Michigan (UM), Kiso, Case, and Second Byurakan (SBS) surveys.  \\par Virtually all of the existing surveys for galaxies which display some form of unusual activity (e.g., blue colors or strong emission lines) have been carried out using Schmidt telescopes and one of three detection methods.  The first, which is commonly referred to as the color-selection technique, uses multiple exposures taken through two or three filters to isolate the bluest galaxies.  This method has been used by Haro (1956), Takase \\& Miyauchi-Isobe (1993; Kiso survey), and Coziol {\\it et al.} (1993, 1997; Montreal survey).  The well known Palomar-Green survey (Green {\\it et al.} 1986; PG survey) for UV-bright stars and QSOs also employed this technique. The other two methods employ objective prisms.  One, the UV-excess technique, was pioneered by Markarian (1967) and Markarian \\etal (1981). It is similar to the color-selection method in that the criterion for inclusion is the presence of a very blue continuum.  The other selects objects via the presence of emission lines in the objective-prism spectra.  Surveys of this type include the Tololo (Smith 1975, Smith {\\it et al.} 1976), Michigan (UM, MacAlpine {\\it et al.} 1977, MacAlpine \\& Williams 1981), Wasilewski (1983), and Universidad Complutense de Madrid (UCM, Zamorano {\\it et al.} 1994, 1996; Alonso {\\it et al.} 1999) surveys. Finally, the Case (Pesch \\& Sanduleak 1983, Stephenson {\\it et al.} 1992) and Second Byurakan surveys (Markarian \\etal 1983, Markarian \\& Stepanian 1983, Stepanian 1994) use a hybrid scheme, selecting galaxies based on both UV excess and line emission.  \\par In recent years various observatories have begun experimenting with the use of CCD detectors on Schmidt telescopes.  While CCDs had become the detector of choice for most observing applications by the mid-1980's, photographic plates have remained in use on Schmidt telescopes due to their vastly superior areal coverage.  However, the advent of large format CCDs (2048 $\\times$ 2048 pixels and larger) now provides sufficiently large fields-of-view (in excess of one degree square) that it has become advantageous to carry out a variety of wide-field survey projects using CCDs on Schmidt telescopes (Armandroff 1995).  These advances have motivated us to initiate a new survey for emission-line galaxies.  \\par Our new survey is called KISS - the KPNO International Spectroscopic Survey.  The technique we employ combines the benefits of a traditional photographic objective-prism survey with the advantages of using a CCD detector, and represents the {\\bf next generation of wide-field objective-prism surveys}.  The goal of the KISS project is to survey a large area of the sky for extragalactic emission-line sources, and to reach a minimum of two magnitudes deeper than any of the previous line-selected Schmidt surveys.  Our initial plan was to be able to detect strong-lined ELGs with continuum magnitudes down to B = 20 - 21, and to be complete to B = 19 - 20.  The use of a CCD detector allows us to achieve these goals.  \\par We stress that the need for this type of survey goes beyond the desire to simply find {\\bf more} AGN and starburst galaxies.  The type of scientific questions we wish to address requires large samples with well defined selection criteria and completeness limits.  Our research cuts across the traditional lines of study of active galaxies, and touches upon nearly all areas of extragalactic astronomy, such as galaxy formation and evolution, chemical evolution, and large-scale structure and cosmology.  Although we fully expect to discover many new objects which possess noteworthy characteristics that will warrant individual study, the prime motivation for this survey is to create a sample of objects that can be used as probes for studying more general questions.  We will discuss some of these applications of the KISS sample of galaxies in Section 5.  \\par In the next section we describe in detail our instrumental setup and observational technique, while Section 3 gives an overview of our image processing and analysis software which is central to the survey effort. Section 4 presents the results from our first two survey lists, evaluates the quality of the survey data (e.g., photometry, astrometry, redshifts), and gives a preliminary analysis of the survey contents. In Section 5 we list some of the possible uses of the KISS sample for attacking specific problems relevant to extragalactic astronomy and cosmology.  Finally, Section 6 summarizes our results. ",
        "conclusions": "We have presented a description of the KPNO International Spectroscopic Survey (KISS), a new objective-prism survey for extragalactic emission-line objects.  KISS is the first large survey to combine CCD detectors with the traditional Schmidt telescope slitless-spectroscopic-survey method, which has proven to be extremely fruitful over the past three decades.  The superior sensitivity of the CCD allows us to survey to fainter flux levels than were previously possible using photographic plates.  In addition, the digital nature of the data will allow us to assess the selection function and completeness limits of the survey far more accurately than with photographic surveys.  KISS is being carried out on the Burrell Schmidt telescope.  The first survey region was observed in both the red and blue spectral regions. Both sets of objective-prism spectra were obtained over a restricted wavelength region in order to minimize the problems of overlapping spectra and to reduce the level of the sky background.  The blue spectral region covers 4800 - 5500 \\AA, and the primary emission line we select by is [\\ion{O}{3}]$\\lambda$5007. The red spectra cover the wavelength range from 6400 - 7200 \\AA, and objects are detected primarily via the H$\\alpha$ line.  Both the red and the blue spectra detect galaxies out to z $\\approx$ 0.095 via the primary line, and are sensitive to higher-z galaxies over restricted redshift intervals as other lines move into the survey bandpasses.  Because of the digital nature of the survey data, all of the candidate selection is carried out using automated software, although the final candidate lists are also checked visually.  In addition to the objective-prism spectra used to select the ELG candidates, we also have direct images taken through B and V filters which provide accurate astrometry, photometry, and morphological information for all objects in each field. These data, plus estimates of the redshift and line flux for each ELG candidate obtained from the objective-prism spectra, yield a fairly complete picture of each candidate galaxy without the need for additional follow-up observations.  This makes the KISS database particularly valuable for statistical studies of galaxian activity.  The results obtained from our first survey lists are extremely encouraging. The first blue survey strip covers 117 deg$^2$ and includes 223 cataloged ELG candidates.  With a surface density of just under 2 ELGs per square degree, the blue KISS sample is substantially deeper than previous surveys of this type. The numbers for the H$\\alpha$ (red) survey are even more impressive.  A total of 1128 ELGs have been found in an area of 62 deg$^2$, for a surface density of 18.1 ELGs per square degree.  This is 32 times the surface density of the H$\\alpha$-selected UCM survey, and 181 times that of the UV-excess Markarian survey.  The first two survey lists will be presented in Salzer \\etal (2000a,b), along with a complete discussion of the characteristics of the two samples. Additional survey lists will be published as the data are acquired, processed, and cataloged.  To date, data covering 200 deg$^2$ of sky have been obtained, and observations are continuing.  Our overall project goal is to cover in excess of 300 deg$^2$ in a series of survey strips in selected areas of both the north and south Galactic caps.  This will yield several thousand ELG candidates, which will be suitable for addressing a wide range of scientific questions that cover nearly the full scope of extragalactic astronomy."
    },
    "0004/astro-ph0004242_arXiv.txt": {
        "abstract": "We present the first measurements of the X-ray size-temperature (ST) relation in intermediate redshift ($z\\sim0.30$) galaxy clusters. We interpret the local ST relation ($z\\sim0.06$) in terms of underlying scaling relations in the cluster dark matter properties, and then we use standard models for the redshift evolution of those dark matter properties to show that the ST relation does not evolve with redshift.  We then use ROSAT HRI observations of 11 clusters to examine the intermediate redshift ST relation; for currently favored cosmological parameters, the intermediate redshift ST relation is consistent with that of local clusters. Finally, we use the ST relation and our evolution model to measure angular diameter distances; with these 11 distances we evaluate constraints on $\\Omega_{M}$ and $\\Omega_{\\Lambda}$ which are consistent with those derived from studies of Type Ia supernovae. The data rule out a model with $\\Omega_{M}=1$ and $\\Omega_{\\Lambda}=0$ with 2.5$\\sigma$ confidence.  When limited to models where $\\Omega_{M}+\\Omega_{\\Lambda}=1$, these data are inconsistent with $\\Omega_{M}=1$ with 3$\\sigma$ confidence. ",
        "introduction": "Nearby galaxy clusters exhibit a tight correlation between X-ray isophotal size and emission weighted intracluster medium (ICM) temperature (\\cite{mohr97}; hereafter ME97).  This correlation is evidence of regularity; it exists in an X-ray flux limited sample of 45 clusters (\\cite{edge90}), where no attempt has been made to use the X-ray morphologies to exclude clusters showing signatures of recent, major mergers.  The scatter around the X-ray size--temperature (ST) relation is approximately 15\\% in size, comparable to the scatter of elliptical and lenticular galaxies around their fundamental plane (\\cite{jorgensen96}).  This small scatter in the galaxy cluster scaling relation is intriguing, because (1) there is overwhelming evidence that galaxy clusters are still accreting mass (e.g. \\cite{mohr95,buote96}) and (2) elliptical galaxies are generally thought to be among the most regular objects in the universe.  ME97 use 48 N-body and hydrodynamical simulations of cluster formation in four different cosmological models to address this apparent contradiction between regularity and ongoing accretion in nearby clusters.  Using simulations from both $\\Omega_{M}=0.3$ and $\\Omega_{M}=1$ cosmologies, they show that a tight ST relation is expected even in cosmogonies where there is significant cluster growth at the present epoch.   The high degree of regularity implied by the ST relation is surprising, because the well known correlation between X-ray luminosity and emission weighted mean temperature (the \\lxt\\ relation) has very large scatter (\\cite{david93}). ME97 show that the same cluster ensemble which exhibits a ~15\\% scatter in the ST relation exhibits a 52\\% scatter in $L_{X}$ around the \\lxt\\ relation.  This higher scatter in the \\lxt\\ relation results from the sensitivity of the X-ray luminosity to the densest regions of the cluster-- a sensitivity to the presence or absence of so-called cooling flows (\\cite{fabian94}).  This interpretation is supported by more recent work where cluster ensembles specially chosen to contain no cooling flow clusters conform to \\lxt\\ relations with significantly reduced scatter of ~25\\% (\\cite{arnaud99}). Additionally, when central parts of cooling flow clusters are excluded, the scatter in the \\lxt\\ relation decreases (\\cite{markevitch98}).  Finally, cluster regularity is also evident in the tight correlation between ICM mass and emission weighted temperature (the \\micmT\\ relation).  When measuring ICM mass within a limiting radius of \\rfive\\ (the radius where the enclosed overdensity is 500 times the critical density) the scatter in mass about the \\micmT\\ relation is 17\\% (\\cite{mohr99}). Observational studies of the ST relation followed by work on the \\lxt\\ and \\micmT\\ relations support a scenario where clusters exhibit regularity similar to that of elliptical galaxies in the properties measured on the scales of their virial regions, but exhibit significant irregularities in the properties of the densest, central regions where physical processes other than gravity and gas dynamics-- such as radiative cooling and magnetic fields-- play significant roles (\\cite{mohr97,arnaud99,mohr99}). The evidence indicating cluster regularity is balanced by evidence for departures from regularity; the scatter in the observed scaling relations is larger than can be accounted for by the measurement uncertainties.  This resolved scatter contains clues about, among other things, cluster peculiar velocities and departures from equilibrium.  Here we examine the ST relation at intermediate redshift ($0.19\\le z\\le0.55$) using ROSAT HRI observations of the Canadian Network for Observational Cosmology (CNOC; e.g. \\cite{yee96,lewis99}) cluster sample. The ST relation provides a potentially powerful tool to study the expansion history of the universe.  As explained in detail below, the ST relation is rather insensitive to cosmological evolution. Thus, armed with an accurate model of cluster evolution, one could use the ST relation to measure distances at intermediate redshift, constraining the deceleration parameter $q_{0}$.  We first present the X-ray ST relation in the nearby cluster sample ($\\S$\\ref{sec:localST}) and then present an interpretation of the ST relation in terms of regularity in the underlying dark matter properties of the cluster.  Section \\ref{sec:midzST} describes observations of the ST relation in intermediate redshift clusters. In $\\S$\\ref{sec:cosmo} we use these observations to constrain cosmological parameters. Section \\ref{sec:conclude}  contains a summary of our conclusions.  Throughout the paper we use $H_{0}=50h_{50}\\,\\rm{km/s/Mpc}$.  \\myputfigure{fig1.eps}{2.8}{0.45}{-20}{-10} \\figcaption{The X-ray isophotal size versus emission weighted mean ICM temperature $T_{X}$ for an X-ray flux limited sample of 45 nearby clusters.  We use the isophote $I=3.0\\times10^{-14}$~erg/s/cm$^{2}$/arcmin$^{2}$ within the cluster rest frame 0.5:2.0~keV band.  The solid line represents the best fit relation, and the RMS scatter about this line is 15\\% in size. \\label{fig:STnearby}} ",
        "conclusions": "\\label{sec:conclude}  We present the local X-ray Size--Temperature (ST) relation for an X-ray flux limited sample of 45 clusters observed with the ROSAT PSPC (\\cite{mohr97}).  We provide an explanation of this scaling relation in terms of underlying scaling relations in the cluster dark matter properties.  The observed ST relation is slightly steeper than the self--similar expectation presented in $\\S$\\ref{sec:STscaling}, but is consistent when the modest variation of the ICM mass fraction  $f_{g}$ with $T_{x}$ is taken into account (\\cite{mohr99}).  We use our theoretical model for the ST relation to explore its evolution with redshift.  Interestingly, for the typical ICM radial distribution observed in nearby and distant clusters, the normalization of the ST relation is not expected to evolve.  The lack of evolution makes the ST relation a plausible source of intermediate redshift angular diameter distances.  Of course, if cluster structure evolution is very different from the current theoretical expectation, our model ($\\S$\\ref{sec:STevolve}) will underestimate ST relation evolution. For example, a shift in the mean ICM mass fraction with redshift would bias ST relation distances; to date there is no compelling evidence that distant clusters have different ICM mass fractions in the mean than nearby clusters (e.g. \\cite{lewis99,grego00}).  We use ROSAT HRI observations of 11 CNOC clusters with measured emission weighted mean ICM temperatures \\TX\\ to make the first measurements of the intermediate redshift ST relation.  Because of the poor image quality, we measure the angular isophotal size $\\theta_{I}$ using the best fit circular $\\beta$ model, rather than measuring it nonparametrically as for the nearby clusters (see Eqn~\\ref{eq:size}). By assuming the cosmological parameters $\\Omega_{M}=0.3$ \\& $\\Omega_{\\Lambda}=0.7$, we show that the slope and zeropoint of this intermediate redshift ST relation is statistically consistent with that of the local ST relation (see Fig~\\ref{fig:STmidz}).  In addition, we examine the relation for $\\Omega_{M}=1$ \\& $\\Omega_{\\Lambda}=0$ (see Fig~\\ref{fig:STmidzSCDM}), showing that although the slope is consistent, the zeropoint is different at greater than 3$\\sigma$ significance.  \\myputfigure{fig5.eps}{2.8}{0.45}{-25}{-10} \\figcaption{Cosmological constraints from X-ray ST relation distance measurements to 11 intermediate redshift clusters.  Contours correspond to 1, 2 and 3 $\\sigma$ confidence regions ($\\Delta\\chi^{2}=2.3$, 6.2 and 11.8).  The solid line marks spatially flat models, and the stars mark models $\\Omega_{M}=0.3$ \\& $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$ and  $\\Omega_{M}=1$. \\label{fig:cosmo}\\vskip5pt}  Finally, we use this cluster sample and our ST relation evolution model to place cosmological constraints.  Given the quality of the cluster images and temperature measurements, it is not surprising that a wide range of cosmological models is consistent with the data. Nevertheless, this sample of 11 intermediate redshift distances is sufficient to rule out $\\Omega_{M}=1$ with between 2$\\sigma$ and 3$\\sigma$ confidence.  Taken together with ICM mass fraction constraints on the cosmological matter density parameter $\\Omega_{M}<0.44$ at 95\\% confidence (\\cite{mohr99}), the cluster ST relation exhibits a slight preference for universes with $\\Omega_{\\Lambda}>0$; models with $\\Omega_{\\Lambda}=0$ are inconsistent with the ST relation at between 1 and 2$\\sigma$.  When considering only models where $\\Omega_{M}+\\Omega_{\\Lambda}=1$, we can rule out $\\Omega_{M}=1$ with 3$\\sigma$ confidence.  With the higher quality X-ray images and ICM temperature measurements available from Chandra and XMM, a significant tightening of these constraints and further tests of the underlying evolution model will be possible. Comparison of local and distant \\micmT\\ relations, which are more sensitive to cluster evolution, would provide important constraints on these models.   In addition, observations with a new generation of Sunyaev-Zel'dovich effect instruments (\\cite{carlstrom99,mohr99b,holder00}) will allow us to more accurately constrain the evolution of cluster structure. With these future observations of intermediate and high redshift clusters, we plan to continue using the ST relation as a tool to provide cosmological constraints independent of those derived from recent high redshift SNe Ia observations (\\cite{schmidt98,perlmutter99})."
    },
    "0004/astro-ph0004148_arXiv.txt": {
        "abstract": "We investigate the rich cluster Abell~2029 ($z \\sim 0.08$) using optical imaging and long-slit spectral observations of 52 disk galaxies distributed throughout the cluster field.  No strong emission-line galaxies are present within $\\sim 400$~kpc of the cluster center, a region largely dominated by the similarly-shaped X-ray and low surface brightness optical envelopes centered on the giant cD galaxy.  However, two-thirds of the galaxies observed outside the cluster core exhibit line emission.  \\hal\\ rotation curves of 14~cluster members are used in conjunction with a deep $I$~band image to study the environmental dependence of the Tully-Fisher relation.  The Tully-Fisher zero-point of Abell~2029 matches that of clusters at lower redshifts, although we do observe a relatively larger scatter about the Tully-Fisher relation.  We do not observe any systematic variation in the data with projected distance to the cluster center:  we see no environmental dependence of Tully-Fisher residuals, $R-I$ color, \\hal\\ equivalent width, and the shape and extent of the rotation curves. ",
        "introduction": "\\label{sec:data}  \\subsection{Optical Spectroscopy}  We obtained long-slit spectroscopy to derive optical rotation curves of galaxies in Abell~2029.  The observations were carried out at the Mt.~Palomar 5~m telescope during the nights of 1998 April 27--29.  We used the red camera of the Double Spectrograph (Oke and Gunn 1982) to observe the \\hal\\ (6563~\\AA), \\NII\\ (6548, 6584 \\AA), and \\SII\\ (6717, 6731 \\AA) emission lines.  The spatial scale of CCD21 (1024$^2$) was 0$\\farcs$468 pixel$^{-1}$. The combination of the 1200~lines~mm$^{-1}$ grating and a 2\\arcsec\\ wide slit yielded a dispersion of 0.65~\\AA~pixel$^{-1}$ and a spectral resolution of 1.7~\\AA\\ (equivalent to 75 \\kms\\ at 6800 \\AA).  The grating angle allowed us to observe \\hal\\ in galaxies with recessional velocities between 7,600 and 38,100~\\kms.  We were fortunate to enjoy extremely mild atmospheric conditions at Mt.~Palomar.  All three nights were photometric and dark.  The seeing was remarkably sharper and more stable than typically encountered at the site; we estimate the median seeing to have been 1\\arcsec, but at times the seeing dropped to 0$\\farcs$6.  Such excellent spatial resolution is important to obtain high sensitivity rotation curves at the redshifts of the target galaxies.  Besides yielding higher signal-to-noise per pixel, a sharper seeing also allows a more accurate placement of the slit.  This is important because slit offsets and incorrect estimations of the position angles of galaxy disks can lead to serious errors in the inferred velocity widths (Bershady 1998, Giovanelli et al. 2000; hereafter G00).  We did not obtain absolute flux calibrations as they were not necessary for the purpose of this paper.  We used deep $R$ and $I$ band images to select candidate galaxies as well as to estimate their position angles.  We discuss these data in the next section. We observed all probable disk-like systems on the reference images that might be members of the cluster, did not appear to be face-on, and were free of contamination from foreground stars.  The limited resolution of the reference images precluded unambiguous identification of appropriate Tully-Fisher candidates.  Our observing strategy began with a five minute test-exposure on each spectroscopic target.  That way we were able to estimate ``on the fly'' the exposure time required in order to sample adequately the outer disk regions.  Furthermore, the test exposure determined whether the galaxy was even useful to our work; a galaxy may lie in the foreground or background of the cluster or it may contain little or no \\hal\\ emission.  If the observation was deemed useful, a second exposure typically ranged between 15 and 45~minutes.  We detected line emission in half of the 52~observed galaxies.  We list the galaxies observed in Table~1, sorting the entries by Right Ascension.  The table contains:  \\noindent Col. 1: Identification names corresponding to a coding number in our database, referred to as the Arecibo General Catalog.  \\noindent Cols. 2 and 3: Right Ascension and Declination in the 1950.0 epoch. Coordinates have been obtained from the Digitized Sky Survey catalog and are accurate to $<$ 2\\arcsec.  \\noindent Col. 4: The galaxy radial velocity as measured in the heliocentric reference frame.  The redshift measurements of the galaxies without emission lines were obtained from the NED\\footnote{The NASA/IPAC Extragalactic Database is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.} database.  They have been previously derived by others using absorption-line spectra.  \\noindent Col. 5: An indication of the usefulness of the optical emission lines in order to apply the Tully-Fisher relation: 0=no lines present; 1=strong emission lines throughout much of the disk; 2=weak or nuclear emission only.  Rotation curves are extracted as discussed in Dale et al. (1997 and 1998; hereafter D97 and D98).  We use the \\hal\\ emission line to map the rotation curve except in the case of the galaxy AGC 251909 where the emission of the \\NII\\ line (6584~\\AA) extends to a larger distance than that of the \\hal\\ emission.  We center the rotation curve kinematically  by assigning the velocity nearest to the average of the 10\\% and 90\\% velocities to be at radius zero, where an N\\% velocity is greater than N\\% of the velocity data points in the rotation curve.  The average of the 10\\% and 90\\% velocities is taken to be the galaxy's recessional velocity.  We define the observed rotational velocity width to be \\Wobs\\ $\\equiv V_{\\rm 90\\%} - V_{\\rm 10\\%}$. We filled-in small portions of the \\hal\\ rotation curve of two galaxies (AGC 251913 and AGC 251912) using data from the \\NII\\ rotation curve in order to provide information on the shape of the inner parts and to ensure consistent estimates of \\Wobs.  The rotation curves in our sample vary in physical extent, and more importantly, they do not all reach the optical radius, \\Ropt, the distance along the major axis to the isophote containing 83\\% of the $I$~band flux. This radius is reported by Persic \\& Salucci (1991) and G00 to be the most useful radius at which to measure the velocity width of rotation curves.  We have extrapolated the rotation curves, and hence made adjustments to \\Wobs, when they did not reach \\Ropt.  The resulting correction, ${\\Delta}_{\\rm sh}$, depends on the shape of the rotation curve and only exceeded 4\\% for AGC 251831 where the correction was large ($\\sim 44$\\%).  To recover the actual velocity widths, a few more corrections are necessary. The first is the factor 1/sin$i$ to convert the width observed when a disk is inclined to the line of sight at an angle $i$ to what would be observed if the disk were edge--on, and the second is the factor 1/(1+$z$) to correct the cosmological broadening of $W$.  A final correction, $f_{\\rm slit} < 1.05$, accounts for the finite width of the slit of the spectrograph (G00).  The corrected optical rotational velocity width is  \\be W_{\\rm cor} = {{W_{\\rm obs} + {\\Delta}_{\\rm sh}} \\over {(1+z)\\sin i}} f_{\\rm slit}. \\ee A discussion of the errors in the velocity widths can be found in D97.  Figure \\ref{fig:RCs} is a display of the rotation curves observed in the field of Abell~2029.  Entries in the figure are sorted by Right Ascension. The name of the galaxy is given along with the CMB radial velocity.  Two dashed lines are drawn: the vertical line is at \\Ropt; the horizontal line indicates the adopted half-velocity width, $W$/2, which in some cases arises from an extrapolation to the rotation curve (see Table 1).  Overlayed are the fits used to infer $W$(\\Ropt).  Details of the fitting procedure can be found in G00.  The error bars include both the uncertainty in the wavelength calibration and the routine used to fit the rotation curve.  Notice that the data are highly correlated due to seeing and guiding jitter.  This is properly taken into account by the fitting routines (see D97 and references therein for details). \\centerline{\\psfig{figure=Dale.fig1a.ps,width=7.0in,bbllx=67pt,bblly=144pt,bburx=515pt,bbury=714pt}} \\begin{figure}[ht] \\centerline{\\psfig{figure=Dale.fig1b.ps,width=7.0in,bbllx=67pt,bblly=485pt,bburx=515pt,bbury=714pt}} \\caption[] {\\ The kinematically folded rotation curves (see text).  The error bars include both the uncertainty in the wavelength calibration and the rotation curve fitting routine used.  Names of the galaxy are given along with the CMB radial velocities.  Two dashed lines are drawn: the horizontal line indicates the adopted half velocity width, $W$/2, which in some cases arises from an extrapolation to the RC; the vertical line is at \\Ropt, the radius containing 83\\% of the $I$ band flux.  A fit to the rotation curve is indicated by a solid line.  Note that the rotation curves are {\\it not} deprojected to an edge-on orientation.} \\label{fig:RCs} \\end{figure}  We list in Table~2 the complete set of spectroscopic data corresponding to the 21~galaxies for which we obtained useful rotation curves, sorting the entries by Right Ascension.  The parameters listed in Table~2 are:  \\noindent Col. 1: Identification names corresponding to a coding number in the Arecibo General Catalog.  \\noindent Col. 2: The spectral exposure time in seconds.  \\noindent Col. 3: The recessional velocity of the galaxy in the CMB reference frame, assuming a Sun-CMB relative velocity of 369.5 \\kms\\ towards $(l,b)=264.4^\\circ,48.4^\\circ)$ (Kogut et al. 1993).  Errors are parenthesized: e.g. 13241(08) means 13241$\\pm$08.  \\noindent Col. 4: The observed velocity width in \\kms.  \\noindent Col. 5: The velocity width in \\kms\\ at \\Ropt\\ after correcting for the shape of the rotation curve, the cosmological broadening, and the smearing effects due to the finite width of the slit of the spectrograph.  \\noindent Col. 6: The corrected velocity width in \\kms\\ converted to an edge-on perspective.  \\noindent Col. 7: The adopted inclination of the plane of the disk to the line of sight, $i$, expressed in degrees, (90$^\\circ$ corresponds to an edge--on perspective); the derivation of $i$ and its associated uncertainty are discussed in Section~4 of D97.  \\noindent Col. 8: The logarithm in base~10 of the corrected velocity width (value in column 6), together with its estimated uncertainty in parentheses.  The uncertainty takes into account both measurement errors and uncertainties arising from the corrections.  The format 2.576(22), for example, is equivalent to 2.576$\\pm$0.022.  \\subsection{Optical Imaging}  $I$~band photometry of Abell~2029 was obtained for a different project by one of us (JMU) in collaboration with S. P. Boughn (Haverford) with the 0.9~m telescope on Kitt Peak National Observatory on 1998~April~19.  They used the T2KA camera mounted at the f:7.5 Cassegrain focus which resulted in square pixels, 0$\\farcs$68 on a side.  The seeing was excellent, between 0$\\farcs$7 and 0$\\farcs$9, which resulted in an effective seeing of $\\sim1\\farcs$2 due to the available pixel size.  Two sets of fifteen partially overlapping frames were used to form a mosaic of about 35$\\arcmin$ (RA) by 58$\\arcmin$ (Dec).  The central three by three mosaic has overlaps of about 3/4 of a frame between immediately adjacent frames, whereas the outer three frames to the north and south overlap by about 1/2 of a frame with the closest of the central ones.  All but four of the outlying frames to the north and south were obtained with air masses between 1.11 and 1.25.  The exposures lasted five minutes.  The data were processed as discussed in Uson, Boughn and Kuhn (1991, hereafter UBK91).  All frames were used to generate a ``sky-flat'' gain calibration frame.  Since the cluster contains a diffuse halo that surrounds the central galaxy, a 12$\\arcmin$ by 12$\\arcmin$ area centered on the cluster was blanked on all the relevant frames before using them to generate the sky flat as discussed in UBK91.  The calibrated frames were used to determine the secant-law extinction which had a slope of 0.06 mag/airmass.  Absolute calibration was done using stars from Landolt's $UBVRI$ secondary calibration list (Landolt 1983).  Details will be given elsewhere.  The $R$~band photometry was obtained from UBK91.  Flux estimation follows from the data reduction methods discussed in D97 and D98 using both standard and customized IRAF\\footnote{IRAF (Image Reduction and Analysis Facility) is distributed by NOAO.} packages.  We will only mention here that the measured fluxes, denoted $m_{\\rm obs}$, include extrapolations of the exponential fits to the surface brightness profiles to eight disk scale lengths and are typically accurate to \\about\\ 0.03 mag (uncertainties at least as large are later included after making corrections for internal extinction).  We apply some corrections to $m_{\\rm obs}$ in order to obtain the final $I$ band fluxes: \\be m_I = m_{\\rm obs} - A_I + k_I - \\Delta m_{\\rm int}. \\ee For the Galactic extinction correction $A_I$, we use the recent work of Schlegel, Finkbeiner, and Davis (1998) who have provided accurate Galactic reddening estimates using COBE/DIRBE and IRAS/ISSA dust maps.  The internal extinction correction, $\\Delta m_{\\rm int}$, is applied using the procedure outlined in G97, \\be \\Delta m_{\\rm int} = - f(T) \\; \\gamma(W_{\\rm cor}) \\; \\log(1-e), \\ee where $\\gamma$ ($\\lesssim$ 1.0) depends on the corrected velocity width $W_{\\rm cor}$ and $e$ is the ellipticity of the spiral disk, corrected for atmospheric seeing effects as described in Section 5 of D97 (the adopted correction $\\Delta m_{\\rm int}$ is slightly smaller for early, less dusty galaxies: $f(T)$=0.85 for types $T$ earlier than Sbc; $f(T)$=1 otherwise). We apply a cosmological k-correction according to Han (1992): $k_I = (0.5876 - 0.1658$T$)z$.  The relevant photometric data are listed in Table~3 with the first column matching that of Table~2.  The remaining parameters are:  \\noindent Col. 2: Morphological type code in the RC3 scheme, where code 1 corresponds to Sa's, code 3 to Sb's, code 5 to Sc's and so on.  We assign these codes after visually inspecting the CCD $I$ band images and after noting the value of $R_{\\rm 75}/R_{\\rm 25}$, where $R_X$ is the radius containing X\\% of the $I$ band flux.  This ratio is a measure of the central concentration of the flux which was computed for a variety of bulge--to--disk ratios.  Given the limited resolution of the images, some of the inferred types are rather uncertain; uncertain types are followed by a colon.  \\noindent Col. 3: The angular distance $\\theta$ in arcminutes from the center of each cluster.  \\noindent Col. 4: Position angle used for spectrograph slit positioning (North: 0$^{\\circ}$, East: 90$^{\\circ}$).  \\noindent Col. 5: Ellipticity of the disk corrected for seeing effects as described in Sec. 5 of D97, along with its corresponding uncertainty expressed using the same convention as in Table~2.  \\noindent Col. 6: The (exponential) disk scale length in arcseconds.  \\noindent Col. 7: The distance along the major axis to the isophote containing 83\\% of the $I$ band flux.  \\noindent Col. 8: The measured $I$ band magnitude, extrapolated to 8 disk scale lengths assuming that the surface brightness profile of the disk is well described by an exponential function.  \\noindent Col. 9: The absolute magnitude, computed assuming that the galaxy is at the distance indicated by the cluster redshift, or by assuming the galaxy is at the distance indicated by the redshift if the galaxy is not deemed to be a member of the cluster.  The calculation assumes $H_\\circ = 100h$ \\kms\\ Mpc$^{-1}$, so the value listed is strictly $M_I - 5\\log h$.  This parameter is calculated after expressing the redshift in the CMB frame and neglecting any peculiar motion.  The uncertainty on the magnitude, parenthetically included in hundredths of a mag, is the sum in quadrature of the measurement errors and the estimate of the uncertainty in the corrections applied to the measured parameter.  \\noindent Col. 10: The difference in the $R$ and $I$ band magnitudes.  When an asterisk appears at the end of the line, we include a detailed comment on that particular object.  Because of the extent of these comments we have not appended them to the table but have included them in the text. Note that a record is flagged in both Tables~2 and~3, whether the comments refer to the photometry, to the spectroscopy, or to both.  \\small \\noindent 251826: Background galaxy.\\\\ \\noindent 251827: Background galaxy; uncertain PA.\\\\ \\noindent 251831: Rising rotation curve; large rotation curve extrapolation.\\\\ \\noindent 251909: Background galaxy; \\NII\\ rotation curve used.\\\\ \\noindent 251913: \\NII\\ patch for radii $<$ 3\\arcsec; uncertain disk ellipticity.\\\\ \\noindent 251838: Asymmetric $I$ band profile.\\\\ \\noindent 251843: Note low $i$.\\\\ \\noindent 251874: Foreground galaxy; 5 minute integration.\\\\ \\noindent 251895: Uncertain disk ellipticity; asymmetric $I$ band profile.\\\\ \\noindent 251897: Center-of-light used for rotation curve spatial and kinematic center.\\\\ \\noindent 251900: Foreground galaxy; 5 minute integration; center-of-light used to determine the center of the rotation curve.\\\\ \\noindent 251908: Tidally interacting with small companion 13\\arcsec\\ to NE (\\S \\ref{sec:tidal}).\\\\ \\noindent 251912: Foreground galaxy; \\NII\\ patch for radii $<$ 3\\arcsec; flux disentanglement with AGC 251911 difficult.\\\\ \\noindent 251911: Flux disentanglement with AGC 251912 difficult; uncertain PA and disk ellipticity.\\\\ \\noindent 251915: Foreground galaxy.\\\\ \\normalsize ",
        "conclusions": "The extended optical envelope centered on the cD galaxy in Abell 2029 was interpreted by UBK91 as the leftover of a period of violent relaxation during the initial cluster collapse. This process would have led as well to the ionization of the gas present in the galaxies that participated in such a collapse and to the hot intracluster gas, observable through its X-ray emission and strongly peaked at the cluster center.  The galaxies that traverse the cluster core will have their interstellar gas stripped by the dense intracluster gas, and is likely the root cause of the lack of emission line signatures from these galaxies.  Indeed, UBK91 show the extremely large extent of the $R$~band profile which they trace to \\about 5$\\arcmin$, a boundary similar to the delineation between emission-poor and relatively emission-rich environs.  Beyond this distance, however, it appears that the projected distance from the cluster core at which a cluster galaxy lies plays only a small role in the \\hal\\ success rate: galaxies with emission lines are spatially evenly mixed with those galaxies lacking emission lines.  A goal of this project was to search for signs of tidal stripping on inner-core galaxies.  Therefore, we observed all disk systems in the cluster core, irrespective of whether or not the galaxy appeared to be a useful ``Tully-Fisher'' galaxy.  This could mean that we inspected a large fraction of early type spirals (S0 to Sab) in the cluster core.  In the Coma cluster, for example, S0 galaxies outnumber spirals in the cluster center by more than a factor of two (Andreon \\& Davoust 1997).  We find that the average (RC3) morphological type for the emission line galaxies that we observed in Abell~2029 is $T=4.7$ (Sbc/Sc),  whereas it is $T=3.2$ (Sb/Sbc) for the non-emission line galaxies.   There is a small difference in the average morphological type of these  two groups, and thus the discrepancy in the spatial distribution of emission line galaxies may not be quite so surprising, as early-type disk galaxies contain fewer \\HII\\ regions than are typically found in late-type spirals.   Outside the cluster core, we find no trends in $R-I$ color, \\hal\\ equivalent width, the shape and physical extent of a rotation curve, nor Tully-Fisher characteristics as a function of projected distance from the cluster center. The latter result is in agreement with the Tully-Fisher work of Biviano et al.~(1990) and G97 at relatively low redshifts, as well as with the work at $0.2<z<0.6$ described in D99.  The dispersion in the Tully-Fisher plot for Abell~2029 is about 50\\% larger than what is seen in clusters at lower redshifts (0.56 mag vs. $0.35-0.38$ mag). The dispersion does not significantly change if each galaxy is placed at the distance indicated by its individual redshift, rather than assuming that all cluster galaxies are at the same distance. Because of the relative difficulty in assessing such galactic properties as disk inclination and major axis position angle, such an increase is not surprising given the larger measurement uncertainties. However, it is also possible that the strong intracluster medium environment of Abell~2029 promotes a wider variety of $M/L$ ratios than what is typically found in other clusters.  In other words, rich clusters of galaxies may have a higher intrinsic contribution to the overall Tully-Fisher scatter.  Using their sample of 522 late-type galaxies in 52 lower redshift Abell clusters, D99 find an intrinsic scatter of 0.25 magnitudes; measurement uncertainties account for the remaining 0.28 magnitudes of the overall scatter.  Similar numbers have been found for the Tully-Fisher samples of G97 and Willick~(1999).  In Abell 2029, we find that measurement uncertainties are responsible for 0.40 magnitudes of the scatter, and thus we infer an intrinsic scatter of 0.39 magnitudes.  In short, the relatively large Tully-Fisher dispersion for Abell 2029 appears to be a consequence of both increased measurement uncertainties and relatively larger variations in the spiral galaxy population.  Finally, the overall offset of the Abell~2029 Tully-Fisher data, with respect to the fiducial template relation for lower $z$ clusters, is indistinguishable from a null offset (at least to within the 0.15~magnitude accuracy of the Abell~2029 zero-point estimation and ignoring the small Malmquist bias correction).  If offsets are interpreted as due to peculiar motions, the Abell~2029 offset implies the cluster is essentially at rest in the CMB frame ($30 \\pm 1600$ \\kms).  Alternatively, this negligible offset may be construed as an indication that little evolution in the Tully-Fisher relation has taken place in clusters up to $z$ \\about 0.08.  However, a much larger sample of clusters at higher redshift is needed to make such a claim statistically significant.  In any case, this null result gives us great confidence in the soundness of our data reduction procedures as it results from the combination of spectroscopic and photometric data that were obtained and reduced independently for different projects."
    },
    "0004/astro-ph0004195_arXiv.txt": {
        "abstract": "We investigate the effect of a hybrid electron population, consisting of both thermal and non-thermal particles, on the synchrotron spectrum, image size, and image shape of a hot accretion flow onto a supermassive black hole. We find two universal features in the emitted synchrotron spectrum: (i) a prominent shoulder at low $(\\lesssim 10^{11}~\\rm{Hz})$ frequencies that is weakly dependent on the shape of the electron energy distribution, and (ii) an extended tail of emission at high $(\\gtrsim 10^{13}~\\rm{Hz})$ frequencies whose spectral slope depends on the slope of the power-law energy distribution of the electrons. In the low-frequency shoulder, the luminosity can be up to two orders of magnitude greater than with a purely thermal plasma even if only a small fraction $(< 1\\%)$ of the steady-state electron energy is in the non-thermal electrons.  We apply the hybrid model to the Galactic center source, Sgr~A$^*$. The observed radio and IR spectra imply that at most $1\\%$ of the steady-state electron energy is present in a power-law tail in this source. This corresponds to no more than $10\\%$ of the electron energy injected into the non-thermal electrons and hence $90 \\%$ into the thermal electrons. We show that such a hybrid distribution can be sustained in the flow because thermalization via Coulomb collisions and synchrotron self-absorption are both inefficient.  The presence of non-thermal electrons enlarges the size of the radio image at low frequencies and alters the frequency dependence of the brightness temperature. A purely thermal electron distributions produces a sharp-edged image while a hybrid distribution causes strong limb brightening. These effects can be seen up to frequencies $\\sim 10^{11}$~Hz and are accessible to radio interferometers. ",
        "introduction": "The mechanisms of particle heating and acceleration, and the emission spectra from the resulting particle energy distributions, are of great importance in the theory of collisionless hot accretion flows onto compact objects.  Discussions in the literature have focused on the physics of electron heating and acceleration (Begelman \\& Chiueh 1988; Bisnovatyi-Kogan \\& Lovelace 1997; Quataert \\& Gruzinov 1999; Gruzinov \\& Quataert 1999, Medvedev 2000), the efficiency of particle thermalization (Ghisellini, Guilbert, \\& Svensson 1988; Ghisellini, Haardt, \\& Fabian 1993; Mahadevan \\& Quataert 1997; Ghisellini, Haardt, \\& Svensson 1998; Nayakshin \\& Melia 1998), and the generation of hybrid thermal-nonthermal electron energy distributions in these plasmas (see, e.g., Coppi 1999 and references therein).  Despite a substantial amount of work, many issues remain unresolved, primarily because of our incomplete understanding of physical processes such as magnetic reconnection, MHD turbulence, and collective plasma modes.  These questions are especially relevant for an optically-thin advection-dominated accretion flow (ADAF). An ADAF is an example of a hot, rarefied, magnetic plasmas with low radiative efficiency (Ichimaru 1977; Narayan \\& Yi 1994, 1995b; Abramowicz et al. 1995; see Narayan, Mahadevan, \\& Quataert 1998b and Kato, Fukue, \\& Mineshige 1998 for reviews). A basic property of the nearly collisionless plasma in an ADAF is that the Coulomb coupling between the electrons and ions is weak so that energy transfer from the ions to the electrons is inefficient. In addition, it is commonly assumed that the viscously generated energy primarily heats the heavier species, the ions, and that the electrons retain a thermal distribution throughout the flow (e.g., Narayan \\& Yi 1995b; Mahadevan 1997).  There are, however, processes such as MHD turbulence, pair production (e.g., through pion decay), and electron-proton coupling which can both heat the electrons and generate non-thermal distributions. Quataert and Gruzinov (1999; see also Gruzinov \\& Quataert 1999) considered two processes specific to MHD turbulence that accelerate particles in magnetic collisionless plasmas: Landau damping by electric fields parallel to the local magnetic field and transit-time damping by time-varying magnetic fields. They found that the assumption of negligible electron heating/acceleration is valid only for weak magnetic fields, i.e., when the ratio of the gas pressure to total pressure $\\beta_{\\rm{ADAF}}$ is larger than a critical value $\\beta_{\\rm{crit}}$. The value of $\\beta_{\\rm{crit}}$ is very uncertain and is around 0.9. For $\\beta_{\\rm{ADAF}} \\gtrsim 0.9$, turbulence primarily accelerates the protons, while for stronger magnetic fields, the results are inconclusive. Shocks and pion decay can also lead to non-thermal electrons in the accretion flow.  Several processes, such as Coulomb collisions and synchrotron self-absorption, can potentially lead to thermalization of particles in accretion flows (Mahadevan \\& Quataert 1997; Ghisellini et al. 1998).  Mahadevan \\& Quataert (1997) showed that Coulomb collisions are ineffective in thermalizing the electrons in an ADAF. However, they argued that, for sufficiently high mass-accretion rates, the electrons in the plasma can be thermalized by synchrotron self-absorption. Nayakshin and Melia (1998) showed that considerable deviations from a Maxwell-Boltzmann distribution can be sustained in a plasma with low source compactness when Coulomb collisions, Comptonization, and pair processes are taken into account.  In view of the difficulty of calculating the heating, cooling and thermalization of particles from first principles, many authors have restricted their models to either purely Maxwellian or purely non-thermal (extended power laws or monoenergetic) electron distributions. Only recently have there been attempts towards explaining spectra of accreting black holes with models including hybrid thermal/non-thermal distributions of electrons. Some models have physically motivated distributions, such as non-thermal electrons produced by decaying pions (Mahadevan 1998), whereas others invoke more {\\it ad hoc} distributions to fit the data (e.g., Beckert \\& Duschl 1997; Falcke \\& Biermann 1999). The reverse process of trying to constrain the energy distributions of particles in accreting plasmas by comparing models to data, however, faces issues of uniqueness which can be addressed only by a more comprehensive study of hybrid models.  In this paper, we consider generalized electron distributions consisting of a dominant Maxwellian plus a small non-thermal power-law component of varying slope and energy content, and study the synchrotron emission from the resulting hybrid plasmas. We identify the characteristic signatures of the non-thermal electrons on the emitted radio synchrotron spectrum of an accretion flow and on its image as observed with a radio telescope. We also discuss to what extent the observed effects could be used to determine the details of the underlying non-thermal electron distribution. This work is relevant for interpreting observations of low luminosity AGNs such as the Galactic Center source, Sagittarius~A$^*$ (Sgr~A$^*$).  In $\\S 2$ we review the basic properties of ADAFs, hybrid plasmas, synchrotron radiation, and radiative transfer. In $\\S 3$ we present a series of models of supermassive black holes with low accretion rates. We apply these models to Sgr~A$^*$ and derive constraints on the fraction of the electron energy that can be present in a non-thermal form.  In $\\S 4$, we study the correspondence between the energy distribution of the electrons and the resulting spectrum. In $\\S5$ we study the energetics of a hybrid flow and calculate the heating, cooling and thermalization rates of the non-thermal electrons. We summarize our conclusions in $\\S 6$. We present in an Appendix approximate analytic expressions for the contribution of a non-thermal particle population to the synchrotron spectrum of an accretion flow. ",
        "conclusions": "In this paper we considered hot accretion flows around supermassive black holes, using the ADAF model as a typical example. Such hot flows are expected to occur at low mass-accretion rates. We assume that a fraction of the viscous dissipation energy in the accretion flow goes into accelerating electrons to a non-thermal power-law distribution. We find that a power-law tail of high energy electrons can be sustained in such flows; neither Coulomb collisions nor synchrotron self-absorption is able to thermalize the power-law electrons. Assuming that there are no other thermalizing mechanisms (e.g., collective plasma modes), we calculate the resulting hybrid thermal/non-thermal spectrum from such a plasma.  The presence of even a population of non-thermal electrons gives rise to two universal and prominent features in the synchrotron spectrum: a low-frequency shoulder and a high-frequency tail (Fig.~2).  These features were identified by Mahadevan (1998) who considered a specific mechanism (via pion decay) for the production of power-law electrons. Even if only a small fraction of the total steady-state electron energy is in the non-thermal power-law component, we find that there is significant super-thermal emission in the low-frequency shoulder and the high frequency tail of the spectrum.  Furthermore, each of these universal features can be associated with a specific range of Lorentz factors of the emitting electrons. The low-frequency shoulder is emitted by electrons with Lorentz factors $\\log \\gamma = 1-2$, while the high-energy tail is emitted by electrons with $\\log \\gamma \\gtrsim 2$ (Fig.8).  Since the power-law electrons cause significant emission at low frequencies, comparing this low-frequency shoulder to data can provide stringent constraints on the fraction of the electron energy that is present in a non-thermal population. In the case of Sgr~A$^*$ (Fig.~4), we conclude that, in steady-state, electrons even with intermediate Lorentz factors ($\\gamma$ as low as 10) make up only a very small fraction of the electron population.  Since Coulomb collisions and synchrotron self-absorption are ineffective in thermalizing these flows, this means that either the energy injected into a non-thermal population, $\\delta_{\\rm{nt}} \\dot{E}_{\\rm{visc}}$, is small as in the calculations described in $\\S 5$ or there is some other efficient thermalization mechanism which plays a role in these flows. The study on Sgr~A$^*$ needs to be extended to other sources before we can judge how universal the result is.  Here it is important to emphasize two points regarding the constraints on non-thermal electrons obtained from the spectrum.  First, for the part of the spectrum below the thermal peak, there is a mapping between each frequency in the spectrum and the thin shell of radii which predominantly emits at that frequency. This is due to the steep radial dependence of the synchrotron absorption that causes a sharp transition between optically thick and optically thin regions in the flow. Therefore, emission at each frequency comes primarily from a specific radius in the flow and can be used to probe the non-thermal content in a shell that extends from that radius out to a somewhat larger radius. (Mathematically, radiation emerging at each frequency comes from the entire range from the radius where the optical depth is unity out to infinity, but because the total emissivity also drops rapidly with radius, the biggest contribution to the intensity integral comes from a relatively narrow shell).  As a result, the constraints on the non-thermal distribution obtained from each frequency of the low-frequency shoulder apply only to electrons within this shell.  Inside the radius corresponding to optical depth unity, the emission is heavily self-absorbed and the electron distribution there is inaccessible to observations.  Second, the constraints on the non-thermal energy content which we have obtained for Sgr~A$^*$ would not be invalidated if the observed low-frequency shoulder originated in a separate, extended region such as a jet or an outflow exterior to the accretion flow. In that case, the observed fluxes may be treated as upper limits to the emission from the inner accretion flow. This would imply that the non-thermal emission from the accretion flow is even smaller than in our models and the resulting constraints on the fraction of non-thermal electrons would become only tighter.  When the non-thermal electron distribution is a power law, three parameters of the model and three corresponding pieces of physics can be extracted by comparing hybrid emission models to spectral data. The first parameter is $\\eta$, or equivalently $\\delta_{\\rm{nt}}$ (cf.\\S5), which measures the fraction of the dissipated energy that goes into the non-thermal electrons.  The second parameter is $p$, the slope of the power-law energy distribution (eq.~4).  This can be uniquely determined by measuring the spectral slope of the high-frequency tail in the spectrum (provided there is no competing emission via Compton scattering at these wavelengths, cf. Fig. 1).  The value of $p$ may provide an indication of the mechanism by which the electrons in a hot accretion flow are accelerated into a power-law tail. The third parameter, which is more significant for distributions with $p<3$, is $\\gamma_{\\rm{max}}$, the maximum Lorentz factor to which electrons are accelerated. If this parameter can be determined by an IR cutoff in the spectrum, it would provide information on the ratio of the acceleration timescale to the synchrotron cooling timescale.  Due to the degeneracy of the low-frequency shoulder to the parameters $\\eta$ and $p$ (Fig.~4), this segment of the spectrum by itself is not sufficient to constrain the two parameters.  However, $\\eta$ and $p$ may be determined uniquely if we have information on both the low-frequency shoulder and the high-frequency tail.  For this to work, there should be no contamination from a jet or outflow to the observed radiation in the low-frequency shoulder or from Compton scattering in the high frequency tail.  We showed in this paper that the presence of a non-thermal population of electrons also has measurable effects on the shape of the image of the source, and the size of the accretion flow, as a function of frequency (Figs.~5-7). Because the frequency dependence of the image size is strongly affected by the fraction of non-thermal electrons in the flow, measuring the size of accretion flows at multiple wavelengths can provide a quantitative constraint on the non-thermal electron energy content.  Moreover, the brightness temperature of the source as a whole, as well as the variation of the brightness temperature across the source, at long wavelengths behave differently for a hybrid plasma compared to a purely thermal population of electrons.  Finally, at the frequencies corresponding to the low-frequency shoulder, we show that there would be limb brightening of the image.  All these image-related signatures may help to identify the presence of a non-thermal population of electrons, and the parameters of the corresponding energy distribution."
    },
    "0004/astro-ph0004256_arXiv.txt": {
        "abstract": "So far all relativistically boosted jets with superluminal motion have only been detected in typical radio galaxies with early type host galaxies. We have now discovered superluminal motion in the Seyfert I galaxy III~Zw~2, classified as a spiral. Superluminal motion was first inferred from the spectral evolution of the source and then confirmed by VLBI observations. The lower limit for the apparent expansion speed is $1.25\\pm0.09$~c. The fact that the spectral and spatial evolution are closely linked demonstrates that we are dealing with real physical expansion. Prior to this rapid expansion we have seen a period of virtually no expansion with an expansion speed less than 0.04 c. Since III~Zw~2 is also part of a sample of so called radio-intermediate quasars (RIQ), it confirms earlier predictions of superluminal motion for this source, based on the argument that RIQs could be relativistically boosted jets in radio-weak quasars and Seyfert galaxies. ",
        "introduction": "The radio properties of quasars with otherwise very similar optical properties can be markedly different. There is a clear dichotomy between radio-loud and radio-quiet quasars in optically selected samples. The radio-loudness is usually characterized by the radio-to-optical flux ratio. In the PG quasar sample, which is probably the best studied quasar sample in the radio and optical (Kellermann et al.~\\cite{kellermann}, Boroson \\& Green \\cite{boroson}), radio-loud and radio-weak quasars separate cleanly in two distinct populations (e.g.~Kellermann et al.~1989).  It is known that radio-loud AGN almost never reside in late type, i.e.~spiral galaxies (e.g.~Kirhakos et al.~\\cite{kirhakos}, Bahcall et al.~\\cite{bahcall}) whereas radio-quiet quasars appear both in spiral and in elliptical host galaxies.  Furthermore, all relativistically boosted jets with superluminal motion and typical blazars have been detected in early type galaxies (e.g. Scarpa et al.~\\cite{scarpa}). It is still unclear, why AGN in spiral galaxies, at the same optical luminosity as their elliptical counterparts, should not be able to produce the powerful, relativistic jets seen in radio galaxies.  However, a few sources with intermediate radio-to-optical ratios appear to be neither radio-loud nor radio-quiet.  They form a distinct subclass with very similar radio morphological and spectral properties.  They all have a compact core at VLA scales and a flat and variable spectrum in common. These properties are very similar to the ones of radio cores in radio-loud quasars, but their low radio-to-optical ratio and their low extended steep-spectrum emission is atypical for radio-loud quasars.  Miller et al.~(\\cite{miller}) and Falcke et al.~(\\cite{falcke95}, \\cite{falcke96a}\\&b) have identified a number of these sources, called ``radio-intermediate quasars'' (RIQs), and suggested that they might be relativistically boosted radio-weak quasars or ``radio-weak blazars''. This would imply that most, if not all, radio-quiet quasars also have relativistic jets. In fact, VLBI observations of radio-quiet quasars already have shown high-brightness temperature radio cores and jets (Blundell \\& Beasly \\cite{blundell}). A crucial test of the relativistic jet hypothesis is the search for apparent superluminal motion in these sources. A prime candidate for detecting this is the brightest radio source in the RIQ sample, III~Zw~2, which we discuss in this paper.  III~Zw~2 (PG 0007+106, Mrk 1501, $z=0.089$) is one of the most extremely variable radio sources and a very unusual AGN. It was discovered by Zwicky (\\cite{zwicky}), classified as a Seyfert I galaxy (e.g., Arp \\cite{arp}; Khachikian \\& Weedman \\cite{khachikian}; Osterbrock \\cite{osterbrock}), and later also included in the PG quasar sample (Schmidt \\& Green \\cite{schmidt}). The host galaxy was classified as a spiral (e.g. Hutchings \\& Campbell \\cite{hutchings_campbell}) and a spiral arm was claimed (Hutchings \\cite{hutchings}). A disk model was later confirmed by fitting of model isophotes to near-IR images (Taylor et al.~\\cite{taylor}).  The most interesting property of III~Zw~2, however,  is its extreme variability at radio and other wavelengths with at least 20-fold increases in radio flux density within 4 years (Aller et al.~\\cite{aller}). The source also shows optical (Lloyd \\cite{lloyd}) and X-ray variability (Kaastra \\& de Korte \\cite{kaastra}; Pounds \\cite{pounds}).  III~Zw~2 is a core-dominated flat-spectrum AGN with only a faint extended structure (see Unger et al.~\\cite{unger}). The weak extended radio emission and the host galaxy is quite typical for a Seyfert galaxy. Its [O{\\sc III}] luminosity is a mere factor three brighter than that of a bright Seyfert galaxy like Mrk~3 (e.g.~Alonso-Herrero et al.~\\cite{alonso}) which explains why it has been classified as either a Seyfert galaxy or a quasar. In this luminosity region a distinction between the two may not be of much significance.   Earlier VLBI observations of the source have only shown a high-brightness temperature core (Falcke et al.~\\cite{falcke96b}, Kellermann et al.~\\cite{kellermann98}) and recent Millimeter-VLBI observations by Falcke et al.~(\\cite{falcke99}) just barely resolved the source into two very compact components.  Based on its average optical-to-radio ratio of $\\sim200$ and its radio properties, it was suggested that III~Zw~2 could be a radio-intermediate quasar and the presence of superluminal motion in this source was predicted. Here we will discuss multi-epoch VLA and millimeter-VLBI observations of this source confirming this hypothesis. ",
        "conclusions": "We found a close connection between the spectral and structural evolution of the radio outburst in III~Zw~2.  While the self-absorption frequency remained constant, we observed no change on VLBA-scales. The excellent consistency of the relative sizes during the first three epochs on the few micro-arcsecond level demonstrates the good quality of our data. The quick drop in peak frequency after 1998 November marked the beginning of a strong structural change. The spectral peak dropped from 43 GHz to 15 GHz within a few months and the VLBA-maps show a rapid expansion with an apparent expansion velocity of 1.25 c. The fact that spectral and spatial evolution are so closely linked also demonstrates that we are dealing with real physical expansion.  With its faint extended radio structure and its spiral host galaxy III~Zw~2 is clearly not a radio-loud quasar, but has properties very typical of luminous Seyfert galaxies or radio-quiet quasars.  The detection of superluminal motion in this galaxy now clearly shows that we are dealing with a relativistic jet on sub-pc scales.  To our knowledge this is the first detection of superluminal motion in a spiral galaxy with a Seyfert nucleus. The maximum aspect angle of $41^{\\circ}$ from the superluminal motion (see Sec.~3) is in good agreement with the orientation based unified scheme (e.g. Antonucci \\cite{antonucci}) for AGN where Seyfert I galaxies are seen under intermediate or small aspect angles, so that the nucleus is not obscured by a dusty torus.  For the question of the nature of the radio-loud/radio-quiet dichotomy this means that radio-weak and radio-loud quasars can indeed have central engines that are in many respects very similar. Their optical properties are almost indistinguishable and both types of quasars can produce relativistic jets in their nuclei. The finding of superluminal motion supports the hypothesis of Miller et al.~(1993) and Falcke et al.~(\\cite{falcke96a}) that RIQs are relativistically boosted intrinsically radio-weak AGN.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{ca037.f3.ps}} \\caption{Component separation from model fitting of point-like components to the closure phases and amplitudes at 43 GHz. The statistic errors for the first four epoch are smaller than the symbols and should be dominated by systematic errors. The separation of the first three epochs is consistent with an expansion speed $\\le0.04~c$ (solid line). The expansion speed between the fourth and fifth epoch is $1.25\\pm0.09~c$.}  \\label{move7} \\end{figure}  While these general conclusions seem to be fairly robust, some characteristics of the outburst need to be modeled in more detail. The initial phase of the flux density rise with its millimeter-peaked spectrum and no detectable expansion perhaps has to be explained similar to the physics of Gigahertz-Peaked-Spectrum (GPS) sources, i.e.~ultra-compact hotspots pumped up and powered by a jet interacting with the interstellar medium or the torus. The rapid expansion thereafter could have marked the phase where the jet breaks free and starts to propagate relativistically into a lower-density medium.  Another explanation of the initial phase could be a fast jet moving through quasi-stationary components as proposed for sources like 4C\\,39.25 (e.g. Alberdi et al.~\\cite{alberdi}). In any case, this dramatic structural change should go together with a change of the polarization vector. This must be checked by future experiments.  The initial slow expansion has also possible implications for the interpretation of other Seyfert galaxies. For example, in observations of the Seyfert galaxies Mrk 348 and Mrk 231 (Ulvestad et al.~\\cite{ulvestad}) only sub-relativistic expansion was found. Hence, one could raise the question whether the region of relativistic expansion in these two sources is at even smaller scales or whether they just happened to be in a `slow' phase, similar to III~Zw~2 early on. If true, it is still possible that Seyfert jets are launched relativistically but are slowed down and disrupted significantly already on the sub-parsec scale. Therefore it would be important to follow the spectral evolution of these sources in more detail."
    },
    "0004/astro-ph0004060_arXiv.txt": {
        "abstract": "To elucidate the effects of the UV background radiation on the collapse of pregalactic clouds, we implement a radiation-hydrodynamical calculation, combining one-dimensional spherical hydrodynamics with an accurate treatment of the radiative transfer of ionizing photons. Both absorption and scattering of UV photons are explicitly taken into account. It turns out that a gas cloud contracting within the dark matter potential does not settle into hydrostatic equilibrium, but undergoes run-away collapse even under the presence of the external UV field. The cloud center is shown to become self-shielded against ionizing photons by radiative transfer effects before shrinking to the rotation barrier. Based on our simulation results, we further discuss the possibility of H$_2$ cooling and subsequent star formation in a run-away collapsing core. The present results are closely relevant to the survival of subgalactic Population III objects as well as to metal injection into intergalactic space. ",
        "introduction": "It is widely recognized that the Ultraviolet (UV) background radiation, inferred from the proximity effect of Ly$\\alpha$ absorption lines in QSO spectra (e.g., Bajtlik, Duncan \\& Ostriker 1988; Bechtold 1994; Giallongo et al. 1996), is likely to exert a significant influence upon the collapse of pregalactic clouds and consequently the formation of galaxies. Several authors have argued that the formation of subgalactic objects is suppressed via photoionization and photoheating caused by the UV background (Umemura \\& Ikeuchi 1984; Ikeuchi 1986; Rees 1986; Bond, Szalay \\& Silk 1988; Efstathiou 1992; Babul \\& Rees 1992; Zhang, Anninos \\& Norman 1995; Thoul \\& Weinberg 1996). Final states of the photoionized clouds, however, are still unclear, because most of these studies assume optically-thin media, which precludes us from the correct assessment of self-shielding against the external UV fields.  Self-shielding not only plays an important role in the thermal and dynamical evolution, but is also essential for the formation of hydrogen molecules (H$_2$) which control the star formation in metal poor environments such as primordial galaxies. Several attempts have recently been made to take account of the radiative transfer of ionizing photons, adopting for instance a pure absorption approximation (Kepner, Babul \\& Spergel 1997; Kitayama \\& Ikeuchi 2000, hereafter KI), the photon conservation method (Abel, Norman \\& Madau 1999), and the full radiative transfer treatment (Tajiri \\& Umemura 1998, hereafter TU; Barkana \\& Loeb 1999; Susa \\& Umemura 2000).  As for the dynamical states, Kepner, Babul \\& Spergel (1997) and Barkana \\& Loeb (1999) have considered {\\it hydrostatic} equilibria of spherical clouds within virialized dark halos. KI have explored the {\\it hydrodynamical} evolution of a spherical system composed of dark matter and baryons.  The full radiative transfer treatment, however, has not hitherto been incorporated with the hydrodynamics of spherical pregalactic clouds.  In this paper, we attempt to implement an accurate radiation-hydrodynamical (RHD) calculation on the evolution of spherical clouds exposed to the UV background, solving simultaneously the radiative transfer of photons and the gas hydrodynamics.  Our goals are to predict at a physically reliable level final states of photoionized clouds and also to assess accurately self-shielding against the UV background that is essential for H$_2$ cooling and the subsequent star formation.  Throughout the present paper, we assume the density parameter $\\Omega_0=1$, the Hubble constant $h=H_0$/(100 km s$^{-1}$ Mpc$^{-1})=0.5$, and the baryon density parameter $\\Omega_{\\rm b} = 0.1$. ",
        "conclusions": ""
    },
    "0004/astro-ph0004382_arXiv.txt": {
        "abstract": "Several facts regarding the structure of NGC 604 are examined here. The three main cavities, produced by the mechanical energy from massive stars which in NGC 604 are spread over a volume of 10$^6$ pc$^3$, are shown here to be undergoing blowout into the halo of M33. High resolution long slit spectroscopy is used to track the impact from massive stars while HST archive data is used to display the asymmetry of the nebula.  NGC 604 is found to be a collection of photoionized filaments and sections of shells in direct contact with the thermalized matter ejected by massive stars. The multiple blowout events presently  drain the energy injected by massive stars and thus the densest photoionized gas is found almost at rest and is expected to suffer a slow evolution. ",
        "introduction": "Upon the formation of massive stellar clusters, the energy input from massive stars leads to a rapid structuring of the surrounding medium. This was originally thought to be due to the high pressure acquired by the constant density photoionized gas which, upon the formation of the H\\,{\\sc ii} region, would then expand with a speed smaller than the sound speed ($c_{HII}$ $\\sim10$ \\kms) onto the surrounding neutral gas to acquire a radius  $R_{HII}(t) = R_{St} (1 + 7/4 (c_{HII}t/R_{St}))^{4/7}$  \\noindent where $R_{St}$ is the Str\\\"omgren radius (Spitzer 1968; Osterbrock 1989). Much larger speeds were envisaged in ''champagne flows'' as a result of the pressure gradient between the photoionized parent cloud and the photoionized inter-cloud medium, suddenly established as the ionization front (either D or R type) overruns the parent cloud edge (\\GTT\\ 1979). In this case the ionization front rushes to infinity and strong isothermal shocks develop within the photoionized region.  The shocks are driven by the ionized cloud matter that supersonically streams into the intercloud medium in an attempt to lift and equalize the pressure between the two regions. The champagne phase would then last until the pressure gradient is smoothed out; fact that is accomplished only once the whole ionized cloud is dispersed.  A completely new paradigm resulted from taking into consideration the mechanical energy produced by massive stars and the evolution of superbubbles. In this case three concentric distinct regions were shown to develop around the exciting source (Castor et al. 1977): The central region ({\\it a}) contains the hypersonic wind from the massive star(s). At a certain radius, $R_a$, this is abruptely thermalized at a reverse shock. Given the speed of the wind, the high temperature ($T$ = 1.4 $\\times$ 10$^7$ ($v_{wind}$/1000 \\kms)$^2$ K $\\sim$ 10$^6$--10$^7$ K) and high pressure acquired upon thermalization, allow the shocked wind to fill a large volume (region {\\it b}), while driving a strong shock into the surrounding gas. The latter shock then sweeps and accelerates the surrounding ISM into a dense expanding shell (region {\\it c}), that grows as a function of time as (Castor et al. 1977; Weaver et al. 1977)   $ R_s(t) = 0.76 (\\dot{E}/\\rho)^{1/5} t^{3/5} $  \\noindent where $\\dot{E}$ is the mechanical energy input rate and $\\rho$ the mass density of the background medium. Castor \\etal (1977) also showed that, given the number of recombinations within the expanding shell and its rapid growth, the ionization front became trapped within the shell early in the evolution of the system. This was shown to happen despite the fact that most of the volume is filled with the hot thermalized wind, which is a transparent medium to the ionizing photon flux due to its high temperature. Consequently, the H\\,{\\sc ii} region gas in this model was shown to be rapidly confined to a narrow 2D structure within the expanding shell.  Other models, accounting only for the mechanical energy deposited by OB associations and starbursts (\\GTT\\ \\& Bodenheimer 1988), have clearly shown how the shells of swept up matter can easily be disrupted upon crossing a large density gradient. This phenomenon, commonly referred to as blowout, occurs as a shell suffers a sudden acceleration, fact that drives it Rayleigh-Taylor unstable and causes its rapid fragmentation. The event then leads to the venting of the hot superbubble interior onto the low density gas, where a new shell of halo swept up matter will develop. More recent studies have also shown that blowout also leads to the escape of the $uv$ photons, which otherwise would be trapped within the decelerating shell. The escape of ionizing radiation has been shown to lead to the development of an extended, low surface brightness conical H\\,{\\sc ii} region ahead of the superbubble leading shock (\\GTT\\ et al. 1999).  Here we center our attention on the nearest northern hemisphere giant \\h2 region NGC 604. The giant nebula is the largest in the neighbouring spiral galaxy M33 (at a distance of 840 kpc; Freedman \\etal 1991), and is located in one of the spiral arms,  2.4 kpc from the galaxy centre. Looking at NGC 604, it is quite evident that it is not a spherical nebula.  It presents a large number of shells, filaments or arches of a wide variety of lengths and intensities.  Its general appearance is, however, very similar to other giant H\\,{\\sc ii} regions, in particular to the 30 Doradus nebula. The large number of shells and networks of shells in 30 Doradus have been identified spectroscopically by Chu \\& Kennicutt (1994). Such a structure, i.e. expanding shells within shells, imply a very inhomogeneous rapidly structured, or pre-structured, parent cloud density distribution. An endless number of tunnels, or cavities, across which the thermalized winds from the massive stars could freely flow causing their expansion and further growth, as to in projection allow for the appearance of a network of nested shells. In such a case, the thermalized ejecta, the hot coronal gas, is to be in direct contact with the H\\,{\\sc ii} region gas, and the latter would be confined to the narrow denser walls of the network of tunnels or to the sections of the expanding shells, presenting thus a three dimensional superposition of multiple 2D structures.   Here we analize the ionization structure, and the velocity field of NGC 604, with the aim of unveiling its true morphology and possible evolution of the nebula. In section 2 we analyse the two data sets used in this work, and in section 3 we present a discussion of the results and our conclusions. ",
        "conclusions": ""
    },
    "0004/hep-ph0004024_arXiv.txt": {
        "abstract": "{Preheating after inflation may lead to nonthermal phase transitions with symmetry restoration. These phase transitions may occur even if the total energy density of fluctuations produced during reheating is relatively small as compared with the vacuum energy in the state with restored symmetry. As a result, in some inflationary models one encounters a secondary, {\\it nonthermal} stage of inflation due to symmetry restoration after preheating. We review the theory of nonthermal phase transitions and make a prediction about the expansion factor during the secondary inflationary stage. We then present the results of lattice simulations which verify these predictions, and discuss possible implications of our results for the theory of formation of topological defects during nonthermal phase transitions. } ",
        "introduction": "The theory of cosmological phase transitions is usually associated with symmetry restoration due to high temperature effects and the subsequent symmetry breaking which occurs as the temperature decreases in an expanding universe \\cite{Kirzhnits,Weinb,Kirzhnits2,book}. A particularly important version of this theory is the theory of first order cosmological phase transitions developed in \\cite{Kirzhnits2}. It served as a basis for the first versions of inflationary cosmology \\cite{infl}, as well as for the theory of electroweak baryogenesis \\cite{Shap}.  Recently it was pointed out that preheating after inflation \\cite{KLS} may rapidly produce a large number of particles that for a long time remain in a state out of thermal equilibrium. These particles may lead to specific {\\em nonthermal} cosmological phase transitions \\cite{ptr1,ptr2}. In some cases these phase transitions are first order \\cite{bubble,Copeland}; they occur by the formation of bubbles of the phase with spontaneously broken symmetry inside the metastable symmetric phase. If the lifetime of the metastable state is large enough for the energy density of fluctuations to be diluted, one may encounter a short secondary stage of inflation {\\it after} preheating \\cite{ptr1}. Such a secondary inflation stage, if it occurs late enough, could be important in solving the moduli and gravitino problems.  In this respect  secondary ``nonthermal''  inflation due to preheating may be an alternative to the ``thermal inflation'' \\cite{LS}, suggested for solving these problems.  In this paper we will briefly present the theory of such phase transitions and then give the results of numerical lattice simulations that directly demonstrated the possibility of such brief inflation. We will also discuss possible implications of our results for the theory of formation of topological defects during nonthermal phase transitions. A detailed description of numerical methods used in our work will be given in   the appendix. ",
        "conclusions": "The results of our lattice simulation confirm our expectations that preheating may lead to nonthermal phase transitions even in those theories where spontaneous symmetry breaking occurs at the GUT scale, $v \\sim 10^{16}$ GeV. Some time ago this question was intensely debated in the literature. Some authors claimed that nonthermal phase transitions induced by preheating are impossible, and the notion of the effective potential after preheating is useless. In our opinion, Figs. 1, 2 and 3 give a clear answer to this question. In particular, Fig. 1 shows that 90\\% of the time from the end of inflation to the moment of symmetry breaking the field $\\phi$ oscillates about $\\phi = 0$ with an amplitude much smaller than $v$. This could happen only because the corrections to the effective potential induced by particles $\\phi$ and $\\chi$ change the shape of $V(\\phi)$ near $\\phi = 0$, turning its maximum into a deep local minimum.   In some theories, this effect may lead to production of superheavy strings, which may have important cosmological implications for the theory of formation of the large scale structure of the universe. In some other theories, these phase transitions may lead to excessive production of monopoles and domain walls. This may rule out a broad class of otherwise acceptable inflationary models.  In this paper we have shown that under certain conditions  a nonthermal phase transition may lead to a short secondary stage of inflation. It would be interesting to study the possibility that a secondary stage of inflation induced by preheating could  help solve the moduli and gravitino problems.  The answer to this question will be strongly model-dependent because gravitinos can be produced by the oscillating scalar field even after the secondary inflation \\cite{KKLP,GTR}. Independently of all practical implications, the possibility of a secondary stage of inflation induced by preheating seems very interesting because it clearly demonstrates the potential importance of nonperturbative effects in post-inflationary cosmology.     The authors are grateful to Julian Borrill for very useful discussions. This work was supported  by  NSERC and  CIAR and by NSF grant AST95-29-225. The work of G.F. and A.L. was also supported   by NSF grant PHY-9870115, and the work of L.K. and A.L. by NATO Linkage Grant 975389.  \\"
    },
    "0004/astro-ph0004088_arXiv.txt": {
        "abstract": "WRA\\,751 is an evolved massive star in our Galaxy closely resembling Luminous Blue Variable stars (LBVs). It is surrounded by a nitrogen enriched nebula of about 23\\arcsec\\ diameter. A comparative study of the nebula's morphology and kinematics is presented, it supports---together with spectroscopical evidence---the classification of WRA\\,751 as a LBV. Images show that the nebula consists of a nearly spherical shell as well as a bipolar-like structure north and south of its main body, the Northern and Southern Caps.  In contrast to the almost spherical appearance of the main body of the nebula, the kinematics shows a deviation even of this part from a classical spherical expansion pattern. From the present data it can be concluded that the main body expands asymmetrically (central expansion velocity $\\sim 26$\\,\\kms), with a thicker shell at the back side. A bump-like structure can be found to the west of the central star. In addition to the main body, bipolar kinematic components can be identified with the morphologically classified Caps. These results put WRA\\,751 into the class of LBVs which are surrounded by a nebula with bipolar components, albeit considerably less pronounced than, for instance, in the classical bipolar LBVs $\\eta$ Car and HR Car. ",
        "introduction": "With today's initial mass function, stars with masses of $M \\sim 50 - 100$\\,M$_{\\sun}$ and luminosities of $L\\sim 10^{5-6}$\\,L$_{\\sun}$ top the {\\it Hertzsprung-Russell diagram\\/} (HRD) and represent the most massive stars known. These massive stars show a remarkable evolutionary behavior (e.g., Schaller et al.\\ 1992, Stothers \\& Chin 1996). After spending a `normal' life as O stars on the main sequence they evolve towards cooler temperatures while entering a phase with very high mass loss (up to 10$^{-4}$M$_{\\sun}$\\,yr$^{-1}$), they become {\\it Luminous Blue Variables\\/} (LBVs). This phase starts when the stars reach the {\\it Humphreys-Davidson limit\\/} (Humphreys \\& Davidson 1979, 1994, Langer 1994) in the HRD. Analyzing HRDs of the Galaxy and the LMC, Humphreys (1978, 1979) and Humphreys \\& Davidson (1979) found a lack of very luminous red supergiants. Apperently the most massive stars do not evolve into red supergiants but instead their evolution is reversed towards the blue supergiant part in the HRD. The turning points form the empirical Humphreys-Davidson limit. Around this Humphreys-Davidson limit in the HRD LBVs are found. One of the most prominent characteristics of the unstable LBV phase are strong stellar winds and possible giant eruptions, which lead to the peeling off of parts of the stellar envelope and the formation of small circumstellar nebulae, the so-called {\\it LBV nebulae\\/} (LBVN, e.g. Nota et al.\\ 1995).  Only very few LBVs are known in our Galaxy (8 classified and candidate objects) and a few in other galaxies, for instance in the LMC, SMC, M31, and M33 (Humphreys \\& Davidson 1994). Altogether nearly 40 LBVs and good candidates are currently known. Not all of them, but many show circumstellar nebulae. In this paper we analyze the LBVN around the galactic star WRA\\,751 and put it into context with the well known and better studied LBVs $\\eta$ Carinae and HR Carinae.  WRA\\,751 ($=$ He 3-591) was first identified as a possible Wolf-Rayet star by Henize (see Roberts 1962) and then appeared in the Carlson and Henize (1979) sample of southern peculiar emission-line stars, where the authors already recognized the strong [Fe\\,{\\sc ii}] lines and classified it as Bep. In addition they proposed a similarity between WRA\\,751 and HR Car (which in the meantime is known to be a LBV). Using spectroscopic and photometric observations Hu et al.\\ (1990) noted for the first time that WRA\\,751 is a good LBV candidate. In addition to the characteristically strong [Fe\\,{\\sc ii}] lines this star shows an irregular light variation typical for LBVs. In their analysis the star was classified as O9.5 with $M_{\\rm bol}= -9.6$\\mag. They derived $T_{\\rm eff} = 30\\,000$\\,K, $E{\\rm (B-V)} = 1.8, M = 50\\,{\\rm M}_{\\sun}$, and a distance of about 5\\,kpc.  One year later Hutsem\\'ekers \\& Van Drom (1991a) found that WRA\\,751 was surrounded by a nebula roughly 22\\arcsec\\ in diameter. They found the nebula to be nearly spherically symmetric and to show a diffuse structure. Their low-dispersion spectra revealed a nebula of low excitation with an electron temperature of $T_{\\rm e,neb} < 15\\,000\\,$K, an electron density of $n_{\\rm e}\\sim 400\\,{\\rm cm}^{-3}$, and an expansion velocity of $v_{\\rm exp} = 26\\,$\\kms . Assuming a distance of 7\\,kpc (derived from radial velocities and making use of the galactic rotation curve) their radius of the nebula is 0.38\\,pc.  Analysis of infrared data (see de Winter 1992) show a NIR excess from which an estimate for the mass loss of $\\dot{M} \\sim 10^{-5.5...-6}\\,M_{\\sun}\\,{\\rm yr}^{-1}$ with a velocity of $v_{\\rm wind} = 500\\,$\\kms\\ were derived. They found a cool dusty circumstellar shell with strong emission in the FIR to surround WRA\\,751. In 1992, from interstellar and circumstellar reddening, van Genderen measured a distance of $4-5\\,$kpc for WRA\\,751 as a lower limit. Garc{\\'\\i}a-Lario et al. (1998) re-determined the characteristic parameters and found $T_{\\rm eff} \\sim 25\\,000\\,$K, $v_{\\rm exp} = 24\\,$\\kms, and $n_{\\rm e} \\sim 200\\,{\\rm cm}^{-3}$.  In summary, the published observations and their interpretations point towards WRA\\,751 being a LBV star with a slowly expanding nebula. ",
        "conclusions": " \\begin{itemize} \\item The LBVN around WRA\\,751 consists of an expanding shell of 23\\arcsec\\ in diameter, with a thicker shell at the back side. The radial velocity distribution deviates from that of a simple spherical expansion pattern. We find bipolar contributions to the morphology as well as radial velocity distribution in approximately north-south direction (Northern and Southern Caps). \\item The maximum expansion velocities are off-centered which most likely can be accounted for by a bump in the back side east of the central star. \\item The [N\\,{\\sc ii}]\\,6583\\AA/H$_\\alpha$ ratio is considerably larger than that of the background material. This is typical for CNO-processed material. \\end{itemize} These three results are fully consistent with WRA\\,751 being a true member of the LBV class of stars. In many of the well investigated ones bipolar components to the morphological as well as kinematic structures of the nebula are found. $\\eta$ Car (Duschl et al.\\ 1995, Morse et al.\\ 1998) and HR Car (Weis et al.\\ 1997, Nota et al.\\ 1997) are the most prominent examples. As stars in late phases of their evolution, LBVs quite naturally show CNO-processed material in their envelopes and ejecta (Garc{\\'\\i}a-Segura et al.\\ 1996, Smith et al.\\ 1978). A large [N\\,{\\sc ii}]\\,6583\\AA/H$_\\alpha$ ratio in a nebula is often used as one of the indicators for a star being a LBV. One finds, for instance, values of [N\\,{\\sc ii}]\\,6583\\AA/H$_\\alpha = 3\\ \\dots\\ 7$ for $\\eta$ Car (Davidson et al.\\ 1982, Meaburn et al.\\ 1987, 1996, and Weis et al.\\ 1999), $= 0.4\\ \\dots\\ 0.9$ for HR Car (Hutsem\\'ekers \\& van Drom 1991b, and Weis et al.\\ 1997), and $= 0.7$ for AG Car (Thackeray 1977 and Smith et al.\\ 1997).  Combining our morphological, kinematic and spectroscopic results, we find mounting evidence that WRA\\,751 is a LBV, indeed. Moreover, there are bipolar components in its nebula, notably the two Caps. While the bipolarity in this object is less pronounced and less obvious than in other LBVs, it still strengthens our suspicion that bipolarity---albeit at different levels---is a genuine property of LBV nebulae. This makes it even more likely that modeling LBVNs as simple windblown spheres is an unjustified oversimplification."
    },
    "0004/astro-ph0004041_arXiv.txt": {
        "abstract": "High angular resolution  far-infrared observations at 143 \\& 185 \\micron, using the TIFR 1-m balloon borne telescope, are presented for two Galactic star forming complexes associated with IRAS 00494+5617 and 05327-0457. The latter map also reveals the cold dust in OMC-3. The HIRES processed IRAS maps at 12, 25, 60 \\& 100 \\micron\\ have also been presented for comparison. Both these regions are illuminated at the edges by high mass stars with substantial UV flux.The present study is aimed at quantifying the role of the nearby stars vis-a-vis embedded young stellar objects in the overall heating of these sources. Based on the FIR observations at 143 \\& 185 \\micron\\ carried out simultaneously with almost identical angular resolution, reliable dust temperature and optical depth maps have been generated for the brighter regions of these sources. Radiative transfer modeling in spherical geometry has been carried out to extract physical parameters of these sources by considering the observational constraints like : spectral energy distribution, angular size at different wavelengths, dust temperature distribution etc. It has been concluded that for both IRAS 00494+5617 and IRAS 05327-0457, the embedded energy sources play the major role in heating them with finite contribution from the nearby stars. The best fit model for IRAS 00494+5617 is consistent with a simple two phase clump-interclump picture with $\\sim$ 5\\% volume filling factor (of clumps) and a density contrast of $\\approx$ 80.    {\\em Subject Headings}: { IRAS 00494+5617 -- IRAS 05327-0457 -- Far Infrared Mapping -- Radiative Transfer } ",
        "introduction": "A long term programme of studying the distribution of cold dust ( down to $\\sim$ 15 K) in and around Galactic star forming regions, is in progress using the 1-m TIFR balloon-borne far-infrared (FIR) telescope. Under this programme, high angular resolution ($\\sim$ 1\\arcmin) mapping is carried out simultaneously in two trans-IRAS FIR wavebands \\citep{ghosh96,bmook98,bmook99a,verma99}. In this paper we present FIR mapping of the sources IRAS 00494+5617 and 05327-0457 in wavebands centered around 143 \\& 185 \\micron. These sources share the similarity that each is heated by a luminous external source in addition to one or more possible embedded sources. The simultaneity of observations in the two wavebands with nearly identical beams is useful in deriving reliable temperature distributions of the interstellar dust in these regions.   The source IRAS 00494+5617 is  part of the western fragment of the molecular cloud NGC281, located in the Perseus arm at a distance of 2.2 kpc \\citep{ces99}. A compact star cluster containing the multiple star HD 5005 as the brightest star excites an ionization front at the edge of NGC 281W. HD 5005 is an O6 V type star  and the cluster is at an angular distance of $\\sim$ 5\\arcmin\\ from  IRAS 00494+5617. Observational evidences in the form of detection of C$^{34}$S (3$\\rightarrow$2) \\citep{meg97}, NH$_{3}$ \\citep{hen94}, 22 GHz H$_{2}$O masers \\citep{tofani95}, CO emission \\citep{carp90} and molecular outflows \\citep{snel90}  together with FIR and millimeter continuum emission, strongly suggest ongoing star formation activity in this region. Existing literature on this source suggests that the ongoing star formation activity has been induced by the compression due to the propagation of the ionization front (energized by HD 5005) into the molecular cloud \\citep{elme78,meg97}. The gas emission features from this source are reasonably well studied. There are not as many high angular resolution observations of the emission from the dust component. Longward of IRAS wavelengths, the only available observations are the high resolution maps at 1.3 mm \\citep{hen94}.   The source IRAS 05327-0457 is associated with the nebulosity NGC 1977, located near the northern star in the sword of Orion.  It is bounded on the south by the northern end of the Orion molecular cloud \\citep{kut76} and is at a distance of 450 pc. The exciting star for the HII region is the B1 V star HD 37018, also known as 42 Ori and is at an angular distance of $\\sim$ 5\\farcm9 from IRAS 05327-0457. Earlier observations of the NGC 1977 region include FIR ($\\leq$ 160 \\micron), near-infrared and radio continuum mapping together with extensive molecular line observations \\citep{mak85,kut85}. High resolution molecular line maps \\citep{kut85} indicate an increase in radial velocity  from south to north implying an expansion of the HII region into the molecular cloud, with a velocity of a few km/sec.   In the  present study the structural details and dust temperature distributions obtained from FIR observations have been used to quantify the relative contributions of the external and internal sources towards heating these regions. In sections 2 \\& 3, the observations and the results are described. The discussion including the radiation transfer modeling of these sources is presented in section 4. ",
        "conclusions": "\\subsection{Radiation transfer scheme \\& Models explored \\label{secrad}}  We have explored the major sources of dust heating based on radiation transfer modeling of the sources. The radiation transfer equations have been solved assuming a 2-point boundary condition  for a spherically symmetric cloud of dust and gas. Based on the boundary conditions, we have constructed 3 types of models : $model~ A$ -- the cloud is  heated by centrally embedded sources and  by an external radiation field due to the average Galactic Interstellar Radiation Field (ISRF) and nearby stars; $model~ B$ \\-- the cloud is heated by internal sources and ISRF only; $model~ C$ -- there is no embedded source, the cloud is heated only externally by the ISRF and radiation due to nearby stars. The contribution of the nearby star in the models $A$ and $C$ has been calculated in the following way: the geometrically diluted stellar radiation intercepted by the cloud surface is estimated and it is then smeared out uniformly over the entire surface of the cloud. This calculated contribution could be a slight overestimate since the absorption in the intervening medium has been neglected. For modeling the observed Spectral Energy Distribution (SED) two types of dust have been considered. The dust properties have been taken from two sources viz., \\citet{draine84} (hereafter DL) and \\citet{mez82} (hereafter MMP). The dust compositions include mixtures of astronomical silicate and graphite. For the DL type of dust grains, size averaged properties are used taking a size distribution of $n(a)da \\sim a^{-3.5}da$ \\citep{mrn77}, with $a$ ranging between 0.01 and 0.25 \\micron. The average Galactic ISRF used for this problem has been taken from \\citet{mmp83}. The cloud is parameterized by the following physical quantities: R$_{\\rm max}$, the outer size of the cloud, R$_{\\rm min}$, the radius of the inner dust cavity, $\\tau_{\\rm 100}$, the total radial optical depth at 100 \\micron\\ and the radial dust (hence gas) density distribution ($r^{\\beta}$, $\\beta$=0, -1, -2). The contributions of ISRF and the external source (if any) are kept fixed, while the luminosity of the embedded source, the dust composition, $\\tau_{\\rm 100}$ value and radial distribution of the dust density and the physical sizes of the cloud are varied to obtain a good match to observations. The gas to dust ratio has been assumed to be 100:1 by mass. Using this scheme a best fit model matching the observed spectral energy distributions (SED), radial profiles at selected wavelengths and the radio continuum flux are obtained. These radiation transfer calculations have been done using a modified version of the code CSDUST3 \\citep{egan88}, which is capable of treating gas and dust in a self-consistent manner \\citep{bmook99a,bmook99b}.  In the following subsections we explore the tenability of the models $A$, $B$ and $C$ in the light of available observational constraints for individual sources.  \\subsection{IRAS 00494+5617}   \\subsubsection{ The main source }  The observed SED for the source IRAS 00494+5617 has been constructed using the flux densities presented in Table~\\ref{ngc281ftab} along with flux densities at 1.3 mm \\citep{hen94} and 2.9 mm \\citep{walk90}. The total observed luminosity is 2.4$\\times$10$^{4}$ \\lsun\\ \\citep{carp90} and  no radio continuum emission is detected at 8.4 GHz \\citep{tofani95}  from this cloud. T(143/185) map shows enhanced temperature towards the boundary closer to HD 5005. In order to explain these, the embedded source in radiation transfer models $A$ and $B$ is assumed to be a cluster of stars. In models $A$ and $C$ an intercepted luminosity of 6.9$\\times$10$^{3}$\\lsun\\  from HD 5005 has been considered. For each type of model it is found that the best fit is obtained for the case of uniform density and grains of DL type. Figure~\\ref{ngc281sed} shows  the observed and the best fit SEDs. Model $A$ fits the SED at all wavelengths very well, while  model $B$ shows some mid-infrared excess. Model $C$ fails to reproduce any of the observed flux densities. The predicted radial profiles (as obtained by convolving the model profiles with instrumental PSFs) for the models $A$ and $B$ are almost similar. Table~\\ref{proftab} compares the FWHMs predicted by model $A$ with those observed at 25, 60, 100, 143 and 185 \\micron. The observed FWHMs presented here refer to the FWHMs for the source along its major and the minor axes. We find that model $A$ not only fits the observed SED better than model $B$ but also achieves an overall consistency by supporting the role of HD 5005 in heating the cloud and fitting the observed FWHMs reasonably well. Table~\\ref{tranparm} presents the parameters of the best fit model $A$ for this source alongwith the uncertainties in the $\\tau_{\\rm 100}$ and the embedded luminosity. These uncertainties have been estimated using only the SED fit and keeping the other parameters fixed. From this radiation transfer modeling and the T(143/185) map we conclude that the source IRAS 00494+5617 is primarily heated by one or more embedded sources with non-negligible contribution from HD 5005.  We compare he FWHM predicted from the best fit spherically symmetric model with the observed FWHM along the minor axis. For IRAS 00494+5617 we find that at 100, 143 and 185 \\micron\\ model predictions for FWHMs match the observed FWHMs (along minor axis) fairly well. However at 25 \\micron\\ (which traces much hotter dust) the observed FWHM is substantially larger than the model prediction. This trend is also seen at 60 \\micron, though not to the same extent as at 25 \\micron. This can be explained by invoking spatially distributed heating sources embedded in clumps. It may be noted that our dust temperature maps also shows the presence of two temperature enhancements. This clumpiness is also substantiated by high resolution C$^{18}$O observations by \\citet{meg97}, which have revealed that this western fragment of NGC 281 is composed of three clumps. The best fit model $A$ obtained above corresponds to a gas density of 2.2$\\times$10$^{4}$ cm$^{-3}$ ( for a gas to dust density of 100:1 by mass) and a total mass of 1890\\msun. \\citet{meg97} obtain 1080\\msun\\ as the lower limit for the total mass in the three clumps. The residual mass (810 \\msun) as predicted by the model can be present in the form of the inter-clump medium. The projected radii (taking the distance to the source to be 2.9 kpc) of the clumps were given by \\citet{meg97} assuming the clumps to be spherical. We assume that all three clumps are spherical and are confined within the boundary of the spherical model $A$ above. After scaling the clump sizes to our assumed distance of 2.2 kpc, we find that the volume filling factor is 4.5 \\%. Using the total mass from our model $A$ and assuming an interclump medium of uniform density we obtain a value of 4.5$\\times$10$^{3}$ cm$^{-3}$  for the number density of the interclump medium. We also calculate the clump-interclump density contrast to be $\\sim$ 80. The resultant clump density of 3$\\times$10$^{5}$ cm$^{-3}$ is consistent with the detection of C$^{34}$S. Although the estimates presented here are based on crude approximations, the values for filling factors of clumps and clump-interclump density contrast obtained here can be called very typical, while comparing with the measured values of such parameters for other clumped sources \\citep{howe91}.   \\subsubsection{Diffuse emission}  The 143 \\micron\\ map of this region shows a considerable amount of diffuse emission particularly towards the west. This emission is to a certain extent positionally coincident and structurally similar to that  in the 185 \\micron\\ map. We have looked into the possible energy sources for this emission. The flux densities (in an aperture of 2\\arcmin\\ dia centered at 0$^{\\rm h}$ 49$^{\\rm m}$ 09$^{\\rm s}$ +56\\arcdeg\\ 17\\arcmin\\ 30\\arcsec\\ (1950)) at 143 and 185 \\micron\\ are 236 and 265 Jy respectively. The bolometric luminosity of the diffuse emission over a diameter of 2\\arcmin\\ has been estimated to be $\\sim$ 870 \\lsun. This luminosity cannot be explained due to heating by either the the average ISRF or the star cluster HD 5005. We explore the possibility of radiation actually leaking from the main source (around 00494+5617) and heating the dust in the neighbourhood. Under this assumption we obtain that about 8\\% of the total luminosity of the embedded source would have to be intercepted by the region where diffuse emission is detected. The separation between the region of diffuse emission and IRAS 00494+5617 is such that the solid angle covered by the diffuse region is adequate to intercept the requisite amount of radiation. This evidence of luminosity leakage further supports the presence of clumpiness discussed in Section 4.2.1.   \\subsection{IRAS 05327-0457}  The observed SED for the region around IRAS 05327-0457 has been constructed using the flux densities presented in Table~\\ref{ngc1977ftab} alongwith the IRAS Low Resolution spectrum (LRS) \\citep{chen95}. VLA observations of this source by \\citet{kut85} have not detected any radio continuum emission at 5 GHz and from the detection limit predict the embedded source to be a ZAMS B2.5 or later type star. Based on their observations \\citet{kut85} have suggested that 42 Ori plays the most significant role in fueling the FIR emission from this source. However the T(143/185) map presented here shows no temperature enhancement towards the boundary illuminated by the star 42 Ori. It rather shows an increase from the boundary towards the position of IRAS 05327-0457.  We have applied the radiation transfer models $A$, $B$ and $C$ (see Section~\\ref{secrad}) to IRAS 05327-0457. The embedded source for models of type  $A$ and $B$ is taken to be a star of type B2. In  models of types $A$ and $C$ the intercepted contribution of 600\\lsun\\ from 42 Ori is considered. Figure~\\ref{ngc1977sed} presents the SEDs predicted from the models $A$, $B$ and $C$ in comparison with the observations. The model outputs shown are the best fits for each type. As in the case of IRAS 00494+5617 uniform density gives the best fit for all three models. Further the MMP type grains lead to better fit than the DL type grains. While the predicted SED from model $A$ matches the observations at all wavelengths reasonably well, model $B$ predicts fluxes less than the observed values longward of 60 \\micron. Predictions from model $C$ do not the fit the observations at all. The predicted FWHMs (post convolution) are similar for models of type $A$ and $B$. Table~\\ref{proftab} presents the comparison of FWHM sizes from model $A$ at different wavelengths with observations. From model $A$ the predicted radio continuum emission at 5 GHz is $\\sim$14 mJy which is well above the sensitivity of the VLA observations by \\citet{kut85}. The reasons for this  non-detection  could be : (i) the embedded source is a cluster of stars (with masses less than that of a B2 star) with combined luminosity equal to that of a B2 star. This is supported by the detection of multiple stars in the near-infrared \\citep{mak85}. They have detected 7 stars having a total luminosity of 2.3$\\times$10$^{3}$\\lsun\\ (model $A$ total luminosity = 2.8$\\times$10$^{3}$\\lsun). Our temperature map also indicates presence of two sources, again close to the boundary. (ii) Since the radio continuum emission has a strong n$_{e}^{2}$ dependence, a smaller gas to dust ratio could explain the reduced emission. (iii) The radio emission could largely be attenuated by free-free absorption. The third possibility proposed would arise, if the gas density distribution very close to the star (r$<<$R$_{\\rm min}$) is r$^{-2}$, a typical situation for ionized wind.  Of the models presented here, model $A$ explains the observed SED, the radial profile, the T(143/185) map and the role of 42 Ori in heating the cloud most satisfactorily. Table~\\ref{tranparm} presents the parameters of this model and the estimated uncertainties in the embedded luminosity and $\\tau_{\\rm 100}$. From radiation transfer modeling we conclude that it is absolutely necessary to have an embedded source to explain the total emission from IRAS 05327-0457. Clearly 42 Ori cannot be the only heating source since the luminosity intercepted by IRAS 05327-0457 is $\\leq$ 20 \\% of the observed luminosity. Further we have verified that the temperature profile from radiation transfer calculation using a slab geometry and 42 Ori as the source, does not fit observations.   The observed angular sizes for IRAS 05327-0457 at 100 \\micron\\ and more prominently at 60 \\micron\\ are larger than the model values. As in IRAS 00494+5617 this is again consistent with a distribution of embedded hot sources as suggested by the T(143/185) map (Figure~\\ref{ngc1977tmptau}) as well. The apparent reduction in the observed FWHMs at 143 and 185 \\micron\\ as compared to the FWHM at 100 \\micron\\ could be due to the effect of sky chopping."
    },
    "0004/astro-ph0004107_arXiv.txt": {
        "abstract": "The Galactic kinematics of Mira variables have been studied using infrared photometry, radial velocities, and Hipparcos parallaxes and proper motions. For Miras in the period range 145 to 200 days (probably corresponding to [Fe/H] in the range --0.8 to --1.3) the major axes of the stellar orbits are concentrated in the first quadrant of Galactic longitude. This is interpreted as a continuation of the bar-like structure of the Galactic Bulge out to the solar circle and beyond. ",
        "introduction": "A somewhat unexpected discovery of microlensing experiments towards the Galactic Bulge was the need to take into account a bar-like structure there (see, e.g.\\ Paczy\\'{n}ski 1996). The precise nature of this bar - its composition, extent and evolution - is only sketchily known. This is an area where progress is likely to come from the combination of microlensing results with other types of observations. In the present paper we summarize the results of an analysis of the Galactic kinematics of Mira variables based on infrared photometry, radial velocities, and Hipparcos parallaxes and proper motions (Whitelock, Feast \\& Marang 2000, Whitelock \\& Feast 2000, Feast \\& Whitelock 2000b). This provides evidence for a rather extended bar-like structure in the Galaxy. ",
        "conclusions": ""
    },
    "0004/astro-ph0004331_arXiv.txt": {
        "abstract": "We study the perturbation modes of rotating superfluid ellipsoidal figures of equilibrium in the framework of the two-fluid superfluid hydrodynamics and Newtonian gravity. Our calculations focus on linear perturbations of background equilibria in which the two fluids move together, the total density is uniform, and the densities of the two components are proportional to one another, with ratios that are independent of position. The motions of the two fluids are coupled by mutual friction, as formulated by Khalatnikov. We show that there are two general classes of modes for small perturbations: one class in which the two fluids move together and the other in which there is relative motion between them. The former are identical to the modes found for a single fluid, except that the rate of viscous dissipation, when computed in the secular (or ``low Reynolds number'') approximation under the assumption of a constant kinematic viscosity, is diminished by a factor $f_N$, the fraction of the total mass in the normal fluid. The relative modes are completely new, and are studied in detail for a  range of values for the phenomenological mutual friction coefficients, relative densities of the superfluid and normal components, and, for Roche ellipsoids, binary mass ratios. We find that there are no new secular instabilities connected with the relative motions of the two fluid components. Moreover, although the new modes are subject to viscous dissipation (a consequence of viscosity of the normal matter), they do not emit gravitational radiation at all. ",
        "introduction": "The problem of the equilibrium and stability of rotating neutron stars is encountered in various astrophysical contexts, ranging from the limiting frequencies of rapidly rotating isolated millisecond pulsars, emission of gravitational waves in neutron star-neutron star or neutron star-black hole binaries, to the generation of  $\\gamma$-rays in the neutron star mergers and $X$-rays in accreting systems\\cite{CONF}. Considerable current interest is attached to the problem of neutron star binary inspiral, which would be the primary source of gravitational wave radiation for detection by future laser interferometers. Such a detection, apart from testing the general theory of relativity,  potentially could provide useful information on the equation of state of superdense matter. Also the stability criteria for rapidly rotating neutron stars are essential for placing firm upper limits on the frequencies to which millisecond pulsars can be spun up thereby constraining the range of admissible equations of states.  High precision, fully relativistic treatments of rapidly rotating isolated neutron stars and binaries comprising two neutron stars have become available in recent years\\cite{CONF}. Nevertheless, the development of simpler models that provide a fast and transparent insight into the underlying physics is needed when  the basic set of equations is modified to include new effects.  A systematic framework for the treatment of the equilibrium and stability of rotating liquid masses bound by self-gravitation in the Newtonian theory is contained in Chandrasekhar's {\\it Ellipsoidal Figures of Equilibrium}\\cite{CHANDRA} (hereafter EFE). The tensor virial method, developed most extensively by Chandrasekhar and co-workers, transforms the local hydrodynamical equations into global virial equations that contain the full information on the structure and stability of the Newtonian self-gravitating system as a whole. The method describes, in a coherent manner, the properties of solitary ellipsoids with and without intrinsic spin, and ellipsoids in binaries subject to Newtonian tidal fields. It is especially useful for studying divergence-free displacements of uniform ellipsoids from equilibrium, in which case the each perturbed virial equation yields (in the absence of viscous dissipation) a different set of normal modes.  Recent alternative formulations of the theory of ellipsoids  are based mainly on either the energy variation method \\cite{ZELDOVICH,ST} or the affine star model\\cite{CARTER} or the (Eulerian) two potential formalism\\cite{IPSER}. A large class of incompressible and compressible  ellipsoidal models has been studied using these methods\\cite{CARTER,IPSER,LRS}. The energy variation method employs the observation that an equilibrium configuration is possible if the energy of the ellipsoid is an extremum for variations of the ellipsoidal semiaxis at a constant volume; the ellipsoidal figure is stable only if the energy is a true local minimum. In the affine star method the figures  are described by a time-dependent Lagrangian as a  function of a deformation matrix and its derivative. The structure of the star at any particular time is related conformally to the initial unperturbed sphere via a quadratic form constructed from the deformation matrix.  The purpose of this paper is to extend previous studies to a treatment of the oscillation modes of ellipsoidal figures of equilibrium that contain a mixture of normal fluid and superfluid. Many-body studies of the pair correlations in neutron star matter show that the baryon fluids in their ground state form superfluid condensates in the bulk of the star. The superfluid phases, in the hydrodynamic limit, can be treated as a mixture of superfluid condensate and normal matter. The superfluid rotation is supported by the Feynman-Onsager vortex lattice state, and on the average leads to quasi-rigid body rotation of the superfluid component. The corresponding time-dependent two-fluid hydrodynamic equations are completely specified by two independent velocities for the superfluid and the normal component and corresponding densities of the constituents. The two fluids are coupled to one another by mutual friction forces, which we model phenomenologically according the prescription given by Khalatnikov \\cite{KHALATNIKOV}. Because of the additional degrees of freedom in this system, there are twice as many modes as for a single fluid. A natural question is whether the new modes affect the stability criteria previously deduced from studying perturbations of a single self-gravitating fluid.  In our treatment of perturbations of a mixture of normal fluid and superfluid, we shall follow most closely Chandrasekhar's formulation. However, since the basic equations of motion for the two fluids will include mutual friction between them, the system we study is inherently dissipative. Nevertheless, since the frictional forces only depend linearly on the relative velocity between the two fluids and vanish in the background, where the two fluids move together, one can still derive relations that resemble Chandrasekhar's tensor virial equations. Because these equations include dissipation, we shall prefer to regard them as moments of the fluid equations, rather than tensor virial equations. In fact, we shall relegate the derivation of perturbation equations from these moment equations to the Appendix of the paper, and instead derive the necessary equations for the fluid displacements directly by taking moments of linearized equations of motion for the two fluid components. In this paper, we concentrate on two-fluid variants of the classical Maclaurin, Jacobi and Roche ellipsoids.  Although our main aim is to access the oscillation modes and instabilities of neutron stars within the ellipsoidal approximation, the results obtained here may be of significance in other contexts (Ref. \\cite{CHANDRA}, the epilogue). One example is the understanding of rapidly rotating nuclei in the spirit of the Bohr-Wheeler model of a charged incompressible liquid droplet\\cite{BOHR}. In this case, the stability is determined by the competition between the attractive surface tension, the repulsive  Coulomb potential,  and the centrifugal stretching due to the rotation\\cite{COHEN,ROSEN}. Another example is the stability of rotating superfluid liquid drops of Bose condensed atomic gases, where the stability is determined through  an interplay among the pressure of the condensate, the confining potential of the magnetic trap and the centrifugal potential~\\cite{FETTER}.  Previous work on the oscillations of superfluid neutron stars concentrated mainly on perturbations of non-rotating or slowly rotating isolated neutron stars \\cite{EPSTEIN,LINDBLOM,UMIN,COMER}, and used methods entirely different from the one adopted here. The propagation of acoustic waves in neutron star interiors, including those related to the relative motion of neutron-proton superfluids, was studied by Epstein\\cite{EPSTEIN}, who found the compressional and shear modes related to short-wavelength oscillations of neutron star matter. The small-amplitude pulsation modes of superfluid neutron stars were derived by Lindblom and Mendell \\cite{LINDBLOM}, who found that the lowest frequency modes were almost indistinguishable from the normal modes of a single fluid star. Their analytical solutions also reveal the existence of a spectrum of modes which are absent in a single fluid star. Subsequent work concentrated on numerical solutions for the radial and non-radial pulsations of the two-fluid stars  and identified distinct superfluid modes in the absence of rotation\\cite{UMIN}. The r modes of slowly rotating two-fluid neutron stars have  been derived  by Lindblom and Mendell \\cite{LINDBLOM}, who find that they are identical to their ordinary-fluid counterparts to the lowest order in their small-angular-velocity expansion. The linear oscillations of general relativistic stars composed of two non-interacting fluids in a non-rotating static background have been studied by Comer {\\it et al.} \\cite{COMER}.   Our calculations allow arbitrary fast rotation, in the context of (incompressible) Newtonian fluid models. One may anticipate that the effects of superfluidity on oscillation modes,  if any, should be affected by the underlying vortex structure of the rotating superfluid. In our treatment, dissipation arises because of the drag forces experienced by the vortex lines as they move through the normal fluid (and there is no dissipation if the drag force is zero). We ignore the motions related to the isospin degrees of freedom in the core of a neutron star and, hence, the mutual entrainment of the neutron and proton condensates, as well as forces arising from deviations from $\\beta$ equilibrium.  The two-fluid equations used in the remainder of this work can adequately describe the mutual friction of a two-condensate fluid in the core of a neutron star, since the entrainment effect renormalizes the effective superfluid densities and the frictional coefficients, i.e. the phenomenological input in the two-fluid equations. While the ellipsoidal  approximation to a superfluid neutron star is restrictive, it allows us to study the effects of vorticity on the oscillation modes of a self-gravitating star in a transparent manner, avoiding complications due to the star's inhomogeneity (multi-layer composition)\\cite{RUDERMAN}. ",
        "conclusions": "Despite a number of simplifying assumptions, the study of the oscillation modes of uniform ellipsoids is useful for understanding the equilibrium and stability of real neutron stars, at least qualitatively. Moreover, the theory is simple enough that it can be extended readily to include modifications to the underlying physics; here, we have considered new features that arise because a neutron star contains a mixture of normal fluid and superfluid coupled by mutual friction. The theory of uniform ellipsoids is interesting from the viewpoint of mathematical physics because it is solvable exactly, and it may also have applications to other physical systems, such as the physics of trapped, rotating Bose-Einstein condensates.  In this paper  we have extended  previous treatments of the oscillation modes of ellipsoidal figures of equilibrium to the case of a mixture of normal fluid and superfluid. The basic equations of motion for the two fluid hydrodynamics include mutual friction exactly, since the frictional forces depend linearly on the relative velocity between the two fluids, and vanish in the background where the two fluids move together. In addition the fluids are coupled via mutual gravitational attraction, which is also treated without further approximations. As a result our relations closely resemble Chandrasekhar's tensor virial equations, even though they are intrinsically dissipative due to the mutual friction. While we have developed these moment equations for general underlying equilibria, they are most useful for perturbations around uniform backgrounds, in which case the moment equations of various orders decouple and yield exact solutions for the normal modes.  Quite generally, there are two classes of modes for small perturbations, one class in which the two fluids move together and the other in which there is relative motion between them. The former are identical to the modes found for a single fluid. As a result our models of superfluid Maclaurin and Roche ellipsoids  undergo dynamical instabilities with respect to these modes  which are indistinguishable from what is found for their classical counterparts. When ordinary viscous dissipation is included they are also subject to secular instabilities related to the modes in which two fluids move together. If the kinematic viscosity is held constant, then the rate of viscous dissipation, when computed in the  ``low Reynolds number'' approximation, is diminished by a factor $f_N$, the fraction of the total mass in the normal fluid (see however below).  The modes involving the relative motion between the fluids are completely new and are shown to be  stable along the entire sequences of the incompressible Maclaurin, Jacobi and Roche ellipsoids independent of the magnitude of the phenomenological mutual friction. These modes also do not become neutral at selected  points along any sequence and the necessary condition for the point of bifurcation is not achieved. Our main conclusion is that mutual friction does not drive secular instabilities in  incompressible and irrotational ellipsoids. In addition we find that even though the new modes are subject to viscous dissipation (a consequence of viscosity of the normal matter), they do not emit gravitational radiation, and are therefore immune to any instabilities associated with gravitational radiation, irrespective of their modal frequencies.  The results summarized above hold within a combined framework of two-fluid superfluid hydrodynamics, Newtonian gravity,  and the ellipsoidal approximation, as formulated in EFE. Each of these elements of our approach contains a number of simplifying approximations which need to be relaxed in realistic applications to neutron stars. For example, to treat the mixture of neutron and proton superfluids in the neutron star cores, the one-constituent two-fluid superfluid hydrodynamics must be replaced by the hydrodynamics of the  multi-constituent superfluid mixtures, in which  case the mutual entrainment of the superfluids and deviations from $\\beta$-equilibrium must be accounted for (see Ref. \\cite{LINDBLOM} for a discussion of these effects and their impact on neutron star oscillations within the real energy functional method). We anticipate that these effects can be incorporated in the tensor virial approach in a perturbative manner and the results of previous sections will hold in leading order of the perturbation expansion. On the other hand, relaxing the incompressible approximation and, hence, including the partial pressures of the superfluid and the normal fluid, will lead to non-perturbative effects as the pressure terms significantly alter the balance between gravitational attraction and  centrifugal stretching. As is well known, the points of the onset of the dynamical and secular instabilities of compressible ellipsoids depend on the adiabatic index of the underlying polytropic equation of state. As noted in Sec. 3, the conclusions reached with respect to the stability of the incompressible ellipsoids should be verified for the compressible models anew. The differences in the pressures (or equations of states) of the normal and superfluid phases, however, are typically small in neutron stars, since the condensation energy is negligible compared to the degeneracy pressures of interacting Fermi-liquids. The coupling between the partial pressures of the noraml fluid and superfluid, therefore, can be treated perturbatively.  Another effect that needs to be included is the density dependence of the mutual friction coefficients and the kinematic viscosity. For example, in neutron stars the kinematic viscosity is density dependent in general, explicitly as a result of  to the density dependence of the phase space of normal quasiparticles undergoing collisions, implicitly because of the density dependence of the in-medium scattering amplitudes (see for further details Ref. \\cite{CUTLER_LINDBLOM}). The ramification for comparison of the secular instabilities of the normal fluid and superfluid ellipsoids is that the modifications of the rate of the viscous dissipation will depend in a non-trivial manner on the fraction of the normal fluid in the system.  The entrainment, $\\beta$-nonequilibrium, compressibility, e. t. c. will couple the relative and center-of-mass modes in general. Such a ``mixing'', as discussed in Sec. 5, implies  $\\xivec_+\\neq 0$ for the relative modes and, similarly, $\\xivec_-\\neq 0$ for the center-of-mass modes. Therefore, the mutual friction might tend to drive the center-of-mass modes secularly unstable; if they emit gravitational radiation, the mutual friction will suppress the gravitational radiation induced instabilities. One of the important issues to be addressed by the future work is the magnitude of the ``mixing'' of the modes for  general equilibria and the corresponding times scales."
    },
    "0004/hep-ex0004007_arXiv.txt": {
        "abstract": "We report on measurements of 11-18~cm wavelength radio emission from interactions of 15.2~MeV pulsed electron bunches at the Argonne Wakefield Accelerator. The electrons were observed both in a configuration where they produced primarily transition radiation from an aluminum foil, and in a configuration designed for the electrons to produce Cherenkov radiation in a silica sand target.   Our aim was to emulate the large electron excess expected to develop during an electromagnetic cascade initiated by an ultra-high-energy particle.  Such charge asymmetries are predicted to produce strong coherent radio pulses, which are the basis for several experiments to detect high-energy neutrinos from the showers they induce in Antarctic ice and in the lunar regolith. We detected coherent emission which we attribute both to transition and possibly Cherenkov radiation at different levels depending on the experimental conditions.  We discuss implications for experiments relying on radio emission for detection of electromagnetic cascades produced by ultra high-energy neutrinos. ",
        "introduction": "In 1962, G. Askaryan\\cite{Ask62} first predicted that large electromagnetic particle cascades would develop an asymmetry of electrons over positrons at the level of 20-30\\% due to the following processes: Compton scattering of atomic electrons in the target material by cascade photons, $\\delta$-ray scattering processes, and annihilation of positrons in flight. Askaryan realized that the propagation of this charge asymmetry in a dielectric medium should result in radio emission due to the processes of Cherenkov and transition radiation.   For wavelengths much larger than the shower dimensions, he noted that the radio emission would be coherent, {\\it i.e.}, the electrons radiate in phase, with a resulting $N_e^2$ increase in the radiated power with the number of excess electrons $N_e$, in contrast with a linear rise expected from incoherent emission.  Recently this suggestion has generated renewed interest both on experimental and theoretical fronts. On the experimental side, it has become the basis for a number of new searches for radio-frequency (RF) emission from cascades induced by ultra-high energy neutrinos---at PeV energies in Antarctic ice \\cite{bol91,fri96,Seck99,Bes99} and at EeV energies from the lunar regolith \\cite{Alv96,Han96,Zhe88,Dag89,Gor99}. On the theoretical side, a number of recent, very detailed simulations\\cite{Alv96,Zas92,Alv97,Alv98} have validated Askaryan's basic result: Particle cascades at energies from $\\sim$1~TeV up to $10^{20}$~eV will develop a net negative charge excess of order 20-30\\% of their total charged-particle number. These simulations also confirm Askaryan's prediction of a corresponding nanosecond pulse of radio Cherenkov emission, with a characteristic frequency and angular dependence. Extending Askaryan's original idea, Markov and Zheleznykh~\\cite{Mark86} have also shown that a number of other coherent radiation processes, such as transition, synchrotron, and bremsstrahlung radiation, are likely to be important contributors to the total emission from cascades depending on the cascade medium, ambient fields, and geometry.  \\subsection{Experimental tests}  Under typical conditions, none of these processes produce radio emission comparable to that of other passbands until the energy of the cascades exceeds $\\sim 10^{15}$~eV. Thus a direct accelerator beam test of these predictions is difficult.   Such a test could be made with existing TeV proton beams where observations of radio emission, though expected to be correspondingly weak and difficult to characterize, would demonstrate the existence of the charge asymmetry. However, a failure to detect radio emission under these circumstances might be attributable to either a failure of the shower to produce the predicted charge asymmetry, or a loss of radio coherence in the emission process after the charge asymmetry has developed.  As an alternative experimental approach, we have utilized lower energy pulsed electron bunches incident on a target material. The electron bunch then emulates the resulting charge asymmetry from the late stages of a high energy cascade, and allows one to focus experimentally on the radio emission process. Simulations have shown\\cite{Zas92} that the bulk of the radio emission arises from low-energy electrons ($\\leq 10$ MeV) produced late in the shower. Thus investigations of the coherent radio emission properties can be done with fairly modest electron energies, provided that the bunch length is at least as small as the longitudinal and transverse shower sizes expected in cascades in solid materials (several~cm).  The total number of charged particles, (which are virtually all electrons and positrons near and after shower maximum) in a cascade is to a good approximation $E_{shower}/\\mbox{GeV}$. Hence, a given cascade energy can be simply related to the number of electrons per pulse at an accelerator.  We have adopted this alternative approach to begin developing a laboratory basis for confirmation of Askaryan's hypothesis. Using the Argonne Wakefield Accelerator (AWA) facility at Argonne National Laboratory, we have employed $20-50$ ps pulsed electron bunches with $N_e \\simeq 10^{10-11}$ electrons per bunch to produce coherent radio emission from several targets.  We looked for radio emission in the 1.7--2.6 GHz microwave band (11-18 cm wavelength).  Our initial measurements show evidence for both coherent transition and perhaps Cherenkov radiation. In the following section we discuss some of the history of related measurements at other facilities.  Section~\\ref{theory} gives a summary of the theoretical basis for the emission.  In section~\\ref{setup} we describe our experimental setup, the measurements made, and the results obtained.  We conclude the paper by considering the applications of these results to coherent radio detection of neutrino-induced cascades.   \\subsection{Related measurements of coherent emission from LINAC beams}  Over the last two decades a significant effort has gone into measurements of coherent Cherenkov radiation (CR) and transition radiation (TR) from electron linear accelerator (LINAC) beams. The primary motivation for these experiments has been the development and characterization of coherent sources of millimeter-wave and far-infrared radiation. Although several early experiments~\\cite{Nei84,Maru86} were done at centimeter wavelengths using continuous-beam devices, most recent work has focussed on the shorter wavelengths. In the early cm-wave work the detected power was attributed to CR from air at the end of the beam pipe. No measurements of the coherence properties were made.  Later work~\\cite{Nak89,Shi91,Ish91} demonstrated the coherence of radiation observed under similar conditions, again attributing it in part to CR. Following this, a number of authors\\cite{Shi91,Tak94,Shi94} realized that the bulk of the observed radiation in such circumstances must in fact be primarily TR, and that the Cherenkov contribution was not easily separable given the experimental configurations used. This is primarily due to two features of these experiments: First, the Cherenkov angular distribution in air is nearly identical to that of TR.  Second, the relatively short effective path-lengths of the electron beam in the air meant that the CR power was suppressed compared to the TR power. ",
        "conclusions": "In the previous section we have presented results which show a clear signature of coherent transition radiation at GHz radio frequencies from electron bunches. Although this result is not surprising given similar results at far infrared and submillimeter wavelengths, it is to our knowledge the first such demonstration at these longer wavelengths.  In addition, we show evidence consistent with a Cherenkov component in the emission from our target. This evidence relies primarily on the design of the target, which produces an angular window where the TR emission is suppressed, and is strengthened by the presence of distinct, linearly polarized component in the emission near the Cherenkov angle, which contrasts sharply with the TR seen at other angles. However, we do not yet observe either Cherenkov or transition radiation that is consistent with the expected power in our experiment.  One possibility is that complete coherence does not fully obtain in either our TR or CR production, due to possible saturation or self-quenching effects in the electron bunch. There is in fact some evidence that a partial loss of coherence is present in the TR runs (see figure~\\ref{trcoh}). This effect appears to set in only at the highest beam charge values. As discussed in the following subsection such a roll-off for the highest energy showers must happen at some point.   \\subsection{Implications for cascade radio emission}  If we take the value of $N_e$ as a measure of the charge excess in a cascade, we can associate our measurements with an equivalent cascade energy $E_e \\simeq N_e/0.2$ GeV, where we have assumed here that the charge excess is 20\\% of the total charge near shower maximum. For the bunch charges in our experiment, $E_e \\simeq 5\\times 10^{20-21}$ eV, corresponding to energies that equal and exceed the highest energy cosmic ray cascades detected to date.  Zas, Halzen, and Stanev~\\cite{Zas92} have shown that high energy cascades in solid ice, under the assumption of full coherence, produce a total radio CR pulse energy $W$ which can be written as \\begin{equation} \\label{cascade-eq} W ~\\simeq~ 5 \\times 10^{-15}~ \\left ( {E_e \\over 1 ~{\\rm TeV} } \\right )^2 \\left ( {\\nu_{max} \\over 1 ~{\\rm GHz} } \\right )^2 ~~{\\rm \\mu J}~ \\end{equation} where $\\nu_{max}$ is the highest frequency observed (in practice loss of coherence sets in for $\\nu_{max} \\geq 2$ GHz). This quadratic dependence of radio power implies a saturation energy: that is, we can equate the total energy in equation \\ref{cascade-eq} with the total energy in the shower to derive an upper limit $E_{max}$ for the maximum energy at which this relation can be valid. Thus we find \\begin{equation} E_{max} ~=~ 3 \\times 10^{24}~{\\rm eV}~. \\end{equation} which is only a factor of $600$ above our highest equivalent cascade energy. However, $E_{max}$ represents the value at which {\\em all} of the shower energy is lost to the RF pulse; in fact a more conservative approach would require some type of equipartition of the shower energy among other energy-loss mechanisms such as ionization, and would require that the radiation reaction of the shower conserve total momentum. Thus we predict that equation \\ref{cascade-eq} must begin to lose validity at some critical energy $E_{crit} = \\epsilon E_{max}$, where we expect that $\\epsilon \\leq 0.1$. Our data in fact suggest $\\epsilon \\sim 0.001$ if the loss of coherence we observe at the highest bunch charges is due to some related effect.  Similar arguments can be applied directly to the theoretical power expected in our experiment from both CR and TR. If we integrate over just the forward angles for our experimental conditions, the total expected power in our passband is $\\simeq 60$ $\\mu$J. The total energy per electron bunch is $N_e E_e \\simeq 2.2 \\times 10^5$ $\\mu$J for 15 nC and 15.2 MeV electrons. Thus the theory predicts that about 1 part in 3600 of the bunch energy to be radiated in our band for fully coherent emission. However, this implies that only a factor of 60 increase in the bunch charge would lead to a complete loss of energy of the bunch to coherent RF emission. Thus it appears that our experimental condition may be more severely affected by possible loss of coherence than equation \\ref{cascade-eq} implies.   \\subsection{Lower limits on cascade power at $E_0 \\geq 10^{20}$ eV }  In spite of the difficulty in finding a satisfactory resolution to the apparent discrepancies in observed vs. theoretical power, we can still use our observed values to derive lower limits on the expected power from both TR and CR from cascades in materials similar to what we have observed. For TR, refractory materials such as sand, lunar regolith, or ice produce nearly the same spectrum as the aluminum window used on our beampipe end. Thus our results are directly applicable to cascades that encounter interfaces in these materials, for example, a cascade that emerges from the material into either air or vacuum is closely analogous to what we have observed for TR. In the CR case, our results are more uncertain but may still provide a preliminary experimental lower limit in estimating cascade emission.  \\subsubsection{Transition radiation}  The highest values in our TR measurements of the pulse energy were $\\sim 6$ $\\mu$J sr$^{-1}$. Converting these to the flux density units commonly used in radio astronomical measurements, we estimate that the TR flux density produced at earth from a cascade of $E_0 = 5 \\times 10^{21}$ eV emerging from the surface of the lunar regolith is \\begin{equation} S (\\theta \\simeq 10^{\\circ}) ~\\geq~ 5000 \\left ( {E_0 \\over 5 \\times 10^{21} ~{\\rm eV}} \\right )^{1.9} ~{\\rm Jy} \\end{equation} where $1$ Jy $\\equiv 10^{-26}$ W m$^{-2}$ Hz$^{-1}$. Since the TR angular distribution appears to be confirmed by our data for these conditions, we can also estimate that the maximum flux density for this case, at an angle of $\\sim 1.5^{\\circ}$ from the cascade axis, is about a factor of 20 higher than at $\\sim 10^{\\circ}$: \\begin{equation} S_{max} (\\theta \\simeq 1.5^{\\circ}) ~\\geq~ 10^5 \\left ( {E_0 \\over 5 \\times 10^{21} ~{\\rm eV}} \\right )^{1.9}~{\\rm Jy}~. \\end{equation} The large antennas used in lunar pulse searches have demonstrated the ability to achieve noise levels of order 400~Jy~\\cite{Gor99} or less. Thus it appears that the energy threshold for detection of events by TR alone is $5 \\times 10^{20}$ eV. Here we have not considered the indications of higher measured pulse energy given by our dipole measurements. If these results are correct, the energy threshold may be an order of magnitude lower.  These results support the claim~\\cite{Mark86} that TR detection of cascades is feasible for the case that a cascade emerges through some surface into a medium transparent to radio emission, and this analysis also applies to detection of cascades on earth. For example, our measurements could be applied to the case of a shower that was initiated in air and then entered ice near a subsurface radio array, such as has been proposed and prototyped in Antarctica~\\cite{fri96,Bes99,Seck99}. In this case, we can write the threshold energy more generally as \\begin{equation} E_{thr}^{TR} ~\\leq~ 5 \\times 10^{20} \\left ( {\\Delta \\nu \\over 900~{\\rm MHz}} \\right )^{-1/2} \\left ( {A_{eff} \\over 2500~{\\rm m^2}} \\right )^{-1/2} \\left ( {  R  \\over 3.8\\times 10^8~{\\rm m}} \\right ) \\left ( {T_{sys} \\over 220~{\\rm K}} \\right )^{1/2} ~{\\rm eV}. \\end{equation} where $A_{eff}$ is the total effective collecting area of the antenna or antenna array (including aperture efficiency), $R$ is the distance to the cascade, and $T_{sys}$ the system temperature. Here we have assumed that full coherence obtains at lower cascade energies; thus the exponent for the energy term above was assumed to be 2.0 rather than 1.9 as we have used at the highest energies.  If we apply this equation to the case of an antenna array in ice comprising 10 m$^2$ effective aperture, operating at $T_{sys}$ of 1000 K with a 900 MHz effective bandwidth, the system will have a TR detection capability for cascades with $E_0 \\simeq 2 \\times 10^{16}$ eV entering the ice from above, out to a distance of 250 m. However, we caution that formation zone effects may be important depending on how rapidly the near-surface refractivity of the ice changes with depth.  \\subsubsection{Cherenkov radiation}  Although we attempted to optimize our experiment in favor or CR detection rather than TR, we cannot yet claim to have yet made measurements of it adequate to allow us to make definite statements about the application of these data to the detection of radio CR from high energy cascades. One thing is much more clear to us after having performed this initial experiment than was evident in the literature: the processes that produce these two forms of radiation are physically very closely aligned.  The problem of separating the two forms of emission is particularly acute when the radiation can be formed in nearly the same physical location, such as at the quartz window entering the target in our case. In fact, the theoretical formulations for TR at this point do not themselves clearly separate the two processes; as we noted in an earlier section, ``Cherenkov-like'' components appear even in the TR equations. Thus, we suggest that a more complete theoretical basis for TR formation under realistic conditions with finite track lengths and boundaries be developed.  An important aspect of the operating regime of our experiment is that the detected Cherenkov power should scale quadratically with both the total track length as well as the number of particles. For solid materials, the mean electron energy near shower maximum for a typical high energy cascade is closer to $\\sim 100 $ MeV rather than the 15 MeV used in our experiment. Thus the total track length is much longer and the Cherenkov production is more directive and therefore more intense near its peak angle. Any scaling from experiments similar to ours must account for this effect.  If we do attribute the observed power of 0.014 $\\mu$J sr$^{-1}$ at our $60^{\\circ}$ angle primarily to Cherenkov radiation, we can make a tentative estimate of the possible energy threshold for CR detection. A cascade of energy $\\sim 10^{20}$ eV in ice or the lunar regolith will have a length of order 10 m\\cite{Zas92} near shower maximum. For this length, the implied angular enhancement in the intensity is of order $10^4$, giving $\\sim 140~\\mu$J sr$^{-1}$, comparable to the observed TR intensity at an angle of $\\sim 1.5^{\\circ}$. The implied cascade energy threshold, using the same arguments as above, is therefore comparable to that implied by our TR results. For the lunar regolith, $E_{thr}~\\approx~ 5 \\times 10^{20}$ eV, similar to values estimated from complete electromagnetic simulations of ultra-high energy cascades~\\cite{Zas92,Alv96,Alv97}. If our assumptions are correct, this value is conservative, and we expect further experiments to yield a potentially much lower cascade energy threshold for coherent Cherenkov detection.  \\subsection{Implications of RF lensing effects}  In concluding this section we note that the lensing effects on the forward-beamed TR emission that we have observed in our target-full configuration have potentially important implications for experiments to detect cascade emission from the lunar regolith or other materials observed through a refractive interface. Under these conditions, our results suggest that total internal reflection of the cascade radio emission can strongly suppress it at certain angles and under certain geometrical configurations.  This effect is potentially helpful in distinguishing particles such as hadronic cosmic rays or photons (which interact immediately upon entering the refractive material) from neutrinos (which can interact near the surface of the material after traversing it over large distances). In the latter case, since the critical angle for total internal reflection is the complement of the Cherenkov angle at a vacuum interface, the CR can exit the surface into the vacuum.  However, in the former case, the CR from the cascade will suffer total-internal reflectance and will not emerge from the surface. Clearly, surface irregularities at scales approaching an RF wavelength will modify this conclusion at some level, but to first order this effect will tend to suppress the detection of cosmic ray events relative to upcoming neutrino cascades in lunar observations."
    },
    "0004/astro-ph0004325_arXiv.txt": {
        "abstract": "Given sufficiently deep optical spectroscopy, the age estimation of high-redshift ($z > 1$) galaxies has been claimed to be a relatively robust process ({\\it e.g.} Dunlop et al. 1996) due to the fact that, for ages $< 5$Gyr, the near-ultraviolet light of a stellar population is expected to be dominated by `well-understood' main-sequence (MS) stars. Recently, however, the reliability of this process has been called into question by Yi \\etal (2000), who claim to have developed models in which the spectrum produced by the main sequence reddens much more rapidly than in the models of Jimenez \\etal (2000a), leading to much younger age estimates for the reddest known high-redshift ellipticals. In support of their revised age estimates, Yi \\etal cite the fact that their models can reproduce the spectrum of the Sun at an age of 5 Gyr, whereas the solar spectrum is not reproduced by the Jimenez \\etal models until $\\simeq 10$ Gyr. Here we confirm this discrepancy, but point out that this is in fact a {\\it strength} of the Jimenez \\etal models and indicative of some flaw in the models of Yi \\etal (which, in effect, imply that the Sun will turn into a red giant any minute now). We have also explored the models of Worthey (1994) (which are known to differ greatly from those of Jimenez \\etal in the treatment of post-MS evolution) and find that the main-sequence component of Worthey's models also cannot reproduce the solar spectrum until an age of 9-10 Gyr. We conclude that either the models of Yi \\etal are not as main-sequence dominated at 4-5 Gyr as claimed, or that the stellar evolutionary timescale in these models is in error by a factor possibly as high as two. Our current best estimate of the age of the oldest galaxies at $z \\simeq 1.5$ thus remains $3-4$ Gyr, as we confirm with a new analysis of the existing data using the updated solar-metallicity models of both Jimenez \\etal and Worthey. Finally, by fitting a mixed metallicity model to the Sun, we demonstrate that, given rest-frame ultraviolet data of sufficient quality, it should be possible to break the age-metallicity degeneracy when analyzing the spectra of high-redshift galaxies. ",
        "introduction": "For over a decade now, astronomers have attempted to estimate the ages of high-redshift galaxies using broad-band optical-infrared photometry ({\\it e.g.} Lilly 1988; Dunlop et al. 1989, Chambers \\& Charlot 1990). Unfortunately, however, the derived ages have been rendered virtually meaningless by disagreements between modellers over post main-sequence evolution ({\\it e.g.} Charlot, Worthey \\& Bressan 1996), and by the extreme susceptibility of such relatively crude broad-band data to dust reddening, emission-line contamination etc.  In contrast, it has long been anticipated that relatively robust age constraints for high-redshift galaxies could be derived given rest-frame near-ultraviolet spectra of sufficient quality. This is because, for the potential ages of interest at $z > 1$ (i.e. ages $<$ 5 Gyr), the ultraviolet light of a stellar population is expected to be dominated by stars close to the turn-off point of the `well-understood' main sequence (MS) ({\\it e.g.} Magris \\& Bruzual 1993).  With the advent of deep optical spectroscopy on 10-m class telescopes, it has now proved possible to put this technique into practice. In particular, Dunlop et al. (1996) were able to use a deep Keck spectrum of the $z = 1.5$ radio galaxy LBDS 53W091 to first confirm that its near-ultraviolet spectrum was indeed dominated by starlight, and then to extract an age constraint of $ > 3$Gyr based primarily on comparison with a main-sequence only model of an evolving stellar population. Spinrad et al. (1997) explored further the reliability of this age estimate, and confirmed that the best agreement between ages derived using alternative evolutionary synthesis models was obtained if fitting was confined to the detailed shape of the near-ultraviolet spectral energy distribution.  Not surprisingly, given its implications for cosmology (for $H_0 = 70 {\\rm km s^{-1} Mpc^{-1}}$, the age of an Einstein-de Sitter universe at $z = 1.5$ is only 2.3 Gyr), this result has been the subject of subsequent close scrutiny, and claims that 53W091 is in fact less than 2 Gyr old have been put forward by, for example,  Bruzual \\& Magris (1997). However, Dunlop (1999) has argued that such young ages are only deduced using some models if the near-infrared photometry is also included in the fitting process, once again placing undesirable emphasis on the reliability of the modelling of post main-sequence evolution (a point previously also explored by Spinrad et al. 1997). Moreover, Dunlop (1999) has shown that, certainly for the slightly redder $z = 1.43$ galaxy 53W069, if fitting is confined to the Keck spectroscopic data (Dey et al. 2000), the models of Bruzual \\& Charlot (1993), Worthey (1994), and Jimenez et al. (2000a) {\\it all} lead to the conclusion that its stellar population is $> 3$ Gyr old (assuming solar metallicity).  Most recently, however, the reliability of even this near-ultraviolet spectroscopic age-dating has been called into question by Yi et al. (2000). Yi \\etal (2000), claim to have derived a much younger age for 53W091, but also claim that this age is not due to differences in post-MS evolution, but rather to the fact that the spectrum produced by the main sequence in their models reddens much more rapidly than in the models of Jimenez \\etal (2000a). In support of their revised age estimates, Yi \\etal cite the fact that their models can reproduce the spectrum of the Sun at an age of 5 Gyr, whereas the solar spectrum is not reproduced by the Jimenez \\etal models until 8-10 Gyr. It is unclear to us why a stellar population should be expected to mimic the spectrum of the Sun at its current age ($\\simeq 5$ Gyr); even if the light from the stellar population is dominated by stars near the main-sequence turnoff the Sun is not expected to leave the main sequence until an age of $\\simeq 10$ Gyr (Jorgensen 1991). Nevertheless this claim has motivated us to explicitly check the calibration of the main-sequence evolution in alternative evolutionary synthesis models of galaxy evolution.  This is the main subject of the present paper. What we have done is to take the 3 alternative and independent models of galaxy evolution developed by Yi et al. (2000), Jimenez et al. (2000a) and Worthey (1994), and to check how rapidly they evolve to mimic the solar spectrum with and (more importantly) without inclusion of their post-MS components. The models are summarized in section 2, and the results of comparison with the solar near-ultraviolet spectrum are presented in section 3. We then proceed, in section 4, to use these models (again with and without post-MS components) to check explicitly the extent to which the age estimates of 53W091 and 53W069 really are affected by different approaches to modelling post-MS evolution. The main remaining uncertainty is the impact of having to assume a value for the metallicity of the stellar population, and in section 5 we explore whether, given near-ultraviolet data of sufficient quality, it may be possible to break the well-known age-metallicity degeneracy. Finally, our conclusions are summarized in section 6. ",
        "conclusions": ""
    },
    "0004/astro-ph0004113_arXiv.txt": {
        "abstract": "New \\ion{H}{1} synthesis data have been obtained for six face--on galaxies with the Very Large Array. These data and reanalyses of three additional data sets make up a sample of nine face--on galaxies analyzed for deviations from axisymmetry in morphology and dynamics. This sample represents a subsample of galaxies already analyzed for morphological symmetry properties in the {\\it R-\\/}band. Four quantitative measures of dynamical nonaxisymmetry are compared to one another and to the quantitative measures of morphological asymmetry in \\ion{H}{1} and {\\it R-\\/}band to investigate the relationships between nonaxisymmetric morphology and dynamics. We find no significant relationship between asymmetric morphology and most of the dynamical measures in our sample. A possible relationship is found, however, between morphology and dynamical position angle differences between approaching and receding sides of the galaxy. ",
        "introduction": "} Despite the fact that most studies of spiral galaxy dynamics concentrate on understanding the properties of axisymmetric disks, evidence is accumulating that many galaxies lack such overall symmetry. Baldwin {\\it et al.\\/} (1980)\\markcite{blbs} were first to seriously examine the asymmetries in galactic disks, pointing out that lopsided galaxies were a common phenomenon not localized to interacting pairs.  More recently, frequency of morphological asymmetry in galactic disks has been quantitatively studied at optical wavelengths. Based on optical morphology, approximately 30\\% of disk galaxies exhibit significant ``lopsidedness'' (Rix \\& Zaritsky 1995\\markcite{rz95}; Kornreich {\\it et al.\\/} 1998\\markcite{khl}, hereafter KHL). These studies are based on data sets containing some 30 targets each.  Dynamical asymmetry, too, has been examined in large samples of single--dish \\ion{H}{1} line profiles. Richter \\& Sancisi (1994)\\markcite{rs94}, followed by the 104 galaxy sample of Haynes {\\it et al.\\/} (1998, hereafter HHMRvZ)\\markcite{h98a}, examined the symmetry properties of single--dish \\ion{H}{1} line profiles and determined that as many as $\\sim50$\\% of spiral galaxies show departures from the expected symmetric two--horned profile. Nevertheless, asymmetries in line profiles are ambiguous evidence at best for disturbed dynamics, since they combine both dynamic and spatial information.  Until recently, however, little work had been done to examine the symmetry properties of neutral hydrogen in synthesis data. As a result, the connection between disturbed morphology and disturbed dynamics in field disk galaxies is only now beginning to be seriously studied. Schoenmakers {\\it et al.\\/} (1997, hereafter SFdZ)\\markcite{SFdZ} outline a method for measuring small deviations from axisymmetry of the potential of a filled gas disk by breaking down the observed velocity field into its harmonic components. Recently, Swaters {\\it et al.\\/} (1999, hereafter S3vA)\\markcite{S3vA} have applied this method to the \\ion{H}{1} synthesis data of two galaxies, DDO~9 and NGC~4395, where the hallmark of dynamical asymmetry is found to be asymmetry in the rotation curve, where one side of the curve rises more steeply than the other. \\ion{H}{1} synthesis observations of several other galaxies, e.g. NGC~3631 (Knapen 1997\\markcite{knapen}), NGC~5474 (Rownd {\\it et al.\\/} 1994, hereafter RDH\\markcite{rownd}), and NGC~7217 (Buta {\\it et al.\\/} 1995\\markcite{buta}), have also curiously revealed asymmetries or offsets between the optical centers of light and the kinematic centers of the neutral gas. While much evidence exists that many field spirals exhibit non--axisymmetric morphology, and that many field spirals exhibit asymmetric dynamics, the question of whether the two phenomena are correlated, or even represent two pictures of a single underlying effect, remains open.  Both optical morphology and gas dynamics provide clues as to the overall structure of a galaxy. For instance, Zaritsky \\& Rix (1997)\\markcite{zr97} proposed that optical lopsidedness arises from tidal interactions, minor mergers, or possibly gradual accretion. Conselice {\\it et al.\\/} (2000)\\markcite{conselice}, on the other hand, find a correlation between optical asymmetry and \\bv\\ color, and are able to use morphological asymmetry to identify whether starbursts in a given galaxy are likely caused by interactions and mergers. Similarly, the neutral hydrogen dynamics of asymmetric galaxies should be able to distinguish between tidally deformed galaxies and those which are dynamically isolated. For instance, while the strongly optically--lopsided galaxy NGC~5474 is well--known to be under the tidal influence of its neighbor M101, the disturbed morphologies of the other relatively isolated objects in the KHL sample are not well--explained by standard tidal interaction models, which require particular dynamics as well as interactions with (unobserved) nearby companions.  Alternatively, asymmetries may arise from the excitation of unstable, one--armed spiral modes possibly triggered by a past interaction or minor merger (Taga \\& Iye 1998a, 1998b; Lovelace {\\it et al.\\/} 1999)\\markcite{love}\\markcite{taga1}\\markcite{taga2}. Librations of the optical galaxy about the minimum of the gravitational potential might also be set in motion by a previous interaction. As density is dependent on radius in a galaxy, one might expect that the natural modes of a galaxy also depend on radius. The resulting differential oscillation might result in both lopsided appearance and kinematic decoupling from the optical light distribution. An understanding of these modes, if they exist, could provide a direct measure of the dark matter distribution (Jog 1997\\markcite{jog}; SFdZ\\markcite{SFdZ}).  Another type of nonaxisymmetry, warping of galactic disks due to non-planar motions, has been proposed as an indicator of inclined flattened halo potentials (e.g. Dekel \\& Shlosman 1983\\markcite{ds83}; Toomre 1983\\markcite{T83}), which are required to stabilize the warp against differential precession, as well as observational indicators of massive dark halos (Tubbs \\& Sanders 1979)\\markcite{sand}. Warps may also be due to tidal interactions and accretion, and have been reported to be related to non--circular motions, particularly $m=1$ modes (Weinberg 1998)\\markcite{W98}. Although ``sloshing'' librations in the plane of the galaxy would not imply a correlation between warps and lopsided dynamics, ``flapping'' librations normal to the galactic plane could contribute to warping.  In this paper, we present HI synthesis data for nine galaxies whose optical asymmetry properties were quantified by KHL. Data obtained from the archives at the Very Large Array (VLA)\\footnote{The VLA is a facility of the National Radio Astronomy Observatory.} have been reanalyzed for the galaxies NGC~5474, NGC~5701, and UGC~12732, and new VLA data are presented for the galaxies NGC~991, NGC~1024, UGC~3685, NGC~3596, UGC~6420, and NGC~4688. These galaxies represent a sample of face--on galaxies selected on the basis of their optical properties.  In \\S \\ref{obs}, we discuss the acquisition and reduction of the new and archived data for our sample of nine galaxies. In \\S\\ref{correlations}, we analyze the data obtained for deviations from morphological and dynamical axisymmetry and describe methods of determining the magnitudes of such deviations. These quantitative measures are useful for quantifying warps in face--on galaxies and for non--circular motions in more inclined galaxies. We then use the sample to draw correlations between the symmetry parameters of the galaxies to one another and to global \\ion{H}{1} properties. Finally, in \\S\\ref{conclusions}, we present our conclusions in the context of other recent work on galaxy asymmetry. ",
        "conclusions": "}  This work represents a concurrent analysis of both optical and 21-cm symmetry properties of a sample of disk galaxies. We have presented methods by which the symmetry properties of a galaxy can be quickly computed over a wide range of dynamic variables. These parameters are sensitive to both non-circular and non-planar motions, depending in large part on the inclinations of the galaxies. In this case, the face--on inclinations of the targets make the parameters most sensitive to non--planar motions.  In our sample presented here, tidal interaction with companion objects dominates the asymmetry in NGC~5474 and cannot be ruled out for NGC~1042 and NGC~4688. Nevertheless, dynamical nonaxisymmetries clearly prosper in more isolated and even morphologically symmetric galaxies such as UGC~12732 and UGC~3685 as well. Pervasive deviations from axisymmetry even outside of tidal interactions seem to indicate that asymmetric galaxies are stable and long--lived phenomena. Our data seem, however, to indicate that there is little or no connection between lopsided morphology and non--planar motions in such galaxies. This lack of clear connection between morphology and warping seems to imply that separate physical mechanisms may be at work in each domain to produce deviations from flat, axisymmetric disks.  The methods presented in this paper can be quickly applied to large samples of galaxies to identify deviations from axisymmetry in dynamics and morphology. The small sample size of this initial study and the kinematic uncertainties produced by the sample galaxies' face--on inclinations can only rule out the strongest of correlations among the various parameters. Larger samples and samples at varying inclinations are required to truly study the relationships between these parameters at more than a cursory level. A future study with galaxies of a wider range of inclinations will also allow for additional parameters such as kinematic offsets which can then be adequately studied. Future numerical work can also benefit from an easily applied method to track and identify departures from symmetry. Such numerical work could, for instance, study the origins and longevity of asymmetries, and predict the course of their evolution in spiral galaxies."
    },
    "0004/astro-ph0004263_arXiv.txt": {
        "abstract": "We present new redshift measurements for 55 galaxies in the vicinity of the rich galaxy cluster Abell 665. When combined with results from the literature, we have good velocity measurements for a sample of 77 confirmed cluster members from which we derive the cluster's redshift $z=0.1829\\pm 0.0005$ and line-of-sight velocity dispersion $\\sigma = 1390^{+120}_{-110}\\,\\rm km\\, s^{-1}$. Our analysis of the kinematical and spatial data for the subset of galaxies located within the central 750 kpc reveals only subtle evidence for substructure and non-Gaussianity in the velocity distribution. We find that the brightest cluster member is not moving significantly relative to the other galaxies near the center of the cluster.  On the other hand, our deep \\rosat\\ high resolution image of A665 shows strong evidence for isophotal twisting and centroid variation, thereby confirming previous suggestions of significant substructure in the hot X-ray--emitting intracluster gas. In light of this evident substructure, we have compared the optical velocity data with N-body simulations of head-on cluster mergers. We find that a merger of two similar mass subclusters (mass ratios of 1:1 or 1:2) seen close to the time of core-crossing produces velocity distributions that are consistent with that observed. ",
        "introduction": "The study of galaxy clusters has revealed them as powerful probes of such cosmological quantities as the baryon fraction of the Universe (e.g., White \\& Fabian 1995), $\\Omega_0$ (Richstone et al.~1992), and the Hubble constant $H_0$ (Gunn 1978; Birkinshaw 1979).  However, our knowledge of the physics of galaxy clusters has not yet reached the same level of understanding that we have gained for, arguably, the most important cosmological probe to date, i.e., Cepheid variables, and this remains as one of the most significant limitations in our use of clusters for cosmological studies. For instance, one can compute $H_0$ by combining measurements of the decrement in the brightness temperature of the cosmic microwave background radiation (CMBR) caused by the inverse Compton scattering of CMBR photons by the hot electrons in the cluster (the Sunyaev-Zel'dovich effect, Sunyaev \\& Zel'dovich 1972) with spectral and imaging observations of the X-ray emission produced by the same hot gas. This technique requires accurate 3-D models of the properties (temperature, density, metallicity, etc.) of the cluster atmosphere, which in turn demands that the physics and astrophysics of clusters be well understood.  However, this understanding presents a problem because even present-day clusters are dynamically young and active, showing evidence for the accretion and merger of other systems. The resulting rich complexity in their internal properties greatly complicates their use for precision cosmology.  Abell 665 was initially classified by Abell (1958) as the richest cluster in his catalog. As such, it has been the subject of considerable study across the wavebands. Evidence for subclustering in the spatial distribution of the galaxies was first presented by Geller \\& Beers (1982). This was later confirmed through BVR photometry of 178 galaxies by Kalloglyan et al.~(1990). The luminosity function has been recently studied by Garilli et al.~(1996), Wilson et al.~(1997), and Trentham (1998) for comparison with other clusters and in an effort to evaluate models for galaxy evolution.  Despite its richness, only 33 cluster members in A665 have published redshifts (Oegerle et al.~1991, hereafter OHFH). OHFH examined the kinematical properties of these galaxies and found that the velocity distribution was well described by a Gaussian with a relativistically-corrected line-of-sight velocity dispersion of $\\sigma = 1201_{-126}^{+183}\\, \\rm km\\, s^{-1}$.  This seemed to indicate a fairly relaxed, but massive cluster. However, these authors did not completely reject the possibility that the cluster could be more complex.  They suggested that the observed spatial substructure in the galaxy distribution and the relatively large peculiar velocity (v$_{pec}$ = 447 km s$^{-1}$) of the brightest cluster member (BCM) argued for a non-relaxed dynamical state for A665. We will have more to say on this point later.  A665 is also a luminous source of X-ray emission, as its optical richness and high velocity dispersion would suggest.  The presence of hot gas in this cluster was demonstrated by early observations performed by the {\\it Einstein Observatory} and the {\\it Ginga} satellite (Birkinshaw, Hughes, \\& Arnaud 1991; Hughes \\& Tanaka 1992). These pioneering observations revealed some evidence that the spatial distribution of the X-ray--emitting gas deviated from circular symmetry. Birkinshaw et al.~modeled the complex gas distribution and found that a combination of two isothermal-$\\beta$ models separated by $\\sim$3$^\\prime$ provided a considerably better fit to the data than a single isothermal-$\\beta$ model did. Hughes \\& Birkinshaw (1996, 2000) have applied this type of modeling in much greater detail to \\rosat\\ PSPC data and came to the conclusion that the properties of the X-ray emission required that a major merger was occurring or had occurred recently in this system. Note that Markevitch (1996) also suggested that a recent merger could explain the temperature gradient and asymmetric X-ray emission detected in \\asca\\ observations of A665.  Further support for a merger scenario comes from Buote \\& Tsai's (1996) study of the X-ray morphologies of a sample of 59 bright X-ray clusters. Their work quantifies cluster substructure by characterizing the X-ray surface brightness in terms of a multipole expansion of the two-dimensional gravitational potential.  They detect a tight correlation between specific multipole power ratios that suggests an evolutionary track for clusters (i.e., the location of a cluster in the $P_2/P_0$--$P_4/P_0$ power ratio plane is a function of its dynamical state). Interestingly, Buote \\& Tsai find that the only cluster in their entire sample that deviates from this correlation is A665. They hypothesized that A665 is undergoing a major merger event and is in a brief period of its evolution when the X-ray emitting hot gas does not follow the dark matter distribution or the gravitational potential.  The goal of this paper is to present new velocity measurements for galaxies in A665 and, by combining them with data from the literature, to study the cluster's dynamics and to investigate the parameters of the cluster merger that we infer to have taken place recently. In section~2 we describe the acquisition and reduction of new optical and \\rosat\\ X-ray data. In the following section we present our analysis of these data. Then in section 4 we present our model for the dynamical state of A665. We summarize our conclusions in section 5. We use $H_0$=75 km s$^{-1}$ Mpc$^{-1}$ and q$_0$=0.5 throughout the paper. ",
        "conclusions": ""
    },
    "0004/astro-ph0004055_arXiv.txt": {
        "abstract": "In the gaseous envelope of protogalaxies, thermal instability leads to the formation of a population of cool fragments which are confined by the pressure of a residual hot background medium. In order to remain in a quasi-hydrostatic equilibrium, the residual gas evolves at approximately the virial temperature of the dark matter halo.  Its density is determined by the requirements of thermal equilibrium. The hot gas is heated by compression and shock dissipation.  The heating is balanced by direct energy loss due to bremsstrahlung emission, and by conductive losses into the cool clouds, which are efficient radiators.  The cool fragments are photoionized and heated by the extragalactic UV background and nearby massive stars.  Several processes interact to determine the size distribution of the cool fragments.  The smallest are evaporated due to conductive heat transfer from the hot gas. All fragments are subject to disruption due to hydrodynamic instabilities. The fragments also gain mass due to collisions and mergers, and condensation from the hot gas due to conduction.  The size distribution of the fragments in term determines the rate and efficiency of star formation during the early phase of galactic evolution.  We have performed one-dimensional hydrodynamic simulations of the evolution of the hot and cool gas.  The cool clouds are assumed to follow a power-law size distribution, and fall into the galactic potential, subject to drag from the hot gas.  The relative amounts of the hot and cool gas is determined by the processes discussed above, and star formation occurs at a rate sufficient to maintain the cool clouds at 10$^4$~K.  We present density distributions for the two phases and also for the stars for several cases, parametrized by the circular speeds of the potentials. Under some conditions, primarily low densities of the hot gas, conduction is more efficient than radiative processes at cooling the hot gas, limiting the x-ray radiation from the halo gas. ",
        "introduction": "According to CDM models of galaxy formation (Blumenthal et al. 1984; Navarro, Frenk, \\& White 1997; Klypin, Nolthenius, \\& Primack 1997), large galaxies emerge through the coalescence of smaller systems, normally identified as dwarf galaxies, which contain mostly dark matter with a fraction of gas. The dark matter components merge with each other in a dissipationless manner to form the halos of larger galaxies.  The crossing of gas streamlines leads to shocks, heating the gas to the virial temperature of the galactic halo (Binney 1977; Rees \\& Ostriker 1977; White \\& Rees 1978).  In order for the heated gas in this extended halo (hereafter, the protogalactic cloud or PGC) to collapse, its cooling timescale $\\tau_c$ must be shorter than its dynamical time scale, $\\tau_d$ during each growth stage of the galaxy.  When large galaxies acquire a mass comparable to the Galaxy, this condition is satisfied once the characteristic length scale of the PGC is $< 100$ kpc (Blumenthal et al. 1984).  For typical values ($\\sim 10^6$ K) of the virial temperature, the cooling timescale increases with temperature, and the PGC's are thermally unstable (Field 1965). Thermal instability leads to the rapid growth of perturbations and fragmentation of PGCs (Murray \\& Lin 1990). The result is that a two-phase medium develops during the initial cooling of the PGC, in which a population of warm fragmentary clouds (WFC's) are confined by the pressure of hot, residual halo gas (RHG) (Burkert \\& Lin 2000).  The RHG is primarily heated by the release of the gravitational energy of the collapsing PGC and cooled by radiative emission and conductive transport.  The WFC's settle into the central region of the halo potential.  They are unable to cool below 10$^4$~K until their density reaches a sufficiently high value that the WFC's become self-shielded from external photodissociating UV radiation (Couchman \\& Rees 1986; Haiman, Rees, \\& Loeb 1997; Haiman, Abel, \\& Rees in press).  Thereafter, molecular hydrogen can form within them as a consequence of non-equilibrium cooling.  Efficient cooling reduces the temperature to $T < 100$~K even in the absence of metals. Consequently, WFC's evolve into cold molecular clouds (CMCs).  The formation of massive stars within the CMC's provides ultraviolet heating, limiting the formation rate of CMC's from WFC's, and leading to the formation of an equilibrium three-phase structure in the galaxy.  In this paper, we discuss the evolution of the hot phase that is left after the initial thermal instability.  Detailed investigations of the warm and cold phases will be made in upcoming papers.  For the current work, we assume a power-law size distribution for the WFC's. The CMC's are not included explicitly, but star formation within them is explicitly assumed to occur at a rate which would provide an adequate flux of UV photons to heat and ionize the WFC's.  In \\S2, we briefly recapitulate the physical processes associated with this scenario.  In \\S~3, we determine the energy balance between the warm and hot phases, while in \\S~4, we discuss the resulting star formation rate in the WFC's.  In \\S~5, we show model results on the evolution of the warm and hot phases.  Finally, we summarize our results and discuss their implications in \\S~6. ",
        "conclusions": "In this paper, we examine the microphysics of gas dynamics in the early epoch of galactic evolution.  Our objective is to provide a description of the dominant physical processes in ordinary matter which may be applied to a large class of galaxy formation models. For example, in the canonical hierarchical galaxy formation scenario, small dwarf-galaxy building blocks, containing non-interactive dark matter and gaseous ordinary matter, form first and subsequently merge to form larger entities such as our Galaxy.  In the present analysis, we neglect the dynamical evolution of the dark matter halo which is undoubtedly important in determining not only the formation process but also the present kinematic properties of galaxies.  This approximation enables us to focus our attention upon the evolution of gas in the early epoch of galactic evolution, which regulates the rate and location of star formation and therefore the light distribution and chemical properties of the emerging galaxies.  We discuss below three possible applications of our results.  \\subsection{Galactic Stellar halo} The results presented here show that thermal instability results in the formation of the residual halo gas (RHG), with density and pressure appropriate for quasi-hydrostatic and energy equilibria, and warm fragmentary clouds (WFC's), which are pressure confined by RHG and heated by ionizing UV photons.  Such systems have been proposed as being the source of observed Lyman-limit systems (Mo \\& Miralda-Escude 1996).  The mass limits of the WFC's are set by the same criteria as in the earlier work, but in this work we include many more details of the interactions between the phases.  Inside $\\sim 10-100$~kpc, WFC's in galaxies with masses comparable to the Galaxy are self-shielded from the extragalactic UV flux (see Figure 1). Unless these WFC's are continually heated by the UV flux from nearby massive stars, further cooling reduces their temperature to $\\leq 10^2$ K.  Gravitational instability in the large WFC's leads to spontaneous formation of stars among which the massive stars radiate UV photons, ionize their surroundings, and quench the formation of additional stars during their lifetime (Lin \\& Murray 1992).  Through such a self-regulating feedback process, the maximum rate at which gas may be converted into stars is determined by the maintenance of an adequate UV flux to photoionize all the WFC's (\\S4.2).  In the regions far from the center of the halo where the density of the RHG is relatively low, the internal density of the WFC's and their average density $<\\rho_w>$ are also low.  The effect of self-regulation limits the star formation timescale $\\langle\\rho_w\\rangle / \\dot{\\rho}_{SF} > \\tau_d$.  But, the star formation efficiency is much higher at smaller galactic distances.  The results in Figures 2-4 show that within a few kpc (depending on the model), the stellar density $\\rho_\\ast$ already exceeds both $<\\rho_w>$ and $\\rho_h$ after 1 Gyr. Once the stars are formed out of gas, they cannot dissipate their orbital energy such that their orbital radii cannot contract further. Thus, the interaction between RHG, WFC's, and the newly formed massive stars essentially determines the asymptotic surface brightness distribution in galaxies.  In Figure~5, we show the computed surface density profiles computed from the stellar distributions of Models~1-3.  As can be seen, the models are fairly well fit by deVaucouleur profiles within the inner 50~kpc, in good agreement with the profiles of spheroidal systems.  The onset of rapid and efficient star formation invalidates the instantaneous mixing and incremental gas-to-star conversion assumptions which are essential to the closed-box models for galactic chemical evolution and enrichment (cf Binney \\& Tremaine 1987), despite the apparent consistency between it and the observed metalicity distribution among population II stars (Chiappini et al. 1999). If these stars are formed in a series of starburst events in which a large fraction of the remaining gas is converted into stars on a timescale shorter than the dynamical timescale of the galaxy or the lifespan of the massive stars, the metallicity distribution would reflect the metallicity inhomogeneity in WFC's.  In this case, the deficiency of extremely metal poor stars (with [Fe/H]$\\lapprox$-3) would be consistent with an evolving initial mass function which gradually becomes less biased towards massive stars as the WFC's are chemically enriched.  A necessary condition for the formation of stars with long-lived, low-mass stars is small $M_{crit}$ (see \\S4.2).  In a self-regulated environment, WFC's with [Fe/H]$\\gapprox$-3 can spontaneously cool from $10^4$ K to $\\sim 10$ K (so that $M_{crit}\\lapprox1M_\\odot$) between successive generations of nearby massive stars (with an interval $\\sim$ a few $10^6$ yr) provided the initial $n_w > 0.1-1$ (Hellsten \\& Lin 2000).  The corresponding external pressure needed to confine such WFC's is $nT\\sim 10^{3-4}$.  In Figures~2-4, we see that these values of $n_h T_h$ are attained outside 10~kpc, similar to the regions where the metal poor population II stars are located.  In the above scenario, we have neglected the effect of magnetic fields which stablize cold clouds against gravitational instability through field-ion coupling and ion-neutral collisions (cf Shu 1985; Shu, Adams, \\& Lizano 1993).  In dense cores of molecular clouds around the solar neighborhood, low-level ionization is maintained through cosmic ray heating and the field strength declines through ambipolar diffusion (Spitzer 1978; Shu 1985). If this process is important in WFC's, the star formation rate would be much reduced from those illustrated in Figures~2-4.  Collisions between WFC's would occur at velocities in excess of 10~km~s$^{-1}$. Shock compression and rapid cooling near the collision interface could lead to a rapid expulsion of magnetic field, rendering its support ineffective.  In the models discussed above, the covering factor of large WFC's is of order unity.  A large cloud would therefore collide with another large cloud only about once during a galactic crossing time. Collisions between large and small clouds would, however, occur much more frequently.  In the limit that the size distribution of the clouds is dominated by relatively small clouds, the area covering factor is large, and small clouds collide with each other much more frequently.  These issues are beyond the scope of the present investigation and they need to be thoroughly investigated in the future.  \\subsection{Formation of globular clusters} The analysis presented here can also be applied to the formation of Galactic globular clusters (Lin \\& Murray 1996). Prior to the conversion of ordinary matter from gas into stars, the progenitors of these clusters were protocluster clouds (PCC's).  The chemical homogeneities within individual clusters and the large metallicity variation among different clusters suggest that PCC's are a distinct entities which must be confined either by their own self-gravity or external pressure. But, if PCC's are entirely bound by their own self-gravity, external UV heating would not be adequate to suppress thermal instability within them (see \\S3.3).  The magnitude of $M_{crit}$ for the WFC's is comparable to the mass of globular clusters.  We identify these warm, marginally self-gravitating and partially pressure-confined WFC's as PCC's.  At Galactic distances $D \\sim 3-30$ kpc, PCC's with mass ($M$) $ < M_{crit} (\\sim 10^6 M_\\odot)$ are confined by the the pressure of the RHG, $n_h T_h \\sim 10^{3-5}$, depending on the Galactic halo structure during the epoch of cluster formation (Figs. 2-4).  If these PCC's are completely ionized and have a $T_w \\sim 10^4$, then $n_w \\sim 0.1-10$~cm$^{-3}$.  For these structural parameters, the UV flux needed is equivalent to that emitted by a few O5 stars at a distance comparable to or greater than their size (typical a few pc) (see Figure~1). (These stars could also reside within the PCC's).  These clouds could persist for a significant fraction of $\\tau_d$ if the accretion of smaller clouds or condensation from the RHG is adequate to compensate for their mass loss due to stripping by the RHG.  On the observational side, in order to verify that the PCC's were pressure-confined, we first estimate their $n_w T_w$ from the current properties of globular clusters, averaged over their half-mass radii ($r_h$) (Murray \\& Lin 1996).  We use these quantities because the stellar density and the velocity dispersion at $r_h$ do not change significantly during post-formation evolution.  Extrapolation to the stage prior to star formation is, however, highly uncertain. If, after their formation, the young stellar objects undergo collapse and virialization from rest, the clouds' initial radii ($r_i$) would be $\\sim 2 r_h$.  Larger $r_i$ would be expected if star formation requires dissipative collisions and coagulation of substellar fragments (Murray \\& Lin 1996). But $r_i$ is unlikely to be larger than the tidal radii of the PCC's, which are typically a few times larger than $r_h$.  Thus, the initial density of the PCC's may be 1-3 orders of magnitude smaller than the average cluster density at $r_h$ today.  Based on the present velocity dispersion of the clusters, we infer the initial temperature of the PCC's to be $\\sim 10^4$~K, comparable to that expected if they were photoionized.  From these estimates, we infer $n_w T_w \\sim 10^{2-5}$ (Murray \\& Lin 1992). The dependence of the pressure upon galactic radius, $D$, is very poorly determined from the observational parameters.  Of more significance is that the magnitude of the pressure inferred from the observations is very similar to that found in the RHG in our models (Figs. 2-4).  From these results, and the cluster metallicities, we can also estimate the cooling time scale ($\\tau_{cc}$) and dynamical time scale ($\\tau_{dc}$), of the PCC's. The ratio $\\tau_{cc}/\\tau_{cd}$ increases from $\\sim 10^{-4}$ near the Galactic bulge to $\\sim 1$ at $\\sim 100$ kpc.  In most PCC's, $\\tau_{cc} < < \\tau_{cd}$ and thermal equilibrium is only possible in the presence of external UV photons with a flux comparable to that required by self-regulated star formation in the halo.  \\subsection{X-ray Luminosity} In the hierarchical galaxy formation process, coalescence of small stellar systems (dwarf galaxies) occurs within a few dynamical timescales.  If the ordinary matter contained initially within these building blocks is to be heated through shock dissipation to the virial temperature of the common halo, the mergers would become luminous X-ray sources (Eke, Navarro, \\& Frenk 1998).  But the results of our models indicate that; 1) thermal instability leads to the formation of WFC's which contain most of the ordinary matter, and 2) even in the RHG, conduction may be a more efficient channel of heat loss than bremsstrahlung emission.  As discussed in \\S~5.2-5.4, the resulting x-ray luminosities are small, ranging from 10$^{11}$~L$_\\odot$ in Model~1 to 10$^{8}$~L$_\\odot$ in Model~3. The spectra of the models are shown in Figure~6.  The spectra of the low luminosity systems may be modified somewhat by supernova emission.  Using the total star formation rates of Models~1-3, we estimate supernova luminosities of 10$^{10}$~L$_\\odot$ for Model~1, 6$\\times$10$^{9}$~L$_\\odot$ for Model~2, and 5$\\times$10$^{8}$~L$_\\odot$ for Model~3.  Based upon these results, we expect relatively little X-ray luminosity to be released from regions with ongoing galaxy merging events, consistent with x-ray observations (Fabbiano \\& Schweizer 1995; Read \\& Ponman 1998).  \\subsection{Limitations} In our first attempt to investigate the complex physics of multi phase gas dynamics during the early epoch of galaxy formation, we have adopted various simplifying assumptions such as 1-D spherical symmetry and power-law size distribution for WFC's. These treatments can be improved with a self-consistent analysis of the WFC's evolution which will be presented in a follow up paper.  We have also neglected the evolution of dark matter which dominates the potential. The results in Figs 2-4 show that the distribution of pressure, density, filling factor of WFC's and the star formation rate depend, though not sensitively on the potential.  A full study of this problem will require the use of multi-dimensional integrated simulations of both dark and ordinary matter.  While such work has been done in a cosmological context, the extreme temperature range spanned by the gas within an individual galaxy (at least four orders of magnitude) places severe demands upon a code to be able to resolve structures over many orders of magnitude in size.  Our basic approach and the prescription provided here can be readily used in such investigations as they begin to be made."
    },
    "0004/astro-ph0004219_arXiv.txt": {
        "abstract": "We present the results from an optical study of the stellar \\& star formation properties of NGC 925 using the WIYN 3.5m telescope.  Images in B,V,R, \\& H$\\alpha$ reveal a galaxy that is fraught with asymmetries.  From isophote fits we discover that the bar center is not coincident with the center of the outer isophotes nor with the dynamical center (from Pisano {\\it{et al.}} 1998).  Cuts across the spiral arms reveal that the northern arms are distinctly different from the southern arm.  The southern arm not only appears more coherent, but the peaks in stellar and H$\\alpha$ emission are found to be coincident with those of the H~I distribution, while no such consistency is present in the northern disk.  We also examine the gas surface density criterion for massive star formation in NGC 925, and find that its behavior is more consistent with that for irregular galaxies, than with late-type spirals.  In particular, star formation persists beyond the radius at which the gas surface density falls below the predicted critical value for star formation for late-type spirals.  Such properties are characteristic of Magellanic spirals, but are present at a less dramatic level in NGC 925, a late-type spiral. ",
        "introduction": "This paper is part of an ongoing series of papers studying the general H~I and optical properties of late-type barred spiral galaxies (Sbc-Sd, hereafter LTBS). We will examine the stellar and star forming properties of NGC 925 based on observations from the WIYN 3.5m telescope{\\footnote{The WIYN observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories.}}.  We previously observed NGC 925 and NGC 1744, both LTBS, in H~I as part of a study of the gaseous properties of this class of galaxies (Pisano {\\it{et al.}}, 1998, hereafter Paper I), and to determine the pattern speed of their bars (Elmegreen {\\it{et al.}} 1998).  NGC 925 was chosen for study because it is a nearby ($\\sim$9.3 Mpc, Silbermann {\\it{et al.}} 1996), apparently prototypical LTBS which is well oriented in inclination for dynamical and morphological study.  This value for the distance will be used throughout this paper.  In paper I, we found that NGC 925 had a weak spiral pattern and bar as indicated by the small streaming motions in the spiral arms and bar.  Asymmetries are quite prevalent in NGC 925.  NGC 925's southern spiral arm has reasonably coherent streaming motions, while the northern arms do not.  The center of the bar is also slightly offset from the dynamical center of the galaxy by $\\sim$950 pc.  These asymmetries could possibly be related to the presence of a small H~I cloud (M$_{H~I}\\sim$10$^7$M$_{\\odot}$) interacting with the main galaxy.  The strength of the interaction, however, is most likely not enough to drive the asymmetries in NGC 925.  The presence of an off-center bar, and a single coherent spiral arm are typically viewed as characteristic properties of barred Magellanic (SBm) galaxies (de Vaucouleurs \\& Freeman 1972), and have been shown to be potentially long-lasting (Levine \\& Sparke 1998, Noordermeer {\\it{et al.}} 2000)  While NGC 925 does not illustrate these properties as dramatically as the LMC, for example, it certainly indicates that these characteristic asymmetries do not suddenly appear in SBm's, but are rather a continuous change in morphology throughout late-type disk galaxies. A summary of basic properties of NGC 925 is given in table 1.  LTBS are an understudied and, hence, a poorly understood class of galaxies.  While bars may play an important role in galaxy evolution by generating gas inflow (e.g., Matsuda \\& Nelson 1977, Athanassoula 1992), driving spiral structure (e.g., Sanders \\& Huntley 1976; James \\& Sellwood 1978), and flattening the radial abundance gradient (Martin \\& Roy 1994), these effects are shown to exist primarily in early-type barred spirals (S0-Sb).  LTBS do not appear to have strong enough bars to cause such effects (see paper I).  Previous studies of LTBS have found that late-type bars tend to have exponential light profiles, as opposed to flat profiles, and end well within corotation, as opposed to just inside it (Elmegreen \\& Elmegreen 1985, Elmegreen {\\it{et al.}} 1998).  Furthermore, late-type bars are not likely to be driving the spiral pattern (Sellwood \\& Sparke 1988), in contrast to early-type bars (Elmegreen \\& Elmegreen 1989).  Star formation properties are also different, with late-type bars tending to be gas-rich and have star formation throughout the bar, while early-type bars have star formation only near the ends of the bar (Phillips 1995).  Our observations in paper I of NGC 925 provided no evidence to contradict any of these properties.  In this paper, we will investigate the optical properties of NGC 925, with a specific eye towards the nature of star formation in the galaxy.  We will see if there are other signatures of asymmetry in NGC 925 that are typically associated with SBm's.  Finally, we will try to devine what role the bar plays in determining the properties of the galaxy as a whole. In section 2 of this paper we discuss the observations and reduction.  In section 3, we discuss the optical properties of galaxy as a whole.  Section 4 discusses the star formation properties of NGC 925 both across the entire galaxy, and in the bar and spiral arms specifically. We conclude in section 5. ",
        "conclusions": "We have observed NGC 925 using the WIYN telescope in B, V, R, and H$\\alpha$ filters to better understand the stellar distribution and star formation properties of the galaxy. We have compared these observations with previously described H~I observations (Paper I) to get a nearly complete picture of the properties of NGC 925.  The global properties of NGC 925 are typical for a late-type spiral galaxy.  It has an obvious bar, and two large spiral arms.  The southern spiral arm is apparently much stronger than the northern arm, which is quite flocculent.  The relative dominance of the southern arm is apparent both optically and in the H~I (paper I).  This difference between the northern and southern spiral arms is not the only asymmetry present in the galaxy.  The center of the galaxy, as derived from the outer isophotes, is coincident with the dynamical center (Paper I), but not with the center of the bar.  These properties are typically associated with barred Magellanic spirals (de Vaucouleurs \\& Freeman 1972), although in the case of SBm's the offset bar and single, dominant spiral arm tend to be much more pronounced than they are in the case of NGC 925.  Nevertheless, we see that such properties are not limited to Magellanic spirals, but are present in more subtle ways in earlier-type spirals.  The optical properties of NGC 925, such as mass-to-light ratio and absolute magnitude, are as would be expected for a typical late-type galaxy.  From our isophote fitting, we find that NGC 925's surface brightness distribution can be characterized by a double exponential; one for the bar and one for the outer disk.  The bar is quite easily distinguished from the rest of the galaxy in the fits, by having a well-defined center, position angle, and ellipticity differing from the outer disk.  The bar also has a exponential brightness distribution with a scale-length much smaller than the rest of the galaxy.  There is little evidence that the bar has any effect on the structure of the rest of the galaxy.  The isophote fits show quite a bit of structure outside of the bar region.  This is almost certainly due to the bright southern spiral arm, which shows up clearly as a surface brightness enhancement in the fits.  At the same radius as this enhancement, we can see distinct changes in the fits for the center, ellipticity, and position angle.  The fits do not seem to be affected by the weaker northern arm.  This is further evidence for NGC 925's similarity to the SBm one-armed spiral pattern.  The fits also show that NGC 925 gets redder with increasing radius, atypical of normal spiral galaxies in the field. This could indicate a recent enhancement of star formation in the inner galaxy.   From our analysis of the critical density for massive star formation in NGC 925, using the technique of Kennicutt (1989), we find that the galaxy has widespread star formation occurring outside of the radius where it should be suppressed.  NGC 925 has a very flat H$\\alpha$ surface brightness distribution, with emission extending all the way to the edge of our image at 15 kpc.  Kennicutt (1989) using a sample of 15 Sc galaxies found that star formation was suppressed when $\\alpha$=$\\Sigma_{gas}$/$\\Sigma_{crit}$ fell below 0.67.  For NGC 925 star formation continues past this radius and is more consistent with an $\\alpha\\sim$0.3, like what was found by Hunter {\\it{et al.}} (1998) and van Zee {\\it{et al.}} (1997) for later-type galaxies.  The exact value of $\\alpha$ is unclear, but probably lower, as the H$\\alpha$ emission could continue beyond our field of view.   As mentioned above, the spiral arms in NGC 925 are quite different.  The north arm is quite weak, lacks coherence, and has only patchy star formation.  The southern arm, on the other hand, can be traced all the way from the bar out to the edge of the galaxy and has plenty of star formation. Aside from a small gap, the arm is quite coherent.  There is a long dust lane associated with the southern arm, which probably obscures some star formation.  If there is dust along the northern arm, it is not nearly so well-behaved.  Taking cuts across the northern and southern arms, we can measure the relative positions of the stars, H~II regions, and H~I.  In the southern arm, the strongest H$\\alpha$ and stellar emission is closely aligned with the peak of the H~I. Large scale coherence breaks down in the northern half of the disk.  At some places (40-50$^{\\prime\\prime}$ in cut f) star formation is well correlated with peaks in the H~I distribution. At other locations ($\\sim$40$^{\\prime\\prime}$ in cut g) star formation is displaced off of the H~I peaks while elsewhere the H~I and H$\\alpha$ are anti-correlated (e.g. $\\sim$50$^{\\prime\\prime}$ in cut e). The difference between the north and south arms suggests that there may be some fundamental difference in the nature of the two arms.  Using the same cuts, we looked for evidence of triggering by the spiral arms.  If there is triggering in the arms, we might expect a color gradient (implying an underlying age gradient) perpendicular to the arms.  Such a color gradient can also be caused by dust extinction.  The lack of a color gradient would imply the lack of an age gradient, and, hence a lack of triggering.  While the cuts of the north arms show that the stellar emission is typically offset from the H$\\alpha$ and H~I emission, there is no offset present between the B \\& R peaks.  Similarly, the southern arm shows no offset between the broadband emission peaks.  The lack of an offset shows that there is no color gradient, and hence it is unlikely that an age gradient due to triggering in the spiral arm is occurring.  This suggests that there is no larger scale organization of the star formation in the northern disk--a result reminiscent of irregular galaxies. Star formation must be driven by local conditions and processes in the northern half, while larger scale phenomena organize star formation in the bar and southern arm.  Finally, we examined the nature of the NGC 925's bar and the star formation occurring within it.  As is expected for late-type bars (Phillips 1995), NGC 925 has star formation occurring all along the bar.  Cuts perpendicular to the major axis of the bar show that the star formation is occurring along the major axis, but the broadband optical emission is systematically offset to the north of the H~I peaks.  While this could be a dust effect, we would expect the H$\\alpha$ emission to be similarly affected, but it is not.  This could also be due to the massive star formation in the bar ionizing or clearing out the H~I in the bar.  The systematic offset may be related to the dynamical center being offset from the bar center, and hence the bar could be a wave pattern moving through the gaseous medium in NGC 925. We might expect to see the peak of the stellar emission downstream of the star formation, as is the case here. It may be that the bar of NGC 925 is not actually a bar, but has more in common with the spiral arms of NGC 925.  Overall, this study has shown that NGC 925 is a prototypical late-type spiral galaxy, with some properties that are characteristic of Magellanic spirals, such as an off-center bar and a dominant spiral arm.  These traits are less pronounced in NGC 925 than in an SBm indicating a probable smooth transition of these properties from Scd galaxies to Sm's.  Massive star formation in NGC 925 persists beyond the radius predicted by Kennicutt (1989) for late-type spirals, but behaves more like dIrr's (Hunter {\\it{et al.}} 1998). This result may be exacerbated by wider field imaging of NGC 925, however, and star formation almost certainly continues at a low level beyond this point (see Ferguson {\\it{et al.}} 1998 for an example).  The north and south arms are not only different in brightness, but in coherence and in the distribution of stars across the arms.  The north arm has stars offset from H~I peaks, while the south arm has both star formation and stellar emission coincident with the neutral gas peaks.  Finally, the bar of NGC 925 appears to be a typical late-type bar being gas-rich and having star formation throughout, but the stars are offset from the H$\\alpha$ and H~I peaks.  This fact coupled with the offset dynamical center of NGC 925 suggests that the bar may be more of a wave phenomenon, similar to a spiral arm."
    },
    "0004/gr-qc0004068_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:Introduction}  Global monopoles are topological defects that arise in certain theories where a global  symmetry is spontaneously broken. The simplest and most studied example is the $O(3)$ model, in which one finds that the static, spherically symmetric solutions, have in general an energy density that decreases for large distances as $1/r^2$\\cite{Vilenkin}. This would lead in a Newtonian analysis to a divergent expression for the total mass. When we turn to the general relativistic analysis, this problem translates in the fact that the resulting spacetime is not asymptotically flat, and thus the standard ADM mass is not well defined. The effect of the $1/r^2$  behavior of the  density on the spacetime is that the latter develops a deficit angle at large distances The fact that, at small distances, the behavior deviates from that,  results in the appearance of a small phenomenological \"core mass\" which turns out to be negative in all cases considered \\cite{Harari}. Moreover, an analysis of the behavior of geodesics in the large distance regime does indeed support such interpretation of this core mass because its effect turns out to be repulsive\\cite{Harari}. This ``core mass\" is then evidently not the standard ADM mass. The question of what exactly  one is talking about when referring to this core mass has been resolved in \\cite{Ulises} through the application of the standard type of Hamiltonian analysis to the class of spacetimes that are Asymptotically-flat-but-for-a-deficit-angle (AFDA$\\alpha$) \\cite{Ulises}.  For these (AFDA$\\alpha$) spacetimes one can also define future and past conformal null infinity, and thus the notion of a black hole and of its horizon. In fact solutions corresponding to global monopoles with such interior horizons have been found in \\cite{Liebling}, \\cite{Maison}.  In this paper we study the behavior of the ADM mass of these AFDA$\\alpha$ spacetimes as a function of the horizon area, concentrating in particular on its sign, which we find changes in the regime where one would interpret as going from a situation that would be naturally described as a ``black hole inside a monopole core\" to that which would be naturally described as a ``black hole with a global monopole inside\".  We shall adhere to the following conventions on index notation in this paper: Greek indices ($\\alpha$, $\\beta$, $\\mu$, $\\nu$,...) range from $0$ to $3$, and denote tensors on (four-dimensional) spacetime. Latin indices, alphabetically located after the letter i (i,j,k,...) range from $1$ to $3$, and denote tensors on a spatial hipersurface $\\Sigma$; whereas Latin indices, from the beginning of the alphabet (a,b,c,d,...) range from $1$ to $3$, and denote  indices in the internal space of the scalar fields. The  metric  for the internal space is just  the flat Euclidean metric $ \\delta_{ab}$. The signature of the spacetime metric {\\bf g} is $ (-,+,+,+)$. Geometrized units, for which $G_N=c=1$ are used in this paper.   \\medskip ",
        "conclusions": ""
    },
    "0004/astro-ph0004169_arXiv.txt": {
        "abstract": "The TGRS experiment on board the {\\em Wind} spacecraft has many advantages as a sky monitor  ---  broad field of view ($\\sim 2 \\pi$ centered on the south ecliptic pole), long life (1994--present), and stable low background and continuous coverage due to {\\em Wind\\/}'s high altitude high eccentricity orbit.  The Ge detector has sufficient energy resolution (3--4 keV at 511 keV) to resolve a cosmic positron annihilation line from the strong background annihilation line from $\\beta$-decays induced by cosmic ray impacts on the instrument, if the cosmic line is Doppler-shifted by this amount.  Such lines (blueshifted) are predicted from nucleosynthesis in classical novae.  We have searched the entire TGRS database for 1995--1997 for this line, with negative results.  In principle such a search could yield an unbiased upper limit on the highly-uncertain Galactic nova rate.  We carefully examined the times around the known nova events during this period, also with negative results.  The upper limit on the nova line flux in a 6-hr interval is typically $<3.8 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$ ($4.6 \\sigma$).  We performed the same analysis for times around the outburst of Nova Vel 1999, obtaining a worse limit due to recent degradation of the detector response caused by cosmic ray induced damage. ",
        "introduction": "Theoretical models of classical novae imply that large quantities of $\\beta$-unstable proton-rich nuclei are formed during a thermonuclear runaway, with half-lives of the order minutes to hours.  The resulting positrons undergo annihilation in the expanding nova envelope, giving rise to a pulse of annihilation $\\gamma$-rays which is blueshifted and broadened by the nova velocity and lasts for up to $\\sim 6$ hr (2).  All space-borne instruments are subject to a strong background emission line at 511 keV, arising from annihilation of cosmic ray induced $\\beta$-decay positrons in the instrument, which hinders detection of the same line from cosmic sources.  However, a spectrometer with sufficient energy resolution could distinguish between cosmic and background lines because the former is blueshifted, typically by 2--5 keV.  The Ge spectrometer TGRS (FWHM resolution 3 keV at the beginning of its mission) fulfills this requirement. Its other advantages as a detector and monitor of the nova 511 keV line are its broad aperture (nearly $2 \\pi$ sterad) and its high altitude high inclination orbit, in which background count rates are rather low, due to the lack of interaction with Earth's trapped radiation belts, and to the virtual absence of Earth albedo background radiation.  The background count rate and spectrum has also been extremely stable throughout the mission.  \\begin{figure} \\centerline{\\epsfig{file=mjh3_fig1.ps,height=2.0in,width=6.0in}} \\vspace{10pt} \\caption{Characteristic TGRS background count spectrum at energies around 511 keV, obtained during the interval 3 June 12h--18h UT with a power law continuum subtracted.  The components of the fit are a line with the width and position of the nova 6-hr line (2, dot-dashed line) and a Gaussian line fitting the blue wing of the 511 keV background line (dashed line).  The total model spectrum is the full line.  (b)  Expansion of Fig. 2a showing the significance of the fitted nova 6-hr line (dot-dashed line of amplitude $5.2 \\pm 0.1 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$; other symbols as in Fig. 2a). Also shown (dotted line) is the theoretically expected level of the nova 6-hr line for a nova at 1 kpc (ref. 6 model HH5).} \\label{Fig. 1} \\end{figure}  A disadvantage of Ge detectors for long-term missions is progressive degradation of performance due to cosmic-ray impact damage. For this reason most of the data used here were obtained during 1995--1997, before this effect became serious.  In an attempt to avoid the effects of degradation, such as low-energy \"tailing\", we fitted only energies on the high-energy (blue) side of the background 511 keV line.  One such fitted background spectrum is shown in Fig. 1.  The dashed line is the strong background 511 keV line.  The dot-dashed line is the fit to a theoretical nova line emitted in a 6-hr interval after the event (3).  Although the amplitude of this line is significant, we will see shortly that it is characteristic of all our background spectra --- in other words, there is a systematic constant positive offset in this measurement. For scale, the dotted line shows how the same line would appear from a model nova at a distance 1 kpc.  This line would be detected at a level $\\sim 8 \\sigma$. ",
        "conclusions": ""
    },
    "0004/astro-ph0004079_arXiv.txt": {
        "abstract": "We present deep imaging in the U, B and I bands obtained under excellent seeing conditions of the double quasar Q0151+048A,B and of the Damped Ly$\\alpha$ (DLA) absorbing galaxy at $z_{\\rm abs} = 1.9342$ named S4.  We analyse the data employing two separate and independent methods. First we deconvolve the images using the MCS algorithm, secondly we decompose the images via an object based iteration process where we fit models to objects without any attempt to improve the resolution of the data. Our detailed analysis of the images reveals, somewhat surprisingly, that extended objects centred on the quasars themselves are much brighter continuum sources than the DLA galaxy.  Due to the complexity caused by the many superimposed objects, we are unable to certify whether or not continuum emission from the DLA galaxy is detected. Continuum emission from the extended objects centred on the positions of the quasars is clearly seen, and the objects are tentatively identified as the ``host galaxies'' of the quasars. The flux of those host galaxies is of order 2--6\\% of the quasar flux, and the light profile of the brighter of the two is clearly best fit with a de Vaucouleurs profile. We discuss two alternative interpretations of the origin of the extended flux: {\\it i)} the early stage of a massive elliptical galaxy in the process of forming the bulk of its stars, and {\\it ii)} quasar light scattered by dust. ",
        "introduction": "The study of the galaxy population at high redshifts has progressed rapidly during the last decade. Through the Lyman-break technique hundreds of normal (i.e. not dominated by active galactic nuclei), star forming galaxies at z=2--4 have been detected and studied with imaging as well as spectroscopy (Steidel et al. 1996). These so called Lyman-Break Galaxies (LBGs) have star formation rates (SFRs) in the range 4--55h$^{-2}$ M$_{\\sun}$ yr$^{-1}$ for $\\Omega$=1.0 or 20 -- 270 M$_{\\sun}$ yr$^{-1}$ for $\\Omega$=0.2 (Pettini et al. 1998). Also, via the study of the class of high column density QSO absorption lines systems known as Damped Ly$\\alpha$ Absorbers (DLAs) a wealth of information on the early chemical evolution of galaxies at z=2--4 has been obtained (e.g. Lu et al. 1996). The DLAs are in general forming stars at a significantly lower rate than the LBGs (M\\o ller \\& Warren 1998, Fynbo et al. 1999).  Independent information about the formation of the brightest galaxies comes from detailed studies of the stellar populations of present day bright cluster ellipticals. These populations seems to have formed early (z$>$2) in strong burst of star formation (Bower et al. 1992). Studies of the fundamental plane for elliptical and lenticular galaxies in rich clusters at intermediate redshifts also indicate early formation times (z$>$5 for $\\Omega$=1, J\\o rgensen et al. 1999), and the fundamental plane for field ellipticals at similar redshifts is consistent with being the same as in clusters (Treu et al. 1999a).  Studies of the globular cluster populations of faint elliptical galaxies also indicate rather early formation times (z$>$1), whereas for bright cluster ellipticals the globular cluster populations do not strongly constrain the possible formation scenarios (Kissler-Patig et al. 1998). Furthermore, the presence of seemingly old stellar populations in elliptical galaxies at z$>$1 proves that at least some elliptical galaxies formed very early in strong bursts of star formation (Spinrad et al. 1997, Treu et al. 1999b, see also Jimenez et al. 1999). For first-rank ellipticals star formation rates as high as SFR$\\sim$10$^{3}$ M$_{\\sun}$ yr$^{-1}$ would then be possible. A reason why such high star formation rates have not been detected in galaxies at high redshift may be that these galaxies are the hosts of powerful QSOs and hence are hidden by the light from the QSOs (e.g. Terlevich \\& Boyle 1993). Support for a connection between QSOs and bright elliptical galaxies comes from the fact that radio quiet QSOs as well as radio loud QSOs and radio galaxies at z=0.1-0.3 are hosted by galaxies for which the light profiles are best fit by de Vaucouleurs profiles indicating that they are early stages of massive ellipticals (McLure et al. 1999). There is increasing evidence that QSOs at redshifts z$\\approx$2 are embedded in extended emission that is consistent with the presence of a stellar population in the QSO host galaxies. In the case of radio loud QSOs host galaxies have been detected in the optical and infrared by Lehnert et al. (1992) and Carballo et al. (1998), and in the case of radio quiet QSOs host galaxies have been detected in the optical and near infrared by Aretxaga et al. (1998a,b). There does not seem to be any systematic differences between the host galaxies of radio loud and radio quiet QSOs. Both populations of host galaxies are extremely bright, R$\\approx$21--22, and have optical-to-infrared colours in the range R-K$\\approx$3--5. However, measured polarisation of the light from some radio galaxies show that scattered QSO light can also contribute significantly to the observed extended emission (e.g. Cimatti et al. 1998).  In 1996  we performed  a  narrow band   study  of the z$_{abs}\\approx$z$_{em}$ Damped Ly$\\alpha$ Absorber (DLA, Wolfe et al. 1986) towards Q0151+048A using the 2.56-m Nordic Optical Telescope (NOT)  (Fynbo et al. 1999). The main result of this study was the detection of extended Ly$\\alpha$ emission from the DLA. The Ly$\\alpha$ emission line had prior to this been detected in a spectroscopic study of Q0151+048A (M\\o ller et al. 1998), but the large extended nature of the DLA absorber was quite unexpected. U band data, also from the 1996 run, hinted at the existence of an extended broad band object, but the signal--to--noise ratio of the object was low. We have therefore obtained deeper imaging of Q0151+048 in broad band U, B and I filters in order to confirm or reject our tentative detection, and to measure the extend and luminosity of the broad band source if real.  In Sect. 2 below we describe our new observations. In Sect. 3 we describe in detail the two independent methods we have used to search for extended objects close to the quasar. First we describe the image-deconvolution, where we used the Magain et al. (1998, hereafter MCS) algorithm, secondly we describe the direct PSF subtraction, and Sect. 4 we discuss our results.  In this paper we adopt H=100 h km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$=1.0 and $\\Omega_{\\Lambda}$=0 unless otherwise stated. ",
        "conclusions": "Our original interest in the field of Q0151+048A was to identify the DLA galaxy in front of it. This identification was accomplished via imaging in Ly$\\alpha$ (Fynbo et al. 1999), but our broad band images left some questions open. The purpose of the deeper broad--band data presented in this paper was to clarify this situation. We shall here first summarise our findings, then briefly consider their implications.  \\subsection{Results summary} Our new data have unambiguously confirmed the presence of extended emission in the field in all three bands I, B and U. The different morphology seen in the three bands strongly suggest that we see three objects superimposed: The quasar, the DLA absorbing galaxy and the quasar host galaxy.  The superposition of three close objects of widely differing brightnesses causes considerable degeneracy for any attempt to determine the brightness of the faintest sources, and it is therefore impossible to find a unique solution for the flux of the faintest object (the DLA galaxy S4). Nevertheless, we find that S4 is clearly detected in the U image. The U--band magnitude of S4 determined via our minimum $\\chi^2$ procedure is fully consistent (to within 1 $\\sigma$) with being caused by the known Ly$\\alpha$ flux at 3565{\\AA} alone. The data are therefore consistent with a zero contribution from any continuum source in the U--band.  It is difficult to determine the exact errors on the I and B magnitudes of S4, but for both images we found a very significant improvement in the reduced $\\chi^2$ of the fit when we included S4. It is therefore likely that S4 is indeed a low surface brightness continuum source, but this question is going to be extremely hard to settle.  The existence of a separate extended continuum source centred on qA is, however, clearly demonstrated independently in all bands. This result was arrived at independently via image deconvolution, and via our iterative object fitting technique.   \\subsection{Discussion: Starburst Galaxy or Dust scattering} The distance modulus (for z=1.93) in the assumed cosmology with h=0.5 is 45.8. Assuming instead $\\Omega$=0.3 and $\\Omega_{\\Lambda}=$0.7 the corresponding distance modulus becomes 46.6. Hence, the absolute AB magnitudes of the host galaxy HGa is $<$-24.0(-24.8) in U (rest frame 1100--1300\\AA), $<$-24.5(-25.3) in B (rest frame 1300--1500\\AA) and $<$-24.0(-24.8) in I (rest frame 2300--3400\\AA). Such extremely bright magnitudes are in the local universe only connected with brightest cluster galaxies (for comparison M87 and Centaurus A both have absolute magnitudes of roughly -23 in the V--band). Brightest cluster members can be as bright as -26 (Oemler 1976). Interestingly we find that the absolute magnitude of HGa is similar to those of the extended `fuzz' that have been detected around other high redshift QSOs by Lehnert et al. (1992), Carballo et al. (1998) and Aretxaga et al. (1998a,b).  The morphology of the host galaxy HGa is best fit by a de Vaucouleurs profile. The fit to an exponential-disc leads to a much poorer fit. A plausible interpretation of the data is therefore that we see the early stage of a massive elliptical galaxy in the process of forming the bulk of its stars. Assuming that all the light is coming from stars, and not e.g. scattered quasar light (see below), we can estimate the star formation rate (SFR) needed to explain the observed fluxes. In the case of continuous star formation we can adopt the relation between the SFR and the luminosity at 1500\\AA \\ SFR = L$_{1500}/(1.3\\times10^{40} erg\\: s^{-1} $\\AA$^{-1}$) commonly used for LBGs (Pettini et al. 1998) and we hence infer a star formation rate of order 100(200) M$_{\\sun}$ \\ yr$^{-1}$ for $\\Omega$=1(0.3) and $\\Omega_{\\Lambda}$=0(0.7). For instantaneous bursts we can use the Starburst99 package (Leitherer et al. 1999) to infer the colours of models calculated with solar metallicity and ages 1, 10 and 100 million years. The colours for these three models are given in Table~\\ref{instan}.  \\begin{table} \\begin{center} \\caption{The colours of instantaneous starbursts with four different ages. The colours of HGa are listed for comparison.} \\begin{tabular}{@{}rrrr} Age   & u-B  & B-I  \\\\ (Myr) &      &      \\\\ \\hline 1     &  -0.2 & -0.7 \\\\ 10    &  -0.1 & -0.3 \\\\ 100   &   0.7 &  0.3 \\\\ \\hline Exta   & 0.0--1.1 & -1.1--0.1 \\\\ HGa   & 0.9$\\pm$0.3 & -0.7$\\pm$0.3 \\\\ \\hline \\end{tabular} \\label{instan} \\end{center} \\end{table} The colours of the host, 0$<$u-B$<$1.1, -1.1$<$B-I$<$0.1 from Deconvolution and 0.9$\\pm$0.3, -0.7$\\pm$0.3 from PSF-subtraction, are roughly consistent with instantaneous bursts with ages in the range 10--100 Myr. The number of stars formed in the burst would be in the range from 10$^8$ to a few times 10$^9$ stars depending on the age of the burst and on the assumed cosmology.  Another interpretation of the extended fuzz frequently seen around quasars, is light from the quasar itself scattered by dust. This mechanism is well known from radio galaxies at high redshifts where scattering off dust grains has revealed the existence of ``hidden'' quasars in the galaxy cores. It is likely that radio quiet QSOs have similar non--isotropic radiation fields (see e.g. the discussion in M{\\o}ller \\& Kj{\\ae}rgaard 1992), and in that case our line of sight is such that we look straight down the emission cone inside of which the scattering is taking place. In this case we therefore expect to see the quasar emission cone ``end on'' via forward scattered quasar light. The scattering process is expected to be essentially grey and recent calculations predict that as much as 10\\% of the quasar light could be scattered in this way (Witt \\& Gordon, 1999; V\\'arosi \\& Dwek, 1999; Vernet et al. in prep.). If considering a clumpy medium, we would expect dust scattered light to be emitted from inside a very large volume in front of the quasar. When taking the cone geometry into account one would expect its total flux to be roughly a few \\% of the quasar flux at any given wavelength (Fosbury, private communication). From Table 3 we find that the flux from HGa is 3, 6 and 2\\% of the flux from Q0151+048A in U, B and I respectively. Similar, but less significant, results are found for HGb. It is not yet known if the light profile of scattered light from a cone will reproduce a de Vaucouleurs profile, but since this seems to be a universally preferred profile it is not unlikely. One thing worth noting in Fig.~\\ref{PSFSUB}e are the negative residuals surrounding the position of qA at a distance of 2-3 arcsec after subtraction of the fitted de Vaucouleurs profile. This indicates that the true profile of HGa in reality falls off steeper than a de Vaucouleurs profile. If model calculations were to show such a steep profile for forward scattered light in a radiation cone, that would be a strong hint towards the nature of the quasar fuzz.  \\begin{figure} \\epsfig{file=h1854_4.eps,width=8cm} \\caption{The B(AB)-I(AB) vs. u(AB)-B(AB) colours of qA and qB (marked by $\\times$), of the instantaneous star burst models of Table 4 (thick, full drawn line). The dotted, dashed and long dashed lines show the colours of reddened bursts assuming A(B$_{Rest}$)=0.5 and MW, LMC and SMC extinction curves respectively. The colours of the extended emission under qA is marked by the error-bars. The dashed error-bar represents the measurement from deconvolution and the full drawn error-bar the measurement from PSF-subtraction. } \\label{colours} \\end{figure}  However, the colours of the extended emission as seen in Table 3 and Fig.~\\ref{colours} are significantly different from those of the two QSOs, which argues against the scattering hypothesis. Hence, we conclude that at least a significant fraction of the observed extended emission must be caused by a star burst."
    },
    "0004/astro-ph0004309_arXiv.txt": {
        "abstract": "We employ an ensemble of 24 hydrodynamic cluster simulations to create spatially and spectrally resolved images of quality comparable to {\\em Chandra}'s expected performance.  Emission from simulation mass elements is represented using the XSPEC {\\texttt mekal} program assuming 0.3 solar metallicity and the resulting spectra are fit with a single-temperature model.  Despite significant departures from isothermality in the cluster gas, single-temperature models produce acceptable fits to 20,000 source photon spectra.  The spectral fit temperature $T_s$ is generally lower than the mass weighted average temperature $T_m$ due to the influence of soft line emission from cooler gas being accreted as part of the hierarchical clustering process.  The nature of this deviation depends on the bandpass used for spectral fitting.  In a {\\em Chandra}-like bandpass of 0.5 to 9.5 keV we find a nearly uniform fractional bias of $(T_m-T_s)/T_s \\simeq 20\\%$, although smaller clusters sometimes demonstrate much greater deviations.  If the minimum energy threshold is raised to 2 keV, however, the effect of line emission on the spectrum is greatly decreased and $T_s$ becomes a nearly unbiased estimator of $T_m$ for smaller clusters. The fractional deviation in $T_s$ relative to $T_m$ is scale-dependent in this bandpass and follows the approximate relation $(T_m-T_s)/T_s = 0.2\\log_{10}T_m$. This results in an observed $M_{ICM}$--$T_s$ relationship for the simulations with slope of about 1.6, intermediate between the virial relation $M \\propto T_m^{3/2}$ and the observed relation $M_{\\rm ICM} \\propto T^2$.  Tracking each cluster in the ensemble at 16 epochs in its evolutionary history, we catalogue merger events with mass ratios exceeding 10\\% in order to investigate the relationship between spectral temperature and proximity to a major merger event.  Clusters that are very cool relative to the mean mass-temperature relationship lie preferentially close to a major merger, suggesting a viable observational method to cull a subset of dynamically young clusters from the general population. ",
        "introduction": "Galaxy clusters are the youngest and largest organized structures in the universe, and as such provide us with a wealth of cosmological information.  The most massive clusters draw their substance from cosmologically significant volumes of linear scale $\\gtrsim 10 h^{-1}$ Mpc.  These scales are large enough that no known coherent process competes against gravity, so rich cluster contents are thought to comprise a fair sample of the universe's ingredients (White \\etal 1993).  Because clusters are rare nonlinear excursions of the cosmic density field, the statistical properties of their population are quite sensitive to both cosmological model and slope of the primordial fluctuation spectrum. Unfortunately, their relative youth can also make interesting physical properties difficult to measure: about\\ 50\\% of the local population bears evidence of ongoing mergers, and the canonical ``relaxed'' cluster is a relatively rare beast.  Nearly all interesting cosmological tests depend on accurate measurement of cluster virial masses.  Observations of the intracluster medium (ICM) have shown promise in this regard: the ICM's high X-ray luminosity and large spatial extent make it possible to probe the content and structure of clusters in great detail. Bulk properties of the ICM such as luminosity, temperature, mass, and gas density profile shape have been found to display highly significant correlations with each other (Edge \\& Stewart 1991; David \\etal 1993; Mohr \\& Evrard 1997; Mushotzky \\& Scharf 1997; Markevitch 1998; Allen \\& Fabian 1998; Mohr, Mathiesen \\& Evrard 1999; Arnaud \\& Evrard 1999).  In contrast to the noisy correlations of early X-ray data (Sarazin 1986 and references therein), many correlations now display scatter at the $10-20\\%$ level, indicating that a high degree of physical uniformity exists even in this structurally diverse population. Hydrodynamic simulations of cluster evolution predict tight relationships between observable quantities and between those quantities and the cluster binding mass (Evrard 1990; Kang \\etal 1994; Navarro, Frenk \\& White 1995; Evrard \\etal 1996; Bryan \\& Norman 1998), even when some members of the sample are far from dynamical equilibrium. The existence of both observed and theoretical correlations implies that the prevalence of cluster substructure is not a fundamental barrier to interpreting the properties of the population.  However, moderate biases caused by the presence of substructure are likely to be present, and we explore the role of substructure in temperature measurements of the ICM in this paper.  A previous paper (Mathiesen, Evrard \\& Mohr 1999) demonstrated that a ICM clumping leads to a modest ($\\sim 15\\%$) overestimate of ICM masses derived under the typical assumptions of spherical symmetry and isothermality. As the X--ray data improve, the limits of simplifying assumptions such as these become clearer.  High-resolution X--ray images reveal secondary peaks and strong asphericities in many clusters and X--ray spectra indicate the presence of multiple temperature components within the cores of many clusters (Fabian \\etal 1994, Holzapfel \\etal 1997, Allen \\etal 2000).  The {\\em Chandra} and {\\em XMM} satellite missions will provide the most detailed maps of the ICM emission and temperature structure yet obtained and will allow more precise definition of the limitations of the current models.  Three-dimensional hydrodynamical simulations of cluster formation can help bridge the gap between the new generation of data and traditional methods and results. While simulated clusters often do not include many processes thought to be important to ICM evolution (e.g. radiative cooling and galactic winds), they excel at the creation of populations with realistic merger histories (Mohr \\etal 1995; Tsai \\& Buote 1996). Ensembles of simulated clusters can therefore be used to investigate the effects of accretion events, major mergers, clumping, and substructure on measurements of the ICM.  In this paper, we analyze the spectral properties of an ensemble of 24 simulated clusters using a realistic plasma emission model and assuming a uniform metallicity $0.3$ times the solar abundance.  We find that even minor accretion events can significantly bias our measurements of the mean, mass-weighted cluster temperature, and that clusters undergoing a major merger can sometimes be identified as extreme examples of this bias.  Section 2 of this paper describes the simulations, the cluster ensemble, and the process of creating our spectral images.  Section 3 discusses the various measures of cluster temperature which have seen frequent use and explores the relationships between them. In particular we explore the relationship between spectrally determined temperatures and the mass-weighted mean temperature.  The latter is found in simulations to follow most closely the virial relationship. Section 4 then delves into cluster dynamics, investigating the effects of a major merger on the ICM and looking for observable signatures of the merging process. Finally, section 5 summarizes our conclusions. ",
        "conclusions": "The first (and perhaps the most suprising) result of this paper is the revelation that the emission spectrum of a realistically complex intracluster medium is barely distinguishable from that of an isothermal gas, even over a broad spectrum. This correspondence arises because we have created spectra with rather coarse (150 eV) bins in an effort to match our observation model to {\\em Chandra's} ACIS instrument.  We provide relationships between spectral, emission, and mass-weighted temperatures in two different bandpasses and within two density constrasts; in particular, we find that spectral temperatures determined in the [2.0,9.5] keV range within $r_{500}$ are very similar in nature to those determined by the {\\em Einstein} MPC, {\\em Ginga}, and {\\em EXOSAT}. The other bandpass, [0.5,9.5] keV, is designed to be similar to {\\em Chandra} observations, and can be used to better interpret its results.  We find that realistically determined spectral temperatures are commonly 10-20\\% lower than the mass-weighted temperatures in these simulations, a fact which has important implications for cluster physics.  The bias arises through the natural and frequent occurence of minor accretion events: small clumps of cool gas which merge into the ICM and produce an excess of soft X-rays, biasing the spectrum towards cooler temperatures until they are fully assimilated. Because the mass-weighted temperature follows the virial relationship, it is a more accurate indicator of the binding mass. Previous measurements of the cluster mass function power spectrum need to account for this source of error if they make use of spectral temperature data. We have calibrated the relationship between spectral and mass-weighted temperature for the [2.0,9.5] keV bandpass within $r_{500}$ for this purpose; this should allow an appropriate correction for clusters observed by satellites with a similar energy range.  The scale-dependent nature of this bias changes the slope of the observed $M_{\\rm ICM}$--$T_s$ relationship in a direction consistent with recent determinations, but it does not account for the whole difference. Additional physics, such as a variation in the ICM mass fraction with temperature, is still needed to explain the observed slope.  Although we are stuck with a discrepance between observed temperatures and mass-weighted temperatures, this bias can be useful in identifying clusters which are dynamically young. Not all of these events will be obvious, even with {\\em Chandra's} high-resolution surface brightness and temperature maps; some of the mergers will be occuring on an axis near our line-of-sight. In such cases the shocks and cool regions will probably be masked by core emission, but the presence of a cool subclump can still produce an unusually large deviation in the spectral temperature.  We should not fall into the trap of assuming that a cluster is relaxed because it looks spherically symmetric; rather, we should do everthing we can to determine whether or not it is truly relaxed. The simulations show that all of the clusters with temperatures falling far ($\\gtrsim 15\\%$) below the mean mass-temperature relation are dynamically young; a subset of merging clusters can thus be identified as outliers in the observed relation.  The results described in this paper are a necessary step towards accurately measuring cluster masses under the assumption of virial equilibrium, as well as towards a reliable technique for determining a cluster's dynamical state observationally. It is not, however, the only step: first the astronomical community must recognize that the various definitions of temperature in common use are {\\em not} equivalent. Further work needs to be done in modeling observed temperatures (e.g. core-excised spectral temperatures and flux-weighted temperatures averaged over an image) and the effects of substructure on observable quantities. Spectral modeling of simulations is computationally more expensive, but not difficult, and the level of systematic error caused by inaccurate modeling is now comparable to the observational constraints on cosmological parameters.  The spectral cubes described in this paper will be made avaiable to the public on Mathiesen's research web site \\footnote{http://redshift.perseus.edu/bfm} along with higher-resolution images for the subset of images used in fitting the mass-temperature relationships, the accretion history of each cluster, and documentation describing the file format and how to model other satellite responses. Our simulations of spatially resolved emission spectra can be used to help interpret {\\em Chandra} and {\\em XMM} observations of real clusters, and we are sure that enterprising researchers can mine these datasets for other interesting results."
    },
    "0004/astro-ph0004373_arXiv.txt": {
        "abstract": "We present a weak-lensing analysis of a region around the galaxy cluster Cl 1604+4304 ($z=0.897$) on the basis of the deep observations with the {\\it Hubble Space Telescope} ({\\sl HST})/Wide Field Planetary Camera 2 (WFPC2). We apply a variant of Schneider's aperture mass technique to the observed WFPC2 field and obtain the distribution of weak-lensing signal-to-noise ratio (S/N) within the field. The resulting S/N map reveals a clear pronounced peak located about $1\\farcm 7$ ($850h_{50}^{-1}$ kpc at $z=0.897$) southwest of the second peak associated with the optical cluster center determined from the dynamical analysis of Postman et al. A non-linear finite-field inversion method has been used to reconstruct the projected mass distribution from the observed shear field. The reconstructed mass map shows a super-critical feature at the location of the S/N peak as well as in the cluster central region. Assuming the redshift distribution of field galaxies, we obtain the total mass in the observed field to be $1.0 \\times 10^{15} h_{50}^{-1}M_{\\odot}$ for $\\left< z \\right>=1.0$. The estimated mass within a circular aperture of radius $280h_{50}^{-1}$ kpc centered on the dark clump is  $2.4\\times 10^{14}h_{50}^{-1}M_{\\odot}$. We have confirmed the existence of the `dark' mass concentration from  another deep {\\sl HST} observation with a slightly different ($\\sim 20\\arcsec$) pointing. ",
        "introduction": "Weak shear fields of high-redshift galaxies are promising, efficient tools to investigate the mass distribution on cluster-supercluster scales (Kaiser \\& Squires 1993; Luppino \\& Kaiser 1997; Kaiser et al. 1998; Bartelmann \\& Schneider 2000) and provide a unique mean to detect dark mass concentrations (Schneider 1996; Erben et al. 2000). Wide-field weak-lensing surveys of projected mass overdensities can probe the statistical clustering properties and underlying cosmology (Bahcall et al. 1999).  Recent observations have revealed the existence of a supercluster at a high redshift of $z\\approx 0.9$ (Lubin et al. 2000). This supercluster contains two massive galaxy clusters, Cl 1604+4304 at $z=0.897$  and Cl 1604+4321 at $z=0.924$. The two clusters are separated by $17\\arcmin$ on the sky, corresponding to a projected separation of $\\sim 9h_{50}^{-1}$ Mpc. Cl1604+4304 is located at ($\\alpha_{{\\rm J}2000}$, $\\delta_{{\\rm J}2000}$) = ( $16{\\rm h}\\,04{\\rm m}\\,19.5{\\rm s}$, $+43\\arcdeg\\,04\\arcmin\\,33\\farcs 9$) and one of the optically-selected high-redshift cluster candidates studied by Oke, Postman, \\& Lubin (1998). Postman, Lubin, \\& Oke(1998) presented a detailed photometric and spectroscopic survey of the cluster and obtained a velocity dispersion of $1226$ km s$^{-1}$ and a dynamical mass estimate of $6.2\\times 10^{15}h_{50}^{-1}M_{\\odot}$.  In this Letter, we present a weak lensing analysis of the Cl 1604+4304 field on the basis of deep images taken with the {\\it Hubble Space Telescope} ({\\sl HST})/Wide Field Planetary Camera 2 (WFPC2). We shall mainly describe the weak-lensing signal-to-noise ratio (S/N) analysis of the observed field; the details of the mass reconstruction procedure will be present elsewhere. Throughout this Letter, we adopt $\\Omega_0=1$ ,$\\Omega_{\\Lambda}=0$, and $H_0=50h_{50}$ km s$^{-1}$ Mpc$^{-1}$; $1\\arcmin$ on the sky corresponds to $0.52h_{50}^{-1}$ Mpc at the cluster redshift. ",
        "conclusions": "We have performed a weak-lensing analysis on the deep images of a region around Cl 1604+4304 taken with the {\\sl HST}/WFPC2. We detected two significant maxima in the resulting weak-lensing S/N map: the second peak associated with the dynamical cluster center and the first peak located about $1\\farcm 7$ south of the dynamical center having no optical counter parts. The identification of high S/N peak from two-different pointing observations provides a robust evidence of dark mass concentration. Assuming that the dark mass clump has the same redshift as Cl 1604+4304, we reconstructed the projected mass distribution in the 1995 data field, from which we have estimated the mass within a circular aperture of radius $280h_{50}^{-1}$ kpc centered on the dark clump to be $2.4\\times 10^{14}h_{50}^{-1}M_{\\odot}$. The obtained results indicate that this dark clump could be a compact, massive bound system associated with the supercluster at $z\\approx 0.9$; the obtained mass will be overestimate if the redshift of the dark clump is much less than $0.897$.  Our results demonstrate that the weak-lensing S/N statistics are powerful and efficient tools even for high-redshift cluster surveys ($z\\sim 1$). For a further confirmation of the dark mass concentration, deep wide-field observations covering the entire supercluster field ($\\sim 20\\arcmin$) are required. Such observations can probe the mass distribution in the supercluster field extending over $\\sim 10h_{50}^{-1}$ Mpc."
    },
    "0004/astro-ph0004145_arXiv.txt": {
        "abstract": "Comparing models of Simple Stellar Populations (SSP) with observed line strengths generally provides a tool to break the age-metallicity degeneracy in elliptical galaxies. Due to the wide range of Balmer line strengths observed, ellipticals have been interpreted to exhibit an appreciable scatter in age. In this paper, we analyze Composite Stellar Population models with a simple mix of an old metal-rich and an old metal-poor component. We show that these models simultaneously produce strong Balmer lines and strong metallic lines without invoking a young population. The key to this result is that our models are based on SSPs that better match the steep increase of \\Hb\\ in metal-poor globular clusters than models in the literature. Hence, the scatter of \\Hb\\ observed in cluster and luminous field elliptical galaxies can be explained by a spread in the metallicity of {\\em old} stellar populations. We check our model with respect to the so-called G-dwarf problem in ellipticals. For a galaxy subsample covering a large range in \\UVV\\ colors we demonstrate that the addition of an old metal-poor subcomponent does not invalidate other observational constraints like colors and the flux in the mid-UV. ",
        "introduction": "More than 20 years ago it has been recognized that the modeling of the spectral energy distribution of ellipticals is affected by an ambiguity in age and metallicity (Faber 1972; O'Connell 1976)\\nocite{Fa72,Oc76}, which has turned out to be a general complication in population synthesis (Renzini 1986)\\nocite{Re86}. However, considering Simple Stellar Populations (SSP) in the two-parameter space of Balmer and metallic lines, the age-metallicity degeneracy can be broken (Gonz\\'alez 1993; Worthey 1994)\\nocite{G93,Wo94}. The major reason for the success of this strategy is that the Balmer line strengths of SSPs are predominantly age sensitive at metallicities above $\\sim 1/3~\\Zsun$ that are supposed to be the only relevant for elliptical galaxies. Strong \\Hb\\ lines are thus taken as an indication for young (intermediate-age) populations, the observed scatter is interpreted as a considerable spread in age (e.g., Faber et al.\\ 1995)\\nocite{Fetal95}.  An alternative path to obtain blue stars and hence strong Balmer lines is to consider old {\\em metal-poor} populations. In this paper we follow this approach and compute composite stellar populations that contain an old metal-poor subcomponent. The principal focus is to check if a combination of only old populations can reproduce strong \\Hb\\ without invalidating further constraints for ellipticals like metallic indices, colors and spectral energy distributions.  The paper is organized as follows. In Section~\\ref{calsec} we calibrate the SSP model indices and spectral energy distributions on globular cluster data. The composite models and their application to galaxy data are presented in Section~\\ref{resultssec}. In Sections~\\ref{discusssec} and~\\ref{sumsec} we discuss and summarize the results. ",
        "conclusions": "\\label{discusssec} \\subsection{Metallicity distributions in ellipticals} The principle idea of this paper is to enhance the Balmer line strengths of old metal-rich stellar populations with a small fraction of metal-poor stars. The real metallicity distribution of the stellar populations in elliptical galaxies is difficult to assess observationally, as the stars cannot be resolved. The {\\em HST} color magnitude diagram of M32 (Grillmair et al.\\ 1996)\\nocite{Getal96} and the spectroscopy of K giants in the Bulge (Rich 1988)\\nocite{Ri88} show that these spheroids contain a tail of low metallicity stars. The closest giant elliptical for which color magnitude diagrams are available is NGC~5128 (Harris, Harris \\& Poole 1999)\\nocite{HHP99}. Analyzing deep {\\em HST} images in the outer halo of NGC~5128, the authors find a differential metallicity distribution that is well reproduced by two closed-box-like chemical enrichment scenarios implying a large number of extremely metal-poor stars. In particular the differential shape of the distribution implies a formation picture in which a metal-poor and a metal-rich population are formed separately from each other.  This picture gets further support from a number of studies that discover bimodal color and metallicity distributions of globular clusters in at least half of the early-type galaxy population (Zepf \\& Ashman 1993; Gebhardt \\& Kissler-Patig 1999)\\nocite{ZA93,GK99}. Spectroscopic and photometric investigations indicate that both populations are old (Cohen, Blakeslee \\& Ryzhov 1998; Kissler-Patig et al.\\ 1998; Kissler-Patig, Forbes \\& Minniti 1998; Kundu et al.\\ 1999; Puzia et al.\\ 1999)\\nocite{Kietal98,CBK98,KFM98,Kuetal99,Puetal99}. More specifically, Puzia et al.\\ (1999)\\nocite{Puetal99} find that the two populations in NGC~4472 are old and coeval, with metallicities $\\sim 0.05~\\Zsun$ and $\\sim\\Zsun$. In the framework of hierarchical structure formation galaxies are built by mergers of smaller objects (e.g., White \\& Rees 1978; White \\& Frenk 1991; Kauffmann, White \\& Guiderdoni 1993)\\nocite{WR78,WF91,KWG93}. A merger of coeval systems without newly induced star formation would result in a coeval composite population. The accretion of a dwarf galaxy by a larger system (minor merger) qualitatively explains the existence of a metal-poor subpopulation.  There is an additional important effect that is independent of the assumed scheme of galaxy formation. As discussed by Greggio (1997)\\nocite{Gre97}, a sizable fraction of metal-poor stars from the outer parts of a galaxy contaminate the light coming from the center due to projection and orbital mixing (Ciotti, Stiavelli \\& Braccesi 1995)\\nocite{CSB95}. The authors find that projection effects lower the actual central metallicity by roughly 10 per cent. The metal-poor components in the composite models of M31, NGC~4472 and NGC~3379 (Table~\\ref{modeltab}) dilute the metallicities of the metal-rich populations by 3, 6, and 4 per cent, respectively. Projection effects alone thus may be sufficient to explain the amount of metal-poor stars considered in the present analysis.   \\subsection{Evidence against recent star formation} There are observational indications from the Fundamental Plane (Djorgovski \\& Davis 1987; Dressler et al.\\ 1987; Bender, Burstein \\& Faber 1992, 1993; Renzini \\& Ciotti 1993)\\nocite{DD87,Dretal87,BBF92,BBF93,RC93} and the color-magnitude relation (Bower, Lucey \\& Ellis 1992)\\nocite{BBF92} that limit the fraction of the young population to less than 10 per cent. As late star formation leads to low $\\alpha$/Fe ratios (Thomas, Greggio \\& Bender 1999; Thomas 1999)\\nocite{TGB99,Th99a}, the addition of a young population provokes inconsistencies with the high $\\alpha$/Fe ratios observed in elliptical galaxies (e.g., Worthey et al.\\ 1992)\\nocite{WFG92}. Also the redshift evolution of colors (Arag\\'on Salamanca et al.\\ 1993)\\nocite{Aetal93}, of the Mg-$\\sigma$ relation (Bender, Ziegler \\& Bruzual 1996)\\nocite{BZB96}, of the Kormendy relation (Ziegler et al.\\ 1999)\\nocite{Zieetal99}, and of the color gradients in cluster ellipticals (Saglia et al.\\ 2000)\\nocite{Saetal00} leave little space for a significant contribution from a young subcomponent. The addition of an old, metal-poor population as discussed in this paper can reconcile these findings with strong Balmer lines.  Disturbed field ellipticals that have strong \\Hb\\ absorption (Schweizer et al.\\ 1990)\\nocite{Schetal90} and blue optical colors (Schweizer \\& Seitzer 1992)\\nocite{SS92} do not show any signature of recent star formation activity in the infrared colors (Silva \\& Bothun 1998a, 1998b)\\nocite{SB98a,SB98b}. As discussed by these authors, recent mergers may not have been accompanied by significant star formation, but metallicity effects are favored to explain the enhanced \\Hb\\ line strengths."
    },
    "0004/astro-ph0004235_arXiv.txt": {
        "abstract": "We present 53 simultaneous photometric (I band) and spectroscopic (6900-9500\\AA) observations of J0422+32, taken during December 1997. From these we determine that J0422+32 was in its lowest state yet observed, at I=20.44$\\pm$0.08.  Using relative spectrophotometry, we show that it is possible to correct very accurately for telluric absorption.  Following this, we use the TiO bands at 7055\\AA\\ and 7589\\AA\\ for a radial velocity study and thereby obtain a semi-amplitude of 378$\\pm$16~km~s$^{-1}$, which yields {\\it f(M)}=1.191$\\pm$0.021M$_\\odot$ and {\\it q}=9.0$^{\\scriptscriptstyle +2.2}_{\\scriptscriptstyle -2.7}$, consistent with previous observations.  We further demonstrate that this little explored method is very powerful for such systems.  We also determine a new orbital ephemeris of HJD=2450274.4156$\\pm$0.0009 + 0.2121600$\\pm$0.0000002 $\\times$ E.  We see some evidence for an ellipsoidal modulation, from which we determine the orbital inclination of J0422+32 to be less than 45${^\\circ}$.  We therefore calculate a minimum mass for the primary of 2.22M$_\\odot$, consistent with a black hole, but not necessarily the super-massive one proposed by Beekman et al (1997).  We obtain an M4-5 spectral type for the secondary star and determine that the secondary contributes 38$\\pm$2\\% of the flux that we observe from J0422+32 over the range 6950-8400\\AA.  From this we calculate the distance to the system to be 1.39$\\pm$0.15kpc. ",
        "introduction": "From the study of X-ray transient systems, it is possible to determine the masses of stellar-mass black-holes, by observing the cool secondary star during quiescence (see e.g. Charles, 1999). In this paper we present new observations of the late-type secondary star in the soft X-ray transient GRO J0422+32 (Nova~Per~1992/ V518~Per). Since the discovery of J0422+32 on August 5th 1992, whilst in outburst, by the `Compton Gamma Ray Observatory' (Paciesas et al, 1992), there have been three subsequent `mini' outbursts: December 1992 (Harmon, Fishman \\& Paciesas, 1992), August 1993 (Filipenko \\& Matheson, 1993) and December-January 1993/4 (Zhao et al, 1993).  The system has been observed at I=20.03 \\cite{Oros95} and I=20.22 \\cite{Casa95}, but never, as we show in this paper, in absolute quiescence.  Therefore, observations in the optical have been dominated by the flux emitted from the accretion disc around the compact object, making observations of the M-star difficult.  Here we present results showing that J0422+32 was fainter still in December 1997, thus observations of the secondary were more accessible, due to less contamination from the disc.  Previously, Beekman et al (1997) have determined a minimum mass of the compact object (black hole) of 15M$_\\odot$, from an I band light-curve.  From calculations by King, Kolb \\& Burderi (1996) based on the `disc instability model', where the mass transfer rate must be below a critical level  in order for low mass X-ray binary systems to become transient, a minimum mass of the compact object can be calculated.  If their assumptions are correct, Beekman et al (1997) calculated a minimum mass of the compact object of 28M$_\\odot$.  This mass is impossibly large for stellar evolution models. From studying the ellipsoidal modulation alone, we constrain the black-hole to have a much lower minimum mass, which is consistent with the evolution of massive stars that form black holes.  We present a new method for conducting a radial velocity study in such systems with a late-type secondary star.  Previously, the strong TiO features seen in M-stars have not been used, as it is very difficult to correct accurately for the telluric absorption that lies on top of them.  Instead, the Na~I doublet at 8183, 8195\\AA\\ is commonly used.  However, by placing two stars along the slit, J0422+32 and a local early-type star, we can correct for the telluric absorption very accurately.  It was also thought that the TiO bands were far too broad to allow accurate cross correlation.  The bands themselves are broad, but have within them, very sharp features.  In this work we show, using the two TiO bands at 7055\\AA\\ and 7589\\AA, how effective cross-correlating these regions can be.  We also exploit the good telluric absorption correction to determine the spectral type of the M-star accurately, through measuring the depths of these two TiO bands.  A 5.1 hour period, assumed to be the orbital period, has been observed in most earlier data sets (e.g. Chevalier \\& Ilovaisky, 1994; Casares et al, 1995; Garcia et al, 1996).  Earlier still, values close to 5.1 hours were observed as the system declined from the initial outburst (e.g. Chevalier \\& Ilovaisky, 1992; Kato, Mineshige \\& Hirata, 1992). This period has not been observed in other datasets, e.g. Shrader et al (1994) and Martin et al (1995). We have used our data to refine the orbital period and to determine a new ephemeris. ",
        "conclusions": "From our data, we determine that J0422+32 was in its lowest state yet observed during December 1997, at I=20.44$\\pm$0.08.  Following a comprehensive search we measure a period of 0.2121600$\\pm$0.0000002 days, consistent with previous analysis.  From the same radial velocity study, we determine a tightly constrained semi-amplitude of 378$\\pm$16 km s$^{-1}$.  Thus we have shown that using TiO bands for radial velocity studies is a very powerful tool.  We also determine a new T$_\\circ$ of 2450274.4156$\\pm$0.0009.  Further to this we see some evidence for ellipsoidal modulation, implying a maximum inclination for J0422+32 of 45$^\\circ$.  We obtain a new mass function of 1.191$\\pm$0.021 M$_\\odot$ and {\\it q}=9.0$^{\\scriptscriptstyle +2.2}_{\\scriptscriptstyle -2.7}$, from which we have calculated a minimum mass for the primary of 2.22M${_\\odot}$, consistent with a black-hole, but not necessarily as a super massive one.  However, this is only a minimum mass and the maximum is not constrained.  We also revise the calculations of Beekman et al (1997) to demonstrate that the minimum possible mass of the black hole is similar to other black hole systems.  We have shown that it is important to perform a slit correction to the spectra, to calculate the correct flux for such a faint system.  We obtain a spectral type of the secondary of M4-5 and determine that the secondary star contributes 38$\\pm$2\\% of the flux that we observe from J0422.  From this we calculate the distance to the system to be 1.39$\\pm$0.15 kpc."
    },
    "0004/astro-ph0004003_arXiv.txt": {
        "abstract": "\\noindent The Universe could be spatially flat, positively curved or negatively curved. Each option has been popular at various times, partly affected by an understanding that models tend to evolve away from flatness. The curvature of the Universe is amenable to measurement, through tests such as the determination of the angles of sufficiently large triangles. The angle subtended by the characteristic scale on the Cosmic Microwave sky provides a direct test, which has now been realised through a combination of exquisite results from a number of CMB experiments.  \\noindent After a long and detailed investigation, with many false clues, it seems that the mystery of the curvature of the Universe is now solved. It's an open and shut case: the Universe is flat! ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/gr-qc0004050_arXiv.txt": {
        "abstract": "We introduce a new method to construct solutions to the constraint equations of general relativity describing binary black holes in quasicircular orbit.  Black hole pairs with arbitrary momenta can be constructed with a simple method recently suggested by Brandt and Br\\\"ugmann, and quasicircular orbits can then be found by locating a minimum in the binding energy along sequences of constant horizon area.  This approach produces binary black holes in a \"three-sheeted\" manifold structure, as opposed to the \"two-sheeted\" structure in the conformal-imaging approach adopted earlier by Cook.  We focus on locating the innermost stable circular orbit and compare with earlier calculations.  Our results confirm those of Cook and imply that the underlying manifold structure has a very small effect on the location of the innermost stable circular orbit. ",
        "introduction": "Binary black holes are among the most promising sources of gravitational radiation for the new generation of gravitational wave detectors such as the Laser Interferometric Gravitational Wave Observatory (LIGO), VIRGO, GEO and TAMA.  This has motivated an intense theoretical effort to predict the gravitational waveform emitted during the inspiral and coalescence of two black holes~\\cite{itp00}.  Because of the circularizing effects of gravitational radiation damping, we expect the orbits of close binary systems to have small eccentricities.  The inspiral of a binary black hole system then proceeds adiabatically along a sequence of quasicircular orbits up to the innermost stable circular orbit (hereafter ISCO), where the evolution is expected to change into a rapid plunge and coalescence\\cite{footnote0}. The ISCO therefore leaves a characteristic signature in the gravitational wave signal, and knowledge of its location and frequency is thus very important for the prospect of future observations.  While various approximations may be adequate to model the adiabatic inspiral up to the ISCO, it is generally expected that only numerical simulations in full general relativity can accurately model the dynamical plunge and merger and predict the gravitational signal from that phase.  It is therefore desirable to construct initial data for numerical evolution calculations describing binary black hole pairs at the ISCO, which adds another motivation for determining the location of the ISCO.  Various approaches have been adopted to locate the ISCO in compact binaries, including first order post-Newtonian approximations~\\cite{ce77}, variational principles~\\cite{bd92}, second order post-Newtonian methods combined with a ``hybrid'' approach~\\cite{kww92}, a Pad\\'{e} approximation~\\cite{dis98} and an effective-one-body approach~\\cite{bd99,bd00}, and numerical solutions to the constraint equations of general relativity~\\cite{c94,betal98}. Unfortunately, however, the results differ significantly and yet have to show any sign of convergence (see Table~II below).  It would clearly be desirable to understand the origin of these differences. In this paper, we revisit binary black hole solutions to the constraint equations, and evaluate how some of the choices which have to be made in this approach affect the location of the ISCO.  Before the constraint equations of general relativity can be solved, a background geometry and topology have to be chosen.  In the conformal-imaging approach adopted by Cook~\\cite{c94}, a conformally flat (spatial) background metric is chosen together with a two-sheeted manifold structure (see Sec.~\\ref{sec2.1}).  It has been suggested that these choices may affect the location of the ISCO, and may explain the difference between these and the more recent post-Newtonian results.  In this paper, we combine the methods of Cook~\\cite{c94} and Brandt and Br\\\"ugmann~\\cite{bb97} to introduce a new approach to constructing binary black holes in quasicircular orbit.  We follow Cook~\\cite{c94} and choose a conformally flat background metric, but do not assume an inversion-symmetry as is done in the conformal-imaging approach.  This considerably simplifies the solution of the momentum constraint (see Sec.~\\ref{sec2.2}), and produces binary black holes in a three-sheeted manifold structure as opposed to the two-sheeted structure in the conformal-imaging approach.  Moreover, adopting the ``puncture'' approach of Brandt and Br\\\"ugmann~\\cite{bb97}, the Hamiltonian constraint can be solved very easily numerically on R$^3$ without having to impose boundary conditions on interior boundaries (see Sec.~\\ref{sec2.3}).  We locate the ISCO, and find that its physical parameters agree very well with those found with the conformal-imaging approach of Cook~\\cite{c94}.  We therefore conclude that the choice of the underlying mani\\-fold structure has a very small effect on the location of the ISCO.  Our new approach, which is significantly simpler than the conformal-imaging approach, may also provide a framework in which the conformal-flatness assumption may be relaxed, and its effect on the ISCO be evaluated.  The paper is organized as follows.  In Sec.~II, we introduce the basic equations and explain how binary black holes in quasicircular orbit can be constructed.  We discuss our numerical implementation in Sec.~III.  In Sec.~IV we present our results and compare with those from other approaches.  We briefly summarize in Sec.~V. ",
        "conclusions": "Since an accurate knowledge of the ISCO is very important for possible future gravitational wave observations, it is very unsettling that different approaches to computing the ISCO lead to very different results (compare Table~II).  One of these approaches, namely constructing binary black hole solutions to the constraint equations of general relativity, involves {\\em choosing} the background geometry and topology, and it would be very desirable to know how much the results depend on these choices.  In this paper, we introduce a new method to construct solutions to the constraint equations of general relativity describing binary black holes in quasicircular orbit.  We combine the approaches of Cook~\\cite{c94} and Brandt and Br\\\"ugmann~\\cite{bb97} to construct binary black holes in a three-sheeted manifold structure, as opposed to the two-sheeted topology in the conformal-imaging approach adopted by Cook~\\cite{c94}.  We locate the ISCO and find that its physical parameters are very similar to those found by Cook~\\cite{c94}.  Our results confirm those earlier results and imply that the underlying manifold structure only has a very small effect on the ISCO.  The latter is perhaps not entirely surprising, since it reflects the fact that the strength of the imaged poles in the conformal-imaging approach is smaller than the strength of the poles themselves~\\cite{m63}.  Our new approach is considerably simpler than the conformal-imaging approach of Cook~\\cite{c94}.  The analytic solution to the momentum constraint simplifies because no inversion-symmetric solutions have to be constructed, and the numerical solution to the Hamiltonian constraint simplifies because we can adopt the ``puncture'' method of Brandt and Br\\\"ugmann~\\cite{bb97}.  In particular, we can solve the Hamiltonian constraint in cartesian coordinates on $R^3$ without having to impose interior boundary conditions.  One disadvantage of our approach is that the apparent horizons have to be located numerically, which we do with the algorithm developed in~\\cite{betal96}.  In this paper, we follow Cook~\\cite{c94} and choose a conformally flat background metric.  Accordingly, we cannot address the dependence of the ISCO on the choice of the background geometry.  However, since our new method is significantly simpler than the conformal-imaging approach, it may provide a useful framework to relax the assumption of conformal flatness and to construct binary black holes in quasicircular orbit for more general background geometries."
    },
    "0004/astro-ph0004017_arXiv.txt": {
        "abstract": "The boundary layer where the accretion disk meets the star is expected to be the dominant source of high-energy radiation in low-mass X-ray binaries which contain weakly magnetized accreting neutron stars.  We present solutions for the structure of the boundary layer in such a system.  We find that the main portion of the boundary layer gas is hot ($\\ga 10^8$ K), has low density, and is radially and vertically extended. It will emit a large luminosity in X-rays, mainly produced by Comptonization of soft photons which pass through the hot gas.  The gas is generally optically thick to scattering but optically thin to absorption.  Energy is transported by viscosity from the rapidly rotating outer part of the boundary layer to the slowly rotating inner part, and this has the important effect of concentrating the energy dissipation in the dense, optically thick zone close to the stellar surface.  Advection of energy also plays an important role in the energy balance.  We explore the dependence of the boundary layer structure on the mass accretion rate and rotation rate of the star. We also examine the effects of changes in the $\\alpha$ viscosity parameter and the viscosity prescription.  Radiation pressure is the dominant source of pressure in the boundary layer.  The radiation flux in the boundary layer is a substantial fraction of the Eddington limiting flux even for luminosities well below ($\\sim 0.01$ times) the Eddington luminosity $L_{Edd}$ for spherically symmetric accretion.  At luminosities near $L_{Edd}$, the boundary layer expands radially, and has a radial extent larger than one stellar radius.  This radial expansion increases the surface area of the boundary layer and allows it to radiate a larger total luminosity.  Based on the temperatures and optical depths which characterize the boundary layer, we expect that Comptonization will produce a power-law spectrum at low source luminosities.  At high luminosities the scattering optical depth is quite large, and bremsstrahlung and Comptonization will produce a Planckian spectrum in the dense region where most of the energy is released.  This spectrum will be altered by Comptonization as the radiation propagates through the lower-density outer boundary layer.  We discuss some implications of our results for standard multi-component fits to X-ray spectra of LMXBs. ",
        "introduction": "\\subsection{The Boundary Layer}  The boundary layer, the region where the rapidly spinning disk material reaches the more slowly spinning accreting star, is a crucial element of an accretion disk.  In a thin accretion disk, the gas rotates at approximately the Keplerian velocity, and so by the time it approaches the surface of the accreting star, half of the gravitational potential energy released in the accretion process has been converted into rotational energy of the gas (in the Newtonian case).  Unless the star is rotating rapidly, most of this energy will be released in the boundary layer.  Furthermore, since this energy comes from a small region close to the star, the boundary layer should be hotter than the disk and should produce harder radiation.  Thus, the boundary layer is expected to be the dominant source of high-energy emission from accretion disks, which for accretion onto neutron stars takes the form of X-rays.  The boundary layer region is quite complex, since the accreting gas must make the transition from a disk to a star, with the accompanying changes in the balance of momentum and energy.  For example, as the rotation rate of the gas drops below Keplerian, rotational support against gravity is replaced by pressure support.  The radiation which cools the disk flows vertically, from the disk midplane to the surface, while near the stellar surface radiation must flow radially outward.  In general, the radial scale over which the disk properties vary becomes comparable to or smaller than the vertical scale of the disk.  As a result, radial transport, particularly of energy, plays an important role in the boundary layer.  Boundary layers around neutron stars are even more complex, due to a number of additional physical processes which become important due to the small size of the neutron star and the resulting strong gravity and enormous luminosity due to accretion.  In the solutions we present here, radiation pressure plays a major role in the dynamics of the flow, and can increase the sound speed to $\\ga 0.1 c$.  Radiation pressure is dominant even for accretion rates where the total luminosity is well below (0.01 times) the Eddington limit, since the local radiation flux still reaches a large fraction of the local Eddington value.  The gas can reach very high temperatures, so Comptonization of soft incident photons can be an important energy loss mechanism and can produce power-law spectra (Sunyaev \\& Titarchuk 1980).  Energy transfer between protons and electrons can become inefficient at high temperatures, producing a two-temperature plasma. Relativistic effects can also be important, since the neutron star radius is comparable to the radius of the last stable particle orbit in the Schwarzschild metric.  The incident radiation could in principle remove angular momentum from the gas (Miller \\& Lamb 1993, 1996), and if the neutron star radius is smaller than the last stable orbit, the gas may spiral toward the surface (Klu\\'zniak \\& Wilson 1991).  \\subsection{Observational Background}  Low-mass X-ray binaries (LMXBs) have been observed in X-rays for almost 40 years, since the discovery of Sco X-1 (Giacconi \\et 1962). Since then, around 100 other LMXBs have been discovered, and many of these have been observed extensively in X-rays and at other wavelengths.  These observations have revealed much about the timing and spectral behavior of LMXBs.  The best-studied LMXBs generally fall into two classes, the atoll sources and the Z sources, based on the paths they trace out in the X-ray color-color diagram as they vary in brightness (Hasinger \\& van der Klis 1989).  These variations are believed to be due to changes in the mass accretion rate $\\dot M$, with the Z sources having luminosities near the Eddington limit and the atoll sources varying over a wider range, down to about 1\\% of the Eddington limit.  If so, then the atoll sources in particular provide an excellent means for directly observing the effects of changing $\\dot M$ and checking the predictions of our models.  As the sources move along the paths in the color-color diagram, their variability properties also change, as indicated by changes in the shape of their power density spectra.  In some regimes, quasi-periodic oscillations (QPOs) appear.  Some of these are at relatively low frequencies of a few Hz or tens of Hz (van der Klis \\et 1985; Middleditch \\& Priedhorsky 1986), but others are the recently discovered kHz QPOs (see van der Klis 1998 for a review).  Sunyaev \\& Revnivtsev (2000) have recently shown that power density spectra of LMXBs in the low/hard state show much more power at high frequencies (> 100 Hz) than those of black hole candidates, probably due to the presence of the neutron star surface and the associated energy release in a boundary layer.  The high-frequency variability and oscillations presumably originate from the innermost portions of the accretion flow, very near the neutron star, and provide strong motivation for studying these inner regions in detail.  Nonetheless, our understanding of the production of X-rays and the formation of the X-ray spectrum in these sources is still very incomplete.  The X-ray spectra of LMXBs are generally modeled in a rather simple way, by combining two or more spectral components to fit the overall spectrum.  The components consist of blackbody, ``disk'' (a sum of blackbodies corresponding to annuli of a disk with a certain temperature and emissivity profile), bremsstrahlung, or power-law spectra.  Thompson (elastic) scattering or Comptonization (where the frequency changes due to scattering) can play important roles in producing or modifying these components.  These model fits are useful in that they have provided a simple picture of the changes in LMXB spectra as a function of luminosity. At high luminosities, above around $10^{37} \\ergs$, the data are fit well by blackbody-type spectra.  Mitsuda \\et (1984) used a two-component model with a disk component and a single-temperature blackbody to fit high-luminosity LMXB spectra.  White \\et (1986) used a similar model in which the disk emission was Compton-scattered to higher energies.  At lower luminosities, the spectra are fit well by a power-law spectrum with an exponential cutoff at high energies (White, Stella, \\& Parmar 1988).  Both a soft component and a hard power law can be present at low luminosities, but the power law seems to disappear at higher luminosities (Barret \\& Vedrenne 1994).  The components used in these fits incorporate much of our current knowledge of the processes by which hot gas emits X-rays.  However, in order to better understand which components should be present and how their temperatures and luminosities relate to each other, it is important to have a model for the accretion flow near the neutron star.  In particular, the boundary layer region, where the accretion disk meets the star, is expected to produce a large portion of the X-ray luminosity.  The size, temperature, and optical depth of this region depend on the dynamics and energetics of the accretion flow near the stellar surface.  In this paper, we study the boundary layer region in detail, as a step toward the eventual goal of being able to directly interpret LMXB spectra in terms of fundamental parameters such as the mass accretion rate and the rotation rate of the accreting neutron star.  \\subsection{Theoretical Background}  The structure of the boundary layer region has been studied in other types of accreting systems, most notably cataclysmic variables (Pringle 1977; Pringle \\& Savonije 1979; Tylenda 1981; Patterson \\& Raymond 1985; Kley 1991; Narayan \\& Popham 1993; Popham \\& Narayan 1995) and pre-main sequence stars such as T Tauri and FU Orionis stars (Popham \\et 1993, 1996).  Cataclysmic variables (CVs) are similar to LMXBs in many respects; the main difference is simply that accretion is onto a white dwarf instead of a neutron star.  CVs also emit X-rays which are believed to originate in the boundary layer.  Narayan \\& Popham (1993, hereafter NP93) showed that the optical depth of the boundary layer region is sensitive to the mass accretion rate.  At high accretion rates, the boundary layer is optically thick and emits approximately as a blackbody with an effective temperature of a $\\few \\times 10^5$.  But at low accretion rates, the boundary layer becomes optically thin to absorption and is unable to cool efficiently, as had been predicted by Tylenda (1981) and King \\& Shaviv (1984).  The accreting gas is heated to $\\sim 10^8$ K by the energy dissipated in the boundary layer, and emits hard X-rays.  While we expect some similarities between boundary layers in CVs and in LMXBs, we also expect a number of major differences.  Because a neutron star is so much smaller than a white dwarf, a much larger luminosity must be emitted from a much smaller area, resulting in much higher radiation fluxes and temperatures, and therefore Comptonization and radiation pressure play critical roles.  Studies of the accretion flow onto neutron stars have largely focused on the case where the neutron star has a very strong magnetic field. This field is believed to truncate the accretion disk at some inner radius and channel the accretion onto magnetic field lines, so that ultimately it falls onto polar caps corresponding to the poles of the magnetic field (Pringle \\& Rees 1972; Basko \\& Sunyaev 1976; Ghosh, Lamb, \\& Pethick 1977).  Evidence for this magnetically-channeled polar accretion is provided by the X-ray pulsations and magnetic cyclotron features first observed in Her X-1 (Tr\\\"umper \\et 1978) and in a number of other X-ray pulsars.  However, the great majority of LMXBs do not show any evidence for periodic pulsations or cyclotron features.  In these systems the magnetic field may be sufficiently small ($B \\la 10^8$ G, as in some millisecond pulsars) to allow the disk to extend all the way in to the stellar surface, resulting in a boundary layer region where the rapidly rotating disk meets the (presumably) more slowly rotating star.  There have only been a few studies of the inner accretion flow onto non-magnetic neutron stars (those where the stellar magnetic field is not strong enough to alter the flow), but for the most part they have not computed the boundary layer structure in detail.  Sunyaev \\& Shakura (1986) computed the relative contributions of the disk and boundary layer to the total accretion luminosity.  In standard Newtonian disk theory each contribute half of the total, but in the Schwarzschild metric the boundary layer contributes more than the disk.  The relative contributions depend on the neutron star radius $R_*$: if $R_*$ equals the radius $R_{ms}=6 G M /c^2$ of the marginally stable particle orbit, the boundary layer luminosity should be about twice that of the disk.  If $R_* < R_{ms}$ and the accreting gas spirals rapidly in from $R_{ms}$ to $R_*$, the relative contribution of the boundary layer is even larger.  The relative luminosities of the disk and boundary layer also depend on the rotation rate of the star, as shown by Sibgatullin \\& Sunyaev (1998), who included the effects of rotation on the shape of the star and its gravitational field.  Shakura \\& Sunyaev (1988) derived analytic estimates for the boundary layer structure at low accretion rates and X-ray luminosities $L_x < 10^{36} \\ergs$, assuming constant temperature and viscosity coefficient.  They did not address the case of higher luminosities, where radiation pressure should dominate. Klu\\'zniak \\& Wilson (1991) computed the structure of an accretion belt on the stellar surface under the assumption that $R_* < R_{ms}$, and the accreting gas impacts the surface at high velocity, and found that high temperatures and hard spectra would be produced.  King \\& Lasota (1987) argued that the boundary layer region would reach high temperatures even if the gas does not experience rapid infall, because as in the CV case, the gas cannot cool efficiently enough to radiate away the dissipated energy.  They show that for luminosities less than $\\sim 10^{35} \\ergs$, where gas pressure is dominant, the boundary layer region should heat up and expand vertically to form a ``corona'' around the neutron star.  Recently, Inogamov \\& Sunyaev (1999) have studied the problem of disk accretion onto neutron stars, using a new approach for modeling the boundary layer region.  The accreting gas arrives at the equator of the star spinning at the Keplerian velocity, and forms a layer on the stellar surface which spreads from the equator toward the poles.  As the gas moves meridionally, it loses angular momentum and dissipates energy, which is radiated away from the surface.  This approach essentially treats the boundary layer as part of the star rather than part of the disk, and the angular velocity decreases with latitude on the stellar surface rather than with radial distance from the surface. This complements the approach used in the current paper, and later we compare the results of the two approaches.  We model the boundary layer as part of the disk, using the slim disk equations (Paczy\\'nski \\& Bisnovatyi-Kogan 1981; Muchotrzeb \\& Paczy\\'nski 1982; Abramowicz \\et 1988), which contain terms which allow for large deviations from the standard thin Keplerian disk with efficient cooling.  A similar approach has been used in most previous studies of boundary layers in CVs and accreting pre-main sequence stars.  This approach has the advantage that it allows one to solve for the structure of the disk and boundary layer together, using a single set of equations throughout.  The interface with the accreting star is treated as a set of boundary conditions implemented at the stellar radius.  We use Newtonian equations throughout, despite the fact that our adopted neutron star radius of 10 km is less than the radius of the marginally stable particle orbit (12.4 km for a neutron star mass of 1.4 $\\msun$).  Since there have been no previous solutions of these equations for neutron star parameters, either in Newtonian or relativistic form, we feel that a Newtonian solution is an important first step.  The effects of relativity will be added in the future.  Figure 1 shows some of the important features of our results.  At the transition from the disk to the boundary layer, the angular velocity $\\Omega$ reaches a maximum and the flow passes through a narrow neck, where the disk height is only $\\sim 40 - 70$ meters in the low-$\\dot M$ solutions.  The radial extent of the boundary layer region is $\\sim 0.1-0.2$ of the stellar radius at low $\\dot M$, and the height is comparable to the radial extent.  The angular velocity drops slowly over most of the boundary layer, and then rapidly at the inner edge. At high $\\dot M$ near the Eddington limit, the situation is quite different: the radial extent and height boundary layer are equal to the stellar radius, and the neck between the disk and boundary layer is much wider.  Here the drop in angular velocity occurs over the whole width of the layer.  In \\S 2, we describe the slim disk equations and radiative transfer scheme which we have used to model boundary layers in LMXBs.  We present expressions for the viscous transport of energy in the disk and boundary layer.  We present the results of our calculations in \\S 3, and show how the transport of energy by viscosity, radiation, and advection play essential roles in determining the boundary layer structure.  We present solutions for a variety of mass accretion rates, stellar rotation rates, and viscosities, and show how the size, temperature, and other properties vary.  In \\S 4 we discuss the energetics of the boundary layer, the behavior near the Eddington limit, and the implications of our results for the spectra of LMXBs. ",
        "conclusions": "\\subsection{Energetics of the Boundary Layer}  We have found that boundary layers around disk-accreting neutron stars will be hot, low in density, optically thin to absorption, and both radially and vertically extended.  We begin by discussing some of the processes which are important in producing this boundary layer structure.  We focus particularly on the important role played by radial energy transfer.  We find that the accreting gas makes a fairly abrupt transition from the disk to the boundary layer.  The disk is thin, relatively cool and dense, and optically thick to absorption, while the boundary layer is geometrically thick, hot and rarefied, and optically thin to absorption for most choices of parameters.  There is a similar abrupt transition in reverse when the hot gas nears the stellar surface. These transitions are related to the thermal instability of the hot, low-density boundary layer gas discussed by King \\& Lasota (1987), in which the gas cannot efficiently radiate away the dissipated energy. However, the situation is more complicated than the one they envisioned, since they confined their analysis to local heating and cooling by dissipation and radiation, in a gas-pressure-dominated disk (corresponding to lower accretion rates and luminosities than those considered here).  In our solutions radiation pressure is dominant and the energy balance is dominated not by local dissipation and radiation, but instead by Comptonization and advection.  In our solutions the gas reaches temperatures of a $\\few \\times 10^8$ K, and nearly reaches $10^9$ K in one case.  One might expect the large energy release to make the boundary layer gas even hotter than this; if all the accretion energy were to go into heating the gas, it would reach the virial temperature of a $\\few \\times 10^{11}$ K. However, the presence of Compton cooling keeps the electron temperature from rising above $\\sim 10^9$ K.  Solutions computed without Compton cooling reached very low optical depths and high temperatures in excess of $10^9$ K even for very small values of $\\alpha = 10^{-4} - 10^{-3}$. For $\\alpha=0.1$ these temperatures would be much higher, and would presumably approach the virial temperature.  This illustrates the dominant role played by Comptonization in transferring energy from the gas to the photons and cooling the boundary layer region.  We have solved separate energy equations for the electrons and ions, allowing for the possibility of a two-temperature plasma.  In most of our solutions, the ion and electron temperatures are essentially the same at all radii; however, in solutions with $\\dot M \\la 10^{-9} \\msyr$ and $\\alpha \\ga 0.1$, the ion temperature is significantly higher than the electron temperature in the boundary layer region. The ion temperature increases rapidly as $\\dot M$ decreases, reaching $\\sim 2 \\times 10^{10}$ K at $\\dot M = 10^{-10} \\msyr$, $\\alpha=0.1$. This increased ion temperature produces an increase in the gas pressure, and together with the smaller luminosity, this will lead to gas pressure becoming the dominant source of pressure at low $\\dot M$.  Radial transport of energy by advection also plays an important role in our solutions.  This is not surprising, since NP93 found that advection plays an important role in hot CV boundary layers, which led to the ``rediscovery'' of advection-dominated accretion by Narayan \\& Yi (1994).  The importance of advection can be seen simply by noting that the disk is not geometrically thin in the boundary layer region, which means that the energy density is significant compared to gravity.  However, the flow in the boundary layer differs in some respects from the standard advection-dominated accretion flow (ADAF), which is heated by viscous dissipation and cooled (very inefficiently) by Compton cooling, so that essentially all the energy goes into heating the gas.  In the boundary layer, as discussed in \\S 3, the energy balance is between Compton heating (or cooling) and advection. In the outer boundary layer, the energetic photons contribute far more energy than viscous dissipation, and the gas cannot radiate this energy away, so it heats up and expands and the energy is carried inward.  In the inner boundary layer, this situation is reversed. Thus, even though the situation differs somewhat from a standard ADAF, advection is nonetheless an important energy term throughout the hot region.  One of the main predictions of the mechanics of the boundary layer is that viscous transport carries energy inward, toward the dense, optically thick region at the surface of the star.  The situation is very similar to that in the standard disk (Shakura \\& Sunyaev 1973), where viscous transport carries energy away from the central regions of the disk and delivers it to larger radii, increasing their energy release by a factor of three compared to the simple estimate based on the the change in the gravitational and rotational energy of the infalling gas.  In the boundary layer, where $d \\Omega/dr$ has a different sign, the viscous torque transports energy toward the star and this leads to important consequences.  This can be seen in Figure 5, which shows the energy release due to viscous dissipation, the viscous transport of energy, and the change in kinetic energy.  In the dense region in the innermost part of the boundary layer, the rotational velocity of the gas is much smaller than Keplerian.  If we consider the deepest layer, where the radiation flux is still much smaller than the local critical Eddington flux, we see immediately that the density of the matter must be very high, because it is not supported strongly by centrifugal forces or by the radiation pressure gradient $dP_{rad}/dR \\propto F_R$.  A large fraction of the energy is dissipated in, and subsequently radiated away from, this high-density region.  Much of the radiation travels radially outward into the hot, low-density region, providing the seed photons for Comptonization.  In the solution shown in Figs. 3--5, the density at the point where the dissipation peaks is nearly $1 \\gcm$, and the radial scattering optical depth $\\tau_s = \\int \\rho \\kappa_s dR$ between that radius and the radius where the temperature peaks is $\\sim 20$, as shown in Fig. 6.  It is crucial for the formation of the spectrum that viscous transport and advection carry the bulk of the energy from the low-density, rapidly rotating outer boundary layer to the dense, slowly rotating region where the gas reaches the stellar surface.  \\subsection{The Eddington Flux Limit}  The Eddington flux limit, where the outward radiation pressure gradient balances gravity, plays an important role in our solutions. For spherical accretion, $\\dot M$ is limited to $\\dot M_{Edd} = 4 \\pi c R_* / \\kappa_s \\simeq 1.73 \\times 10^{-8} \\msyr$ for $R_* = 10^6 \\cm$.  In a disk geometry, where the luminosity is not emitted isotropically, the crucial quantity is the critical Eddington flux which balances the local gravitational force.  The radial flux through the disk is limited to the local Eddington value $F_{Edd} = (G M / R^2) c/\\kappa_s$.  In the vertical direction, the downward component of gravity increases with distance $z$ from the midplane as $\\sim z/R$.  Thus at the disk surface $z \\sim H$, the vertical flux is limited to $F_{Edd,V} = (H/R) (G M / R^2) c/\\kappa_s$.  For high values of $\\dot M$, the radial and vertical radiation fluxes approach their respective Eddington limiting fluxes.  In Fig. 8 we showed that for $\\dot M = 10^{-8} \\msyr$, the radial flux comes very near the Eddington value, but only over a small range of radius, and drops steadily as the radiation moves outward through the boundary layer.  This reflects the fact that more of the radial support against gravity comes from the centrifugal force and less from radiation pressure as one moves outward through the boundary layer.  The vertical flux, on the other hand, stays within 2\\% of the Eddington value throughout the entire boundary layer.  Even at $\\dot M = 10^{-9} \\msyr$, the vertical flux stays between 80\\% and 90\\% of $F_{Edd,V}$. This close correspondence between $F_V$ and $F_{Edd,V}$ means that the radiation flux from the boundary layer has a radial profile varying as $H/R^3$.  Physically, it means that at high $\\dot M$ the boundary layer is radiating as much flux as it can.  In order to increase $\\dot M$, the boundary layer must radiate more energy, and in order to do this it must expand either radially or vertically, or both.  By expanding radially, it increases the surface area through which the energy can escape, while vertical expansion increases the gravity and the limiting flux.  In the outer portion of the expanded boundary layer, both the radiation flux and centrifugal forces are important in supporting matter against gravitational attraction to the neutron star.  The rotational velocity decreases significantly in this region as the gas flows inward, and is replaced by radiation pressure support; however, the energy release due to viscous dissipation is very small, since as we have seen, viscous transport carries most of the energy to be dissipated farther in. Thus this region serves mostly to radiate the energy which is dissipated farther in.  We have shown that the boundary layer is both radially and vertically extended, with the radial width reaching $\\sim R_*$ and $H/R \\simeq 0.8$ at $\\dot M \\sim 10^{-8} \\msyr$ (Fig. 7).  If we continue to increase $\\dot M$, the boundary layer continues to expand radially and increase its emitting area (Fig. 9).  The radial expansion is quite rapid, with the boundary layer size doubling for a 40\\% increase in $\\dot M$.  By increasing the emitting area, the radial expansion makes it possible for the boundary layer to radiate away the very high luminosities associated with these high values of $\\dot M$, while keeping the vertical flux below the Eddington limit.  This illustrates an important difference between disk accretion and spherical accretion.  In spherical accretion, the Eddington limit is global, in the sense that any luminosity produced inside a given radius must contribute to the total outward flux at that radius.  Thus the total luminosity of the system is constrained by the Eddington limit.  In a disk system, the Eddington limit is local; the local radial and vertical fluxes cannot exceed the local gravity.  However, there is no global limit on the luminosity, since it can be radiated away through the surface area of the boundary layer, and as we have seen, the boundary layer can expand to radiate additional luminosity as needed.  \\subsection{Radiation Spectrum}  The structure of the boundary layer described above permits simple modeling of the formation of the radiation spectrum which leaves the boundary layer.  We plan to calculate detailed spectra based on our solutions in a future paper; however, we can make some general statements.  It is clear that Compton scattering will play a dominant role in the formation of the X-ray spectrum of the radiation which travels through the hot, low-density region.  The seed photons which are Compton-scattered are primarily emitted from the denser gas which piles up at the inside edge of the hot region.  The Comptonized spectrum will be characterized by $T_e$ and $\\tau_s$, and especially by the Compton $y$-parameter $y = (T/T_{com}) max(\\tau_s, \\tau_s^2)$.  These parameters vary with radius in our solutions, and of course $T_e$ and $\\rho$ also vary with vertical position in the disk, and this variation is not included directly in our solutions.  Also, $\\tau_s = \\kappa_s \\rho H$ in the vertical optical depth for our solutions, but the photons are not travelling only vertically, but also radially and azimuthally, and being scattered in all directions, so the path traveled by a given photon may be much longer than $H$. Therefore, although we can get a rough sense of what the X-ray spectrum will look like based on the characteristic values of $T_e$ and $\\tau_s$ for a given solution, the formation of the real spectrum will be considerably more complex.  In the optically thick region, where Thompson scattering nevertheless strongly dominates the opacity, free-free processes easily produce many photons at low frequencies where they ($K_\\nu \\sim \\nu^{-2}$) are more effective than Comptonization.  Comptonization increases the energy of low frequency photons due to the Doppler effect and leads to the diffusion of the photons towards higher frequencies.  At some frequency $x_0 = {{h\\nu_0\\over kT_e}}$ the rate at which Comptonization can take photons and bring them to higher frequencies equals the rate of photon absorption due to bremsstrahlung. Practically every photon born with a frequency higher than $x_0$ will be transported toward the frequency $x=3$ due to the Comptonization process.  This picture is very similar to the processes occurring in the early universe and is described in detail by Illarionov \\& Sunyaev (1975).  Comptonization is very effective because the parameter $y = {{kT_e}\\over{m_ec^2}} \\tau_s^2 \\gg 1$.  Under these circumstances a Wien-type spectrum must be formed when $x_0$ is very small and $y \\ga 1$ but not very high.  At higher $\\tau_s$ and $y$ a Bose-Einstein spectrum is formed and the formation of the black body spectrum is possible inside very deep regions.  Radiation with this spectrum diffuses out from the dense regions toward more and more rarified regions, where the production of new soft photons becomes more and more difficult due to the low density. Comptonization continues to dominate, and therefore the spectrum tends to be close to a Wien spectrum, with a strong increase of intensity at low frequencies where the spectrum becomes a Rayleigh-Jeans spectrum. The temperature of electrons in the outer regions is determined by the radiation: high energy photons heat electrons due to the recoil effect, and low energy photons cool them down.  (See Levich \\& Sunyaev 1971).  This process is illustrated by Fig. 5. We see that Comptonization takes energy from the plasma in the inner part of the boundary layer and gives energy back to the plasma in the outer part, heating the electrons.  As a result, we are producing a quasi-Wien-type radiation spectrum with strong low frequency excess, and radiation with such a spectrum escapes from every point of the surface of the boundary layer.  Different regions of the boundary layer have a range of temperatures, depending on the distance from both the stellar surface and from the midplane of the disk.  Therefore we will observe a spectrum which is a sum of Wien spectra with different temperatures.  The picture described above follows the results of extended calculations by Grebenev \\& Sunyaev (2000) for the spreading layer model (Inogamov \\& Sunyaev 1999) of the surface of the neutron star. It is interesting to note that this ``spreading layer'' picture of the boundary layer, in which the gas loses angular momentum as it spreads over the stellar surface, gives rather similar results to the ones presented here.  In both cases, the size of the layer increases as the mass accretion rate and luminosity increase.  The ``spreading layer'' has a meridional extent of about 0.45 km at $L/L_E = 0.01$ (which should correspond approximately to our $\\dot M = 10^{-10} \\msyr$ solutions), increasing to $\\sim 17$ km at $L/L_E = 0.8$.  In our solutions the radial extent of the boundary layer is $\\sim 1, 0.5$ km for $\\dot M = 10^{-10} \\msyr$ and $\\alpha=0.1,0.01$, respectively, and $\\sim 10$ km for $\\dot M = 10^{-8} \\msyr$.  In addition, the optical depths and $y$-parameters are quite similar in the two types of solutions.  Inogamov \\& Sunyaev (1999) found $\\tau_s \\sim 3$ at $L/L_E = 0.01$ and $\\tau_s \\sim 1000, y \\sim 10^5$ at $L/L_E = 0.8$, quite similar to the values presented above.  Thus the two treatments should result in rather similar spectra.  Most importantly, in both cases the majority of the energy release occurs in a dense region which is covered by a levitating low density region where the final spectrum is formed.  The agreement is remarkable considering that the accretion flow is treated quite differently in the two approaches.  The question of which approach is the correct one is difficult to answer at present.  The main difference between them is the assumed geometry of the boundary layer region.  The ``spreading layer'' treatment assumes that the disk material enters the spreading layer at nearly Keplerian rotation velocity and small disk height; i.e., very little angular momentum is lost in the disk and there is no disk boundary layer of the type calculated in the present paper. Conversely, the approach taken here assumes that the drop from Keplerian to the stellar angular velocity takes place during the radial inflow of the gas, rather than during the spreading of the accreted gas over the stellar surface.  A multi-dimensional treatment will be required to distinguish between these two possibilities.  The treatment of viscosity in the disk and surface layer may have an important impact on the results.  Note that the viscosity prescription adopted by Inogamov \\& Sunyaev (1999) is also quite different from the one used here; the viscosity decreases as the gas approaches the neutron star surface in analogy with the behavior of fluid near a wall in laboratory experiments.  Yet despite these differences their results are rather similar to ours.  \\subsection{Comparison to Observed Spectra of LMXBs}  Most observed LMXBs have been in one of two spectral states: a low state characterized by low luminosity and a hard, power-law spectrum, or a high state characterized by high luminosity and a softer spectrum.  For example, the four LMXBs recently observed by Barret \\et (1999) included three in the low state and one in the high state.  The three low state sources 1E1724-3045, GS1826-238, and SLX1735-269, all had $1-200 \\keV$ luminosities of $\\sim 1 - 1.5 \\times 10^{37} \\ergs$, and their spectra were fitted by thermal Comptonization by gas with electron temperatures of $\\sim 25-30$ \\keV and optical depths of a few.  The high state source KS1731-260 had $L (1-200 \\keV) \\simeq 8-9 \\times 10^{37} \\ergs$, and was fitted by a much softer Comptonized spectrum with $T_e \\sim 2.6-2.8$ keV and $\\tau \\sim 10$.  Our $\\dot M = 10^{-10}$ solution with $\\alpha=0.1$ and regular viscosity (shown in Fig. 7) has $T_e$ in the hot boundary layer varying from $\\sim 10^8 - 10^9$ K and $\\tau_s \\sim 0.5 - 1$, giving a Compton $y$-parameter of less than 1.  The large variations in $T_e$ and $\\tau_s$ make it difficult to predict what the spectrum will look like, but probably it will have a general power-law shape, and may extend to rather high energies due to the very high $T_e$ in the hottest part of the boundary layer.  At moderate $\\dot M \\sim 10^{-9} \\msyr$, $T_e \\sim 2-3 \\times 10^8$ K and both $\\tau_s$ and $y \\sim$ a few.  Unsaturated Comptonization in the hot boundary layer should produce a power-law spectrum with a cutoff at $\\sim 20-30$ keV.  This solution has a total luminosity of $\\sim 10^{37} \\ergs$, which corresponds approximately to that of the low-state LMXBs oberved by Barret \\et (1999), and the values of $T_e$ and $\\tau_s$ agree well with those from their fits.  At high $\\dot M$, $\\tau_s$ is in the hundreds (for $\\alpha=0.1$) or thousands (for $\\alpha=0.01$) and $y \\sim 10^5$.  The optical depth for combined absorption and scattering is $\\tau_* \\sim 0.1$ for $\\alpha=0.01$ but only $\\tau_* \\sim 0.001$ for $\\alpha=0.1$.  Saturated Comptonization will then produce a Wien spectrum which peaks at $\\sim 3 k T$ (Illarionov \\& Sunyaev 1975).  $T$ varies from $\\sim 10^8$ K for $\\alpha=0.01$ to $\\sim 3 \\times 10^8$ K for $\\alpha=0.1$.  This will produce spectra with $kT \\sim 10-30$ keV.  However, these are the spectra that will be produced by the gas near the disk midplane, but these solutions have such large values of $\\tau_s$ and $y$ that we expect a substantial temperature gradient, with the gas near the surface much cooler than at the midplane.  Thus the spectrum will strongly resemble a blackbody spectrum with a characteristic temperature which is close to the effective temperature of the gas rather than the midplane temperature, i.e. around 1.5--2 keV. Blackbody fits to luminous LMXBs (e.g. Mitsuda \\et 1984) give $kT \\sim 2$ keV, while White \\et (1986) used a $1-2$ keV blackbody plus an unsaturated Comptonized component with a cutoff at $3-8$ keV.  Overall, it appears that our solutions should fit the spectral data reasonably well.  The variation of boundary layer temperature with $\\dot M$ depends on $\\alpha$; for low $\\dot M$, the boundary layer is much hotter for $\\alpha=0.1$ than for $\\alpha=0.01$.  Since low-luminosity LMXBs generally are observed to have rather hard power-law tails which are fit using electron temperatures of $\\sim 30 \\keV$, it seems that our models will agree with the observations better if a large value of $\\alpha$ is used.  X-ray spectra of LMXBs frequently show evidence for iron line emission at about 6.4 keV, which is believed to result from X-ray irradiation of the surface of the disk.  In our solutions, the hot boundary layer is much thicker than the inner disk, so the disk surface should intercept a reasonable fraction of the X-ray emission.  This fraction will depend on the boundary layer thickness and radial extent and the disk height profile.  Fig. 1 shows the height profile of the boundary layer and disk for four solutions.  The solution with $\\dot M = 10^{-8} \\msyr$ has a thicker boundary layer and disk than the $\\dot M = 10^{-9} \\msyr$ or $10^{-10} \\msyr$ solutions, so a larger fraction of the boundary layer emission should be intercepted by the disk.  The solution for a rotating neutron star has a thinner boundary layer than for the nonrotating star.  Note that the axes are chosen to emphasize the differences between the three solutions and give a distorted impression of the shape of the disk; in fact the disk is quite thin, with $H/R \\sim 0.01-0.02$ at all radii.  Thus the disk is basically flat, and the fraction of the boundary layer emission intercepted by the disk will be around 25\\% (Lapidus \\& Sunyaev 1985).  The stellar surface will also intercept a substantial fraction of the X-ray emission (Popham 1997), but most of this will be reradiated back into the hot boundary layer gas.  The X-ray flux incident on the disk and star will result in a number of interesting effects, including polarization (Lapidus \\& Sunyaev 1985), line emission, and a ``Compton reflection'' spectrum.  \\subsection{Implications for Multicomponent Spectral Fits}  With a self-consistent picture of the dynamics and energetics of the boundary layer region, we are now in a position to assess the multi-component models commonly used to fit LMXB spectra.  In particular, we can discuss the emitting regions which are present. The first component, present in most of our solutions, is a hot, low-density boundary layer region which cools by inverse-Compton scattering of photons emitted from the cooler, optically thick zone near the stellar surface.  The hot region also emits bremsstrahlung radiation, but in our solutions this is an insignificantly small fraction of the total emission.  The temperature and density of the gas vary across the boundary layer, so the use of a single-temperature Comptonizing cloud will only approximate the true emitted spectrum. This temperature variation is quite pronounced at low $\\dot M$ for $\\alpha=0.1$; for other choices of parameters the single-temperature approximation may not be so bad.  The second component is the disk, which is optically thick due to the combined effects of absorption and scattering.  The scattering opacity is much larger than the absorptive opacity in our solutions, so the disk should emit a modified blackbody spectrum.  In general the effective temperature of the inner disk is in the range $3-6 \\times 10^6$ K, so $kT \\simeq 0.25-0.5$ keV, but the color temperature of the disk radiation will be higher due to the modified blackbody spectrum. In multi-component models, the ``multi-color disk'' component is often parameterized using an inner radius $r_{in}$; in this context it is useful to note that the inner disk radius in our solutions is set by the radial extent of the boundary layer, and it varies by a factor of two as $\\dot M$ changes.  Two-component models which consist of a Comptonized blackbody spectrum plus a modified blackbody multi-color disk would provide the best approximation to the solutions shown here.  The two components would be constrained to having the appropriate luminosities.  For a narrow boundary layer around a non-rotating star, the disk and boundary layer each contribute half of the total accretion luminosity $L_{acc} \\equiv G M \\dot M / R_*$.  The disk luminosity can be substantially lower than this, since the inner radius of the disk can be as much as $\\sim 2 R_*$.  The boundary layer luminosity is much less than $L_{acc}/2$ if the accreting star is rotating in the same direction as the disk, which it should be due to accretion spinup.  The total luminosity of the boundary layer and disk varies as $1 - jf + 0.5f^2$, where $f \\equiv \\Omega_* / \\Omega_K(R_*)$ is the spin rate of the star as a fraction of the breakup rate (PN95).  If the outer edge of the boundary layer is at a radius $b R_*$, then $j \\simeq b^{1/2}$.  The disk and boundary layer luminosities will then be \\begin{equation} L_{disk} \\simeq L_{acc} \\left( {1 \\over 2b} \\right) \\qquad L_{bl} \\simeq L_{acc} \\left(1 - b^{1/2} f + {f^2 \\over 2} - {1 \\over 2b} \\right). \\end{equation} By using two components with the correct temperatures and optical depths and these luminosities, it should be possible to produce a reasonable approximation to the spectrum that would be emitted by our solutions.  \\subsection{Comparison with Boundary Layers in CVs and Implications for Oscillations}  In many respects the boundary layer solutions for accreting neutron stars in this paper resemble the solutions for very hot boundary layers in CVs presented by NP93.  In both CVs and LMXBs, the radial extent of the boundary layer can be comparable to the stellar radius. One difference between the two types of solutions is that the radial extent of the boundary layer generally increases with increasing $\\dot M$ for LMXB solutions, while it increases with {\\it decreasing} $\\dot M$ for CV solutions.  Both types of boundary layers reach high temperatures $T \\sim 10^8$, but in some cases the LMXB solutions get quite a bit hotter than this.  One might expect the LMXB boundary layers to be much hotter than the CV ones, but as discussed above, Compton cooling limits the temperature.  The CV boundary layer solutions of NP93 have $\\tau_s < 1$ and $T_e \\sim 10^8$ K, so the effects of Compton cooling should be small.  One important observational characteristic of LMXBs is the presence of kHz QPOs.  One of us has argued that kHz QPOs in LMXBs are very similar to dwarf nova oscillations (DNOs) observed in CVs, and that DNOs could arise at the boundary between the disk and the hot, low-density boundary layer (Popham 1999, 2000).  Since we have shown here that LMXBs should have hot, low-density boundary layers similar to those in CVs, the logical next step is that kHz QPOs could arise at this same location.  In our solutions the disk--boundary layer transition occurs at $\\sim 1.1 - 2.0 R_*$, moving outward as $\\dot M$ increases from $10^{-10}$ to $10^{-8} \\msyr$.  The Keplerian rotation frequencies for this range of radii are $\\sim 770-1880$ Hz for our choice of the neutron star mass ($1.4 \\msun$) and radius (10 km).  For a larger neutron star radius, e.g. 13 km, the range would be $\\sim 520-1270$ Hz, which matches the observed range of frequencies quite well.  One difficulty with this picture is that unlike the CV solutions (NP93), our LMXB solutions have this transition radius increasing with increasing $\\dot M$.  If the oscillation period is just the Keplerian period at the transition radius, it should also increase with increasing $\\dot M$.  However, in the oscillations observed thus far, as the oscillation period increases, $\\dot M$ is inferred to decrease. Thus, if kHz QPOs are formed at the disk--boundary layer transition, either $\\dot M$ must change in the opposite sense to that inferred from the observations, or the change in the transition radius with $\\dot M$ must be opposite to what our models predict.  Note that the transition radius reaches a minimum for $\\dot M < 10^{-9.8} \\msyr$ (for $\\alpha=0.1$), and then begins to move back out again as $\\dot M$ decreases.  Thus for low values of $\\dot M$, the sense of the variation of the Keplerian period with $\\dot M$ would agree with the kHz QPO observations; however, these values of $\\dot M$ correspond to lower luminosities than are observed from the systems which show kHz QPOs.  For higher values of $\\alpha$, the transition radius might turn around at a higher $\\dot M$, and when we add additional physical effects to our model, the dependence of the transition radius on $\\dot M$ may change.  \\subsection{Limitations of the Model and Future Improvements}  One obvious limitation of the model used here is the use of one-dimensional, vertically-averaged equations to model the boundary layer region.  This region is inherently two-dimensional, and by using one-dimensional equations, we are unable to simulate a number of aspects of the flow of matter and radiation.  Most of the important physical quantities are assumed to be constant with height $z$ above the midplane, but in reality they vary with both $R$ and $z$.  The viscous dissipation rate varies with $z$, and the processes which produce radial energy transport will also produce vertical energy transport.  As discussed in \\S 2, the abrupt drop in the disk height $H$ just before the gas reaches the stellar surface may also be a consequence of our use of simplified one-dimensional equations. Finally, the main differences between the model presented here and the ``spreading layer'' model of the flow of the gas over the neutron star surface are different assumptions about the geometry of the gas as it reaches the star.  By constructing a two-dimensional model, we could eliminate a number of these problems.  Another important improvement to be made to the model is the inclusion of general relativistic effects.  We have found infall velocities $v_R$ which are at most $\\sim 0.01 c$, but the accreting gas may reach much larger infall velocities if it falls inside the marginally stable orbit.  In addition to purely dynamical effects, radiation drag can remove angular momentum from the gas (Miller \\& Lamb 1993, 1996).  Our current equations for radiative transfer are also rather crude. In particular, the assumption of frequency independence makes our treatment of absorption and of Compton scattering very approximate. Also, in our current form of the radiative transfer equations, all radial flux stays inside the disk, except that which is scattered or absorbed and reemitted as vertical flux.  This is reasonable when the disk height varies slowly with radius, but in our boundary layers the disk height varies rapidly (see Fig. 1).  This should be taken into account; for instance, some of the outward radial flux should escape from the outer side of the boundary layer as ``vertical'' flux when the disk height drops rapidly there.  Since radiation pressure is important in supporting the gas, this could affect the size of the boundary layer region."
    },
    "0004/astro-ph0004221_arXiv.txt": {
        "abstract": "Color gradients in elliptical galaxies in distant clusters ($z=0.37-0.56$) are examined by using the archival deep imaging data of Wide Field Planetary Camera 2 (WFPC2) on-board the Hubble Space Telescope (HST). Obtained color gradients are compared with the two model gradients to examine the origin of the color gradients. In one model, a color gradient is assumed to be caused by a metallicity gradient of stellar populations, while in the other one, it is caused by an age gradient. Both of these model color gradients reproduce the average color gradient seen in nearby ellipticals, but predict significantly different gradients at a redshift larger than $\\sim$0.3. Comparison between the observed gradients and the model gradients reveals that the metallicity gradient is much more favorable as the primary origin of color gradients in elliptical galaxies in clusters. The same conclusion has been obtained for field ellipticals by using those at the redshift from 0.1 to 1.0 in the Hubble Deep Field-North by Tamura et al. (2000). Thus, it is also suggested that the primary origin of the color gradients in elliptical galaxies does not depend on galaxy environment. ",
        "introduction": "It has been known that nearby elliptical galaxies have color gradients; colors in an elliptical galaxy gradually become bluer with increasing radius (e.g., Vader et al. 1988; Franx, Illingworth, \\& Heckman 1989; Peletier et al. 1990a; Peletier, Valentijn, \\& Jameson 1990b, Goudfrooij et al. 1994; Michard 1999). Since many of elliptical galaxies show radial gradients in metal absorption line strengths such as Mg$_{2}$, Fe$_{1}$(5270 \\AA) and Fe$_{2}$(5335 \\AA) (e.g., Carollo, Danziger, \\& Buson 1993; Davies, Sadler, \\& Peletier 1993; Gonzalez 1993; Kobayashi \\& Arimoto 1999), the origin of the color gradients has been naively interpreted to be the metallicity gradients.  However, such an interpretation for the origin of the color gradient is premature, because both metallicity gradient and age gradient in stellar population can cause the same color gradient, and we cannot distinguish the cause for the gradient.  This is called {\\it age-metallicity degeneracy}, which was originally pointed out by Worthey, Trager, \\& Faber (1996) in terms of the origin of the color-magnitude relation of nearby elliptical galaxies (see also Arimoto 1996). In order to break this degeneracy and to know the primary origin of the color gradients in elliptical galaxies, comparing the observed color gradients in distant ellipticals with predicted model gradients caused by either the metallicity gradient or the age gradient is a very effective approach, as was successful for examining the origin of the color-magnitude (CM) relation (Kodama \\& Arimoto, 1997).  Tamura et al. (2000; hereafter called Paper I) constructed the two models both of which reproduce a typical color gradient of elliptical galaxies at $z=0$ using a population synthesis model.  In one model, the mean metallicity of the stellar population decreases with increasing radius at a fixed old mean age. While in the other one, the mean age decreases with a radius at a fixed mean metallicity.  These models were then made evolve back in time.  The evolution of color gradients thus predicted are confronted with the observed ones in distant ($z = 0.1 \\sim 1.0$) ellipticals sampled from the {\\it Hubble Deep Field-North} (HDF-N; Williams et al. 1996).  As a result, Paper I found that the metallicity gradient is the primary origin of color gradients and the age gradient model cannot reproduce the observed gradient at such redshift.  The elliptical galaxies in the HDF-N, however, are only those in field environment. It has never been obvious that ellipticals in clusters evolve similarly as those in field.  In rich clusters, it has been found that the color-magnitude relation still holds even at around $z \\sim 1$ (e.g., Stanford, Eisenhardt, \\& Dickinson 1998) and these observational results seem to favor the classical monolithic collapse scenario associated with the galactic wind and high-$z$ formation (e.g., $z > 3$) of elliptical galaxies (e.g., Kodama et al. 1998).  However, this kind of evolution has not been established for ellipticals in lower density environment (but see Kodama, Bower, \\& Bell 1998).  Some predictions either theoretically or observationally show that field ellipticals formed by recent (at $z \\leq 1$) merging processes (e.g., Baugh, Cole, \\& Frenk 1996; Barger et al. 1999).  An internal structure of a galaxy such as a metallicity gradient and an age gradient must depend on its formation process.  If cluster ellipticals pass different formation histories from those for field ellipticals, their internal structures, thus the origin of the color gradients, may not be the same. Or some environmental effects on color gradients may exist. Thus, the same approach is needed for cluster ellipticals to clarify the origin of their color gradients.  It is noted that dust extinction in elliptical galaxies may also have some effects on the color gradients (Goudfrooij \\& de Jong 1995; Wise \\& Silva 1996; Silva \\& Wise 1996). However, about half of the detection towards ellipticals in far infrared with IRAS are around $3\\sigma$ threshold and confirmation is needed to be definitive (Bregman et al. 1998). In addition, spatial distribution of dust in an elliptical galaxy as well as dust mass which could affect a color gradient are not established yet. These are still open problems and will be examined in detail in our forthcoming papers.  Therefore, in this paper, we have chosen to focus on age and metallicity effects only.  This paper is organized as follows. The sample selection and data analysis of elliptical galaxies are described in \\S~2. Histograms of color gradients are presented in \\S~3 together with the representative color profiles of the sample ellipticals.  Discussion is given in \\S~4. The cosmological parameters adopted throughout this paper are the same as those in Paper I; $H_{0} = 50$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{0}=0.2$ and $\\Lambda = 0$. ",
        "conclusions": ""
    },
    "0004/astro-ph0004198_arXiv.txt": {
        "abstract": "The detection of a pulsed X-ray counterpart (RX J1824.2-2R52P) of the 3.05 ms pulsar PSR 1821-24, suggests the possibility of  a part of the rotational energy loss of this high spindown rate pulsar being in the optical band. Archival HST data for M28 is used here to set upper limits on the optical V-band magnitude of PSR 1821-24. The optical limit extends the multiwavelength observations for this source and  provides a constraint for theoretical models of pulsar emission. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004151_arXiv.txt": {
        "abstract": "Weak gravitational lensing observations probe the spectrum and evolution of density fluctuations and the cosmological parameters which govern them. The non-linear evolution of large scale structure produces a non-Gaussian signal which is potentially observable in galaxy shear data. We study the three-point statistics of the convergence, specifically the bispectrum, using the dark matter halo approach which describes the density field in terms of correlations between and within dark matter halos. Our approach allows us to study the effect of the mass distribution in observed fields, in particular the bias induced by the lack of rare massive halos (clusters) in observed fields.   We show the convergence skewness is primarily due to rare and massive dark matter halos with skewness converging to its mean value only if halos of mass $M > 10^{15} M_{\\sun}$ are present.  This calculational method can in principle be used to correct for such a bias as well as to search for more robust statistics related to the two and three point correlations. ",
        "introduction": "Weak gravitational lensing of faint galaxies probes the distribution of matter along the line of sight.  Lensing by large-scale structure (LSS) induces correlation in the galaxy ellipticities at the percent level (e.g., \\cite{Blaetal91} 1991; \\cite{Mir91} 1991; \\cite{Kai92} 1992).  Though challenging to measure, these correlations provide important cosmological information that is complementary to that supplied by the cosmic microwave background and potentially as precise (e.g., \\cite{JaiSel97} 1997; \\cite{Beretal97} 1997; \\cite{Kai98} 1998; \\cite{Schetal98} 1998; \\cite{HuTeg99} 1999; \\cite{Coo99} 1999; \\cite{Vanetal99} 1999; see \\cite{BarSch00} 2000 for a recent review). Indeed several recent studies have provided the first clear evidence for weak lensing in so-called blank fields (e.g., \\cite{Vanetal00} 2000; \\cite{Bacetal00} 2000; \\cite{Witetal00} 2000; \\cite{Kaietal00} 2000), though more work is clearly needed to understand even the statistical errors (e.g. \\cite{Cooetal00b} 2000b).  Given that weak gravitational lensing results from the projected mass distribution, the statistical properties of weak lensing convergence reflect those of the dark matter. Non-linearities in the mass distribution induce non-Gaussianity in the convergence distribution. With the growing observational and theoretical interest in weak gravitational lensing, statistics such as the skewness have been suggested as probes of cosmological parameters and the non-linear evolution of large scale structure (e.g., \\cite{Beretal97} 1997; \\cite{JaiSelWhi00} 2000; \\cite{Hui99} 1999; \\cite{MunJai99} 1999; \\cite{Vanetal99} 1999).  Here, we extend previous studies by considering the full convergence bispectrum, the Fourier space analog of three-point function. The bispectrum contains all the information present at the three point level, whereas conventional statistics, such as skewness, do not.  The calculation of the convergence bispectrum requires detailed knowledge of the dark matter density bispectrum, which can be obtain analytically through perturbation theory (e.g., \\cite{Beretal97} 1997) or numerically through simulations (e.g., \\cite{JaiSelWhi00} 2000; \\cite{WhiHu99} 1999). Perturbation theory, however, is not applicable at all scales of interest, while numerical simulations are limited by computational expense to a handful of realizations of cosmological models with modest dynamical range. Here, we use a new approach to obtain the density field bispectrum analytically by describing the underlying three point correlations as due to contributions from (and correlations between) individual dark matter halos.  Techniques for studying the dark matter density field through halo contributions have recently been developed (\\cite{Sel00} 2000; \\cite{MaFry00b} 2000b; \\cite{Scoetal00} 2000) and applied to two-point lensing statistics (\\cite{Cooetal00b} 2000b). The critical ingredients are: the Press-Schechter formalism (PS; \\cite{PreSch74} 1974) for the mass function; the NFW profile of \\cite{Navetal96} (1996), and the halo bias model of \\cite{Moetal97} (1997). The dark matter halo approach provides a physically motivated method to calculate the bispectrum.   By calibrating the halo parameters with N-body simulations, it can be made accurate across the scales of interest. Since lensing probes scales ranging from linear to deeply non-linear, this is an important advantage over perturbation-theory calculations.  Throughout this paper, we will take $\\Lambda$CDM as our fiducial cosmology with parameters $\\Omega_c=0.30$ for the CDM density, $\\Omega_b=0.05$ for the baryon density, $\\Omega_\\Lambda=0.65$ for the cosmological constant, $h=0.65$ for the dimensionless Hubble constant and a scale invariant spectrum of primordial fluctuations, normalized to  galaxy  cluster abundances ($\\sigma_8=0.9$ see \\cite{ViaLid99} 1999) and consistent with COBE (\\cite{BunWhi97} 1997). For the linear power spectrum, we take the fitting  formula for the transfer function given in \\cite{EisHu99} (1999).   In \\S \\ref{sec:density}, we review the dark matter halo approach to modeling the density field.  In \\S \\ref{sec:convergence} we apply the formalism to the convergence power spectrum, skewness, and bispectrum.  We summarize our results in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented an efficient method to calculate the non-Gaussian statistics of lensing convergence at the three point level based on a description of the underlying density field in terms of dark matter halos. The bispectrum contains all of the three point information, including the skewness. The prior attempts at calculating lensing bispectrum and skewness were limited by the accuracy of perturbative approximations and the dynamic range and sample variance of simulations.  Though the present technique provides a clear and an efficient method to calculate the statistics of the convergence field, it has its own shortcomings.  Halos are not all spherical, which can to some extent affect the configuration dependence in moments higher than the two point level. Substructures due to mergers of halos can also introduce scatter. Though such effects unlikely to dominate our calculations, further work using numerical simulations will be necessary to determine to what extent present method can be used as a precise tool to study the higher order statistics associated with weak gravitational lensing.  The dark matter halo approach also allows one to study possible selection effects that may be present in weak lensing observations due to the presence or absence of rare massive halos in the small fields that are observed. We have shown that the weak lensing skewness is mostly due to the most massive and rarest dark matter halos in the universe. The effect of such halos is stronger at the three point level than the two point level. The absence of massive halos, with masses greater than $10^{14}$ M$_{\\sun}$, leads to a strong decrease in skewness, suggesting that a straightforward use of measured skewness values as a test of cosmological models may not be appropriate unless prior observations are available on the distribution of masses in observed lensing fields.  One can correct for such biases using the halo approach, however. To implement such a correction in practice, further work will be needed to calibrate the technique precisely against simulations across a wide range of cosmologies. Efficient techniques to correct for mass biases both in the lensing power spectrum and bispectrum will be needed. Alternatively, this technique can be used to search for generalized three point statistics that are more robust to sampling issues. Given the great potential to study the dark matter distribution through weak lensing, this issues merit further study."
    },
    "0004/astro-ph0004367_arXiv.txt": {
        "abstract": "We compute theoretical predictions for surface brightness fluctuations (SBFs) of single-burst stellar populations (SSPs) using models optimized for this purpose.  We present results over a wide range of ages (from 1 to 17~Gyr) and metallicities (from 1/200 to 2.5~times solar) and for a comprehensive set of ground-based and space-based optical and infrared bandpasses. Our models agree well with existing SBF observations of Milky Way globular clusters and elliptical galaxies.  Our results provide refined theoretical calibrations and $k$-corrections that are needed to use SBFs as standard candles.  We suggest that SBF distance measurements can be improved by (1) using a filter around 1~\\micron\\ to minimize the influence of stellar population variations, and (2) using the integrated $V-K$ galaxy color instead of $V-\\Ic$ to calibrate $I$-band SBF distances.  We show that available SBF observations set useful constraints on current population synthesis models, and we suggest SBF-based tests for future models.  The existing SBF data favor particular choices of stellar evolutionary tracks and spectral libraries among the several choices allowed by comparisons based on only the integrated properties of galaxies.  Also, the tightness of the empirical $I$-band SBF calibration as a function of $V-I_c$ galaxy color is a useful constraint. It suggests that the model uncertainties in the lifetimes of the post-main sequence evolutionary phases are probably less than $\\pm$50\\% and that the initial mass function in elliptical galaxies is probably not much steeper than that in the solar neighborhood.  Finally, we analyze the potential of SBFs for probing unresolved stellar populations in elliptical galaxies. Since SBFs depend on the second moment of the stellar luminosity function, they are sensitive to the brightest giant stars and provide complementary information to commonly-used integrated light and spectra. In particular, we find that optical/near-infrared SBFs are much more sensitive to the metallicity than the age of a stellar population.  Therefore, in combination with age-sensitive observables, SBF magnitudes and colors are a valuable complement to metal-line indices to break the age/metallicity degeneracy in elliptical galaxy studies.  Our preliminary results suggest that the most luminous stellar populations of bright galaxies in nearby clusters have roughly solar metallicites and about a factor of three spread in age. ",
        "introduction": "When observing the inner regions of a nearby elliptical galaxy or the bulge of a nearby spiral galaxy, there are two noticeable characteristics of the surface brightness structure of these spheroidal stellar systems.  The first characteristic is that the galaxy is brightest in the center with the surface brightness falling off gradually with increasing radial distance.  The second characteristic is only apparent in good seeing conditions: on small scales, the galaxy has a clumpy appearance on the spatial scale of the seeing disk.  The clumpiness arises from Poisson statistical variations in the number of stars within each resolution element.  This effect can be easily recognized by visual inspection of images of nearby galaxies like M~31 and M~32, in which the small-scale clumpiness can be a few percent of the mean surface brightness.  Historically, this effect was called ``incipient resolution.''  In the modern context, it is known as surface brightness fluctuations (SBFs).  \\citet{1988AJ.....96..807T} devised a technique to quantify SBFs for use as an extragalactic distance indicator for undisturbed early-type galaxies. (See also the reviews by \\citealp{1992PASP..104..599J} and \\citealp{1999phcc.conf..181B}.) This method relies on using the ratio of the second moment to the first moment of the stellar luminosity function (LF) of the galaxy as a standard candle: \\begin{equation} \\Lbar \\equiv \\frac{\\sum_i n_i L_i^2}{\\sum_i n_i L_i} \\label{Lbar} \\end {equation} where $n_i$ is the number of stars of type $i$ and luminosity $L_i$. The quantity \\Lbar\\ has units of luminosity and is referred to as \\Mbar\\ when represented as an absolute magnitude.  The apparent SBF magnitude \\mbar\\ can be determined observationally, and if the distance to the galaxy is known, \\Mbar\\ can also be determined.  Of course, as in the case of ordinary photometry, SBF colors are distance-independent (provided that the $k$-correction is negligible).  Using SBFs as a distance indicator requires that (1) the bright end of the stellar LF in elliptical galaxies and spiral bulges is universal, or (2) variations in the LF from galaxy to galaxy can be measured and corrected so that \\Mbar\\ remains a standard candle.  SBFs are an intrinsic property of a stellar population as a whole. Therefore, in addition to their utility as a distance indicator, SBFs offer much promise in adding to our knowledge of the stellar content of elliptical galaxies.  In fact, the use of SBFs for stellar population studies arguably preceded its use as a distance indicator: \\citet{1955AJ.....60..247B} used the \"count-brightness ratio,\" the ratio of the number of resolved stars to the integrated light, to study the populations of M~31 and M~32.  Hence, the idea of using observations near or at the limit of resolution to explore stellar populations has a long history.  Furthermore, SBFs provide information about stellar populations {\\em unique} from ordinary integrated light, which is the first moment of the stellar LF.  Since SBFs also depend on the second moment of the stellar luminosity function, they are especially sensitive to the most luminous stars in elliptical galaxies, the evolved cool giant stars.  Thus, SBFs can put stronger constraints on the evolution of these stars than integrated light alone. Both the interior structure and emergent spectral energy distributions of cool giant stars are poorly understood, especially near the tip of the red giant branch (RGB) and the asymptotic giant branch (AGB) populated by low and intermediate mass ($\\lesssim5-7\\Msun$) stars.  Ideally, we would study these stars using Local Group star clusters, which comprise populations of homogenous composition and age. However, because cool giants evolve rapidly, only a handful are present in any cluster; therefore, small number statistics and stochastic fluctuations are undesireable factors \\citep[e.g.,][]{1997ApJ...479..764S}.  In this context, SBF analyses of entire galaxies can complement star cluster studies, since galaxian light arises from several orders of magnitude more stars. Also, while nearby globular clusters are mostly metal-poor, the dominant stellar populations in ellipticals are thought to be generally old and metal-rich. There is a dearth of such systems in the Local Group --- the best examples are the bulges of the Milky Way and M~31, but these may be imperfect analogs. In order to study metal-rich stellar evolution, one naturally turns to elliptical galaxies.  There are two basic motivations for modeling SBFs: (1) as pointed out by \\citet{1988AJ.....96..807T}, one can derive the calibration of SBF absolute magnitudes purely from models provided that stellar populations in galaxies can be modeled accurately, and alternatively, (2) one can use the observations of SBF magnitudes and colors to test and improve the models.  The first attempt at deriving a purely theoretical SBF zeropoint was that of \\citet{1990AJ....100.1416T}. They used stellar evolutionary models from the Revised Yale Isochrones \\citep[][hereinafter RYI]{gre87} supplemented with simple prescriptions for the horizontal branch and AGB; the resulting \\Ibar\\ zeropoint and especially its dependence on the integrated $V-\\Ic$ galaxy color disagreed significantly with observations \\citep{1991ApJ...373L...1T}.  This was due to the fact that the RYI giant branches failed to turn over in the optical at high metallicity, probably because of inaccurate bolometric corrections for the coolest giant stars \\citep{1992IAUS..149..181M, 1994ApJ...429..557A}.  A subsequent study was made by \\citet{1993ApJ...409..530W}, who computed SBF magnitudes using his own population synthesis models \\citep{1994ApJS...95..107W}.  For the main-sequence and RGB stars, these models used an amalgamation of isochrones from Vandenberg and collaborators with the RYI; post-RGB evolution was added using the fuel consumption theorem, including ``schematic'' treatments of the HB as a single red clump and of AGB evolution using a variety of theoretical prescriptions. The resulting SBF predictions agreed well with the observed optical SBF colors (\\Vbar, \\Rbar, and \\Ibar) and with the empirical calibration of \\Ibar\\ versus $V-\\Ic$ \\citep[see also][]{1997ApJ...475..399T}. \\citet{1993A&A...275..433B} also computed predictions for SBF magnitudes which were consistent with the existing optical SBF data at the time, though there were some uncertainties in transforming from the Johnson filters used in his models to the Kron-Cousins ones used for the observations. Since his models used older theoretical spectra \\citep{1978A&AS...34..229B}, Buzzoni had to extrapolate the spectra for wavelengths longward of 1.08~\\micron\\ and also for stars with $T_{eff}<4000$~K.  For both of these reasons, the \\citet{1993A&A...275..433B} models are expected to be less accurate for SBF predictions in the IR. For example, their predictions are at least several tenths of a magnitude fainter in the $K$-band than the observations.  The major observational effort on SBFs has been focused on $I$-band measurements for distance determinations.  As mentioned above, comparisons of stellar population models with observed SBF magnitudes and colors can also help us calibrate the colors (\\ie, stellar spectral energy distributions [SEDs]) and numbers (\\ie, evolutionary lifetimes) of cool luminous giant stars in old stellar populations. SBF stellar population studies have been less explored than distance measurements, partly because of the lack of suitable datasets.  Multicolor optical ($VRI$) SBF measurements for Virgo cluster galaxies from \\citet{1990AJ....100.1416T} were analyzed by both \\citet{1993ApJ...409..530W} and \\citet{1993A&A...275..433B}, who reached opposite conclusions on whether the optical SBF colors indicated that the galaxies contained a significant metal-poor component. \\citet{1994ApJ...429..557A} measured $I$-band SBFs for a sample of Galactic globular clusters, the only observations to date for these systems.  They sought to understand the empirical $I$-band SBF zeropoint and its correlation with integrated $V-\\Ic$ galaxy colors, as well as the conflict between the data and the RYI models; they also addressed the use of optical SBF colors for stellar population studies in galaxies and for disentangling age from metallicity effects in globular clusters. \\citet{1995AJ....110..179S, 1996AJ....111..208S} performed a detailed study of the optical SBF gradients within the nearby ellipticals M~32 and NGC~3379. Finally, \\citet{1998ApJ...505..111J} found reasonable agreement between their \\Kp-band SBF data for 11 nearby galaxies and the \\citet{1993ApJ...409..530W} models, with most of the galaxies lying around the [Fe/H]~=~--0.25 models with a spread in ages.  Now is a ripe opportunity to revisit the issue of stellar population modeling of surface brightness fluctuations.  There have been significant recent improvements in the stellar evolution calculations and spectral libraries used by population synthesis models. For example, the latest stellar evolution calculations include updated input radiative opacities \\citep[e.g.,][]{1992ApJ...397..717I}.  There has been even more progress on the observational front. The amount of $I$-band SBF data has increased by nearly tenfold since the early modeling of \\citet{1993ApJ...409..530W}, and new data in the near-infrared have extended the spectral range of SBF measurements \\citep{1994ApJ...433..567P, 1998ApJ...505..111J, liu2000, mei2000}.  In this paper, we present new models for optical/infrared SBFs of intermediate-age and old single-burst stellar populations (SSPs) and discuss their implications for SBF distance measurements and stellar population studies.  Though interesting issues remain to be addressed by blue/near-UV SBF measurements \\citep{1993ApJ...415L..91W}, we focus on the optical and near-infrared (\\ie, $V$-band to $K$-band) SBFs, since these constitute the bulk of past and ongoing observations. In \\S~2, we describe our models for computing SBFs. The models cover a wider range of ages and metallicities than in previous SBF studies. Furthermore, we have optimized the models for this work by refining the prescription for the luminous cool stars, which are important contributors to the SBF signal.  In \\S~3, we present the predictions of our models, including SBF magnitudes, integrated colors, and the fractional contribution of different stellar evolutionary phases to the SBFs.  We also derive theoretical calibrations and $k$-corrections that are needed to use SBFs standard candles.  In \\S~4, we compare our results with current SBF observations and discuss implications for the stellar content of elliptical galaxies.  In \\S~5, we review the uncertainties in our results and explore the potential of SBFs for breaking the age-metallicity degeneracy in studies of elliptical galaxies.  Finally in \\S~6, we summarize our findings and offer some future directions for SBF studies. ",
        "conclusions": "}  We have presented theoretical predictions for SBFs of single-burst stellar populations (SSPs) spanning a wide range of ages (from 1 to 17~Gyr) and metallicities (from 1/200 to 2.5~times solar).  Our calculations are based on the population synthesis models of Bruzual \\& Charlot (2000), in which the stellar evolution prescription and spectral libraries are improved over the models used in previous SBF studies.  In particular, our models have been optimized during the course of this work by refining the prescription for the latest phases of stellar evolution, which are important contributors to the optical and infrared SBF signal. Our standard predictions are based on multi-metallicity evolutionary tracks from the Padova school and semi-empirical stellar spectra designed to match the observed color-temperature relations of solar-neighborhood stars at solar metallicity (LCB97).  Using our models, we generate several {\\bf basic predictions as a function of age and metallicity}.  \\begin {enumerate}  \\item We compute SBF magnitudes and integrated colors for a large set of ground-based and space-based (\\HST) optical and infrared bandpasses. These are supplemented with the strengths of several optical absorption-line indices on the Lick/IDS system.  \\item We provide results for solar-metallicity models using several combinations of stellar evolutionary tracks and spectral libraries. These can be used to assess the systematic effects of different model inputs on the results.  \\item We predict the fractional contribution of different stellar evolutionary phases to the SBFs. Since this information cannot be easily derived from observations, the models provide insight into which phases are important contributors to the SBF signal for a given bandpass.  \\end{enumerate}   \\noindent Our model results directly {\\bf benefit SBF distance determinations}, specifically:   \\begin{enumerate} \\setcounter{enumi}{3}  \\item We use the models to determine purely theoretical calibrations for SBFs in many bandpasses. These are independent of any systematic errors in Cepheid distances or reddening corrections, which affect only empirical calibrations.  \\item We tabulate $k$-corrections out to $cz\\leq15,000$~\\kms\\ ($z=0.04$), which are required for accurate determinations of \\Ho.  We find that the $k$-corrections are roughly linear in this redshift range. Metallicity has a stronger effect on $k$-corrections in the near-infrared than in optical, but the amplitudes of the corrections are also generally smaller in the near-infrared. We conclude that systematic errors from uncertainties in the $k$-corrections are not important sources of error for $H_0$ determinations.  \\item We suggest that the scatter in $I$-band SBF distances can be further reduced by using the integrated $V-K$ galaxy color instead of $V-\\Ic$ to correct for stellar population variations between galaxies. The reason for this improvement is that the $V-K$ color is more sensitive to metallicity, which also drives the \\Ibar\\ signal.  \\item Our models predict that the fluctuation magnitudes should be independent of population age and metallicity around 1~\\micron. A similar conclusion was reached by \\citet{1993ApJ...409..530W} using very different population synthesis models --- this suggests that the prediction is robust. Therefore, observations taken with a $Z$-band filter from a large (8--10~m) ground-based telescope or with the $F110W$ filter in \\HST\\ NICMOS should allow SBF distance measurements which are more robust against galaxian population variations.  \\end{enumerate}   \\noindent We have compared our model results with nearly all the SBF measurements available to date.  Since SBFs are especially sensitive to the cool, luminous stars on the upper RGB and AGB, they provide important tests for population synthesis models.  The existing dataset comprises Galactic globular clusters, M~31, M~32, and early-type galaxies in nearby clusters. The $I$-band dataset is by far the most extensive; there are some $K$-band SBF magnitudes and optical/IR SBF colors, but more measurements are needed for further testing.  We find generally {\\bf good agreement between models and data} and also suggest some {\\bf new tests for the models}. Specifically:   \\begin{enumerate} \\setcounter{enumi}{7}  \\item Our models reproduce \\Vbar\\ and \\Ibar\\ observations of Galactic globular clusters. This test is complementary to those based on galaxy data, since the globular clusters have much lower metallicities.  Models with [Fe/H]~$\\approx$~--0.7 might be $\\approx$0.2~mag too red in \\Vbar--\\Ibar, although more data are needed to verify this.  \\item Our standard models provide the best agreement to date with the tight empirical calibration of \\Ibar\\ over the entire observed range of $V-\\Ic$ galaxy color.  The zeropoint and slope of the calibration predicted by our models agree remarkably well with those derived from the data. Moreover, the models indicate a saturation of \\Ibar\\ for $V-\\Ic\\lesssim1.0$, which is also seen in the observations. The reason for this flattening is most likely the constancy of the $I$-band tip of the RGB for metal-poor stellar populations. The small scatter in the empirical calibration as a function of $V-\\Ic$ galaxy color is also reproduced by the models; this arises because of the partial age/metallicity degeneracy in the \\hbox{\\{\\Ibar, $V-\\Ic$\\}} parameter space. This degeneracy is a boon for distance measurements.  \\item Our standard models also agree with \\Kpbar\\ observations, although this is based on a much smaller sample of galaxies. In the \\hbox{\\{\\Kpbar, $V-\\Ic$\\}} parameter space, changes in age and metallicity are roughly orthogonal.  \\Kpbar\\ brightens for populations of higher metallicities and younger ages, as expected from observations of the RGB and AGB of star clusters in the Galaxy and the Magellanic Clouds.  \\item The optical/IR fluctuation colors predicted by our models agree with the observations, although some discrepancies exist at the $\\approx0.1-0.2$~mag level. An advantage of testing the models against measurements of SBF colors is that the data are immune to errors in the galaxy distances.  \\item The semi-empirical SEDs of LCB97 provide better agreement with SBF observations at all metallicities than their theoretical SEDs. Observations of \\Ibar--\\Jbar\\ and \\Ibar--\\Hbar\\ colors would help to verify this result.  For solar metallicity, the results obtained from the empirical spectral library of \\citet{1998PASP..110..863P} agree closely with those from the semi-empirical library of LCB97.  \\item For solar metallicity, the Padova evolutionary tracks seem to provide better agreement with SBF observations than the Geneva tracks. The integrated spectral properties from the two sets of tracks are very comparable at solar metallicities. However, the SBF data are a sensitive test for deciding between the Geneva and Padova tracks, since the differences are larger in the SBF predictions than in those for the integrated spectra.  \\item From the tightness of the empirical \\Ibar\\ calibration, we conclude that the lifetimes of post-main sequence phases (RGB, core-He burning, and AGB) in the evolutionary tracks are probably accurate to within better than $\\pm$50\\%.  \\end{enumerate}  \\noindent By comparing our single-burst models with the available dataset, mostly composed of luminous galaxies in nearby clusters, our preliminary findings on {\\bf the stellar populations dominating the SBFs} are:   \\begin{enumerate} \\setcounter{enumi}{14}  \\item The metallicities inferred from SBF magnitudes and SBF colors show little spread.  The metallicities favored by the optical/SBF colors are slightly sub-solar, while those favored by the \\Kpbar\\ data are around solar.  \\item SBF color measurements show no obvious differences for galaxies observed with different linear aperture sizes, though the available dataset is small.  The implication is that SBF distance measurements should be relatively insensitive to systematic errors due to aperture effects.  \\item The ages inferred from comparisons of both \\Kpbar\\ and \\Vbar--\\Ibar\\ with the $V-\\Ic$ integrated color span a range of about a factor of three, with the youngest ones near 3~Gyr. Note that estimates based on combinations of SBF colors and integrated colors are independent of the galaxy distances.  \\item For old populations, the tightness of the empirical $I$-band SBF calibration also indicates that the IMF in elliptical galaxies cannot be significantly steeper than that in the solar neighborhood.  \\end{enumerate}   \\noindent Finally, we suggest that SBF measurements can offer {\\bf useful new tools for stellar population studies}:   \\begin{enumerate} \\setcounter{enumi}{18}  \\item In old populations, the SBF magnitudes and colors are predicted to be very sensitive to metallicity, especially at near-IR ($JHK$) wavelengths. This may offer a potent means of breaking the age/metallicity degeneracy inherent in studies based on integrated spectral properties.  \\item We find that the \\Ibar--\\Kbar\\ SBF color is very sensitive to metallicity because of the decreasing temperature of the giant branch with increasing metallicity. Thus, \\Ibar--\\Kbar\\ might be used in combination with age-sensitive observables such as Balmer absorption lines to constrain the ages and metallicities of elliptical galaxies. SBF colors may also present advantages over metal absorption lines such as Mg$_2$ and C4668, which are affected by uncertainties in the patterns of $\\alpha$-element enhancement in elliptical galaxies.  \\item Our models suggest that the $L$-band and $M$-band SBFs are very sensitive to age, although our predictions are not optimized in this wavelength range. This potentially interesting result should be further investigated using more appropriate models.  \\item Observations of \\Bbar--\\Ibar\\ with \\Ibar--\\Kbar\\ may be useful to identify stellar populations of different metallicities in elliptical galaxies. \\end{enumerate}   The single-burst models we have investigated can account for the full observed ranges of SBF magnitudes, SBF colors, and integrated colors for bright elliptical galaxies in nearby clusters.  It is important to realize that, although the SBF observations can be most simply reproduced by models with around solar metallicity and a significant spread in age, a more refined analysis is required to interpret these measurements in terms of the star formation history of elliptical galaxies.  In particular, there are multiple lines of evidence that both cluster and field elliptical galaxies have experienced more than one episode of star formation \\citep[e.g.,][]{1992AJ....104.1039S, 1996MNRAS.279....1B, 1999ApJ...518..576P}.  To constrain the ages and metallicities of different stellar generations in elliptical galaxies, we then require a combination of various age and metallicity indicators. Unfortunately, the published SBF and absorption-line studies contain few galaxies in common. In a future paper (M. Liu \\etal, in preparation), we exploit a more extensive set of new SBF measurements to investigate the stellar content of elliptical galaxies."
    },
    "0004/astro-ph0004401_arXiv.txt": {
        "abstract": "Spectra taken with the {\\it Hubble Space Telescope} (\\hst) reveal that NGC 4450 emits Balmer emission lines with displaced double peaks and extremely high-velocity wings.  This characteristic line profile, previously seen in a few nearby LINERs and in a small fraction of broad-line radio galaxies, can be interpreted as a kinematic signature of a relativistic accretion disk.  We can reproduce the observed profile with a model for a disk with a radial range of 1000--2000 gravitational radii and inclined by 27 degrees along the line of sight. The small-aperture \\hst\\ data also allow us to detect, for the first time, the featureless continuum at optical wavelengths in NGC 4450; the nonstellar nucleus is intrinsically very faint, with $M_B$ = --11.2 mag for $D$ = 16.8 Mpc.  We have examined the multiwavelength properties of NGC 4450 collectively with those of other low-luminosity active nuclei which possess double-peaked broad lines and find a number of common features.  These objects are all classified spectroscopically as ``type 1'' LINERs or closely related objects.  The nuclear luminosities are low, both in absolute terms and relative to the Eddington rates.  All of them have compact radio cores, whose strength relative to the optical nuclear emission places them in the league of radio-loud active nuclei.  The broad-band spectral energy distributions of these sources are most notable for their deficit of ultraviolet emission compared to those observed in luminous Seyfert~1 nuclei and quasars.  The double-peaked broad-line radio galaxies Arp 102B and Pictor~A have very similar attributes.  We discuss how these characteristics can be understood in the context of advection-dominated accretion onto massive black holes. ",
        "introduction": "A minority of active galactic nuclei (AGNs) exhibit double-peaked broad emission lines --- permitted lines whose profile shows two displaced maxima, offset from the line center by several thousand \\kms.  According to Eracleous \\& Halpern (1994), approximately 10\\% of broad-line radio galaxies show double-peaked lines.  A variety of mechanisms have been proposed to explain this unique kinematic signature, including relativistic motions in an accretion disk (Chen, Halpern, \\& Filippenko 1989; Chen \\& Halpern 1989), two separate broad-line regions due to a binary black hole (Gaskell 1983), biconical outflow (Zheng, Binette, \\& Sulentic 1990), and anisotropic illumination of the broad-line region (Goad \\& Wanders 1996). A combination of basic physical arguments and recent observational results, however, has shown that the most plausible origin of double-peaked emission lines is the accretion disk around the central black hole. Alternative interpretations are slowly being ruled out (Eracleous 1999).  A notable characteristic of double-peaked broad-line AGNs is that their distinctive line profile can be transient or highly variable.  In recent years, double-peaked broad emission lines have been found in several galaxies that previously had none.  These are NGC 1097 (Storchi-Bergmann, Baldwin, \\& Wilson 1993), Pictor~A (Halpern \\& Eracleous 1994; Sulentic et al. 1995), and M81 (Bower et al. 1996).  That so many cases turn up serendipitously among nearby galaxies suggests that this is not an uncommon phenomenon.  An intriguing, possibly important, connection may exist with low-ionization nuclear emission-line regions (LINERs; Heckman 1980), a class of emission-line nuclei often found in nearby galaxies, but one whose nature remains controversial (Filippenko 1996; Ho 1999a).  The transient emission sources listed above all qualify spectroscopically as LINERs according to the low-ionization state of their optical spectra, as do the majority of the broad-line radio galaxies that contain double-peaked lines (e.g., Eracleous \\& Halpern 1994).  This paper reports the discovery of double-peaked broad emission lines in yet another LINER galaxy, NGC 4450, based on observations made with the {\\it Hubble Space Telescope (HST)}.  NGC 4450 is a bright ($B_T$ = 10.9 mag) Sab galaxy located in the Virgo cluster.  The nucleus is spectroscopically classified by Ho, Filippenko, \\& Sargent (1997a) as a LINER of ``type 1.9'' based on detection of weak broad H\\al\\ emission. ",
        "conclusions": "\\subsection{Evidence for an Accretion Disk in NGC 4450}  The double-shouldered H\\al\\ line in NGC 4450 strongly resembles the double-peaked profiles observed in some broad-line radio galaxies.   In those objects the distinctive line profile has often been modeled in the context of a relativistic accretion disk (e.g., Chen et al. 1989; Eracleous \\& Halpern 1994; Halpern et al. 1996).  In this interpretation, the tremendous breadth of the line reflects the high rotation speed of the gas; the blue peak has greater intensity than the red peak because of relativistic beaming; and the overall asymmetry of the profile arises from the combined effects of transverse and gravitational redshift.  In NGC 4450, the observed profile of H\\bet\\ differs slightly from that of H\\al, but its signal-to-noise ratio is considerably lower because of its reduced strength, and the difference between H\\al\\ and H\\bet\\ is not inconsistent with that seen in broad-line radio galaxies.  Figure 3 illustrates explicitly the viability of the disk interpretation for NGC 4450.  After subtracting the continuum, the H$\\alpha$ profile was fitted with a model according to which the line originates in a relativistic Keplerian disk (Chen et al. 1989; Chen \\& Halpern 1989). In this model the line photons are emitted from a thin, photoionized layer at the surface of a circular accretion disk, which is illuminated by a source of ionizing radiation located at the center of the disk. The calculation of the line profile includes, in addition to Doppler broadening, the effects of transverse and gravitational redshift and light bending. The free parameters of the model are the inner and outer radii of the line-emitting portion of the disk, $\\xi_1$ and $\\xi_2$ (expressed in units of the gravitational radius, $r_{\\rm g}\\equiv GM/c^2$, where $M$ is the mass of the black hole), the inclination angle of the disk axis to the line of sight, $i$, and the broadening parameter, $\\sigma$, which represents the velocity dispersion in a parcel of gas in the disk, presumably due to turbulence. The emissivity of the disk is assumed to vary with radius as $\\epsilon\\propto\\xi^{-q}$, with $q$ set to 3, following the results of photoionization calculations by Dumont \\& Collin-Souffrin (1990a,b,c). The model parameters that yield the best fit are $\\xi_1=1010^{+610}_{-440}$, $\\xi_2=2030^{+1200}_{-990}$, $i=27^{+8}_{-7}$~degrees, and $\\sigma=1000^{+300}_{-200}~{\\rm km~s}^{-1}$. These parameters are comparable to those required to fit the profiles of double-peaked emission lines found in broad-line radio galaxies (Eracleous \\& Halpern 1994) and in the LINER NGC 1097 (Eracleous et al. 1995; Storchi-Bergmann et al. 1995, 1997).  The bottom panel of Figure 3 shows the residual spectrum obtained by subtracting the disk model from the observed spectrum.  Note that the residual spectrum looks very similar to the starlight-subtracted ground-based spectrum shown in Figure 1{\\it b}.  NGC 4203, the other double-peaked broad-line object discovered in our survey (Shields et al. 2000), is qualitatively very similar to NGC 4450 in its line profile.  Notably, the line profiles of both of these objects differ from those of NGC 1097 and M81, the two spirals that previously were found to emit transient double-peaked lines.  At the time of discovery by Storchi-Bergmann et al. (1993), the H\\al\\ line in NGC 1097 displayed a red peak stronger than the blue peak, and the blue shoulder was the more extended of the two.  The sense of the asymmetries, however, varied during subsequent monitoring, with the overall pattern interpretable in the context of a model of a precessing elliptical disk (Eracleous et al. 1995; Storchi-Bergmann et al. 1995, 1997).  When M81 was observed by Bower et al. (1996), it, too, had the initial appearance of NGC 1097.  \\begin{figure} \\psfig{file=fig3_new.ps,width=6.5truein,angle=270} \\caption{ The top panel shows an expanded view of the spectrum near the broad H\\al\\ line; the continuum has been removed.  The disk model described in the text is overplotted as a continuous curve.  The bottom panel shows the residual spectrum obtained by subtracting the disk model from the observed spectrum. } \\end{figure}  \\vbox{ \\hbox{ \\hskip -0.5truein \\psfig{file=table1_v7.ps,width=7.5truein} } }   \\subsection{The Incidence and Detectability of Double-Peaked Broad Lines}  Double-peaked broad lines have recently been found serendipitously in three cases, namely NGC 1097, M81, and Pictor~A (\\S\\ 1).  Our STIS study adds two new members to the roster, NGC 4450 and NGC 4203. Detection of such emission in low-luminosity sources is complicated by its weak contrast with the underlying stellar continuum.  In the case of the first three objects, detection was made possible by a transient brightening of the broad feature.  For the last two sources, we do not know whether variability contributed to their detection, but the use of the STIS aperture probably did.  Indeed, the detection of two such objects out of a sample of only 19 weak AGNs (Rix et al. 2000) is perhaps an indication of both the prevalence of such emission and the difficulty of discerning it from the ground, a point reinforced by the scarcity of such objects in ground-based surveys.  Disk-like line profiles are found preferentially in LINERs.  All the previously known transient cases, along with the two found in our survey, are spectroscopically classified as LINERs (see Table 1), and the detection rate among LINERs in the STIS survey is 25\\% (2/8).  The association with {\\it type 1} LINERs --- those that have detectable emission from a broad-line region --- is especially striking.  With the exception of NGC 1097, all of the objects in Table 1 were known to be LINER~1s prior to the discovery of the double-peaked component, and two out of the four LINER~1s in the STIS sample turned out have disk-like profiles.  Consistent with this trend, Eracleous \\& Halpern (1994) find that broad-line radio galaxies that show disk-like emission preferentially emit a stronger low-ionization, LINER-like spectrum.  \\subsection{Spectral Properties and Nature of the Accretion Flow}  Despite the limited number of known nearby low-luminosity AGNs with double-peaked broad emission lines, it is instructive to examine their spectral properties as a group.  Table 1 summarizes multiwavelength data available for NGC 1097, M81, NGC 4203, and NGC 4450, and, for comparison, the radio galaxies Pictor~A and Arp 102B.  The spectral energy distributions (SEDs) of these objects will be presented in greater detail by Ho et al. (2000a).  It is important to stress that the SED measurements pertain to the galaxy {\\it nucleus}, which at virtually all wavelengths is much fainter than the host galaxy itself.  We draw attention to a set of common spectral properties shared among the first four objects, all of which have nuclei with very low luminosities.  The absolute magnitudes of the nuclei range from $M_B^{\\rm nuc}$  = --11.3 to --12.6.  The most salient feature of the SEDs is their conspicuous deficit of UV emission, which manifests itself as continuum slopes in the optical--UV ($\\alpha_{\\rm ou}$ = 1.4--2.5; see Table 1 for specific definitions) and in the UV--X-ray ($\\alpha_{\\rm ox}$ = 0.9--1.1) regions that are steeper and flatter, respectively, than is found in luminous AGNs ($\\alpha_{\\rm ou}\\, \\approx$ 0.2--0.5, Malkan 1988; $\\alpha_{\\rm ox}\\,\\approx$ 1.2--1.6, Mushotzky \\& Wandel 1989).  Following the arguments of Ho (1999b), we believe that the UV deficit is intrinsic to the source and not an artifact of dust extinction.  This suggests that the ``big blue bump'' (Shields 1978; Malkan \\& Sargent 1982), and by implication the optically thick, geometrically thin accretion disk thought to give rise to this feature, is weak or absent.  The strength of the radio band is also notable.  All four objects, among them three spirals, contain compact, flat-spectrum radio cores.  The absolute radio power of the nuclei is tiny by traditional standards, with $P_{\\rm 6~cm}\\, \\approx\\,10^{20}$ W Hz$^{-1}$, but their optical output is correspondingly low, such that the normalized radio power remains relatively high.  The ratio $\\Re$, defined as $f_{\\nu}({\\rm 6~cm})/f_{\\nu}(B)$, spans $\\sim$20--100, to be compared with $\\Re\\,>$ 10 for radio-loud quasars (Kellermann \\etal 1989; Visnovsky \\etal 1992).  By this conventional criterion, therefore, all four sources formally qualify as radio-loud, a finding that refutes the popular prejudice that radio-loud AGNs are confined to luminous early-type hosts.  We note, however, that none of these low-luminosity objects possess well-collimated, large-scale radio jets.  The accretion flow in the systems highlighted here may be systematically different from the standard Shakura \\& Sunyaev (1973) disk normally assumed to exist in luminous AGNs.  Recently there has been increasing theoretical and observational evidence that black holes fed at a low rate accrete through an optically thin, advective mode (see review by Narayan, Mahadevan, \\& Quataert 1998).  ``Advection-dominated'' accretion flows (ADAFs) have several characteristics that make them an attractive framework in which to interpret our observations.  First, ADAFs are typified by a low radiative efficiency ($L\\,\\propto\\,\\dot{m}^2$ instead of $\\dot{m}$), and so they provide a natural explanation for the very low accretion luminosities actually measured.  The nuclei of the four nearby sources in Table 1, for instance, have bolometric luminosities of $L_{\\rm bol}\\,\\approx\\, 10^{41}-10^{42}$ \\lum. Second, the ADAF phase can only exist in sub-Eddington systems, when $L_{\\rm bol}/L_{\\rm Edd}$ \\lax $10^{-2}$--10$^{-1}$ (Narayan et al. 1998); this condition is fulfilled by our sources, which have $L_{\\rm bol}/L_{\\rm Edd},\\approx$ 1\\e{-4} -- 3\\e{-5} (Table 1).  Third, ADAFs do not emit a thermal UV bump, in qualitative agreement with the observed SEDs.  Fourth, thermal synchrotron radio emission contributes significantly to the ADAF spectrum, and consequently we expect the SED to be prominent in the radio.  Finally, we note that the characteristically harder ionizing spectrum of an ADAF lowers the effective ionization parameter and hence favors the production of a LINER-like emission-line spectrum (Halpern \\& Steiner 1983; Ferland \\& Netzer 1983).   The picture we are advocating for low-luminosity AGNs with double-peaked broad lines is likely to be more widely applicable to low-luminosity AGNs as a class. The characteristics of the SEDs discussed above are not restricted to the specific sample emphasized in this study; they were noted by Ho (1999b) in a number of other low-luminosity systems, among them M81.  Quataert et al. (1999) have successfully applied ADAF models to fit the SEDs of M81 and NGC 4579, the latter a LINER not known to possess disk-like broad lines but otherwise very similar to M81.  The basic elements of the above scenario may also be appropriate in general for radio galaxies with disk-like broad emission lines.  We illustrate this connection in Table 1 by comparing the first four nearby, low-luminosity systems with two well-studied, double-peaked broad-line radio galaxies, Arp 102B and Pictor~A.  Although both objects are substantially more luminous than the rest ($M_B^{\\rm nuc}$ = --17.9 mag for Arp 102B; $M_B^{\\rm nuc}$ = --18.3 mag for Pictor~A), they still rank among the least luminous classical Seyferts known.  More germane to the present discussion is the Eddington ratio, which is $\\sim$1\\e{-3} for Arp 102B and 2\\e{-2} for Pictor~A, both within the threshold in which the ADAF framework holds. In their study of Arp 102B, Chen et al. (1989) already suggested that the structure of its accretion flow may take the form of an ``ion-supported torus'' (Rees et al. 1982), the physical concept of which is very similar to that of an ADAF.  Eracleous \\& Halpern (1994) subsequently extended this idea to interpret a number of statistical properties observed in disk-like emitters.  However, one of the key pieces of evidence --- the absence of the UV bump --- has, until very recently, remained elusive.  Halpern \\etal (1996) found that the \\hst\\ spectrum of Arp 102B indeed shows weak UV emission.  As shown in Table 1, its continuum has a steep optical--UV slope ($\\alpha_{\\rm ou}$ = 2.04) and a relatively hard UV to X-ray slope ($\\alpha_{\\rm ox}$ = 1.08).  Pictor~A, too, has an exceptionally flat UV--X-ray slope ($\\alpha_{\\rm ox}$ = 0.88), although its optical--UV slope ($\\alpha_{\\rm ou}$ = 0.60) does not differ markedly from those typically seen in luminous AGNs.  This may reflect the fact that its Eddington ratio ($\\sim$2\\e{-2}), the largest in the sample, is only marginally within the ADAF limit."
    },
    "0004/astro-ph0004292_arXiv.txt": {
        "abstract": "{ROSEBUD (Rare Objects SEarch with Bolometers UndergrounD) is an experiment which attempts to detect low mass Weak Interacting Massive Particles (WIMPs) through their elastic scattering off Al and O nuclei. It consists of three small sapphire bolometers (of a total mass of 100 g) with NTD-Ge sensors in a dilution refrigerator operating at 20 mK in the Canfranc Underground Laboratory. We report in this paper the results of several runs (of about 10 days each) with successively improved energy thresholds, and the progressive background reduction obtained by improvement of the radiopurity of the components and subsequent modifications in the experimental assembly, including the addition of old lead shields. Mid-term plans and perspectives of the experiment are also presented.} \\\\ \\\\ PACS: 95.35.+d; 07.57.Kp; 07.62.+s \\\\ {\\it{Keywords:}} Dark matter; WIMPs; Underground detectors; Bolometers ",
        "introduction": "Among the particle candidates to the non-baryonic dark matter of the universe, the neutral, weak interacting massive particles (WIMPs)---supposedly forming a significant part of the galactic haloes---enjoy a prominent position. A distinguished WIMP is the neutralino, the lightest stable particle of the minimal supersymmetric extension of the Standard Model \\cite{Gri}. Accelerator results settle a lower bound of 20--30 GeV to the neutralino mass in most of the models, but in unconstrained SUSY models masses as low as 2 GeV could be allowed \\cite{Gab96}. WIMPs of m $\\sim 1-1000$ GeV with non-relativistic velocity, can interact with the nuclei of a detector target producing a nuclear recoil of a few keV, a fraction of which is visible in the detector (depending on the nuclear target, the detector and the mechanism of energy deposition). Because of the low interaction rate and small energy deposition the WIMP detectors require very low backgrounds and energy thresholds \\cite{Mor3}. A further contribution to the direct detection rate---in the very low energy region $\\leq 1$ keV--- could be provided by a non-thermal, low velocity WIMP population recently proposed \\cite{Dam98} and discussed \\cite{Bot00}.  Thermal detectors \\cite{Cryo}, which measure the temperature increase produced by the WIMP interaction in an absorber crystal, use more efficiently the energy deposition of WIMPs than the conventional ionization detectors because most of the nuclear recoil energy in WIMP scattering goes to heat. They are true low energy detectors, where the visible energy is practically the whole recoil energy (quenching factor close to one). Moreover, the mechanisms and quanta involved in the physics of the detection imply that they should have better energy threshold and energy resolution than the conventional ionization detectors. In particular, the energy resolution on nuclear recoils achieved by the sapphire bolometers of the experiment reported in this paper is significantly better than that obtained with HPGe diodes, like COSME \\cite{Jmor92} (which features a full width at half maximum (FWHM) energy resolution $\\Gamma (10 \\rm keV)=0.4 \\rm keV$). In much the same way the energy threshold of the sapphire bolometers of this work is also lower than that obtained in Ge-diodes looking for dark matter \\cite{Mor3,Jmor92,Reu91,Mor00}. Such low energy thresholds, hopefully achieved by thermal detectors, make them valuable tools in the search for phenomena \\cite{Dam98} leaving very small energy in the absorbers. So the cryogenic detectors are very well suited to explore either low mass WIMPs or low velocity WIMPs \\cite{Dam98}. They offer moreover the possibility of using different targets to tune up the sensitivity for different WIMP masses. Finally, they can discriminate the background (electron recoils) from nuclear recoils by collecting not only the phonon energy but also the charge (or light) produced by the ionizing component of the deposited energy (hybrid detection) \\cite{Cha00,Gai00}. The resulting background reduction permits to increase significantly the sensitivity of WIMP searches allowing to explore regions below $\\sigma^{p} \\sim 10^{-8}$ nbarn for medium mass ($\\sim 50$ GeV) WIMPs---where the neutralino configuration \\cite{Bot98} relevant to interpret an annual modulation effect \\cite{Ber99} is located. Although the radioactivity content of bolometers and their environment is still higher than that of HPGe, significant progresses are being accomplished.  A WIMP search with small sapphire bolometers, ROSEBUD (Rare Objects Search with Bolometers Underground), is being carried out in the Canfranc Underground Laboratory (at 2450 m.w.e.) \\cite{rosebud1,rosebud2,Ceb00} with the purpose of exploring the WIMPs scattering off Al and O nuclei. It features three small sapphire bolometers (with Ge thermistors) placed in a small, mobile dilution refrigerator in the ultralow radioactivity environment of the Canfranc underground site. ",
        "conclusions": ""
    },
    "0004/astro-ph0004200_arXiv.txt": {
        "abstract": "Inverse Problem techniques offer powerful tools which deal naturally with marginal data and asymmetric or strongly smoothing kernels, in cases where parameter-fitting methods may be used only with some caution.  Although they are typically subject to some bias, they can invert data without requiring one to assume a particular model for the source.  The Backus-Gilbert method in particular concentrates on the tradeoff between resolution and stability, and allows one to select an optimal compromise between them.  We use these tools to analyse the problem of reconstructing features of the source star in a microlensing event, show that it should be possible to obtain useful information about the star with reasonably obtainable data, and note that the quality of the reconstruction is more sensitive to the number of data points than to the quality of individual ones. ",
        "introduction": "Where once all the interest in microlensing was in the details of the lensing population, there is now an increasing interest in lensing events as probes of the stellar sources.  From this point of view, once the lens' geometrical details have been worked out, the event can be used as a `super-telescope', providing otherwise completely unattainable resolution of the surfaces of distant stellar disks.  Initially, analyses assumed that the microlensing source star could be taken to be a point source, and the first discussion of `finite source effects' was in the context of a problem -- Witt \\& Mao (1994) asked at what point the point-source approximation would break down; {Nemiroff} \\& {Wickramasinghe} (1994) put the question more positively, asking what information could be obtained from the distortions to the light curve which finite-source effects would cause.  It is not merely intensity information which can be obtained from events.  Simmons, Newsam, \\& Willis (1995) discuss the information which can be extracted when polarization measurements are made of microlensing events, and in (Newsam {et~al.} 1998) show how even relatively poor polarization data can substantially improve fits of source parameters.  Although the basic microlensing effect is achromatic, the fact that stars have different limb-darkening profiles in different colours means that a lens differentially amplifying the disk will produce a chromatic effect.  Other workers ({Valls-Gabaud} 1995; Sasselov 1996; Valls-Gabaud 1998) have discussed how one might obtain such chromaticity information.  It is even possible to discuss how one might observe the signatures of stellar spots (Heyrovsk{\\'y} \\& Sasselov 1999; Bryce \\& Hendry 2000).  The usual way in which source structure is detected is by applying a model-fitting (equivalently, parameter-fitting) algorithm to the observed data, to obtain the best-fit parameters of a suitable limb-darkening model; this is the approach used, for example, by the MACHO collaboration ({Alcock} {et~al.} 1997) and the PLANET collaboration ({Albrow} {et~al.} 1999) to make the first detections of limb-darkening in microlensing events.  It is also the approach which underlies the insightful error analysis by {Gaudi} \\& {Gould} (1999).  A parameter-fitting algorithm essentially consists of a mechanism for systematically moving through parameter space, repeatedly solving the `forward problem' -- calculating the data to be predicted from a given limb-darkening profile -- until the predicted data is optimally close to the data actually observed.  Here we want to suggest that, because of the fact that the underlying source function is convolved through a broad and asymmetric amplification kernel, a model-fitting approach is potentially problematic, and that this recovery problem is more naturally addressed using the well-established technology of inverse problems.  We plan to discuss the merits of inverse problem techniques in general, and the Backus-Gilbert method in particular, and exemplify the possibilities by inverting simulated microlensing data to recover limb-darkening and limb-polarization effects.  We will see that we are able to discuss explicitly and robustly the tradeoff between resolution and stability which is implicit in any such inversion, including recoveries obtained by model-fitting. ",
        "conclusions": "One of the aims of this paper is to emphasise the seriousness of the ill-conditioning of the source reconstruction problem, and hence the desirability of using an analytical technique which starts by examining that ill-conditioning, goes on to discuss what information is nonetheless recoverable, and only then produces the numerical information which is the point of the exercise.  Of course, the same questions can be asked using a parameter-fit approach, but less naturally, since such approaches assume, in a sense, that the information is recoverable, with the result that problems can only be uncovered \\textit{post hoc}, by an intelligent examination of goodness-of-fit measures, or by tracing the propagation of errors through a calculation.  The second aim is to use this inverse problem approach to analyse the kernel which turns the underlying limb-darkening and limb-polarization functions into microlensing intensity and polarization data.  It turns out that the information is indeed recoverable with adequate uncertainties but, depending on the quality of the data available, one many have to make significant compromises over the resolution one is prepared to accept.  It is also possible to use this approach to analyse the effect of modifications in the way the data is collected, and discover that this can indeed significantly improve the data's effective quality.   \\let\\JApJ\\apj \\let\\JMNRAS\\mnras \\let\\JPASJ\\pasj \\let\\JAAS\\aaps"
    },
    "0004/astro-ph0004036_arXiv.txt": {
        "abstract": "We detect stars from the Hipparcos and Tycho Catalogues in the field of view of observations  with the ROSAT HRI of three globular clusters. We use the positions of these stars to reduce the systematic error in the positions of X-ray sources in the clusters to $\\sim2''$ for \\omcen\\ and NGC\\,6752, and $1''$ for NGC\\,6397. We detect three X-ray sources in the core of \\omcen, and four in the core of NGC\\,6752; the data for the center of NGC\\,6397 may be fitted with five or six sources. Outside the cores, but within the half-mass radius of the clusters, we detect two sources in \\omcen, one in NGC\\,6397 and two in NGC\\,6752; these may or may not be cluster members. A ROSAT HRI observation of Liller\\,1 does not detect a low-luminosity source, at a limit below a detection with ASCA. We discuss the nature of the low-luminosity X-ray sources in globular clusters in the light of these new results. ",
        "introduction": "Globular clusters contain many X-ray sources at lower luminosities, $L_{\\rm x}\\ltap 10^{34}$\\,\\ergs. These sources were first discovered with the Einstein satellite (Hertz \\&\\ Grindlay 1983), and many more were found with ROSAT (for a compilation, see Johnston \\&\\ Verbunt 1996). The nature of these low-luminosity sources is the subject of debate, because various types of objects can emit X-rays at such luminosities, such as soft X-ray transients in quiescence, cataclysmic variables, RS~CVn binaries, and recycled neutron stars (see e.g.\\ Fig.~8 in Verbunt et al.\\ 1997). The most compelling identification of a dim X-ray source with an object observed at other wavelengths is the recycled radio pulsar in M\\,28: the X-ray flux varies on the pulse period (Danner et al.\\ 1994). Plausible identifications with cataclysmic variables have been suggested for dim X-ray sources in NGC\\,6397, NGC\\,6752, NGC\\,5904 and 47\\,Tuc (Cool et al.\\ 1995b, Grindlay 1993, Hakala et al.\\ 1997, Verbunt \\&\\ Hasinger 1998). These identifications are based on the proximity of the X-ray position to that of a cataclysmic variable, and thus their probability depends on the accuracy of the X-ray position. \\nocite{hg83}\\nocite{jv96}\\nocite{vbrp97}\\nocite{dkt94}\\nocite{cgc+95} \\nocite{gri93}\\nocite{hcjv97}\\nocite{vh98}  The accuracy of the ROSAT position of a detected X-ray source is determined by two factors: the statistical accuracy of the position of the source on the detector, and the accuracy with which the position of the detector as a whole is projected on the sky. For a sufficient number of photons the statistical error is less than an arcsecond, but the projection in general has a typical error of $\\sim5''$. Secure identification of a source in the detector field reduces the error in the projection to the statistical error of the identified source, provided that the optical (or radio) position has better accuracy. Only one identification is necessary, because the roll angle of the detector (i.e.\\ the North-South direction) is accurately known; nonetheless, identification of more than one source is preferable to allow checks on internal consistency. In a globular cluster the surface density of possible counterparts is so high that chance coincidence usually cannot be excluded; a secure identification can usually be made only for X-ray sources detected well outside the cluster. This method has been used by Verbunt \\&\\ Hasinger (1998) to improve the positional accuracy of the sources in the core of 47\\,Tuc from $5''$ to $2''$, whereby the area in which the source is expected to lie is reduced sufficiently to exclude several proposed identifications, and increase the probability of others, including two possible cataclysmic variables.  In this paper we investigate three clusters known to contain multiple dim X-ray sources in their core, which have been observed in long exposures with the ROSAT HRI, and one cluster known to harbour a transient. We analyse hitherto unpublished observations and detect both previously published and new X-ray sources. All source positions are checked in the SIMBAD data base versus positions of other objects, and we find objects in the Hipparcos or Tycho Catalogues (ESA 1997, Perryman et al.\\ 1997, H{\\o}g et al.\\ 1997) with each cluster, i.e.\\ counterparts with very accurate positions. In Sect.\\,2 we describe the observations and our data reduction procedures; Sections 3 to 5 describe the results for \\object{$\\omega$\\,Cen}, \\object{NGC\\,6397}, and \\object{NGC\\,6752}, respectively. In Sect.\\,6 we discuss an observation of \\object{Liller\\,1}. A discussion of our results is given in Sect.\\,7. \\nocite{ESA97}\\nocite{plk+97}\\nocite{hbb+97} ",
        "conclusions": "In the three low-reddened clusters \\omcen, NGC\\,6397 and NGC\\,6752 we have detected a total of 17 dim X-ray sources, of which 5 are well outside the core. The X-ray luminosities of these sources are listed in Table\\,\\ref{txlum}, and plotted in Fig.\\,\\ref{fxlum}. The interpretation of Fig.\\,\\ref{fxlum} must be made with some care. First, sources outside the core may not belong to the cluster; the faintest core source in \\omcen\\ may be a fore- or background source. Second, the conversion of observed countrate to luminosity depends on the assumed spectrum, and from PSPC observations we know that different sources have different spectral parameters (Johnston et al.\\ 1994). For example, the 0.6\\,keV black body spectrum used for the sources in \\omcen\\ gives a 40\\%\\ higher flux for the same countrate than an assumed 0.6\\,keV bremsstrahlung spectrum would give. The bremsstrahlung spectrum is used for the three other clusters. Third, the detection limits in NGC\\,6397, NGC\\,6752 and \\object{47~Tuc} are higher in the cores, where the point spread functions of sources overlap, than outside the core. Such a difference is not present in \\omcen. Fourth, we show the average luminosity, and several sources are known to be variable.  \\begin{table} \\caption{X-ray luminosities in \\ergs\\ in the 0.5-2.5 keV band of the dim X-ray sources in globular clusters described in this paper. For sources in \\omcen\\ we assume a 0.6\\,keV black body spectrum; for those in NGC\\,6397 and NGC\\,6752 a 0.6\\,keV bremsstrahlung spectrum. For the same countrate, the blackbody spectrum corresponds to a flux higher by about 40\\%\\ than the bremsstrahlung spectrum. \\label{txlum}} \\begin{tabular}{rl@{\\hspace{1.cm}}rl@{\\hspace{1.cm}}rl} \\multicolumn{2}{l}{\\omcen} & \\multicolumn{2}{l}{NGC\\,6397} & \\multicolumn{2}{l}{NGC\\,6752} \\\\ X & $\\log L_{\\rm x}$ & X & $\\log L_{\\rm x}$ & X & $\\log L_{\\rm x}$ \\\\ \\multicolumn{2}{c}{core} & \\multicolumn{2}{c}{core} & \\multicolumn{2}{c}{core} \\\\ 9a & 32.14 &          13 & 31.84      & 7a & 31.86 \\\\ 9b & 32.22 &          4a & 31.73      & 7b & 31.89 \\\\ 20 & 31.96 &          4b & 31.75      & 21 & 31.35 \\\\ \\multicolumn{2}{c}{outside} & 4c & 31.43 & 22 & 31.26 \\\\ 7 & 32.33 &           4d & 31.53     & \\multicolumn{2}{c}{outside} \\\\ 21 & 31.88 & \\multicolumn{2}{c}{outside} & 6 & 31.66  \\\\ &       &          12 & 30.95       & 14 & 31.50 \\\\ \\end{tabular} \\end{table}  With these points in mind, we note from Fig.\\,\\ref{fxlum} that in all clusters except possibly \\omcen\\ the most luminous sources appear to be in the cluster core. The main difference between \\omcen\\ and the other clusters is that the collision frequency in \\omcen\\ is so low that one expects no low-mass X-ray binaries in it, and that most cataclysmic variables in it will be evolved from primordial binaries (Verbunt \\&\\ Meylan 1988, Davies 1997). In addition, the mass segregation in this cluster is very low. Thus in \\omcen\\ there is no marked difference between the core and the regions outside the core. \\nocite{dav97}\\nocite{vm88}  In each cluster we detect sources down to the detection limit; this suggests that more sensitive observations will detect more sources. In the cores of NGC\\,6397 and NGC\\,6752 the detection of more source will also require better imaging, so that the faint sources can be detected against the brighter ones. We do not detect a difference between the luminosities of sources in the collapsed globular cluster NGC\\,6397 and the much less concentrated globular cluster NGC\\,6752. On the other hand, the highly concentrated cluster 47~Tuc contains three sources which are an order of magnitude brighter than the brightest sources in NGC\\,6397 and NGC\\,6752.  \\begin{figure} \\centerline{\\psfig{figure=compa.ps,width=\\columnwidth,clip=t} {\\hfil}} \\caption{X-ray countrates of the dim sources in globular clusters as a function of their visual magnitude, compared with the ROSAT PSPC countrates and visual magnitudes of various types of cataclysmic variables (data from Verbunt et al.\\ 1997; filled symbols represent systems first discovered in X-rays and only subsequently identified as cataclysmic variables, i.e.\\ X-ray selected systems) and with RS CVn systems (data from Dempsey et al.\\ 1993) respectively. PSPC countrates of the dim cluster sources have been computed for an assumed 0.6\\,keV bremsstrahlung spectrum, corrected for absorption, from the observed HRI countrates. Visual magnitudes are also corrected for absorption. T indicate sources in 47\\,Tuc (\\x9 and \\x19, $V$ as estimated by Verbunt \\&\\ Hasinger 1998), A in NGC\\,6397 (\\x4b and \\x4c, $V$ from Cool et al.\\ 1998), B in NGC\\,6752 (\\x7a, $V$ from Bailyn et al.\\ 1996). The dotted lines indicate a constant ratio of X-ray to optical flux. \\label{fcompa}} \\end{figure} \\nocite{dlfs93}  Viable optical counterparts have been suggested for only five among the 26 sources shown in Fig.\\,\\ref{fxlum}, all of them probable cataclysmic variables. We compare the ratio of X-ray flux to optical flux of these sources with the ratios measured for cataclysmic variables and for RS CVn systems in the Galactic Disk in Fig.\\,\\ref{fcompa}. It is seen that the suggested optical counterparts for the sources in NGC\\,6397 and NGC\\,6752 lead to ratios which are compatible with those of cataclysmic variables, whereas those in 47~Tuc are too bright in X-rays, in agreement with Fig.\\,\\ref{fxlum}. If these sources are indeed cataclysmic variables, their excessive X-ray luminosity needs to be explained; alternatively, the suggested identifications may be chance coincidences (as discussed by Verbunt \\&\\ Hasinger 1998). All suggested counterparts lead to higher X-ray to optical flux ratios than those of RS CVn binaries.  The accurate positions that we determine for individual sources are valid for separately detected sources in particular. In the case of overlapping sources, we do not have unique solutions. Thus, in the core of NGC\\,6397 fits with 5 and 6 sources are both acceptable, at similar quality; and we cannot exclude that more sources contribute to the observed flux, which would invalidate our derived positions.  Binaries may reside away from the core either because the cluster has undergone little mass segregation, or because a three-body interaction (i.e.\\ a close encounter of a binary with a single star) in the core has expelled the binary from the core (e.g.\\ Hut et al.\\ 1992). \\nocite{hmr92} In the latter case the binary is expected to be eccentric immediately after being expelled; tidal forces may in time circularize the orbit again. Such binaries are only a minority of the overall binary population of a cluster; however, X-ray observations may preferably select such binaries if tidal forces act in them. Since sources away from the core can be fore- or background sources, optical identification of them is required to settle whether they belong to the cluster or not. Our accurate positions should help in finding such counterparts."
    },
    "0004/astro-ph0004346_arXiv.txt": {
        "abstract": "This paper presents four color narrow-band photometry of clusters A115 ($z=0.191$) and A2283 ($z=0.182$) in order to follow the star formation history of various galaxy types.  Although located at similar redshifts, the two clusters display very different fractions of blue galaxies (i.e. the Butcher-Oemler effect, $f_B = 0.13$ for A115, $f_B = 0.30$ for A2283).  A system of photometric classification is applied to the cluster members that divides the cluster population into four classes based on their recent levels of star formation.  It is shown that the blue population of each cluster is primarily composed of normal starforming (SFR $< 1 M_{\\sun}$ yrs$^{-1}$) galaxies at the high luminosity end, but with an increasing contribution from a dwarf starburst population below $M_{5500}= -20$.  This dwarf starburst population appears to be the same population of low mass galaxies identified in recent HST imaging (Koo \\etal 1997), possible progenitors to present-day cluster dwarf ellipticals, irregulars and BCD's.  Deviations in the color-magnitude relationship for the red galaxies in each cluster suggest that a population of blue S0's is evolving into present-day S0 colors at this epoch.  The radial distribution of the blue population supports the prediction of galaxy harassment mechanisms for tidally induced star formation operating on an infalling set of gas-rich galaxies. ",
        "introduction": "The investigation of the blue and red populations in rich clusters have been our most lucrative glimpse into the evolution of galaxies.  The search for changes in the stellar population of galaxies has mostly focused on clusters of galaxies due to their high visibility and ease of cataloging rich clusters even at distant redshifts.  However, the gravitational clumping of a cluster also minimizes the effort required for distance determination where the measurement of a few of the brightest galaxies provides the redshift for the entire cluster.  Attention concerning recent and rapid evolution has been diverted in the past decade to field galaxies (Tyson 1988), and the dilemma of the blue field population.  However, clusters are the sites of numerous dramatic evolutionary effects, such as galaxy cannibalism (Moore \\etal 1996) and the Butcher-Oemler population (Oemler, Dressler and Butcher 1997).  In addition, cluster populations are key to understanding galaxy characteristics since they cover not only a full range of galaxy masses (from giants to dwarfs) but also the full range of Hubble types and intrinsic density (e.g. surface brightness) that are missing from field populations.  In a series of papers extending over the last 12 years (Rakos \\etal 1988, 1991, 1995, 1996, 1997, 1999), we have used a narrow band filter system to perform photometry of galaxies in rich clusters for redshifts ranging from 0.2 to 1.  Our studies have differed from previous photometry of distant clusters by the use of a color system surrounding the 4000\\AA\\ break (the Str\\\"omgren $uvby$ system) and modified such that the filters are ``redshifted'' to the cluster of galaxies in consideration.  This method results in effectively no k-corrections and allows discrimination between cluster membership based on spectrophotometric criteria.  We call our modified system $uz,vz,bz,yz$ to distinguish it from the original $uvby$ Str\\\"omgren system and we believe we have demonstrated that these color indices are a profitable tool for investigating color evolution of both the red and blue populations in clusters of galaxies.  In an previous exploratory paper (Rakos, Maindl and Schombert 1996, hereafter RMS96), the modified $uz,vz,bz,yz$ Str\\\"omgren system was used to develop a photometric classification scheme based on the recent star formation history of a galaxy.  From the $mz$ color index ($mz=(vz-bz)-(bz-yz)$), the classification technique was shown to be extremely successful at discriminating normal starforming galaxies (spirals) and starburst galaxies despite the presence of heavy reddening. Rakos \\etal (1996, 1997) applied this technique to the members of the blue population in several clusters of intermediate redshift ($0.2 < z < 0.6$) and demonstrated that many of the cluster members have strong signatures of star formation activity.  An extreme example is CL0317+1521 at $z=0.583$ which has a blue fraction of $f_B=0.60$ and which 42\\% of the blue population have $mz<-0.2$, the photometric signature for a starburst.  Deep photometry of the cluster A2317 ($z=0.211$, Rakos, Odell and Schombert 1997) demonstrated that the ratio of the blue population to red population has a strong dependence on luminosity, such that blue galaxies dominate the very brightest and very faintest galaxies in the cluster.  In contrast to the bright blue galaxies, the fraction of galaxies displaying the signatures of a starburst increases towards the faint end of the luminosity function, a dwarf starburst population first suggested by Koo \\etal (1997).  The origin of this dwarf starburst population remains an enigma.  Tidal interactions are frequently invoked as an explanation for the high fraction of starburst galaxies in Butcher-Oemler clusters (Dressler \\etal 1994, Couch \\etal 1994).  If the same phenomenon acts on the low mass galaxies, then these starburst systems would have their origin as gas-rich dwarf galaxies who have had a short, but intense, tidally induced episode of star formation which would quickly exhaust their limited gas supply. It should be noted, however, that the orbits of cluster galaxies are primarily radial (Bothun and Schombert 1990), and the typical velocities into the dense cluster core are high.  This makes any encounter extremely short-lived, with little impulse being transfered as is required to shock the incumbent molecular clouds into a nuclear starburst.  The galaxy harassment mechanism (Moore \\etal 1996) emphasizes the influence of the cluster tidal field and the more powerful impulse encounters with individual central galaxies at the cluster edges.  These two processes can then conspire raise the luminosity of cluster dwarfs, increase their visibility and, thus, their detectability.  To further explore the behavior of the red E/S0's, blue Butcher-Oemler galaxies and the newly detected dwarf starburst population, we began a program of obtaining deeper photometry then our previous work on intermediate redshift clusters.  Our goal in this study is to 1) illuminate the types of galaxies involved in the Butcher-Oemler effect, 2) determine the dominance of the blue galaxies to the cluster luminosity function, 3) examine the spatial extent of the blue and red population, and 4) determine the existence and characteristics of the dwarf starburst population in other blue clusters.  In order to maintain a correct timescale to stellar population models, calibrated to globular cluster ages, values of $H_o=50$ km sec$^{-1}$ Mpc$^{-1}$ and $q_o=0$ are used throughout this paper. ",
        "conclusions": "The trends with respect to cluster populations for A115 (southern subcluster) and A2283 (combined with the data for A2218 and A2317) are shown in Figure 9.  The data for A2317 are published in Rakos, Odell and Schombert (1997) and the data for A2218 will be published in Rakos, Dominis and Steindling (2000).  Four parameters are plotted with respect to the elliptical/S0 fraction; the fraction of the blue population ($f_B$), the fraction of starburst galaxies, the richness of the cluster ($C$) and the radius of where the maximum value of $f_B$ occurs.  The fraction of the blue population drops inversely with the E/S0 percentage.  This is not too surprising since a majority of the red population is classified photometrically as E/S0's.  However, the relation is nearly 1-to-1, suggesting that it is conversion of the E/S0 objects into the blue population that comprises the Butcher-Oemler effect.  Since our previous work on present-day ellipticals and the color evolution of ellipticals (Schombert \\etal 1993, Rakos and Schombert 1995) demonstrates that their colors have a narrow range and closely match the models for a passively evolving population with a formation redshift of 5, we must conclude that it is primarily the S0 galaxies that are the normal starforming objects of the Butcher-Oemler population.  This scenario matches well to our expectations of gas depletion, where large bulge S0's exhaust their supply first, followed by Sa's, then Sb's in the near future.  We can assume that a gas-rich, proto-S0 undergoing star formation rates of around 1 $M_{\\sun}$ yrs$^{-1}$ will produce spiral structure and have morphologies similar to early-type spirals as revealed by HST imaging.  The fraction of starbursts is only weakly related to the E/S0 fraction confirming the split impact that the reddened starburst galaxies have on the color distribution of cluster populations.  This is a new phenomenon for distant clusters and, while the same mechanisms may be at work for the dwarf starbursts, we believe this is a separate event from the Butcher-Oemler effect.  The radius of the peak density of the blue population is also uncorrelated with any population fraction or color fraction of the clusters.  This fact is illuminating in the sense that distribution of the blue population in the red clusters (A115 and A2218) is similar to that of the blue clusters (A2283 and A2317).  If tidal forces are responsible for the blue population, then there will be a mix of effects due to close encounters with other galaxies and the mean cluster tidal field which will only be weakly dependent on the cluster morphology.  If ram pressure stripping is important, then this radius should be correlated with the x-ray profile of the cluster.  One of the primary results of this study is that there appears to exist a {\\it duality} with respect to cluster population studies by morphology versus photometric classification.  Morphological studies of intermediate redshift clusters (Oemler, Dressler and Butcher 1997, Dressler \\etal 1997) find the fraction of ellipticals to be similar to present-day clusters with a sharp drop in the number of S0's and a proportional increase in the number of spirals and irregulars.  There is also an increase in the number of disturbed galaxies with redshift (Couch \\etal 1998), galaxies with tidal signatures of past encounters.  On the other hand, photometric classification (such as this study and Dressler \\etal 1999) find a population of bright blue galaxies with normal star formation rates (the Butcher-Oemler effect) and a starburst population, dominated by low mass galaxies. But, the characteristics of these populations differ in their luminosity and geographical distributions compared to their morphological counterparts.  This is particularly evident in magnitude diagrams such as Figure 5 compared to similar diagrams for nearby, rich clusters.  {\\it The underlying consequence of this duality is the decoupling of the presence of star formation of the properties of Hubble types with redshift} (i.e. the existence of blue S0's, starbursting dwarfs, disturbed spirals at high redshift).  How this decoupling takes place is due to a separation of formation events (such as the density of the protogalaxy, the fraction of dark to baryonic matter) and the later effects of environment.  Clearly, local density plays an important role in the evolution of a galaxy. As shown in Hashimoto \\etal 1998, galaxies at intermediate densities have higher star formation rates than galaxies located in high dense regions (such as the core of a rich cluster).  If the dominate process were galaxy encounters that induce star formation, then this correlation would be expected as intermediate density environments (such as loose groups) have the appropriate mix of galaxy density and low velocities to maximize tidal effects (Hashimoto \\etal 1998).  However, a competing process, that a pure density analysis does not take into account, is the possibly of ram pressure effects from the hot intracluster gas identified with every rich cluster of galaxies.  There is clear evidence of such stripping ongoing in nearby clusters, such as Virgo, based on the HI and CO images of cluster spirals (Kenney and Young 1989).  We conclude that the Butcher-Oemler effect in rich clusters is the phenomenon of starforming S0's being ram pressure stripped as they encounter the hot gas in the core of the cluster.  The elimination of the atomic gas prematurely ceases star formation (although the question remains of the interplay with the molecular gas) and the S0's disks age to the population we see today.  On the other hand, the dwarf galaxy interactions in the outer regions of the cluster (where the velocities are lower) produce the starburst population.  Both the blue and starburst population avoid the core of the cluster, but each for different reasons, one being intrinsic, the other environmental."
    },
    "0004/astro-ph0004170_arXiv.txt": {
        "abstract": "We present a catalog of extremely red objects discovered using the NICMOS/HST parallel imaging database and ground-based optical follow-up observations. Within an area of 16 square arc-minutes, we detect 15 objects with $\\rm R - F160W > 5$ and $\\rm F160W < 21.5$. We have also obtained K-band photometry for a subset of the 15 EROs. All of the $\\rm R - F160W$ selected EROs imaged at K-band have $\\rm R - K > 6$. Our objects have $\\rm F110W - F160W$ colors in the range of 1.3 $-$ 2.1, redder than the cluster ellipticals at $z \\sim 0.8$ and nearly 1 magnitude redder than the average population selected from the F160W images at the same depth. In addition, among only 22 NICMOS pointings, we detected two groups or clusters in two fields, each contains 3 or more EROs, suggesting that extremely red galaxies may be strongly clustered. At bright magnitudes with $\\rm F160W < 19.5$, the ERO surface density is similar to what has been measured by other surveys. At the limit of our sample, F160W $= 21.5$, our measured surface density is 0.94$\\pm 0.24$ arcmin$^{-2}$. Excluding the two possible groups/clusters and the one apparently stellar object, reduces the surface density to 0.38$\\pm 0.15$~arcmin$^{-2}$. ",
        "introduction": "Deep near-IR imaging surveys have revealed a population of extremely red objects (``EROs''; Elston, Rieke \\&\\ Rieke 1988; McCarthy, Persson \\&\\ West 1992; Graham \\&\\ Dey 1996; Hu \\&\\ Ridgeway 1994; Soifer et al. 1994; Dey, Spinrad \\&\\ Dickinson 1995; Thompson et al. 1999).  The nature of the extremely red population remains unclear. As it is defined largely by a single color, primarily $\\rm R - K$, there is no certainty that it represents a uniform class of object, and it may contain contributions from galaxies, cool stars or substellar objects and active nuclei.  The precise definition of an ERO varies among the different surveys and depends on the particular bandpasses employed. Most samples were defined by $\\rm R - K \\gta 5 - 6$ or $\\rm I - K \\gta 4 - 5$.  The present work is based on a somewhat different color system, that defined by the NICMOS F160W bandpass and conventional Kron-Cousins R magnitudes. The NICMOS F160W band is similar to the Johnson H band filter and the $\\rm H - F160W$ color term is negligable for a flat spectral energy distribution (in $f_{\\nu}$ units) (M. Rieke, 1999, private communication). Among the resolved objects there are reasonable expectations that there should exist stellar systems at redshifts such that the  K-correction applied to an old or intermediate age population will produce very red optical to near-IR colors. Alternatively even fairly modest extinction, when observed in the same redshift range, can produce very steep spectral energy distributions in the rest-frame near-UV and there are local examples of such objects among the dusty starburst population.  The central issue regarding the nature of the resolved EROs is to understand to what degree these two classes of objects contribute to the overall population.  The earliest interpretation of the colors of EROs were centered around the old stellar population hypothesis (e.g. McCarthy, Persson, \\& West 1991; Hu \\& Ridgeway 1994) and redshifts of 1 - 2 were inferred from their multi-band photometry. There are clear examples which support this interpretation, such as weak radio source LBDS 53W091 at $z = 1.55$ (Dunlop et al. 1996) and near-IR selected object CL0939+4713B at z$=1.58$ (Soifer et al. 1999), and a concentration of EROs at $z = 1.3$ (Liu et al. 2000). Graham and Dey (1996) argued that the spectral energy distribution of HR10 (Hu \\& Ridgeway 1994) is well matched by a dusty star-forming galaxy at $z = 1.5$.  The detection of a strong sub-mm continuum from HR10 (Cimatti et al. 1998; Dey et al. 1999) provided conclusive evidence that some of EROs, if not all, are dust enshrouded starburst galaxies with inferred star formation rates of $\\rm 500-2000 h_{50}^{-2}M_\\odot~yr^{-1}$ at moderate redshifts ($z \\sim 1 -2$).  Recent deep near-IR follow-up observations of the sub-mm sources detected with the SCUBA (Smail et al. 1998) have suggested that two faint sub-mm sources may also be EROs with $\\rm I - K > 6$ (Smail et al. 1999). Liu et al. have measured redshifts for several other EROs and most of these are also in the $0.8 < z < 1.5$ range and some have moderately strong emission lines (Liu et al. 2000).  Two important open issues concern the surface density of EROs and the relative contribution of different classes of objects as a function of both color and apparent magnitude. We continue to lack statistically large samples of EROs. The most recent systematic large survey, covering an area of 154~square arcminutes by Thompson et al. (1999), has yielded six objects with $K \\le 19.0$ and $ R - K > 6$. To quantify the fraction of different classes of EROs as a function of colors and magnitude, larger samples are required. The ERO surface density inferred from the Thompson et al. survey is 0.04$\\pm 0.016$ arcmin$^{-2}$ for $K \\le 19.0$ and $R - K > 6$. Depending on the K-band magnitude limit, $\\rm R - K$ color and possibly environment, the reported surface densities of EROs range from $0.01 - 0.7$ arcmin$^{-2}$, derived from several serendipitous surveys over small areas (Hu \\&\\ Ridgway 1994; Cowie et al. 1994; Beckwith et al. 1998; Thompson et al. 1999). There have been suggestions that EROs tend to cluster, particularly, in regions around high redshift AGNs compared with blank fields.  The large luminosities and (admittedly uncertain) space densities imply that these objects represent a singificant consitutent of the overall galaxy population and that their contribution to the overall rate of star formation is non-negligable (e.g. Liu et al. 2000).   In this paper, we present a sample of EROs discovered using NICMOS on HST while operating in the parallel mode. Our combined NICMOS/optical survey covers only 16 square arcminutes but it provides high spatial resolution and better signal-to-noise in the near-IR than most ground-based surveys.  \\medskip ",
        "conclusions": ""
    },
    "0004/astro-ph0004058_arXiv.txt": {
        "abstract": "Planet formation is accompanied by the formation of comet clouds. In systems where planets migrate on rapid timescales, the diffusive evolution of comet orbits may stall, resulting in a comet cloud intermediate between a flattened Kuiper Belt and a spherical Oort cloud. These `failed Oort clouds' may provide a `smoking gun', indicating that planetary migration has taken place. If some fraction of the scattered component consists of planetary embryos, it may be possible to observe transits of such bodies even when the planetary system is not edge-on to the line of sight. ",
        "introduction": "Planet formation is a messy business. The process of accumulation of small bodies into large bodies results not only in the formation of large planetary-mass objects but also a significant population of scattered comets and asteroids. Studies of short and long period comets in our solar system (Oort 1950; Fernandez 1980; Duncan, Quinn \\& Tremaine 1988) point to the existence of two repositories of cometary material, the spherically distributed Oort cloud (Oort 1950) and a low inclination population of trans-Neptunian objects, the Edgeworth/Kuiper Belt (Edgeworth 1949; Kuiper 1951). These cometary reservoirs convey valuable information regarding the primordial conditions in the solar system and their manifestations around other stars offer insights into the formation of extrasolar planetary systems. In particular, the detection of infra-red emission from dust in such systems (Backman \\& Paresce 1993; Trilling \\& Brown 1998) is thought to supplied by ongoing evaporation or ablation of comet-like objects (Weissman 1984) .  The properties of recently detected extrasolar planets (e.g. Marcy, Cochran \\& Mayor 1999) suggest that the formation and  evolution of planetary systems may be a much more dynamic and violent process than previously envisaged. The purpose of this letter is to examine the impact of this new paradigm on the configuration of cometary reservoirs in such systems. In particular, we will demonstrate the existence of a cometary component intermediate between a disk and isotropic component that will result in systems where significant planetary migration has taken place.  Section~\\ref{Scatter} reviews the process of formation of an Oort cloud and examines the implications of planetary migration for such a scenario. In section~\\ref{Dust} we describe the implications for observations of dust disks around extrasolar planetary systems. ",
        "conclusions": ""
    },
    "0004/astro-ph0004091_arXiv.txt": {
        "abstract": "We present the first dynamical study of the optically selected Palomar Distant Cluster Survey (PDCS).  We have measured redshifts for seventeen clusters of galaxies in the PDCS and velocity dispersions for a subset of eleven.  Using our new cluster redshifts, we re-determine the X-ray luminosities and upper limits.  We show that eleven out of twelve PDCS clusters we observed are real over-densities of galaxies.  Most clusters have velocity dispersions appropriate for clusters of galaxies.  However, we find a fraction ($\\sim$ \\slantfrac{1}{3}) of objects in the PDCS which have velocity dispersions in the range of groups of galaxies (200 km s$^{-1}$ $\\pm$ 100 km s$^{-1}$) but have richnesses appropriate for clusters of galaxies.   Within our survey volume of $31.7^{+0.5}_{-0.8}\\times 10^4$ h$^{-3}$ Mpc$^3$ ($q_o = 0.1$) for Richness Class 2 and greater clusters, we measure the richness function, X-ray luminosity function (using both the detections and upper limits), and the mass function derived from our velocity dispersions.  We confirm that the space density, as a function of richness, of clusters of galaxies in the PDCS is $\\sim$ 5 times that of the Abell catalog.  Excluding the above fraction of \\slantfrac{1}{3} of objects with low velocity dispersions, we measure a space density $\\sim$ 3 times that of the Abell catalog for equivalent mass clusters of galaxies, raising the possibility that the Abell catalog is incomplete.  However, our space density estimates are in agreement with other low-redshift, optically-selected cluster surveys such as the EDCC, APM and EDCC2.  Our X-ray luminosity function agrees with other measurements based on both X-ray and optically selected samples, so we find that the PDCS does not miss clusters of galaxies that would be found in an X-ray selected survey. Our resulting mass function, centered around $10^{14}$ M$_{\\sun}$ $h^{-1}$, agrees with the expectations from such surveys as the Canadian Network for Observational Cosmology cluster survey, though errors on our mass measurements are too large to constrain cosmological parameters.  We do show that future machine-based, optically-selected surveys can be used to constrain cosmological parameters. ",
        "introduction": "The space density of virialized clusters of galaxies as a function of mass and redshift are primary predictions of theories on the formation and evolution of structure.  In general, three parameters control the shape, normalization and evolution of the mass function of clusters of galaxies.  These parameters are the density of matter in the universe, $\\Omega_m$, the variance of the distribution of mass density fluctuations at cluster scales, $\\sigma_8$, and the shape of the spectrum of density fluctuations, commonly quantified as $\\Gamma$ \\citep{press74,efstathiou92,lacey93,viana96, Eke96, Kitayama97}. We can directly test these theoretical predictions and, therefore, constrain these parameters, by measuring masses of clusters of galaxies in a sample with a known survey volume.  Most efforts to measure the space density of clusters of galaxies as a function of mass have focused on X-ray selected samples.  X-ray selection has two advantages.  First, X-ray luminosity is strongly correlated with mass.  Second, X-ray selected surveys have a selection function that is easy to quantify.  Therefore, previous work focused on the mass function derived from velocity dispersions of X-ray selected clusters \\citep[as examples]{carlberg97b, borgani99}, the temperature function \\citep[for an example]{henry97} or the luminosity function \\citep[as examples]{reichart99, borgani99b}.  There is a large spread in the above results.  Yet, most of the papers mentioned above use the sample of clusters of galaxies from the Extended Medium Sensitivity Survey or EMSS \\citep{gioia90b,gioia90a,henry92,nichol97}.  For example, \\citet{reichart99} finds the most likely value of $\\Omega_m$ from the mass function of the EMSS to be around 1, while the CNOC survey \\citep[for example]{carlberg96} finds a the most likely value to be to 0.2.  Recently, \\citet{borgani99} finds that $\\Omega_m$ from the CNOC survey can be constrained to the range $0.35 < \\Omega_m < 1.0$. \\citet{borgani99} finds that the large uncertainty in the best fitting values of $\\Omega_m$ is partly the result of the uncertainty in the mass function of low redshift clusters of galaxies.  However, the high redshift sample of the EMSS only contains the highest mass clusters and, furthermore, is incomplete even at those masses (see Borgani \\etal 1999a\\nocite{borgani99} for a discussion of the completeness of the sample used by Carlberg \\etal 1997\\nocite{carlberg97b}, Borgani \\etal 1999a\\nocite{borgani99} estimate that the CNOC sample contains only 25\\% of all clusters with $\\sigma_v > 800$ km s$^{-1}$ because of the X-ray threshold of $ >10^{45}$ erg s$^{-1}$).  Though there is great potential in the CNOC sample, both a larger number of clusters and a larger range in masses at both high and low redshift is needed to constrain cosmological parameters.  To increase the range in mass at high redshift, we have used optically selected clusters of galaxies.  Specifically, we have collected a number of redshift measurements towards clusters of galaxies in the Palomar Distant Cluster Survey \\citep[PDCS]{postman96} so we can measure velocity dispersions, the first such velocity dispersion measurements of this catalog.  With our sample of $0.2 \\le z \\le 0.6$ clusters of galaxies, we could potentially improve the sample of \\citet{carlberg97b} and make a better measurement of the value of $\\Omega_m$.  The PDCS contains a much larger space density of clusters of galaxies, between $\\sim 10^{-5}$ and $\\sim 10^{-6}$ h$^{-3}$ Mpc$^{3}$, in the same redshift range as the EMSS, and, therefore, should contain lower mass clusters.  We should then have a sample that, when combined with the EMSS, increases the dynamic range of the cluster mass function making a more robust measurement of $\\Omega_m$.  The approach of the PDCS was to create a model of what clusters of galaxies look like, called a matched-filter.  The model for the galaxy distribution is a Schechter function for the luminosity distribution of the cluster galaxies and a profile of \\(\\frac{1}{\\sqrt{1 + (r/r_c)^2}} \\).  The core radius of the radial profile were fixed as were the slope and ``knee'' of the Schechter luminosity function.  The total cluster size was also restricted to 1 $h^{-1}$ Mpc.  The strength of any observed correlation of the observed galaxy catalog with this matched-filter can be used to measure how well the model matches the data, with the strongest correlations being assigned to cluster candidates.  In this way, the PDCS also generated an estimated redshift, a galaxy richness based on the normalization of the luminosity function and other parameters for each cluster candidate based on the best-fitting model.  The advantage of this approach has lead other groups to use similar techniques, see for example \\citet{dalton97}, \\citet{kepner99}, \\citet{olsen99} and \\citet{bramel99}.  Using only the derived quantities from the cluster catalog, the authors of the PDCS found that the space density of clusters of galaxies was $5\\ \\pm\\ 2$ times that of the space density in the Abell catalog (though the space density is consistent with that found in low redshift automated catalogs such as the Edinburgh-Durham Cluster Catalog and APM cluster catalogs).  Secondly, the authors found no evidence for evolution in the space density with redshift.  Both of these results rely on the cluster catalog alone, a catalog based entirely on imaging data.  We have completed a program of obtaining redshifts and velocity dispersions of PDCS clusters to measure the space density as a function of mass as well as richness and X-ray luminosity.  This will allow us to both test the original results of the PDCS and possibly provide a complementary sample to that of \\citet{carlberg97b}.  We began this program with the X-ray survey of \\citet{holden97}, or H97. Our first spectroscopic observations are described in \\citet{holden99}, H99, but the majority of spectra are from the CFHT Optical PDCS survey which is described in \\citet{adami99}.  We summarized our data in \\S 2, with an emphasis on what changes we have made from the previously mentioned papers.  We then derive cluster redshifts, velocity dispersions, X-ray luminosities and masses in \\S 3.  Using the survey volume estimated in \\S 4, we find the richness and X-ray luminosity function of PDCS clusters in \\S 5.  We compute the mass function of PDCS clusters, computed in \\S 6, and find it is consistent with the mass function found for the sample of \\citet{carlberg97b} given specific choices for the relation between $\\sigma_8$ and $\\Omega_m$, the value of $\\Gamma$ and the form of the richness-mass relation.  Finally, in \\S 7 we summarize our results and discuss the future prospects of using optically selected clusters of galaxies to probe theoretical models of cluster formation and evolution.  Unless otherwise noted, we use $q_o = 0.1$ ($\\Omega_m = 0.2$ and $\\Lambda = 0$) and $H_o = 100\\ h\\ {\\rm km\\ s^{-1}}$. ",
        "conclusions": "In this paper we presented redshift measurements, velocity dispersions and X-ray imaging data for a subset of the PDCS.  We aimed to construct the first dataset of velocity dispersions of a subsample of the PDCS. This dataset was constructed to be complementary to the CNOC cluster survey and to be easily compared with lower redshift samples.  Our goal was to use this dataset to determine the space density of clusters of galaxies in the PDCS catalog as a function of richness, X-ray luminosity, and mass.  We could then compare these space densities with low redshift samples to test the apparent lack of evolution in the PDCS cluster catalog and to compare the distribution of masses with the expectations of cluster formation theory.  We began this process by collecting X-ray imaging data for a total of thirty-one cluster candidates in the PDCS.  We followed up the X-ray imaging data with spectroscopy of PDCS cluster candidate galaxies.  We successfully measured redshifts for seventeen PDCS clusters (fifteen of which we have X-ray imaging data for).  The PDCS clusters of galaxies for which we have redshifts are part of two samples with different selection criteria.  One sample, the ARC/KPNO sample, was selected on the number of $V-I \\ge 1$, $V \\le 21$ galaxies in a 2\\farcm 5 circle centered on the PDCS cluster candidate position. This provided a sample selected independently of many of the derived cluster parameters in the PDCS (such as richness, estimated redshift, {\\em etc}.)  Our second sample, the CFHT sample, was selected using cluster parameters from the catalog, specifically all clusters in the X-ray imaging data with $\\Lambda_V \\ge 40$ and $0.3 \\le z_{estimated} \\le 0.5$ where $\\Lambda_V$ is the richness estimate of the PDCS in the V filter.  This second sample was constructed to mimic the selection of the ESO Nearby Abell Cluster Survey in that we selected clusters of Richness Class 1 or greater \\citep{katgert96}.  Given these two samples of redshifts, we then derived redshifts for the clusters, and, when possible, velocity dispersions based on the galaxy redshifts.  We fit a Gaussian plus flat background using a maximum likelihood estimator to measure the velocity dispersions.  This approach has the advantage of using the galaxy velocities directly and is completely automated.  However, as we made an explicit choice on the shape of the velocity distribution, we are biased towards over-estimating the velocity dispersion of clusters with complicated dynamics.  We tested our approach using the published CNOC velocities and found that for seven out of the eight clusters we measured the same velocity dispersions and the same size errors, $\\sim$ 10\\%.  The one exception was MS 1512+36, a dynamically complicated cluster with two secondary peaks in the velocity distribution.  Most of the clusters in our sample had a much smaller number of cluster members than the CNOC sample.  Therefore, our velocity dispersions had much larger errors, up to $\\sim$ 50\\%.  Using the redshifts, we measured X-ray luminosities or upper limits for thirty-one PDCS cluster candidates.  We chose to measure our X-ray luminosities within an aperture that contains a fixed fraction of the total flux, using a $\\beta$ model to calculate the size of that aperture.  For those sixteen PDCS cluster candidates for which we did not have spectroscopic redshifts, we used the estimated redshifts of the PDCS to compute the apertures.  Twenty-four of the candidates in our sample were not X-ray detections.  For those, we assigned $3\\sigma$ upper limits on the flux and luminosity.  When we compare our masses derived from the velocity dispersions to the richnesses in the PDCS catalog in Figure \\ref{richvmass}, we find that \\slantfrac{1}{3} of our masses are lower than the expected values based on the richnesses.  For example, we have in our sample two Richness Class 3 clusters with velocity dispersions of only $\\sim$500 km s$^{-1}$, with implied masses around $1-3 \\times 10^{14}$ $h^{-1}$ M$_{\\sun}$.  We find that the relation between masses and richnesses in the PDCS sample differs from the Abell mass-richness relation from \\cite{girardi98b} at the 99.5\\% confidence limit.  This observed difference comes from these low mass systems with high richnesses.  In \\S 4.2 and 4.3 we derived the volumes surveyed for our two samples. We find that these volumes are accurate to around $\\sim$ 20\\% except for changes in $q_o$.  These are much smaller errors than our errors based on our sample size.  We used these volumes to derive the space density of PDCS clusters as a function of richness, X-ray luminosity and mass.  If we examine the richness function of the clusters in our sample, we find space densities in line with those of the whole of the PDCS.  As we mentioned before, this result is not surprising given how accurate the estimated redshifts are for the clusters in our sample and that most of the candidates in the PDCS are true over-densities of galaxies.  Based on this result, we find that the PDCS agrees with the space densities for lower redshift catalogs of clusters of galaxies like the EDCC2 of \\citet{bramel99} and the APM results from \\citet{croft99}.  However, this raises the question of why all of these catalogs find a much higher space density of clusters of galaxies than the Abell catalog at ostensibly the same richnesses?  To try to answer this question, we first look at a different distribution function, the X-ray luminosity function.  What is interesting is that the X-ray luminosity function of PDCS clusters of galaxies is quite in line with those based on X-ray selected samples, such as \\citet{burke97} and \\citet{rosati98} as well as those based on other optically selected samples such as \\citet{burns96}.  This means that it is unlikely that the PDCS missed any X-ray emitting clusters of galaxies.  However, considering that most of the clusters of galaxies are not detected in our X-ray data, those clusters must have quite faint luminosities ($ \\le 10^{43}$ erg s$^{-1}$, see Figure \\ref{nvl}).  The clusters at those luminosities in the optically-selected, low redshift sample of \\cite{burns96} are quite poor, being either Richness Class 0 Abell clusters or groups from the catalog of \\citet{white99}.  If we examine the mass function, once again we find that the distribution of masses in our sample matches that of other samples. We chose to compare our mass function with those of other samples by finding which cosmological parameters are most likely to produce the distribution of masses in our sample.  We find the best fitting $\\Omega_m$ to be quite in line with those from the CNOC cluster survey if we assume both the $\\sigma_8 - \\Omega_m$ relation and the richness-mass relation from \\citet{girardi98b}. By the theoretical mass function that best fits our data, we see that most of our clusters should have masses on the order of $\\sim 10^{14}$ h$^{-1}$ M$_{\\odot}$, which is in agreement with our data, and is the mass range found for Richness Class 0 and 1 clusters in Figure \\ref{richvmass}.  Based on our data, we find that the PDCS finds clusters of galaxies have distributions of mass and X-ray luminosity that match those of other samples.  This leaves the richness function as the discrepant function, but only when compared with the Abell catalog.  Other optically selected catalogs of clusters of galaxies find a much higher space density of objects than the Abell catalog, such as the APM cluster catalog \\citep{dalton97} or the Edinburgh-Durham Cluster Catalog 2 \\citep{bramel99}.  Therefore, it is likely that either the Abell catalog is strongly incomplete by a factor of three or four, a rather unlikely occurrence, or that there is a mismatch between the various richness measures used by the automated catalogs and the Abell Richness.  These ideas will be tested with future cluster catalogs that cover a large area of the sky, such as the DPOSS \\citep{gal99} and SDSS cluster catalogs.  The promise of optically selected catalogs has always been that, given the relative ease of collecting large amounts of imaging data, catalogs with literally thousands of clusters of galaxies could be constructed.  Our investigations of the PDCS shows that modern sophisticated cluster finding algorithms do an excellent job of finding over-densities of galaxies and correctly estimating their apparent redshift.  From Figure \\ref{massf}, we can see that it will be possible to use optically selected clusters of galaxies to study cluster formation and evolution.  The challenge will lie in improving the contrast for truly massive systems, to decrease the contamination we see in Figure \\ref{richvmass}.  Machine based richnesses appear to be only an incremental improvement over counts by eye \\citep[for a comparison of richness to mass for the APM and Abell catalogs]{alonso99}.  A possible solution could be in using such techniques as a color-based or photometric redshift selection, see for example \\citet{kepner99} or \\citet{gal99}.  The approaches have the promise of reducing the number of contaminating field galaxies in cluster candidates.  A second solution would be to replace a simple count of the number of galaxies with new measures of the total galaxy content such as the total luminosity from \\citet{adami98}, \\citet{fritsch99} or \\citet{miller00}.  Catalogs using the combination of these two techniques should have the observed optical galaxy signal much more strongly correlated with mass.  With such improvements and the new large area surveys such as the SDSS and DPOSS, optically selected clusters can be used to probe the mass function over 4 or 5 orders of magnitude in density and, thus, be used to explore the formation of the most massive gravitationally bound objects observed.  This paper represents the submitted version of BH's dissertation.  I would like to thank a number of people for reading over various versions of this work.  This list includes Rich Kron, who read the whole thing at least twice, and my collaborators: Christophe Adami, Francisco Javier Castander, Lori Lubin, Robert Nichol, and A. Katherine Romer.  I would also like to thank the anonymous referee for improving this paper.  For help in understanding statistics, I thank Erik Reese for many useful pointers and Carlo Graziani who spent hours explaining Maximum Likelihood parameter estimation and the theory behind confidence limits.  I would like to credit Gil Holder for correcting many of my misunderstandings of the Press-Schechter formulation and for testing my predictions of the cluster mass function.  Daniel Reichart also helped with my understanding of Press-Schechter theory as well as all of the pitfalls of converting observational quantities like richness into the masses predicted from theory.  Finally, I would like to thank Chris Metzler and Martin White for long explanations during the Cluster Eating Group about cluster formation and the spherical collapse model, I would like to thank Chris especially for his long patient explanations of relating theoretical predictions to observations.  For scientific guidance, I would like to thank Rich Kron for teaching so well his particular way of looking at observational data.  His intuition and ability to quickly get to the heart of a complicated collection of numbers served me well, despite my writing a thesis in a field that is not his specialty.  I owe a great debt to Bob Nichol for teaching me, on a daily basis, how to start, progress, and complete a scientific project.  Without his ideas and guidance, I would have been a graduate student much longer.  I only hope I can learn how to form the ideas that he does.  Kath Romer has taught me almost all of the X-ray astronomy I know, and good deal of the optical astronomy as well.  She also patiently waded through the horrid language of my first paper, molding it into something that could be read and understood.  Francisco Castander had many long conversations with me about cluster X-ray emission, stellar population synthesis, and all sorts of other fields of astronomy.  I would like to thank Lori Lubin taking the time to repeatedly dig through all of her old notes to answer my obscure questions about the PDCS.  I would also like to thank her for providing very insightful comments on earlier drafts of this dissertation and my other papers.  In addition, I would like to thank Mel Ulmer for scientific and financial support all throughout my graduate career, as well as a fresh perspective on many different scientific problems.  This project was funded in part by NASA grant NAG5-3202.  BH was partially supported by the Center for Astrophysical Research in Antarctica, a National Science Foundation Science and Technology Center, by NASA GO-06838.01-95A and by NSF AST-9256606.  This research was supported through NASA ADP grant NAG5-2432 (at Northwestern University) and NASA LTSA grant NAG5-6548 (at Carnegie Mellon University).  \\newpage"
    },
    "0004/astro-ph0004328_arXiv.txt": {
        "abstract": "For a flat universe presently dominated by smooth energy, either cosmological constant (LCDM) or quintessence (QCDM), we calculate the asymptotic collapsed mass fraction as function of the present ratio of smooth energy to matter energy $\\rat$.  Identifying the normalized collapsed fraction as a conditional probability for habitable galaxies, we observe that the observed present ratio $\\rat \\sim 2$ is likely in LCDM, but more likely in QCDM. Inverse application of Bayes' Theorem makes the Anthropic Principle a predictive scientific principle: the data implies that the prior probability for $\\rat$ must be essentially flat over the anthropically allowed range. Interpreting this prior as a distribution over {\\em theories} lets us predict that any future theory of initial conditions must be indifferent to $\\rat$.  This application of the Anthropic Principle does not demand the existence of other universes. ",
        "introduction": "In the absence of a recognized symmetry principle protecting its value, no theoretical reason for making the cosmological constant zero or small has been found.  Inflation makes the universe appear flat, so that, at present, the vacuum or smooth energy density $\\Omega_{Q0}=1-\\omm < 1$, is $10^{120}$ times smaller than would be expected on current particle theories.  To explain this small but non-vanishing present value, a dynamic vacuum energy, quintessence, has been invoked: a background scalar field whose potential energy dominates its kinetic energy, so that the fluid pressure $P$ and its ratio to energy density $w_Q \\equiv P/\\rho <0$.  (When we speak of a static vacuum energy or cosmological constant, we mean the limiting case, $w_Q=-1$, which is homogeneous on all scales.)  With any positive cosmological constant or quintessence, an expanding universe starts out radiation or matter dominated, but ultimately becomes dominated by smooth energy and goes into exponential expansion (Fig. 1).  The best evidence for a flat low-density universe comes from the location (at $l \\sim 200$) of the first Doppler peak in the CBR anisotroy in the combined BOOMERANG-98, MAXIMA-1 and COBE-DMR measurements: $\\omm+\\omq=1.11 \\pm 0.07~^{+0.13}_{-0.12}$ \\citep{Jaffe}.  Supporting evidence \\citep{WCOS,RHOR} comes from the slow evolution of rich clusters, the mass power spectrum, the curvature in the SNIa Hubble diagram, and the dynamic age.  The cosmic flow implies $\\omm=0.3\\pm0.05$.  The height of the first Doppler peak, and gravitational lensing imply $\\omq=1-\\omm \\sim 2/3$.  Of these, the SNIa evidence is most subject to systematic errors due to precursor intrinsic evolution.  A large set of such observational data have been combined \\citep{RHOR} in a two-step constrained fit. Firstly, ten independent constraints in the ($\\omm,\\omlam$)-plane yielded the result $\\omm + \\omlam = 0.99 \\pm 0.14$, which clearly supports the view of a flat universe. Secondly, assuming exact flatness, five more constraints were included in the fit with the result $\\omm = 1 - \\omlam = 0.33 \\pm 0.04$, or equivalently, $\\rat=\\omq/\\omm = 2.03\\pm 0.25$. We can interpret this as evidence that we live in a low-density universe with a smooth energy component with present density $\\omq \\sim 2/3$ and negative pressure $ -1 \\leq w_Q < -1/3 $.  Accepting this small but non-vanishing value for static or dynamic vacuum energy, a flat Friedmann cosmology (CDM) is characterized by $\\omm,~\\Omega_{Q0}=1-\\omm$ or the present ratio $$\\rat \\equiv \\Omega_{Q0}/\\omm=(1-\\omm)/\\omm~, $$ and by the equation of state for the smooth energy component. The {\\em Cosmic Coincidence} problem now becomes pressing: Why do we live when the clustered matter density $\\Omega(a)$, which is diluting as $a^{-3}$ with cosmic scale $a$, is just now comparable to the static vacuum energy or present value of the smooth energy i.e. when the ratio $\\rat \\sim 2$ ?  In this paper, we study the {quintessence} range $ -1 \\le w_Q < -1/3 $ for the smooth energy component, distinguishing in particular the two cases \\\\ LCDM: cosmological constant with $w_Q=-1$, and\\\\ QCDM: quintessence with the specific choice $ w_Q =-1/2$ .  The next section compares the cosmic expansion and the freeze-out of structure formation, in these two models for smooth energy. Section 3 extends to QCDM the calculation of asymptotic mass fraction as function of a hypothetical continuous variable $\\omm$ for LCDM, presented by Martel {\\it et al} \\citep{MSW,MS}.  In Section 4, identifying these collapsed mass fractions with anthropic probabilities for $\\rat$, we show that the presently observed ratio, while reasonable in an LCDM universe, is more likely in a QCDM universe.  This confirms empirically that the prior probability for our universe is flat in $\\omm$, as is expected in a large class of fundamental theories \\citep{W}.  The anthropic answer to the cosmic coincidence problem is: ``If not now, then when?'' \\citep{Hil}    \\section {Expansion of a Low Density Flat Universe}  The Friedmann equation in a flat universe with clustered matter and smooth energy density is $$ H^2(x) \\equiv (\\dot{a}/a)^2=(8 \\pi G/3)(\\rho_m+\\rho_Q), $$ or, in units of $\\rho_{cr}(x)=3H^2(x)/8\\pi G$, $$1=\\Omega_m(x)+\\Omega_Q(x),$$ where the reciprocal scale factor $x \\equiv a_0/a \\equiv 1+z \\rightarrow \\infty$ in the far past, $\\rightarrow 0$ in the far future.  With the effective equation of state $w \\equiv P/\\rho=$ constant, different kinds of energy density dilute at different rates $\\rho \\sim a^{-n},~n \\equiv 3(1+w)$, and contribute to the deceleration at different rates $(1+3w)/2$ shown in the table:\\\\ \\begin{table} \\centering \\begin{tabular*}{115mm}{@{\\extracolsep{\\fill}}l|ccc@{}}  \\hline  {\\em substance}              &{\\em w}       &{\\em n}       &{\\em (1+3w)/2}   \\\\  \\hline radiation      &  1/3   &  4     &1    \\\\ NR matter      &   0    &  3     &1/2      \\\\ quintessence   & -1/2   &  3/2   &-1/4   \\\\ cosmolconst    & -1     &  0     &-1     \\\\ \\hline \\end{tabular*} \\caption{Energy Dilution for Various Equations of State} \\end{table}\\\\ The expansion rate in present Hubble units is $$ H(x)/H_0=(\\omm x^3+ (1-\\omm) x^n_Q)^{1/2}. $$ The Friedmann equation has an unstable fixed point in the far past and a stable attractor in the far future.  (Note the tacit application of the anthropic principle: Why does our universe expand, rather than contract?)  \\begin{figure} \\begin{center} \\epsfig{file=RecpresPlot.eps,width=12cm,height=13cm} \\caption{Scale evolution of LCDM and QCDM low-density flat universes in the recent past and near future. The lower curve shows the SCDM universe from which both LCDM and QCDM evolved in the far past. The upper curve shows the flat de Sitter universe towards which both LCDM and QCDM will evolve in the future. The inflection points marked (O) show where first LCDM and later QCDM change over from decelerating to accelerating universes.} \\end{center} \\end{figure} \\nopagebreak The second Friedmann equation is $$q(a)\\equiv -\\ddot{a} a/{\\dot{a}^2}=\\sum_{i} \\Omega_i (1+3w_i)/2= (1+\\Omega_Q(a))/2 .$$ The ratio of smooth energy to matter energy, $\\mathcal R(a)=\\mathcal R_0 (a_0/a)^{3w_Q}$, increases as the cosmic expansion dilutes the matter density. A flat universe, characterized by $\\rat,~w_Q$, evolves out of an SCDM universe in the remote past towards a flat de Sitter universe in the future. As shown by the inflection points (O) on the middle curves of Figure 1, for fixed $\\rat$, QCDM  expands faster than LCDM, but begins accelerating only at the present epoch.  The top and bottom curves refer respectively to a de Sitter universe ($\\Omega_m=0$), which is always accelerating, and an SCDM universe ($\\Omega_m=1$), which is always decelerating.  The matter-smooth energy transition (``freeze-out'') $\\Omega_Q/\\Omega_m=1$ took place only recently at $(x^*)^{-w_Q}=\\rat^{-1/3}$ or at $x^*=\\rat^{2/3}=1.59$ for QCDM and, even later, at $x^*=1.26$ for LCDM. Because, for the same value of $\\rat$, a matter-QCDM transition would take place earlier and more slowly than a matter-LCDM transition, it imposes a stronger constraint on structure evolution. As summarized in the table below, quintessence dominance begins 3.6 Gyr earlier and more gradually than cosmological constant dominance. (In this table, the deceleration $q(x) \\equiv -\\ddot{a}/aH_0^2$ is measured in {\\em present} Hubble units.)  The recent lookback time is $$H_0t_L(z)=z-(1+q_0)z^2+...,\\quad z<1 ,$$ where $q_0=0$ for QCDM and $q_0=- 1/2$ for LCDM.   \\begin{table} \\centering \\noindent \\begin{tabular*}{125mm}{@{\\extracolsep{\\fill}}l|cc@{}}  \\hline {\\em event}                 &  {\\em LCDM}      &  {\\em QCDM} \\\\ \\hline \\hline {\\bf Cross-Over to Smooth Energy Dominance}&               &             \\\\ reciprocal scale $x^*=a_0/a=1+z$                &$\\rat^{1/3}$=1.260           &$\\rat ^{2/3}$=1.587     \\\\ age   $t(x^*)/H_0^{-1}$    &0.720             &0.478           \\\\ in units $h_{65}^{-1}$Gyr      &10.8              &7.2            \\\\ \\hline horizon size in units $cH_0^{-1}$   &2.39              &1.58           \\\\ in units $h_{65}^{-1}$Gpc      &11.0              &7.24            \\\\ \\hline deceleration $q(x^*)$ at freeze-out               &-0.333            &0.333 \\\\  \\hline \\hline {\\bf Present Epoch}         &                  &              \\\\ age $t_0/H_0^{-1}$         &0.936             &0.845          \\\\ $h_{65}^{-1}$Gyr      &14.0              &12.7           \\\\ \\hline horizon in units $cH_0^{-1}$                      &3.26              &2.96          \\\\ in units $h_{65}^{-1}$Gpc &15.0              &13.6            \\\\ \\hline present deceleration $q_0$  &-0.500            &0              \\\\ \\hline   \\end{tabular*} \\caption{ Comparative Evolution of LCDM and QCDM Universes} \\end{table} ",
        "conclusions": ""
    },
    "0004/astro-ph0004164_arXiv.txt": {
        "abstract": "The Transient Gamma Ray Spectrometer (TGRS) on board the {\\em WIND\\/} spacecraft has spent most of the interval 1995--1997 in a high-altitude orbit where $\\gamma$-ray backgrounds are low. Its high-resolution Ge spectrometer is thus able to detect weak lines which are slightly offset from stronger background features.  One such line is predicted from nucleosynthesis in classical novae, where $\\beta$-decays on a time-scale of a few hours in an expanding envelope produce positrons that annihilate to generate a line which is blueshifted by a few keV away from the background annihilation line at 511 keV.  The broad TGRS field of view contained five known Galactic novae during 1995 January -- 1997 June, and we have searched the spectra taken around the times of these events for the blueshifted nova annihilation line.  Although no definite detections were made, the method is shown to be sensitive enough to detect novae occurring on ONeMg-rich white dwarfs out to about 2.5 kpc. ",
        "introduction": "In the standard thermonuclear runaway model of classical nova outbursts the energy source is the explosive burning of H in a degenerate layer on the surface of a white dwarf, composed of H-rich material accreted from a companion contaminated by material taken up from the white dwarf by diffusion.  The nuclear burning time-scale is rapid compared to the lifetimes of certain key $\\beta$-unstable isotopes in the reaction chain, and substantial abundances of these nuclei are therefore built up. The isotopes involved may be $^{7}$Be, $^{13}$N, $^{14,15}$O, $^{17,18}$F, $^{22}$Na or $^{26}$Al, depending on the temperature and the chemical composition of the contaminating material from the white dwarf (CO-rich or ONeMg-rich).  Convection in the accreted layer carries the longer-lived ($>100$ s) of these species to the surface before they decay.  The $\\gamma$-ray lines resulting from decays at the surface are in principle observable, and would provide a rather direct test of the model.  The $\\beta$-unstable nuclei involved are all proton-rich and therefore decay predominantly by emission of positrons.  Thus in addition to $\\gamma$-ray lines from de-excitation of the daughter nuclei, we also expect a line at 511 keV from the annihilation of the positrons with the ambient material (Clayton \\& Hoyle 1974, Leising \\& Clayton 1987).  Several previous experiments have searched for the decay lines from the relatively long-lived isotopes $^{22}$Na and $^{7}$Be, either from individual nearby novae (Leising et al. 1988, Iyudin et al. 1995) or from the accumulated production of many novae in a wide field of view (Leising et al. 1988, Harris, Leising \\& Share 1991, Harris et al. 1997).  Here we report first results of the first search for the 511 keV line, which reflects primarily the production of shorter-lived isotopes such as $^{13}$N and $^{18}$F.  The Transient Gamma Ray Spectrometer experiment (TGRS) is very well suited to the search for the 511 keV line, for several reasons.  First, it is located on board a spacecraft whose orbit is so elliptical that it has spent virtually all of its mission (1994 November--present) in interplanetary space (the {\\em WIND\\/} mission).  In this environment the $\\gamma$-ray background level is relatively low, and interruptions due to passages of Earth's trapped radiation belts are minimal (typically lasting a few hours at intervals of several months). This is important because the bursts of 511 keV radiation are predicted to occur only during a $\\sim 6$ hr period at an uncertain interval before the nova event (Leising \\& Clayton 1987, Hernanz et al. 1997, and \\S 2). Second, the TGRS instrument is attached to the south-facing surface of the rotating cylindrical {\\em WIND\\/} body, which points permanently toward the south ecliptic pole.  The detector is unshielded, and TGRS therefore has an unobstructed view of the entire southern (ecliptic) hemisphere.  Third, and most importantly, the TGRS Ge detector has sufficient spectral resolution to detect a 511 keV line which is slightly Doppler-shifted away from the background 511 keV line which is always present in space experiments, which arises from annihilation following $\\beta^{+}$ decays of cosmic-ray spallation product nuclei.  This background line is always very close to the rest energy, whereas the most recent nova models predict the source line to be blueshifted by 2--5 keV (Hernanz 1997, private communication); by comparison, the TGRS energy resolution at 511 keV has varied from 3--4 keV FWHM during the mission (Harris et al. 1998). No previous experiment has possessed the same advantages of long and continuous temporal coverage, broad spatial coverage, and fine spectral resolution.  Our analysis procedure (described in \\S 3) relies heavily on the theoretically predicted properties of the 511 keV line.  There has not been much incentive for theoretical work upon this line because of the difficulty of resolving the nova line from the background line in the previous generation of low-resolution experiments. The early work of Clayton \\& Hoyle (1974) and Leising \\& Clayton (1987) assumed white dwarf models which were CO-rich, and treated the explosion parametrically. Much fuller hydrodynamic models which can be applied at all stages from accretion through explosion and nucleosynthesis have recently been developed (Starrfield et al. 1992, Hernanz et al. 1996, Jose et al. 1997) and applied to the positron annihilation problem (Hernanz et al. 1997, G\\'{o}mez-Gomar et al. 1998).  In this series of models the largest 511 keV line fluxes come from the more massive ONeMg-rich white dwarfs; indeed, fluxes as high as $10^{-2}$ photon cm$^{-2}$ s$^{-2}$ are expected in favorable cases. For the purpose of analysis we shall in general assume the line properties of the Hernanz et al. (1997) ONe2 model, which generates this level of flux over a $\\sim 6$ hr period (\\S 3).\\footnote{ This model is referred to as A2 in G\\'{o}mez-Gomar et al. (1998).}  In \\S 2 we describe the difficulties which are involved in a search for the 511 keV line around the times of those novae which are known to have been in TGRS's field of view during the mission.  Our results for the individual novae are presented in \\S 4, and we will show in \\S 5 that for the ONe2 and other nova models TGRS sensitivity is good enough to justify a large-scale search for the line in the full 1995--1997 TGRS data. ",
        "conclusions": "\\subsection{Results for Individual Novae}  As emphasized in \\S 2, it is very difficult to use our results for the individual novae to constrain the ONe2 model because of our lack of knowledge of two related nova parameters, the epoch of explosion and the distance.  In Figs. 4c and 7--10 we use the limited information given in Table 2 to plot the predicted fluxes $\\phi_{pred}$ as dotted lines, which depend on the time relative to discovery according to \\begin{equation} \\phi_{pred}(t) = 0.016/r(t)^2 \\end{equation} where the distance $r(t)$ in kpc is obtained from Equation (3). Note that $\\phi_{pred}$ is {\\em not\\/} a light curve; as shown in Table 1, there is usually a wide range of allowed explosion times, and $\\phi_{pred}$ is the predicted flux {\\em at any one time\\/} of explosion.  In two cases (N Sco 1997 and N Sgr 1996; Figs. 4c and 9a) the predicted flux is always below our $3 \\sigma$ upper limit, for any of the allowed explosion times; these novae must be too distant for our results to be used to constrain even the most optimistic nova model (Hernanz et al. 1997, model ONe2).  The other three novae (N Cir 1995, N Cen 1995 and N Cru 1996) are predicted to yield a detectable $\\phi_{pred}$ only if the explosion occurred early in the allowed range of times [prior to $\\sim 12$ d before discovery for N Cir (Fig. 7), 7 d before discovery for N Cen (Fig. 8), and 9 d before discovery for N Cru (Fig. 10)].  In other words, the allowed range of distances for these novae includes the TGRS detection limit $\\sim 2.5$ kpc.  The ONe2 model can be tested using our results for these novae if further investigations yield either the appropriate constraints on explosion time, or an independent measurement of the distance to the nova.  \\subsection{Implications for Nova Searches}  The data analyzed in this paper covers about 70 d, representing only a small fraction of the whole TGRS database.  An obvious extension of this work would be to analyze the entire database in order to detect, or place upper limits upon, the occurrence of all neon novae lying within the 2.5 kpc radius estimated in \\S 4.  There are good reasons for believing that many novae within this radius escape detection.  As noted in \\S 2, visual searches for novae are usually episodic, containing gaps in temporal coverage within which fast-rising and decaying novae are likely to be missed; coverage is particularly uneven in the southern hemisphere. Further, the southern Galactic plane contains regions of very strong interstellar absorption; $A_{V}$ values up to 4 are found within 2.5 kpc in some directions (Neckel \\& Klare 1980).  There is therefore considerable uncertainty, not only about the global Galactic nova rate, but also about its spatial distribution; novae in a disk population may be more common than the spheroidal population modeled by Leising \\& Clayton (1985), but less observable due to obscuration (Hatano et al. 1997).  There is evidence that the neon subclass may be more concentrated toward the plane, and thus more under-represented in global nova statistics (Della Valle et al. 1992).  Estimates of the success rate of a $\\gamma$-ray line search must be very uncertain, for the above reasons; in fact, the best justification for such a search is precisely that it is the only way to obtain an unbiased census of novae on which to base global nova rates.  Estimated rates for the Galaxy as a whole range from 11--260 yr$^{-1}$, tending to cluster about values 35--50 (Shafter 1997), with about one-third occurring on ONeMg white dwarfs (Gehrz et al. 1998).  To estimate TGRS's rate of detections, we note that a full-scale search of almost 3 years of TGRS data will yield a much higher chance detection rate than that found in the present work ($< 1$ detection in 70 d), so that the threshold for detection must be set higher than the value $3 \\sigma$ used here. Since the distribution of our null measurements is very close to the normal distribution expected by chance, it is easy to show that a threshold level $\\simeq 4.6 \\sigma$ yields a probability $< 1$\\% of a single false detection by chance (Abramowitz \\& Stegun 1964). This threshold level corresponds to a distance $\\simeq 2$ kpc for a neon nova.  A simple model of nova distribution following star counts (Bahcall 1986, integrating the disk over $z$) suggests that $\\sim 2$\\% of Galactic novae occur within 2 kpc of the Sun in the southern hemisphere.  Moderate estimates of the global nova rate $\\sim 50$ yr$^{-1}$ then imply that TGRS would detect one neon nova during $\\sim 3$ yr; higher global rates, or a higher fraction of neon novae, would imply more detections.  Further papers in this series will utilize this search strategy, and will follow up the possibilities arising from detailed study of the already known novae as described in the previous section."
    },
    "0004/astro-ph0004352_arXiv.txt": {
        "abstract": "The central densities of dark matter (DM) halos are much lower than predicted in cold DM models of structure formation.  Confirmation that they have cores with a finite central density would allow us to rule out many popular types of collisionless particle as candidates for DM.  Any model that leads to cusped halos (such as cold DM) is already facing serious difficulties on small scales and hot DM models have been excluded.  Here I show that fermionic warm DM is inconsistent with the wide range of phase space densities in the DM halos of well-observed nearby galaxies. ",
        "introduction": "The well-known mass discrepancies in galaxies (Bosma 1978; Broeils 1992; Verheijen 1997) and in clusters of galaxies (Zwicky 1937; Carlberg, Yee \\& Ellingson 1997; Tyson, Kochanski \\& Dell'Antonio 1998) are usually taken to imply the existence of a large fraction of invisible, or ``dark,'' matter (DM) in the universe.  A popular candidate is Cold Dark Matter, for which the DM particles, whatever they are, are almost at rest with respect to the Hubble flow in the early universe.  CDM is often imagined to be a heavy, non-baryonic relic particle from the early universe which has essentially only gravitational interactions with itself and with normal, or baryonic, matter.  The CDM model has been studied intensively for twenty years and now has many well-worked out predictions for the formation of structure in the universe (\\eg\\ Bertschinger 1998).  The broad-brush impression one now has is that the currently favored $\\Lambda$CDM model boasts a considerable degree of success in predicting large-scale structure (\\eg\\ Pearce \\etal\\ 1999; Bahcall \\etal\\ 1999). But it has become apparent in recent years that the predictions of almost any flavor of CDM are seriously at variance with the observed properties of galaxies because the central densities of collapsed objects and fragments are predicted to be too high (\\S\\ref{sec:cusps}).  A number of authors have therefore begun to explore variations of the CDM model; most favor a modification to the properties of the DM particles rather than the alternative -- a change to the law of gravity.  The simplest is the warm DM matter model (\\eg\\ Colombi, Dodelson \\& Widrow 1996; Sommer-Larsen \\& Dolgov 1999; Hogan 1999) in which streaming of the DM particles in the early universe suppresses small-scale power in the fluctuation spectrum.  In addition, WDM particles in halos of greater volume density must have larger velocity spreads, because of Liouville's theorem, thereby precluding strong density gradients.  Simulators of the WDM model suppress small-scale power in the fluctuation spectrum but generally ignore the initial finite velocity spread, which is difficult to include without wrecking the quiet start.  Since the DM in these simulations still has infinite phase space density, the resulting halo profiles have density cusps resembling those which form in CDM (\\eg\\ Moore \\etal\\ 1999; Col\\'\\i n \\etal\\ 2000). ",
        "conclusions": "\\label{sec:concl} Fermionic WDM halos should have well-defined cores with a characteristic $\\fm$. This prediction is inconsistent with the range of values shown in Figure 1 which rules out this particular DM candidate.  Halos having finite central densities are also inconsistent with the density cusps predicted by CDM, and possibly also bosonic WDM.  Other work has excluded HDM.  Thus the existence of finite central density cores in halos would rule out most forms of simple, collisionless relic particle as DM candidates.  It is therefore of great importance to obtain tighter observational constraints on the inner density profiles of halos.  It is suddenly popular to hypothesize extra properties for the CDM particle in order to soften the collapsed halos (Spergel \\& Steinhardt 1999; Peebles 2000; Goodman 2000; Hu \\etal\\ 2000; Kaplinghat \\etal\\ 2000; \\etc).  The simple CDM model would become much less attractive if an extra {\\it ad hoc\\/} property were needed to rescue it.  The nature of DM is increasingly constrained by the observed properties of galaxy halos.  In particular, any successful theory of galaxy formation will need to account for the wide range of phase space densities for DM shown in Figure 1."
    },
    "0004/astro-ph0004214_arXiv.txt": {
        "abstract": "We use numerical simulations to investigate the cusp at the centre of elliptical galaxies, due to the slow growth of a super-massive black hole. We study this problem for axisymmetric models of galaxies, with or without rotation. The numerical simulations are based on the `Perturbation Particles' method, and use GRAPEs to compute the force due to the cusp. We study how the density cusp is affected by the initial flattening of the model, as well as the role played by initial rotation. The logarithmic slope of the density cusp is found to be very much insensitive to flattening; as a consequence, we deduce that tangential velocity anisotropy-- which supports the flattening-- is also of little influence on the final cusp. We investigate {\\it via} two different kinds of rotating models the efficiency with which a rotation velocity component builds within the cusp. A cusp in rotation develops only for models where a net rotation component is initially present at high energy levels. The eventual observation of a central rotational velocity peak in E galaxies has therefore some implications for the galaxy dynamical history. ",
        "introduction": "\\label{sec_intro} \\indent  HST observations have shown that a central density cusp is present in most, if not all, elliptical galaxies (Lauer {\\it et al.} 1995, and companion papers). Furthermore the slope of the density cusp has shown a dichotomy between  high mass (luminosity) systems and low mass (luminosity) ones, analogous to what has been found for other observational properties (Bender {\\it et al.} 1989). Low mass ellipticals tend to host steep cusps, the radial profile of their luminosity $\\rho(r)$ having a logarithmic slope $\\gamma\\equiv - {\\rm d } log(\\rho)/{\\rm d} log(r)  \\simeq 1.9 \\pm 0.5 $. On the other hand, luminous ellipticals have shallower cusps, with $\\gamma \\simeq 0.8 \\pm 0.5$ (Gebhard \\etal 1996).  The most popular explanation for the formation of a density cusp is a central super-massive black hole (BH)-- see Richstone {\\it et al.} (1998), Kormendy \\& Richstone (1995) for reviews. Such BHs are believed to be present in a large fraction (possibly close to 1, Haenelt \\& Rees 1993, Tremaine 1997) of present day galaxies. Their influence on a collisionless galactic nucleus was first addressed by Peebles (1972), Young (1980), Goodman \\& Binney (1984) using semi-analytic models.  All these models suppose an isolated galaxy, at the center of which a BH grows by gas accretion. The slope of the cusp produced in such models is within the range observed for the less massive E galaxies. In this paper we shall also consider the formation of a cusp in an isolated galaxy, as we are interested in the origin for the steep cusp ($\\gamma \\simeq 1.9$) of low luminosity ellipticals. On the other hand, numerical simulations (Makino \\& Ebisuzaki, 1996, Quinlan \\& Hernquist 1997, Nakano \\& Makino 1999) have shown that a BH sinking within the core of an elliptical galaxy, or a binary of BHs produce a shallow cusp, similar to that observed for massive E galaxies. Such a BH binary could result from a merger. This picture is therefore compatible with the general belief (e.g. Nieto \\& Bender 1989) that the apparent existence of two classes of ellipticals is related to a more important role played by merging and interactions in the history of massive E galaxies.  All the semi-analytic models of cusp formation due to a central BH use the adiabatic invariance of actions to derive the distribution function after a BH has slowly grown (see Young 1980). Since actions are explicitly known for spherical systems, analytical models have been derived in spherical symmetry. Based on such quasi-analytic computations for spherical models, Quinlan \\etal (1995) investigated the consequences of the initial density profile on the slope of the cusp, and addressed the influence of a velocity anisotropy.  For non-spherical systems -- except for St\\\"ackel potentials -- an expression for the actions is unknown, and therefore models based on the conservation of actions have not been computed hitherto. In this paper, we propose to use numerical simulations with high central resolution, based on the PP method (Leeuwin \\etal 1993), to investigate the cusp formed in axisymmetric systems by the growth of a BH.  Many low luminosity galaxies display little evidence for triaxiality, but are close to axisymmetry; they can be modeled as two-integral models, with $f(E,L_z)$. Such a distribution function (d.f.) implies velocity dispersions obeying $\\sigma_R=\\sigma_z$, meaning that flattening would be due to an excess of azimuthal motion (which increases the net rotation, or the tangential anisotropy). This does not preclude that three-integral models would not do as well, or better ({\\it e.g.} review by Merritt, 1998). Many of those galaxies however exhibit kinematic features consistent with isotropic rotators models: at least for them, an excess of tangential motion is plausibly the main support for their flattening.  The BH itself may have perturbed orbits sufficiently to erase the part of memory for initial conditions that corresponds to conservation of a third integral (Norman \\etal 1985, Gerhard \\& Binney 1985, Merritt \\& Quinlan 1998). In this case axisymmetry is a consequence of the BH growth. Nevertheless, one still ought to investigate the scenario in which, for some of these galaxies, the presently observed axisymmetry already existed before the BH growth.  This is the goal of this paper, where we will study the cusp generated by the growth of a central BH within axisymmetric systems having various degrees of rotation. If a BH grows in a roughly axisymmetric galaxy, do the initial flattening and rotation have a sizeable effect on the evolution, so that it deviates from that of the well known spherical case? Could we in certain cases be able, eventually with more detailed observational data, to infer from the properties of presently observed cusps any information about the pre-black hole galaxy? Should one be cautious in certain cases about deriving the BH mass using the spherical adiabatic model?  This paper is organized as follows: first, the existing models are briefly recalled in section 2. We then specify the initial conditions (section 3) and the numerical techniques used for this work (section 4). Results for non-rotating and rotating models are given in sections 5 and 6, respectively. We conclude in section 7. ",
        "conclusions": ""
    },
    "0004/astro-ph0004028_arXiv.txt": {
        "abstract": "{early universe --- galaxies: evolution --- galaxies: morphology -- galaxies: stellar content --- infrared: galaxies}  The Hubble Deep Fields continue to be a valuable resource for studying the distant universe, particularly at $z > 2$ where their comoving volume becomes large enough to encompass several hundred $L^\\ast$ galaxies or their progenitors.  Here I present recent results from a near--infrared imaging survey of the HDF--North with NICMOS, which provides structural and photometric information in the optical rest frame ($\\lambda\\lambda_0$4000--5500\\AA) for hundreds of ordinary galaxies at $2 < z < 3$, and which offers the means to search for still more distant objects at $z \\gg 5$.  Lyman break galaxies at $2 < z < 3$ are compact and often irregular in the NICMOS images;   ordinary Hubble sequence spirals and ellipticals seem to be largely absent at these redshifts, and apparently reached maturity at $1 < z < 2$.   The Lyman break galaxies have UV--optical spectral energy distributions like those of local starburst galaxies. Population synthesis models suggest typical ages $\\sim$few$\\times 10^8$ years and moderate UV extinction ($\\sim 1.2$~mag at 1700\\AA), but the constraints are fairly weak and there may be considerable variety. Considering a near--IR selected galaxy sample, there is little evidence for a significant number of galaxies at $z \\sim 3$ that have been missed by UV--based Lyman break selection.  Using the well--characterized $z \\sim 3$ galaxy population as a point of reference, I consider Lyman break galaxy candidates at $4.5 < z < 9$, as well as one remarkable object which might (or might not) be at $z > 12$.  The space density of UV--bright galaxies in the HDF appears to thin out toward larger redshifts, although surface brightness selection effects may play an important role. ",
        "introduction": "\\subsection{The first galaxies?}  The past five years have seen remarkable breakthroughs in our ability to identify and systematically study ordinary galaxies at very large redshifts, not just as isolated case studies, but {\\it en masse} as a galaxy {\\it population.}  To date, nearly 1000 galaxies have been spectroscopically confirmed at $z > 2$, mostly identified via broad--band color selection techniques (Steidel \\etal 1996, i.e., the Lyman break galaxies, or LBGs), but with other important and complementary methods also coming into play (sub--mm and radio surveys, emission line searches, QSO absorption systems, etc.).  A broad census of galaxy properties at $z \\sim 3$ now seems within reach, covering star formation rates (SFRs), dust content, morphologies, spatial clustering, and perhaps even chemical abundances and internal kinematics.  Although I have retained my assigned title, `The First Galaxies...,' it is far from clear that we know when, where, or how to find the `first' galaxies.  The $z \\sim 3$ LBGs may or may not be the first major wave of galaxy formation.  If anything, current data favors a roughly constant global SFR (as traced by cosmic UV luminosity density, at least) from $2 \\simlt z \\simlt 4$, with no certain evidence for a decline at higher redshifts (Steidel \\etal 1999).  The sub--mm population detected by SCUBA (cf.\\ Cowie, this volume) may or may not represent the bulk of early star formation, and the upper redshift bound to SCUBA sources remains unknown. The reionization of the intergalactic medium at $z > 5$ and the presence of metals in the Ly~$\\alpha$ forest at $z \\sim 3$ point toward earlier epochs of star and galaxy formation, at least in trace amounts. A handful of galaxies now have plausible spectroscopic confirmation at $z > 5$, but too few for any systematic census.  For this reason, I cannot promise to live up to my title:  I do not know what the first galaxies are or what they look like.  Given this ignorance, I will focus the first part of this article on the structure and stellar populations of the most distant {\\it well studied} galaxies, the $z \\sim 3$ Lyman break objects.  This is not meant as a comprehensive review, but will instead concentrate on new imaging and photometric data from NICMOS on {\\it HST} that extend our knowledge of LBG properties to the optical rest frame.   In his contribution to this volume, Max Pettini provides a complementary discussion of recent efforts to measure chemical abundances and internal kinematics for these same galaxies.  In the second part of my article, I will describe efforts to extend Lyman break color selection to still larger redshifts, approaching (or perhaps even exceeding) $z \\approx 10$.  In this way I hope to at least provide a look into the epoch when `the first galaxies' might plausibly have been formed, and to catalog what we can find right now, in the pre--{\\it NGST} era, given the best available survey data.  \\subsection{Infrared observations of the Hubble Deep Field}  For the past five years, the Hubble Deep Fields (HDFs) have provided the most exquisitely deep, high angular resolution optical census of the distant universe.  It is important (if somewhat pedantic) to consider what an HDF is actually good for.  One WFPC2 field covers 5 arcmin$^2$, and probes a very small co--moving volume at $z < 1$, enough to hold only $\\sim 12$--30 $L^\\ast$ galaxies, depending on the cosmology.   Given small number statistics and concerns about clustering, the central HDF is therefore not the best place study massive galaxies in the `low' redshift universe, despite the fact that most of cosmic time and most bright galaxies with spectroscopic redshifts are at $z \\simlt 1$.  There is far more volume at high redshift: 10.5 to 40$\\times$ more at $2 < z < 10$ than at $z < 1$ for plausible cosmologies, room enough for several hundred $L^\\ast$ galaxies or their progenitors.  At $z > 1$, the optical light emitted from galaxies shifts into the near--infrared.  Thus in order to compare $z > 2$ galaxies to their local counterparts, and to search for still more distant objects at $z \\gg 5$, it is important to extend the wavelength baseline.  The HDF--North was observed in the near--IR from the ground in several different programs (e.g., Hogg \\etal 1997; Barger \\etal 1998; and our own KPNO 4m $JHK_s$ imaging, cf.\\ Dickinson 1998).  The depth and angular resolution (typically $\\sim 1\\arcsec$) of these data are a poor match to that of the optical WFPC2 HDF images.  Two programs therefore targeted the HDF--North with NICMOS on board {\\it HST}, providing much deeper images with high angular resolution.   The NICMOS GTOs (Thompson \\etal 1999) imaged one NICMOS Camera 3 field ($\\sim 51\\arcsec\\times 51\\arcsec$) for 49 orbits each at F110W (1.1$\\mu$m) and F160W (1.6$\\mu$m).  Our own program mosaiced the complete HDF with a mean exposure time of 12600s per filter in F110W and F160W.  Sensitivity varies over the field of view, but the average depth is AB~$\\approx 26.1$ at $S/N=10$ in an $0\\secspt7$ diameter aperture.  The drizzled PSF has FWHM~=~$0\\secspt22$, primarily limited by the NIC3 pixel scale.   Because most galaxies have spectral energy distributions (SEDs) which brighten (in $f_\\nu$ units) at redder wavelengths, our images detect roughly half of the galaxies from the WFPC2 HDF,  despite their short exposure times.  We have also reanalyzed our KPNO $K_s$ images to optimally extract photometry matched to the WFPC2+NICMOS data.   These data are not as deep as one would like, which is unfortunate because they provide the only access to rest frame optical wavelengths for objects at $3 < z < 4.4$, but they are the best presently available.  Thanks to the dedicated efforts of the observers, a remarkably high density of spectroscopic redshifts is available in the HDF--N:  $\\sim 150$ galaxies (plus a few stars) in the central WFPC2+NICMOS field alone, with 33 objects at $2 < z < 5.6$.  Taking advantage of the high quality photometric data, many investigators have used multicolor selection (e.g.\\ the two color Lyman break criteria of Steidel \\etal 1996, Madau \\etal 1996, and others) or photometric redshifts (e.g., Fern\\'andez--Soto \\etal 1999) to identify high redshift galaxy candidates.  There are advantages and drawbacks to both approaches, but both have demonstrated remarkable successes.  In this discussion, I will make use of both methods.  For the photometric redshifts, I will use fits to our 7--band HDF photometry by Budav\\'ari \\etal (2000), whose `adaptive template' method is a modification of an otherwise straightforward spectral template fitting scheme.  I will use AB magnitudes here throughout, and notate the six WFPC2 + NICMOS bandpasses by \\U300, \\B450, \\V606, \\I814, \\J110 and \\H160.  Unless stated otherwise, I will assume a cosmology with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$~\\kmsmpc. ",
        "conclusions": "Overall, the HDF/NICMOS data demonstrate both the promise and the challenges which lie ahead for finding and studying the `first' galaxies.   The rest--frame optical view of Lyman break galaxies presented in \\S\\S2 and 3 strongly suggests that the galaxy population at $2 < z < 3$ had not yet achieved maturity.  The giant, Hubble sequence spirals and ellipticals that dominate the high--mass end of the galaxy population today are not seen at $z > 2$. In a sample of HDF galaxies selected in the near--infrared, nearly all galaxies with spectroscopic or plausible photometric redshifts $2 < z < 3.5$ are evidently forming stars quite rapidly and can also be identified via their emitted--frame UV light. The evidence from SCUBA shows that there are occasional `monsters' whose obscured star formation may be quite important to the global emissive energy budget from galaxies.  The identification of these objects, relatively rare by number, remains an important dilemma.  Broad band color selection has been the most successful means for identifying high redshift galaxies, but we seem to be pushing the limits of what can be accomplished at $z > 6$ with present--day capabilities.  The NICMOS HDF images are the deepest near--IR data now available, and they do include plausible candidates for galaxies at $6 < z < 9$, but they are relatively few, and most are quite probably too faint for spectroscopic confirmation.  We probably should not expect to find galaxies much {\\it brighter} than these candidates unless some of the `first' galaxies were significantly more luminous than the boring old `later' galaxies that we have now surveyed extensively at $z \\approx 3$.  This is not impossible of course:  the Chen \\etal object and HDFN--JD1 are both possible (but unconfirmed) $z > 6$ candidates more luminous than any normal Lyman break galaxy at $z < 5$.  Perhaps indeed there are very luminous, relatively unobscured proto--galaxies out there waiting to be found,  a hope that was once quite widespread, but which seems to have gradually faded in the modern era of 25th magnitude LBGs and optically invisible SCUBA sources.  Perhaps it will still make a comeback...  These few rather speculative candidates aside, the evidence from the HDF alone would suggest that the population of UV--bright LBGs may be thinning out at $z > 5$, at least for objects comparable to those at the bright end of the $z \\approx 3$ luminosity function. It should be remembered, however, that this was also the conclusion reached by Madau \\etal (1996) at $z > 3.5$, a result that has since been challenged by larger surveys with extensive spectroscopy. It is undoubtedly dangerous to draw conclusions too strongly from one 5~arcmin$^2$ field.  However, extending this work to larger areas and more sightlines will be an expensive effort. Surface brightness dimming and limited solid angle coverage may limit our ability to see much more with NICMOS (assuming that it is successfully revived in 2001), and ground--based near--IR imaging may never go deep enough to detect any but the most luminous objects at $z > 5$.   Wider fields imaged with the {\\it HST} WFC3 near--IR channel (coming circa 2004) may offer the best survey opportunity until NGST, but a substantial investment of observing time will be needed to survey adequate solid angles to sufficient depth.  Alternatively, we may turn to other observing strategies, e.g., by taking advantage of gravitational lensing from foreground galaxy clusters to boost very distant objects to detectable magnitudes. Narrow band and blind multislit emission line searches are being carried out through airglow windows (e.g., at $\\lambda 9150$\\AA, corresponding to $z \\approx 6.5$;  cf.\\ Crampton \\& Lilly 1999; Stockton 1999).  Or perhaps concerted efforts to identify SCUBA sources will indeed turn up objects at $z \\gg 5$, where the advantage of the negative sub--mm $k$--correction is enormous.  These data are offering a first glimpse into the so--called `dark ages,' and giving hope that there may be luminous things there to find and study.  In some sense, we may not know that we've found the first galaxies until we can find no more beyond them.  Holding to that standard will ensure that a more challenging (and hence, more rewarding) threshold of proof always lies ahead."
    },
    "0004/astro-ph0004391_arXiv.txt": {
        "abstract": "We present a new numerical code which is designed to derive a spectral energy distribution (SED) for an arbitrary spatial distribution of stellar and gaseous components in a dusty starburst galaxy. We apply a  ray tracing method to  numerical simulations and thereby  estimate extinction and  reemission of stellar light by dusty gas in an explicitly self-consistent manner. By using this code, we can investigate simultaneously dynamical and photometric evolution of a dusty galaxy based on stellar and gaseous dynamical simulations. As an example, we demonstrate when and how a major galaxy merger with dusty starburst becomes an ultra-luminous infrared galaxy owing to strong internal dust extinction. We furthermore discuss advantages and disadvantages of the present new code in clarifying the nature and the origin of low and high redshift dusty starburst galaxies. ",
        "introduction": "Recent observational studies have discovered possible dusty starburst candidates in various classes of galactic objects such as low z ultra-luminous infrared galaxies (ULIRGs) (e.g., Sanders et al. 1988;  Sanders \\& Mirabel 1996), intermediate z ones (e.g., Tran et al. 1999), faint submillimeter sources detected by  the Sub-millimeter Common-User Bolometer Array (SCUBA) (Holland et al. 1999) on the James Clerk Maxwell Telescope (Smail, Ivison, \\& Blain 1997; Hughes et al. 1998; Barger et al. 1998; Smail et al. 1998; Ivison et al. 1999; Lilly et al.  1999), extremely red objects (EROs) (Elston, Rieke, \\& Rieke 1988; Graham \\& Dey 1996; Cimatti et al. 1998; Dey et al. 1999; Smail  et al. 1999), optically faint radio sources detected by the Very Large Array (VLA) (e.g., Richards et al. 1999), Lyman-break galaxies (Steidel et al. 1996; Lowenthal et al. 1997), and near-infrared emission line galaxies (e.g., Mannucci et al. 1998). These dusty starburst candidates are generally considered to provide valuable information on the formation and the evolution of galaxies, accordingly the origin and the nature of these galaxies have been extensively discussed in variously different contexts. These discussions include, for example, the importance of dissipative dynamics in the formation of elliptical galaxies at low and high redshift (Kormendy \\& Sanders 1992; Barnes \\& Hernquist 1992), an evolutionary link between starburst and active galactic nuclei (Norman \\& Scoville 1988), fueling of dusty gas to the central $\\sim$ 100 pc for nuclear starburst in major galaxy mergers (e.g., Mihos \\& Hernquist 1996), physical correlations between morphological and photometric properties in ULIRGs (Bekki, Shioya, \\& Tanaka 1999; Bekki \\& Shioya 2000a), merging and clustering process of high $z$ dusty mergers within a hierarchical clustering scenario (e.g., Somerville \\& Primack 1998), and cosmic star formation history (e.g., Pei \\& Fall 1995; Madau et al. 1996; Meurer et al. 1997; Madau, Pozzetti, \\& Dickinson 1998; Pascarelle, Lanzetta, \\& Fernandez-Soto 1998; Blain et al. 1999; Steidel et al. 1999). One of important problems in extra-galactic astronomy is to settle the above discussions on high redshift  dusty starburst galaxies, though observational analysis still has some difficulties in determining  precisely  redshifts (Sanders 1999), identifying  optical counterparts (Richards 1999), and estimating the degree of dust extinction (e.g., Meurer, Heckman, \\& Calzetti 1999) for high $z$ dusty starburst galaxies.    Spectral energy distributions  (SEDs) are generally considered to be one of essentially important factors that can determine the nature of dusty starburst galaxies. Accordingly several theoretical attempts have been made to derive the SEDs of galaxies with dusty starburst regions. Rowan-Robinson \\& Crawford (1989) discussed how the far-infrared spectra of a sample of galaxies observed by $Infrared$ $Astronomical$ $Satellite$ (IRAS) can be reproduced by a theoretical model with  a given temperature of stars embedded by dust, the optical depth, and the density distribution (or geometry) of dust. Witt et al. (1991) investigated transfer processes of stellar light for models with variously different spherical geometries of dust with  special emphasis on the effects of scattered light  on photometric properties of galaxies and SEDs. Calzetti, Kinney, \\& Storchi-Bergmann (1994) analyzed ultraviolet (UV) and optical spectra of 39 starburst and blue compact galaxies and thereby provided an analytical formulation of the effects of dust extinction in galaxies and an effective extinction law for correcting the observed UV and optical spectral continua. Witt \\& Gordon (1996) investigated the radiative transfer processes of a central stellar source surrounded by a spherical, statistically homogeneous but clumpy two-phase scattering medium and found that the structure of dusty medium can greatly affect the conversion of UV, optical, and near-infrared radiation into thermal far-IR dust radiation in a dusty system. By comparing the observed SEDs of 30 starburst galaxies with theoretical radiative transfer models  of dusty systems, Gordon, Calzetti, and Witt (1997) discussed the importance of geometry of stellar and gaseous components in determining the SED of a dusty starburst galaxy. Takagi, Arimoto, \\& Vansevi$\\rm \\check{c}$ius (1999) investigated radiative transfer models with variously different ages of secondary starburst components and optical depths for dusty starburst galaxies and proposed a new method for estimating  precisely the effect of ages of young starburst populations  and that of dust attenuation on the shape of SED in UV and near$-$infrared bands. Efstathiou, Rowan-Robinson, \\& Siebenmorgen (2000) treated starburst galaxies as an ensemble of optically thick giant molecular clouds (GMCs) centrally illuminated by recently formed stars and thereby constructed a new radiative transfer model for calculating SEDs from UV to millimeter band of dusty starburst galaxies. By using this model, they discussed how the age and the star formation history of a dusty starburst galaxy can control the SED particularly for the prototypical starburst galaxy M82 and NGC 6090.    Although the above previous studies have succeeded in clarifying important dependences of SEDs on physical parameters such as spatial distribution of dust, geometries of stellar and gaseous components, and dust properties in  dusty starburst galaxies, many of them  did not discuss $the$ $time$ $evolution$ $of$ $SEDs$. Main reasons for $some$ previous theoretical studies' not discussing the time evolution of SEDs are the following three. Firstly previous models did not follow the time evolution of stellar and gaseous distribution and accordingly could not derive the time evolution of SEDs in dusty starburst galaxies. Secondly,  hydro-dynamical evolution of interstellar gas (e.g., time evolution of gaseous density) was not included in previous studies, and accordingly the time evolution of optical depth could  not be derived. Thirdly, some previous studies did not consider chemical evolution  of  dusty interstellar medium and that of stellar components, they could not follow  time evolution of metallicity and that of dust properties. Considering that low z infrared luminous galaxies with dusty starburst and high $z$ faint SCUBA sources with possible dusty starburst show very peculiar morphology (Sanders et al. 1988;  Sanders \\& Mirabel 1996; Smail et al. 1998), spherical symmetric approximation  or   axisymmetric one adopted in previous theoretical and numerical studies for  the stellar and gaseous distribution of a dusty galaxy should be also relaxed in order that the nature and the origin of low and high z dusty galaxies can be  extensively investigated by theoretical studies.   The purpose of the present paper and our future papers is to investigate in detail the formation and the evolution of dusty starburst galaxies by using a new numerical code for deriving the time evolution of SEDs of these galaxies. We first perform numerical simulations that can follow dynamical evolution of stellar and gaseous components, star formation history, and chemical evolution for dusty starburst galaxies in an explicitly self-consistent manner. We then derive stellar and gaseous distribution and age and metallicity distribution of stellar populations and thereby calculate galactic SEDs. Furthermore  we  describe how the present  numerical code is useful and helpful for investigating variously different physical properties of low and high z dusty galaxies. As an example,  we here present the results of a major merger with dusty interstellar gas. We particularly demonstrate how dynamical evolution of stellar and gaseous components controls the time evolution of SEDs in dusty galaxy mergers  and thus their photometric evolution.    The layout of this paper is as follows. In \\S    2, we summarize  numerical models used in the present study and describe   in detail the methods for deriving the SEDs corrected by internal dust extinction. In this section, we also point out the limitations of the present numerical code. In \\S 3, we present numerical results on the time evolution of morphology, SED, and photometric properties in a gas-rich major merger. In \\S 4, we discuss the origin of high $z$ dusty starburst galaxies such as faint SCUBA sources and EROs. The conclusions of the preset study are  given in \\S 5. ",
        "conclusions": "We present a new code which is designed to derive a SED for an arbitrary spatial distribution of stellar and gaseous components for a dusty starburst galaxy. By using this code, we can calculate SEDs based on numerical simulations that can analyze simultaneously dynamical and chemical evolution, structural and kinematical properties, morphology, star formation history, and transfer of metals and dust in interstellar medium for a starburst galaxy. Accordingly, we can investigate variously different properties of starburst galaxies, such as effects of dynamical evolution on galactic SEDs, physical correlations between morphology and SEDs, photometric evolution from UV to submillimeter  wavelength, two dimensional distribution of $A_{\\rm V}$, and dependence of SEDs on line-of-site of an observer. Thanks to this code, we can furthermore try to clarify the origin of possible  candidates of starburst galaxies, such as low $z$ ULIRGs (e.g., Sanders et al. 1988;  Sanders \\& Mirabel 1996), intermediate z ones (e.g., Tran et al. 1999), faint SCUBA sources (Smail, Ivison, \\& Blain 1997; Barger et al. 1998; Hughes et al. 1998; Smail et al. 1998; Ivison et al. 1999; Lilly et al.  1999), EROs (Elston, Rieke, \\& Rieke 1988; Dey et al. 1999; Smail  et al. 1999), optically faint radii sources detected by VLA (e.g., Richards et al. 1999), Lyman-break galaxies (Steidel et al. 1996; Lowenthal et al. 1997), and emission line galaxies (e.g., Mannucci et al. 1998). By using a new  code developed in the present study, we try to answer the following seven questions in our  forthcoming papers: (1)  When and how a gas-rich major galaxy merger becomes an ULIRG during the dynamical evolution of the merger ? (2) How high z faint SCUBA sources with dusty starburst form and evolve? (3) What are physical conditions for high z dusty galaxies to become EROs ? (4) Are there any evolutionary links between high z possible dusty starburst galaxies, such as faint SCUBA sources, EROs, optically faint radio sources recently detected by VLA, emission line galaxies, and Lyman-break ones ? (5) When does a forming disk galaxy show the strongest  submillimeter flux? (6) What physical processes can determine the shapes of SEDs observed in Lyman-break galaxies? (7) How dusty interstellar gas affects apparent morphology of intermediate and high $z$ dusty galaxies ? Although this code has some disadvantages in deriving very precise SEDs of dusty starburst galaxies, we believe that this code enables us to grasp some essential ingredients of physical processes related to galaxy formation with starburst at low and high z universe."
    },
    "0004/astro-ph0004358_arXiv.txt": {
        "abstract": "We present new observations of $H\\alpha $ emission in the Hickson Compact Group 18 (HCG 18) obtained with a scanning Fabry-Perot interferometer. The velocity field does not show motions of individual group members but, instead, a complex common velocity field for the whole group.  The gas distribution is very asymmetric with clumps of maximum intensity coinciding with the optically brightest knots. Comparing $H\\alpha $ and HI data we conclude that HCG 18 is not a compact group but instead a large irregular galaxy with several clumps of star formation. ",
        "introduction": "Compact groups of galaxies have been known for over 30 years (Vorontsov-Vel`yaminov 1959, Arp 1966, Rose 1977, Hickson et al. 1977, Hickson 1982).  A spectroscopic survey confirmed that 92 of the 100 groups catalogued by Hickson (1982) have at least three accordant redshift members with 69 groups having at least four (Hickson et al. 1992).  With a median galaxy-galaxy separation of 40 h$^{-1}$ kpc  and  a low typical velocity dispersion ($\\sigma$ $\\sim$ 200 km s$^{-1}$, Hickson et al. 1992), compact groups are usually considered to be ideal laboratories for studying galaxy interactions. The most direct way to determine if interactions have occurred among group galaxies is to measure their kinematics in order to check if they are disturbed or normal. Study of the ionized gas kinematics for several of the Hickson group galaxies has shown that they are in different evolutionary stages (Mendes de Oliveira et al. 1998, Plana et al.  1998).  HCG 18 (H18, Arp 258 or VV 143) was cataloged by Hickson (1982) as a group of three irregular galaxies (H18b, c and d with radial velocities between 4080 and 4163 km $s^{-1}$) plus one discordant-redshift S0 galaxy (H18a with a velocity of 10000 km $s^{-1}$; Hickson 1993). In this paper the H18 group refers to the triplet H18b, c and d.  Verdes-Montenegro et al. (1998) determined the IRAS flux for H18d and upper limits on the flux for H18b and H18c.  They also determined the H$_2$ mass for H18b and H18c, from their CO observations. Allam et al. (1996) also give IRAS flux but for H18b, c and d together. We eventually use Allam's values because H18 IRAS emission can not be spatially resolved. HI data is also available for this group, besides the optical measures (Hickson 1993). From a study of the HI velocity field of the group Williams and van Gorkom (1988, hereafter W\\&vG) concluded that the HI gas is concentrated in a single cloud with m$_{HI}$ = 10$^{10}$ M$_{\\odot}$ encompassing all of the optical structures.   \\begin{deluxetable}{lll} \\tablenum{1} \\tablewidth{0pt} \\tablecaption{Journal of Perot-Fabry observations} \\tablehead{ \\colhead{} & \\colhead{Compact Group of Galaxies Hickson 18 } & \\colhead{}} \\startdata Observations & Telescope & CFHT 3.6m  \\nl & Equipment & MOS/FP @ Cassegrain focus  \\nl & Date & August, 22th 1996  \\nl & Seeing & $<$ 1\" \\nl Interference Filter &   Central Wavelength & 6653 \\AA \\tablenotemark{1} \\nl &   FWHM & 23 \\AA  \\tablenotemark{2} \\nl &   Transmission (maximum) & 0.6  \\nl Calibration & Neon Comparison light & $\\lambda$ 6598.95 \\AA \\nl Perot--Fabry & Interference Order & 1162 @ 6562.78 \\AA \\nl & Free Spectral Range at H$\\alpha$ & 265 km s$^{-1}$ \\nl & Finesse at H$\\alpha$ & 12 \\nl & Spectral resolution at H$\\alpha$ & 27344 at the sample step \\nl Sampling & Number of Scanning Steps & 24 \\nl & Sampling Step & 0.24 \\AA\\ (11 km s$^{-1}$)\\nl & Total Field & 430''$\\times $430'' (500$\\times $500 px$^2$) \\nl & Pixel Size & 0.86'' \\nl Detector & & STIS 2 CCD  \\nl Exposures times & Total exposure & 2 hours \\nl & Total exposure time per channel & 300s \\nl   \\enddata \\tablenotetext{} { $^1$ For a temperature de 5$^o$. $^2$ For a mean beam inclination of 2.7$^o$. }  \\end{deluxetable}  We observed the H18 group  in the H$\\alpha$ emission line with a scanning Fabry-Perot instrument and we obtained velocity and H$\\alpha$ integrated flux maps for the system. This observation is part of a larger program that has the goal of unveiling kinematic evidences of interactions in compact groups of galaxies in order to determine their evolutionary stages.  H18 appears to be one of the most complex groups of the Hickson sample. A key question raised by the HI study of W\\&vG was the nature of the three components b, c and d, namely if they are individual entities or form one single galaxy. Based on our observations and comparisons between our data and the HI results, we conclude that H18 is likely to be a single large irregular galaxy. The H$\\alpha$ maps presented in this study permit the determination of the H$\\alpha$ luminosity (L(H$\\alpha$)) and star formation rate (SFR) of the system, which can be compared with typical values for samples of irregular galaxies studied by Hunter et al. (1986, 1989 and 1993). We also compare the total B luminosity (L$_B$) of the group and its far-infrared luminosity (L($_{FIR}$)) with values given in the literature for typical irregular galaxies.  The paper is organized as follows. Section 2 gives details about the reduction of the Fabry-Perot data. Section 3 describes the results. Section 4 has a comparison of our results with the HI observations. Section 5 presents  a discussion about the nature of H18 in light of our new H$\\alpha$ observations and section 6 contains our summary and final remarks. ",
        "conclusions": "The nature of H18 in light of HI observations was discussed by W\\&vG. They considered two possibilities for the nature of this system: 1) it is a knotty irregular galaxy (hereafter referred to as ``the irregular-galaxy scenario''), or 2) it is an interacting group and the observed HI cloud is a remnant of an interpenetrating collision which stripped the gas from the colliding galaxies. Our data and analysis confirm the first scenario.  \\subsection{Arguments in favour of the irregular-galaxy scenario}  The strongest argument in favour of H18 being an irregular galaxy comes from the kinematics of the HI and H$\\alpha$ gas.  The velocity maps show velocity amplitudes of $\\pm$ 70 km s$^{-1}$ and gradients of 15km s$^{-1}$ kpc$^{-1}$ for the ionized gas and 10km s$^{-1}$ kpc$^{-1}$ for the HI.  These confirm that H18 is a slow rotator.  The few detailed studies concerning the kinematics concerning irregular galaxies (Hunter 1982, Tomita et al. 1998) confirm that irregular galaxies are indeed slow rotators.  Hunter (1982) studied a sample of 15 irregular galaxies, five of which showed velocity gradients between 60 and 80 km s$^{-1}$ kpc$^{-1}$, consistent with results found for H18.  Tomita et al. (1998) produced position--velocity diagrams for four dwarf galaxies along several slit orientations.  They did not detect disk-like rotation but velocity gradients were clearly seen.  They also found a velocity difference of 10  to 20 km s$^{-1}$ between the HII regions and the HI gas disk.  Saito et al. (1992) reported a kinematic study of the ionized gas in IC 10 and, particularly, a comparison between the ionized and HI gas kinematics for this galaxy giving the same kind of velocity differences that we found for H18.  Two other works have studied a few irregular galaxies with a Fabry Perot interferometer. Sasaki et al. (1997) reported observations of NGC 4449 and its H$\\alpha$ velocity field that show a kpc-scale mosaic structure of blueshift and redshift components with a slow global rotation. They confirmed the counter rotation between ionized gas and the HI halo for this galaxy. Rosado et al. (1998) and Valdez \\& Rosado (1998) present an optical velocity field for NGC 4449 showing a decreasing gradient along the optical bar and an anticorrelation with respect to the HI velocity field.  The LOS velocity curve of the southern part (Fig. 4a) suggests an independent disk rotation. The velocity curve along the major axis of H18 (Fig 4b) shows that H18b and c lie on the curve, suggesting that they are gravitationaly bound.  The ratio L$_B$/L(H$\\alpha$) of H18 is consistent with the irregular-galaxy hypothesis. Hunter et al. (1989) find a value for the ratio L$_B$ /L(H$\\alpha$) $\\sim$ 44, for giant irregular galaxies while we find 50 for H18.  One last piece of evidence for the irregular-galaxy scenario comes from the HI data.  The size of the large HI cloud around H18 measured by W\\&vG and the total HI mass they found (M(HI)=10$^{10}$ M$_{\\odot}$) are consistent with values found for HI clouds around irregular galaxies such as IC 10 (Shostak \\& Skillman 1989) and NGC 4449 (Hunter et al. 1998). Normalized by the total B luminosity the HI total mass also compares with the typical values found for the Hunter (1993) sample.   \\subsection{Differences between H18 and irregular galaxies}  The main difference between H18 and irregular galaxies concerns the FIR properties and star formation rates.  Hunter et al. (1989, 1993) present properties (in FIR, H$\\alpha$ and broad band imaging) for a sample of 43 irregular galaxies of different types (dwarf, giants, distant, amorphous). We used their sample of dwarf and giant irregular galaxies as a control sample for comparison with the properties of H18.  We measured ratios of L(FIR)/L$_B$ = 0.22 and L(FIR)/L(H$\\alpha$)=11 for H18.  The L(FIR)/L$_B$ ratio is significantly lower than the mean value of 1.9  and 1.0 found for Hunter et al.'s (1989) sample of giant and dwarf irregulars respectively (the  two other classes, ``amorphous'' and ``distant'', show much larger ratios).  The L(FIR)/L(H$\\alpha$) ratio is also lower than the mean values of 90 and 71 for the giant and dwarf irregulars respectively. However, within each subclass, giant, dwarf, amorphous and distant, there is significant scatter in the ratios.  We also derived a dust temperature for H18 (from the S(100$\\mu$) and S(60$\\mu$) IRAS data) of T$_d$=27K, which is cooler than typical dust temperatures derived for other irregular galaxies.  The recent SFR (SFR(B)) and the current SFR ( SFR(FIR) or SFR(H$\\alpha$) ) are much higher (by one or two orders of magnitude) compared with values from Hunter et al's sample.  If, however, we calculate a SFR per area, using the SFR(H$\\alpha$) and taking the area to be that enclosed within the ellipse of major and minor axes defined by the H$\\alpha$ integrated flux map (Fig.5), we find that the SFR/area is 5 $\\times$ 10$^{-9}$ M$_{\\odot}$ yr$^{-1}$ pc$^{-2}$. Hunter \\& Gallagher (1986) have a similar average for the SFR/area for the giant irregular galaxies of their sample.  The other subclasses of irregular galaxies have average values for the SFR/area ratio that are slightly lower than that found for H18."
    },
    "0004/astro-ph0004322_arXiv.txt": {
        "abstract": "s{ We present an analysis of the Gaussianity of the 4-year COBE-DMR data (in HEALPix pixelisation) based on spherical wavelets. The skewness, kurtosis and scale-scale correlation spectra are computed from the detail wavelet coefficients at each scale. The sensitivity of the method to the orientation of the data is also taken into account. We find a single detection of non-Gaussianity at the $>99\\%$ confidence level in one of our statistics. We use Monte-Carlo simulations to assess the statistical significance of this detection and find that the probability of obtaining such a detection by chance for an underlying Gaussian field is as high as $0.69$. Therefore, our analysis does not show evidence of non-Gaussianity in the COBE-DMR data.} ",
        "introduction": "Testing the Gaussianity of the cosmic microwave background (CMB) fluctuations has become of great interest since it would make it possible to distinguish between competing theories of structure formation in the early Universe such as inflation and topological defects. With this aim, a large number of techniques have already been proposed, many of them being applied to the 4-year COBE-DMR data (see e.g. Barreiro 2000). Although most of these analyses did not find evidence for non-Gaussianity, methods based on bispectrum analyses (Ferreira et al. 1998, Magueijo 1999; see also Zaroubi et al. 1999) and on statistics of wavelet coefficients (Pando et al. 1998) have yielded detection of non-Gaussianity in the COBE-DMR data. In particular, Pando et al. have found a significant non-Gaussian signal at the $99\\%$ confidence level on computing the scale-scale correlation of the wavelet coefficients. This analysis has recently been revised by Mukherjee et al. (2000) (hereinafter MHL) who take into account that the results depend critically on the orientation of the signal and that a large number of the computed statistics do not show deviation from the Gaussian case. According to MHL, Gaussianity can be ruled out only at the $41\\%$ confidence level in the DSMB data and at the $72\\%$ level in the 53+90 GHz coadded data. The above wavelet analyses have been performed applying planar wavelets to Face 0 and Face 5 of the COBE-DMR QuadCube pixelisation. Therefore only one-third of the available data is being considered and, at the same time, undesirable projection effects may be present. In this work, we have performed a similar analysis to those of Pando et al and MHL, but applying orthogonal spherical Haar wavelets (SHW) (see Sweldens 1995 and references therein) to the COBE data in HEALPix pixelisation (G\\'orski et al 1999) with the customised Galactic cut (Banday et al. 1997). On the one hand, this hierarchical pixelisation scheme is particularly well-suited to the application of such a wavelet decomposition. On the other hand, SHW are more appropriate to study data over a large region of the sky and also allow an easy identification of those coefficients affected by the Galaxy. Therefore {\\em all} data lying outside the Galactic cut is used in the analysis, i.e., approximately two-thirds of the total number of COBE-DMR pixels.  As an illustration, we present our results for three different orientations of the data to account for the sensitivity of these estimators to the orientation of the input signal. ",
        "conclusions": "We have investigated the Gaussianity of the 4-year COBE data (in HEALPix pixelisation) with an analysis based on SHW. We have taken into account the sensitivity of our method with respect to the orientation of the input signal, presenting the results for three different orientations.  We have found a single detection of non-Gaussianity at the $>99\\%$ confidence level out of our 126 computed statistics, corresponding to the value of the kurtosis at $j=2,m=3$ for one of the chosen orientations. Using Monte-Carlo simulations we estimate that the probability of having such a detection in one of our statistics is as high as $0.69$ for the case of an underlying Gaussian field. Therefore, we conclude that an analysis based on SHW of the 4-year COBE-DMR data show no evidence of non-Gaussianity in the CMB."
    },
    "0004/astro-ph0004114_arXiv.txt": {
        "abstract": "We first briefly review how we investigate the modes of oscillation trapped within the inner region of accretion disks by the strong-field gravitational properties of a black hole (or a compact, weakly-magnetized neutron star). Then we focus on the `corrugation'(c)--modes, nearly incompressible perturbations of the inner disk. The fundamental c--modes have eigenfrequencies (ordered by radial mode number) which correspond to the Lense--Thirring frequency, evaluated at the outer trapping radius of the mode, in the slow rotation limit. This trapping radius is a decreasing function of the black hole angular momentum, so a significant portion of the disk is modulated only for slowly rotating black holes. The eigenfrequencies are thus strongly increasing functions of black hole angular momentum. The dependence of the eigenfrequencies on the speed of sound within (or the luminosity) within the disk is very weak, except for slowly rotating black holes. ",
        "introduction": "For twenty years [beginning with \\citet{kf}], it has been known that general relativity can trap normal modes of oscillation near the inner edge of accretion disks around black holes. The strong gravitational fields that are required can also be produced by neutron stars that are sufficiently compact (with a soft equation of state) and weakly magnetized to produce a gap between the surface of the star and the innermost stable orbit of the accretion disk. Although we shall not explicitly consider such neutron stars here, the results obtained will also apply to them to first order in the dimensionless angular momentum parameter $a=cJ/GM^2$, since their exterior metric is identical to that of a black hole to that order.  These modes of oscillation provide a potentially powerful probe of both strong gravitational fields and the physics of accretion disks, since they do not exist in Newtonian gravity. In addition, their frequencies depend upon the angular momentum as well as the mass of the black hole. The fractional frequency spread of each mode depends upon the elusive viscosity parameter of the accretion disk.  The subject of `relativistic diskoseismology' has recently been reviewed by \\citet{w}. In this paper we shall focus on the `corrugation'(c) modes, previously studied (less generally) by \\citet{k90,k93}, \\citet{i94,i96}, and \\citet{p}. We shall briefly compare these modes with the `gravity'(g) modes and the pressure (p) modes. ",
        "conclusions": "We now apply the above analysis to models of black hole accretion disks. The numerical results discussed below were obtained by assuming that gas pressure dominates within the disk, so that $\\Gamma = 5/3$. However, this is only true near the inner edge $r_i$ [justifying our choice $\\mu=2/5$ in equation (\\ref{eq:33}) consistent with the results of \\citet{pt}], at radii $r\\gg r_i$, and at all radii for luminosities $L\\ll L_{Edd}$.  Recall that the interior structure of the (zero buoyancy) accretion disk enters our formulation only through the constant $\\Gamma$ and the function $\\alpha(r)$ (not to be confused with the usual viscosity parameter, here denoted by $\\alpha_*$), inversely proportional to the speed of sound on the midplane and modeled by equation (\\ref{eq:33}). In this equation we also chose $\\gamma(r)$ to be constant, and took $\\nu=0$ unless otherwise indicated. In section 4.3, we found that the properties of the modes do not change greatly when small amounts of buoyancy are introduced.  For all of the tables (but not the figures), we used $M=10^8M_\\sun$, $\\alpha_*=0.01$, and $L=0.1L_{Edd}$ so we could compare our results with those of \\citet{p}. For a standard thin accretion disk \\citep{ss}, at large radii (where the disk is nonrelativistic and gas pressure dominates) the speed of sound is given by \\be c_s/c = 2.21\\times 10^{-2}\\frac{(L/L_{Edd})^{1/5}}{(\\alpha_*M/M_\\sun)^{1/10}} (rc^2/GM)^{-9/20} \\qquad (r\\gg r_i) \\; , \\label{eq:75} \\ee corresponding to $\\nu=9/20$ \\citep{n}. (We are here reverting to ordinary units.) Since the physical conditions within the disk are more uncertain near its inner edge (but where gas pressure should also dominate), we have also used the nonrelativistic expression \\be c_s/c = 4.32\\times 10^{-3}\\frac{(L/L_{Edd})^{1/5}}{(\\alpha_*M/M_\\sun)^{1/10}}[c^2(r-r_i)/GM]^{2/5} \\qquad (r\\cong r_i)  \\label{eq:76} \\ee there \\citep{n}, as indicated previously. We employ this relation (\\ref{eq:76}) when using $\\nu=0$ in equation (\\ref{eq:33}). We obtain the (constant) value of $\\gamma$ used in equation (\\ref{eq:33}) by setting $\\alpha(r)=1/c_s(r)$ there (again neglecting the relativistic corrections near the inner edge). The various relativistic corrections only become significant when $a\\gtrsim 0.5$. We have only studied the fundamental c--mode ($m^2=j=1$), but have employed various values of the radial mode number $n$.  In Table 1 we present the cutoff values of $a$ [$a_n$, from equation (\\ref{eq:63}); for $a<a_n$, the mode does not exist] for radial mode numbers $n=0,1,2$ and two choices of the boundary condition parameter. There is no c--mode for $n=0$ and $\\theta=0$; no solution is also indicated by no entry in the subsequent tables. The results for $\\theta=\\pi/4$ are essentially identical to those shown for $\\theta=\\pi/2$. The results agree with the analytic expressions (\\ref{eq:64}) valid for values of $a_n\\ll 1$. These results for $\\theta=\\pi/2$ are extended to larger values of $n$ in Table 2. The maximum $n=175$ included corresponds to the requirement $a<a_\\perp\\cong 0.953$ for the validity of our analysis.  Table 3 presents values of our basic expansion parameter $\\epsilon(r_i,r_c)$, obtained from equation (\\ref{eq:20}), as a function of $a$ for the lower radial mode numbers and two choices of boundary condition. We see that it is indeed much less than unity, as also required for our analysis. Recall that $\\epsilon(r_i,r_c) > \\epsilon(r,r_c) > 0$ for all $r_c>r>r_i$. The same format is used for Tables 4 and 5. The eigenfrequency $\\sigma$ (in units of $c^3/GM$, as usual) is presented in Table 4. We see that it is a decreasing function of the radial mode number $n$, more strongly for lower values of $a$. The corresponding radial extent $r_c-r_i$ of the mode (in units of $GM/c^2$, as usual) is shown in Table 5. It increases significantly with radial mode number. We note that the sensitivity of the results in Tables 3--5 to the boundary condition parameter $\\theta$ is modest.  Table 6 shows the dependence of the outer radius $r_c$ of the capture zone on values of $a$ near its cutoff $a_n$, for $n=0$ and $\\theta=\\pi/2$ (corresponding to $a_0=3.22\\times 10^{-5}$). From these data, one can see that when $\\log(a-a_0)\\lesssim -6$, the asymptotic formula (\\ref{eq:69}) holds: $r_c\\propto(a-a_0)^{-q}$. The exponent $q=0.56$ is very close to that predicted, $q=4/7$, for $\\nu=0$.  The effect of changing the behavior of the speed of sound at large radii is illustrated in Table 7. The radial size of the mode for $\\nu=0$ (from Table 5) is compared to that obtained with the choice $\\nu=9/20$ of the standard accretion disk model, equation (\\ref{eq:75}). This shows that the decrease in the speed of sound with radius also reduces the size of the trapping region, as expected. Correspondingly, the values of $a_n$ shown in Table 1 are reduced by about a factor of 7.5, and the values of $\\epsilon(r_i,r_c)$ in Table 3 are reduced by about a factor of 2. The eigenfrequencies in Table 4 are increased by factors of 1.2--2.8 (with $n=0,1$) for $a=10^{-3}$, but were unaffected for $a=0.5$, as expected.  We present our major observationally relevant results in the following figures. For them, we have chosen the radial mode number $n=0$ and the boundary condition parameter $\\theta=\\pi/2$. (One might expect that the lowest radial mode could be the one most easily excited and with the largest net modulation.) The speed of sound parameter is taken to be $\\nu=0$ for Figures 2 and 3, and $\\nu=9/20$ for Figures 4(a,b) (which differed very little when $\\nu=0$ was used).  In Figure 2, we plot the relation between the fundamental c--mode and g--mode eigenfrequencies (scaled by the black hole mass) and the black hole angular momentum. This illustrates the dramatically different dependence of the c--mode frequency. Also shown is the orbital frequency $\\Omega_{max}=\\Omega(r_i)$ of a (commonly invoked) `blob' at the inner disk radius.  As expected, the c--mode frequency approaches $\\Omega_{max}/2\\pi$ as $a\\rightarrow 1$. The case $\\nu=9/20$ lies almost entirely within the band containing the range of masses and luminosities indicated. It is significant that the c--mode results of \\citet{p}, obtained by numerical integration of equations (\\ref{eq:4}) and (\\ref{eq:5}), agree within the band shown in this figure (obtained via a further radial WKB approximation).  In Figure 3, we present our numerical results (points) for the size of the trapping region. The curves shown are the fits of these results to the formula \\be r_c-r_i = K_0(GM/c^2)a^{-K_1}(1-a)^{K_2} \\; , \\label{eq:77} \\ee giving \\begin{eqnarray*} K_0 & = & 0.058\\; , \\quad K_1=0.66\\; , \\quad K_2=0.31\\; \\quad (M=10M_\\sun)\\; , \\\\ K_0 & = & 0.021\\; , \\quad K_1=0.55\\; , \\quad K_2=0.39\\; \\quad (M=10^8M_\\sun)\\; . \\end{eqnarray*} This illustrates the fact that the radial extent of the mode is only appreciable for slowly rotating black holes ($a\\ll 1$). To obtain the corresponding eigenfrequency, one can then use equation (\\ref{eq:77}) and the known function $r_i=Mf(a)$ to obtain the value of the trapping radius $r_c$. Our fundamental result, equation (\\ref{eq:17}) [or equation (\\ref{eq:19}) for $a\\ll 1$] with $m=-1$, then gives $|\\sigma|$.  The c--mode frequency (scaled by mass) depends only on the black hole angular momentum and the accretion disk speed of sound. The first dependence has been shown in Figure 2. Rather than showing the second dependence directly, it is more relevant to instead use an observable, the luminosity. For fixed $\\alpha_*$ and $M$, we use equation (\\ref{eq:75}) to relate it to the speed of sound. We want to know how the frequency changes as the luminosity (proportional to the mass accretion rate) varies with time. That dependence is known to be very weak for the g--modes \\citep{per}. In contrast, the $m=0$ fundamental p--mode has a frequency $|\\sigma|\\propto c_s^{1/3}$ and a radial size $(r_- - r_i) \\propto c_s^{2/3}$ \\citep{kf,nw91,p}.  We fit the dependences of the eigenfrequency on accretion disk luminosity shown by the points in Figures 4(a) and 4(b) by the form $|\\sigma|\\propto L^{-K(M,a)}$. The results shown correspond to \\begin{eqnarray*} K(10M_\\sun,10^{-3}) & = & 0.30\\phn\\; , \\quad K(10^8M_\\sun,10^{-3}) = 0.039\\phn \\; ;\\\\ K(10M_\\sun,10^{-1}) & = & 0.020 \\; , \\quad K(10^8M_\\sun,10^{-1}) = 0.0044 \\; . \\end{eqnarray*} Note that the dependence is weaker for the larger value of $a$, again as expected since the properties of the disk (except its inner radius) become irrelevant as $a\\rightarrow 1$.  Although the fundamental c-mode is almost incompressible, the changing projected area of the mode could modulate the luminosity via reflection of radiation from the postulated `corona' surrounding the accretion disk. Of course, this requires that the disk not be viewed close to face-on. The observability of a c--mode induced modulation of the detected flux would seem to require large values of $r_c-r_i$. This in turn would imply small values of $a$, from Figure 3, and correspondingly small values of frequency, from Figure 2. We then see from Figure 4(a) that the dependence of the frequency on luminosity is relatively weak but might be detectable for the stellar mass black holes. Issues such as the excitation and damping of the c--modes, including their leakage into the black hole via accretion from the inner edge of the  disk, are beyond the scope of this paper.  Finally, we emphasize that of the fundamental (g, p, c) modes, only the c--mode can (generically) have a frequency $|\\sigma|\\ll\\Omega$. The c--modes are candidates for those low frequency features in the power spectra of accreting black holes whose frequency varies only weakly with changes in luminosity. An example of a candidate is the 9 Hz modulation in the `microquasar' GRO J1655-40 observed by the RXTE satellite \\citep{rem}. Since the mass of the presumed black hole in this binary has been determined to be $M\\cong 7M_\\sun$ \\citep{sha}, this frequency requires a black hole angular momentum $a\\cong 0.18$ (Figure 2) if it is produced by a low $n$ c--mode.  However, we see from Figure 3 that the lowest radial mode would have a radial extent of only $0.2GM/c^2$. (But from the results in Table 5, we expect the radial extent of the $n=1$ mode to be about 3 times greater.) In addition, the energy spectrum of the X-rays when this modulation was present was `softer' (more thermal) than typical. This may complicate the above proposal that the photons are modulated via disk reflection from the corona. Clearly, one should search for other radial (or vertical) modes to confirm any identification."
    },
    "0004/astro-ph0004264_arXiv.txt": {
        "abstract": "We present a detailed analysis of the number count and photometric redshift distribution of faint galaxies in the Hubble Deep Field (HDF), paying a special attention to the selection effects including the cosmological dimming of surface brightness of galaxies, under the observational condition employed in this field.  We find a considerably different result from previous studies ignoring the selection effects, and these effects should therefore be taken into account in the analysis. We find that the model of pure luminosity evolution (PLE) of galaxies in the Einstein-de Sitter (EdS) universe predicts much smaller counts than those observed at faint magnitude limits by a factor of more than 10, so that a very strong number evolution of galaxies with $\\eta \\gtilde$ 3--4 must be invoked to reproduce the $I_{814}$ counts, when parametrized as $\\phi^* \\propto (1+z)^\\eta$.  However we show that such a strong number evolution under realistic merging processes of galaxies can not explain the steep slope of the $B_{450}$ and $V_{606}$ counts, and it is seriously inconsistent with their photometric redshift distribution.  We find that these difficulties still persist in an open universe with $\\Omega_0 \\gtilde 0.2$, but are resolved only when we invoke a $\\Lambda$-dominated flat universe, after examining various systematic uncertainties in modeling the formation and evolution of galaxies.  The present analysis revitalizes the practice of using faint number counts as an important cosmological test, giving one of the arguments against the EdS universe and suggests acceleration of the cosmic expansion by vacuum energy density.  While a modest number evolution of galaxies with $\\eta \\ltilde 1$ is still necessary even in a $\\Lambda$-dominated universe, a stronger number evolution with $\\eta > 1$ is rejected from the HDF data, giving a strong constraint on the merger history of galaxies. ",
        "introduction": "Number counting of faint galaxies is one of the most fundamental observational tests with which the formation/evolution of galaxies as well as the geometry of the universe is probed. The best view to date of the optical sky to faint flux levels is given by the Hubble Deep Field (HDF, Williams et al. 1996), and it provides a valuable information to a wide range of studies on galaxies and cosmology.  A comprehensive study of the HDF galaxy counts has been performed by Pozzetti et al. (1998), and they found that a simple model of pure luminosity evolution (PLE), in which galaxies evolve passively due to star formation histories without mergers or number evolution, gives a reasonable fit to the HDF counts in all the four passbands of $U$, $B$, $V$, and $I$, when an open universe with $\\Omega_0 = 0.1$ is assumed.  The increase of the number of galaxies with their apparent magnitude was originally proposed as a measure of the geometry of the universe (Sandage 1961), and considerable efforts have been made along this line (e.g., Yoshii \\& Takahara 1988; Fukugita, Takahara, Yamashita \\& Yoshii 1990; Yoshii \\& Peterson 1991). However, the obtained constraints on cosmological parameters based on the PLE model have not been thought deterministic because of possible number evolution of galaxies by mergers.  Particularly, when the PLE model is used, the Einstein-de Sitter (EdS) universe ($\\Omega_0=1$) underpredicts the observed galaxy counts at faint magnitudes, but a simple model of galaxy number evolution can reproduce the observed counts and save the EdS universe (Rocca-Volmerange \\& Guiderdoni 1990; Pozzetti et al. 1996).  This degeneracy between the effects of galaxy evolution and cosmology has been a major problem when one uses the galaxy number count to determine the geometry of the universe.  The information of redshifts is able to break such a degeneracy, because luminous galaxies at great distance are distinguishable from dwarf galaxies in a local universe.  Although most of the HDF galaxies are too faint to measure the spectroscopic redshifts, several catalogs of their photometric redshifts have been published (Sawicki, Lin, \\& Yee 1997; Wang, Bahcall, \\& Turner 1998; Fern\\'andez-Soto, Lanzetta, \\& Yahil 1999).  The follow-up studies based on these catalogs show that the photometric redshifts give reasonably reliable estimates of spectroscopic redshifts and are useful for a statistical study of the HDF galaxies.  Here we give a combined analysis for the HDF counts and redshifts and constrain the cosmological parameters separately from the merger history of galaxies.  Both the number count of faint galaxies and their redshift distribution are significantly affected by the selection effects inherent in the method of detecting galaxies in faint surveys, but these important effects have been ignored in almost all previous studies except for Yoshii \\& Fukugita (1991) and Yoshii (1993).  It is well known that the surface brightness of galaxies rapidly becomes dimmer with increasing redshift as $\\propto (1+z)^{-4}$ (Tolman 1934), and this cosmological dimming makes many high-redshift galaxies remain undetected below the threshold value of surface brightness adopted in a galaxy survey (Pritchet \\& Kline 1981; Tyson 1984; Ellis, Sievers, \\& Perry 1984). The seeing or smoothing of an image furthermore lowers its surface brightness, and the photometry scheme used in a survey heavily affects a magnitude estimate of the faintest galaxies.  Some observers apply corrections to raw counts of faint galaxies for those undetected, but it is in principle difficult and heavily model-dependent to estimate the number of undetected galaxies. Rather, the best way is to make theoretical predictions with the selection effects taken into account and then compare them directly to raw counts (Yoshii 1993).  This paper is the first analysis of the HDF galaxies in which the above selection bias against high-redshift galaxies is explicitly incorporated.  We use a standard PLE model of galaxies including the effects of internal dust obscuration and intergalactic HI absorption.  Number evolution of galaxies is also allowed for with simple modifications to the PLE model.  Throughout this paper, we use the AB photometry system with the notation of $U_{300}$, $B_{450}$, $V_{606}$, and $I_{814}$ (Williams et al. 1996). In \\S \\ref{section:model}, we present a detailed description for models of galaxy evolution and formulations to calculate galaxy counts and redshift distribution with the selection effects taken into account. Extensive calculations of number count predictions and comparison to the HDF counts are given in  \\S \\ref{section:counts}, checking in great detail the uncertainties arising from the prescribed properties of local galaxies and their evolution. We will give the comparison of the model predictions with the observed photometric redshift distribution in \\S \\ref{section:redshifts}. We discuss the results in \\S \\ref{section:discussion}.  The summary and conclusion of this paper are given in \\S \\ref{section:conclusions}. ",
        "conclusions": "\\label{section:conclusions} We have modeled the number count and redshift distribution of faint HDF galaxies, with the observational selection effects properly taken into account.  As a consequence of the selection effects in the theoretical modeling, predicted counts from the PLE model are smaller than previously considered, and they are more than 10 times smaller than the observed HDF counts at the faintest magnitudes in the EdS universe.  A strong number evolution with $\\eta \\gtilde$ 3--4 under the assumption of conserved luminosity density is required to explain the faintest counts in this EdS universe, when the number evolution is parametrized as $\\phi^* \\propto (1+z)^\\eta$ and $L^* \\propto (1+z)^{-\\eta}$.  However, such a strong number evolution is not consistent with the overall $N$-$m$ slope or the photometric redshift distribution.  These discrepancies become even worse when one considers a more realistic merger process, i.e., enhanced star formation following by gas-rich mergers.  In addition, such a strong evolution is rejected at least for average $L^*$ galaxies at $z<1$ from the data of spectroscopic redshift surveys (Totani \\& Yoshii 1998; Shade et al. 1999; Le F\\'evre et al. 2000).  Therefore, we conclude that it is almost impossible to explain the HDF galaxies in the EdS universe, unless we invoke ultra-exotic galaxy populations such as galaxies forming only massive stars at high redshifts to escape from local galaxy surveys due to the complete lack of long-lived stars.  The present work revitalizes the practice of using faint number counts as an important cosmological test, which gives one of the arguments against the EdS universe by its outstanding statistics compared with other cosmological tests. An open universe with $\\Omega_0 > 0.2$ does not fit to the HDF data either, for the similar reasons for rejecting the EdS universe.  An open universe with $\\Omega_0 \\sim 0.1$ might be consistent with the HDF data, but such a low value of $\\Omega_0 \\sim 0.1$ would not be reconciled with other constraints on $\\Omega_0$, such as the baryon-gas to dark-matter mass ratio in poor clusters of galaxies combined with the standard big-bang nucleosynthesis prediction of baryon mass density in the universe (e.g., Pedersen, Yoshii, \\& Sommer-Larsen 1997).  We have extensively checked systematic uncertainties in our theoretical modeling of galaxy formation and evolution, and found that they are unlikely to resolve the above discrepancies emerged in the EdS universe and also in an open universe.  On the other hand, such discrepancies are naturally resolved if we invoke a $\\Lambda$-dominated flat universe.  This suggests that the existence of the cosmological constant or an exotic form of the vacuum energy density of the universe which is now accelerating the expansion of the universe.  The PLE model in a $\\Lambda$-dominated flat universe with $\\Omega_0 \\sim 0.2$ gives a reasonable fit to the HDF data, and a modest number evolution with $\\eta \\ltilde 1$ is also suggested by the HDF counts at the faintest magnitudes.  It should be noted that this number evolution does not necessarily mean mergers of galaxies, but may suggest strongly clumpy star-forming regions within an individual galaxy system becoming visible at high redshifts (Colley et al. 1996; Bunker, Spinrad, \\& Thompson 1999).  On the other hand, it is interesting to note that this indication of mild number evolution is consistent with the merger rate evolution of $L^*$ galaxies at $z<1$ recently inferred from a high-resolution image study for galaxies in the CFRS survey (Le F\\'evre et al. 2000). This result is consistent with some models of galaxy formation based on the hierarchical structure formation in the CDM universe (Le F\\'evre et al. 2000), although there are considerable uncertainties in the theoretical calculations for the merging history of baryonic component.  A stronger number evolution with $\\eta \\gtilde 1$ is, however, strongly disfavored by the observed HDF galaxy counts.  This will give an important constraint when galaxy formation is modeled in the framework of the structure formation in a cold dark matter universe.  Inclusion of the selection effects in this paper leads to a considerably different result from previous studies on galaxy number count and redshift distribution.  This means that any cosmological interpretations will be seriously misled if the selection effects are ignored.  All future studies related to the detection and statistics of high-redshift galaxies should take into account these effects.  The selection effects give a bias against high-redshift galaxies, reducing a problem of overprediction of such galaxies by the PLE model, which has been claimed by several studies ignoring the selection effects (e.g., Ellis 1997). In fact, we have shown that the PLE model is in overall agreement with the HDF galaxies, even if a modest number evolution of galaxies ($\\eta \\ltilde 1$) may be required.  A strong number evolution, however, predicts too small a number of high-redshift galaxies to be consistent with the photometric redshift distribution of the HDF galaxies.  The authors would like to thank K. Shimasaku for providing numerical data for the filter functions of HST photometry bands, and T. Tsujimoto and C. Kobayashi for providing their models of galaxy evolution in a tabular form.  We also thank an anonymous referee for many useful comments which have considerably improved this manuscript. This work has been supported in part by a Grand-in-Aid for Conter-of-Excellence Research (07CE2002) of the Ministry of Education, Science, and Culture in Japan."
    },
    "0004/astro-ph0004052_arXiv.txt": {
        "abstract": "The spectral energy distribution (SED) of gamma-ray loud BL~Lac objects typically has a double-humped appearance usually interpreted in terms of synchrotron self-Compton models.  In proton blazar models, the SED is instead explained in terms of acceleration of protons and subsequent cascading.  We discuss a variation of the Synchrotron Proton Blazar model, first proposed by M\\\"ucke \\& Protheroe (1999), in which the low energy part of the SED is mainly synchrotron radiation by electrons co-accelerated with protons which produce the high energy part of the SED mainly as proton synchrotron radiation.  As an approximation, we assume non-relativistic shock acceleration which could apply if the bulk of the plasma in the jet frame were non-relativistic. Our results may therefore change if a relativistic equation of state were used. We consider the case where the maximum energy of the accelerated protons is above the threshold for pion photoproduction interactions on the synchrotron photons of the low energy part of the SED. Using a Monte Carlo/numerical technique to simulate the interactions and subsequent cascading of the accelerated protons, we are able to fit the high-energy gamma-ray portion of the observed SED of Markarian 501 during the April 1997 flare.  We find that the emerging cascade spectra initiated by gamma-rays from $\\pi^0$ decay and by $e^\\pm$ from $\\mu^\\pm$ decay turn out to be relatively featureless.  Synchrotron radiation produced by $\\mu^\\pm$ from $\\pi^\\pm$ decay, and even more importantly by protons, and subsequent synchrotron-pair cascading, is able to reproduce well the high energy part of the SED.  For this fit we find that synchrotron radiation by protons dominates the TeV emission, pion photoproduction being less important with the consequence that we predict a lower neutrino flux than in other proton blazar models. ",
        "introduction": "During its giant outburst in April 1997, the nearby BL~Lac object Mkn~501 (at redshift z=0.034) emitted photons up to $24$~TeV and $0.5$~MeV in the $\\gamma $-ray and X-ray bands, respectively, and has proved to be the most extreme TeV-blazar observed so far (e.g. Catanese et al 1997, Pian et al 1998, Protheroe et al 1998, Quinn et al 1999, Aharonian et al 1999). This energy is the highest so-far observed for any BL~Lac object, and the flux is approximately 2 orders of magnitude higher than the synchrotron peak at its quiescent level. BeppoSAX and OSSE observations (Maraschi 1999) suggest that the X-ray spectrum is curved at all epochs, and the spectrum during flaring has been fitted by a multiply-broken power-law (Bednarek \\& Protheroe 1999). COMPTEL has not seen any significant signal from Mkn~501 at any time (Collmar 1999), while a $3\\sigma $ upper limit of $F(>100\\,{\\rm MeV})<3.6\\times 10^{-7}$cm$^{-2}$ s$^{-1}$ has been derived for the April 1997 EGRET viewing period (Catanese et al 1997).  A flux increase at TeV-energies was also observed with the Whipple, HEGRA and CAT telescopes (Catanese et al 1997), with the most intense flare peaking on April 16 at a level $\\sim100 $ times higher than during its quiescent flux. The non-detection of Mkn~501 by EGRET indicates that most of the power output of the high energy component is in the GeV-TeV range. The TeV-observations revealed a power-law spectrum with photon index $\\sim 2$ up to $\\sim 10$~TeV and a gradual steepening up to 24~TeV. The extragalactic diffuse infrared background leads to significant extinction of $\\gamma $-rays through $\\gamma \\gamma $-pair production above 10~TeV. The extinction-corrected TeV-spectrum (e.g. Bednarek \\& Protheroe 1999), shows the spectral energy distribution (SED) peaking at $\\sim 2$~TeV. Optical observations did not show any significant variations (Buckley \\& McEnery 1997), indicating that the change in the low energy part of the SED was mainly confined to the X-ray band above $0.1$ keV.  Various models have been proposed to explain the observed $\\gamma $-ray emission from TeV-blazars, all of which are identified as high-frequency peaked BL~Lac objects. Leptonic models, in which electrons inverse-Compton scatter a population of low energy photons to high energies, currently dominate the thinking of the scientific community. Because of the low luminosity of accretion disks in BL~Lacs, the main target photons for the relativistic electrons would be the synchrotron photons produced by the same relativistic electron population, as in the synchrotron self-Compton (SSC) model. An alternative scenario for the production of the observed $\\gamma $-ray flux has been proposed involving pion photoproduction by energetic protons with subsequent synchrotron-pair cascades initiated by decay products (photons and $e^\\pm$) of the mesons (e.g. Mannheim et al 1991, Mannheim 1993). These proton-initiated cascade (PIC) models could, in principle, be distinguished by the observation of high energy neutrinos produced as a result of photoproduction.  In this paper, we consider the April 1997 flare of Mkn~501 in the light of a modified Synchrotron Proton Blazar (SPB) model. We assume that electrons ($e^{-}$) and protons ($p$) are accelerated by 1st order Fermi acceleration at the same shock. The relativistic $e^{-}$ radiate synchrotron photons which serve as the target radiation field for proton-photon interactions, and for the subsequent pair-synchrotron cascade which develops as a result of photon-photon pair production. This cascade redistributes the photon power to lower energies where the photons escape from the emission region, or ``blob,'' which moves relativistically in a direction closely aligned with our line-of-sight.  Until recently, this model was not able to reproduce the general features of the double-humped blazar spectral energy distribution (SED), but produced a rather featureless spectrum (see e.g. Mannheim 1993), nor could it explain correlated X-ray/TeV-variability. Here, we present a comprehensive description of our Monte-Carlo simulations of a stationary SPB model, including all relevant emission processes, and show that this model is indeed capable of reproducing a double-humped SED as observed. Here, the origin of the TeV-photons are proton synchrotron radiation, as first proposed by M\\\"ucke \\& Protheroe (1999); a similar model has also been proposed by Aharonian (2000), and Rachen (1999) presented speculations about $\\mu^\\pm$- and proton-synchrotron radiation leading to narrow cascade spectra during flares, which might explain correlated X-ray/TeV-variability. Jet energetics and limits from particle shock acceleration, however, put severe constraints on this scenario. The goal of this paper is to discuss the physical processes included in our SPB model Monte-Carlo code, and give the results of applying this code, as an example, to reproduce the SED of the giant flare from Mkn~501 which occurred in April 1997. A comprehensive study of the whole parameter-space (magnetic field, Doppler factor, etc.) for this model will be the subject of a subsequent paper.  In Section 2, we discuss constraints on the maximum particle energies imposed by the co-acceleration scenario, and by the pion production threshold. Section 3 is devoted to the emission processes in the present model. Energy losses and particle production are treated in Sect. 3.1, while the cascade calculations, including a brief description of our code, are outlined in Sect. 3.2. In Sect. 4 we apply our model to the April 1997 flare of Mkn~501. The multifrequency photon spectrum is shown in Sect. 4.1, while in Sect. 4.2 the predicted neutrino spectrum is discussed. We conclude with a discussion and summary in Section 5. ",
        "conclusions": "This paper describes an application of our newly-developed Monte Carlo program which simulates a modified version of the stationary SPB-model.  The Monte Carlo technique allows us to use exact cross sections, and all important emission processes are considered here.  As an example, we have used our code to model the giant April 1997 TeV-flare from Mkn~501. Here, the TeV-photons are due to synchrotron radiation of the relativistic protons in the highly magnetized emission region. This proton synchrotron model was first proposed by M\\\"ucke \\& Protheroe (1999); a similar model for TeV emission in Mrk~501 has just been proposed by Aharonian (2000), and his conclusions regarding the required Doppler factor and magnetic field are very similar to ours. Our model departs from the standard SPB model as introduced first by Mannheim and co-workers mainly in two areas: (i) we use the {\\it{observed}} synchrotron radiation as the target photon field for $p\\gamma$-interactions and pair-synchrotron cascades assuming it to be produced by co-accelerated electrons, and (ii) our model takes into account synchrotron radiation from muons and protons.  The model parameters derived assuming diffusive shock acceleration of $e^-$ and $p$ in a Kolmogorov turbulence spectrum are consistent with the X-ray to TeV-data in the flare state. However, the total jet power we obtain is too large to comply with the steady-state jet--disk symbiosis scenario, but then we are not dealing with a stready-state phenomenon.    While the emerging cascade SED initiated by $\\pi^0$ decay and $\\pi^\\pm$ synchrotron photons turns out to be relatively featureless, as was also found by, e.g., Mannheim (1993), the $\\mu^\\pm$ (see also Rachen 1999, and Rachen \\& Mannheim 2000) and, more importantly, the proton synchrotron radiation and its cascade produces a double-humped SED as is commonly observed in flaring blazars. For the present model, we find that proton synchrotron radiation dominates the TeV emission, while the contribution of the synchrotron radiation from the pairs, produced by photon-photon interactions of gamma-rays from the high energy hump is only minor.  Our model considers the emission region to be homogeneous.  Inhomogeneities in particle density, magnetic field, etc. within the source would result in a broader X-ray and TeV-peak in the SED. This indicates that for Mkn~501 a homogeneous model of the emission region seems to be appropriate.  Being a hadronic model, our model predicts neutrino emission and we give the expected neutrino flux of Mrk~501 during flaring. Comparing our predicted neutrino flux with that for previous proton blazar models (e.g. Mannheim 1993), we find that the neutrino output in our model is significantly less than was previously estimated due to the synchrotron losses dominating the energy losses of protons, producing synchrotron $\\gamma$-rays at the expense of $\\pi^0$ $\\gamma$-rays and neutrinos."
    },
    "0004/gr-qc0004001_arXiv.txt": {
        "abstract": "\\baselineskip 16pt \\noindent One of the remarkable features of black holes is that they possess a thermodynamic description, even though they do not appear to be statistical systems.  We use self-gravitating magnetic monopole solutions as tools for understanding the emergence of this description as one goes from an ordinary spacetime to one containing a black hole. We describe how causally distinct regions emerge as a monopole solution develops a horizon.  We define an entropy that is naturally associated with these regions and that has a clear connection with the Hawking-Bekenstein entropy in the critical black hole limit. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004270_arXiv.txt": {
        "abstract": "We are using the Multibeam 21cm receiver on the Parkes Telescope combined with the optical Two degree Field spectrograph (2dF) of the Anglo-Australian Telescope to obtain the first complete spectroscopic sample of the Fornax cluster.  In the optical the survey is unique in that all objects (both ``stars'' and ``galaxies'') within our magnitude limits ($16.5\\leq B_J \\leq19.7$) are measured, producing the most complete survey of cluster members irrespective of surface brightness. We have detected two new classes of high surface brightness dwarf galaxy in the cluster. With 2dF we have discovered a population of very low luminosity ($M_B\\approx -12$) objects which are unresolved from the ground and may be the stripped nuclei of dwarf galaxies; they are unlike any known galaxies. In a survey of the brighter ($16.5\\leq B_J \\leq18$) galaxies with the FLAIR-II spectrograph we have found a number of new high surface brightness dwarf galaxies and show that the fraction of star-forming dwarf galaxies in the cluster is about 30\\%, about twice that implied by earlier morphological classifications.  Our radio observations have greatly improved upon the sensitivity of the standard Multibeam survey by using a new ``basket weave'' scanning pattern. Our initial analysis shows that we are detecting new cluster members with HI masses of order $10^8$M$_\\odot$ and HI mass-to-light ratios of 1--2 M$_\\odot$/L$_\\odot$. ",
        "introduction": "It has long been suggested that optical selection effects limit the galaxies in optical surveys to a narrow range of surface brightness (Disney \\& Phillipps 1983 and references therein). The idea is that very low surface brightness (LSB) galaxies are lost in the sky noise and compact, high surface brightness galaxies would be confused with stars.  It now seems unlikely that there are large numbers of undetected giant LSB galaxies (Driver, these proceedings) and the number of unresolved giant galaxies missed in photographic surveys is small (Drinkwater et al.\\ 1999). However the situation for dwarf galaxies may be different.  Most flux-limited galaxy surveys are dominated by giant galaxies, so the only way to study significant samples of dwarf galaxies is to observe nearby galaxy clusters where the galaxy density is so elevated that many dwarfs can be detected. The advantage of cluster samples is that cluster membership can often be assigned on morphological grounds, avoiding the need for spectroscopy. This has been done very effectively in surveys of the Virgo and Fornax Clusters which now form the basis of our knowledge of dwarf galaxies (Binggeli, Sandage \\& Tammann 1985; Ferguson 1989=FCC).  However the lack of spectroscopy becomes a serious handicap when it comes to unusual types of dwarf galaxy: these may not be included on morphological grounds.   \\begin{figure} \\plotone{drinkwater1.eps} \\caption{New galaxies detected in the Fornax Cluster. These Bj-band photographic images (from the DSS: see acknowledgements) are all 3 arcminutes across with North at the top and East to the left.} \\end{figure}  In this paper we present results from an extensive multi-wavelength {\\em spectroscopic} survey of the Fornax Cluster. We are obtaining the most complete spectroscopic sample of cluster galaxies ever made, detecting new types of galaxy missed in the previous morphological surveys. In the optical we are using the Two degree Field spectrograph (2dF) of the Anglo-Australian Telescope (AAT) to make a complete survey of faint objects in the core of the cluster and the FLAIR-II spectrograph of the UK Schmidt Telescope (UKST) to measure brighter compact galaxies over a six degree field. In the radio we are making a blind spectroscopic survey of an even larger ten degree field with the Multibeam 21cm receiver on the Parkes Telescope. In Fig.~1 we show images of a selection of the cluster galaxies observed; these are described in more detail below. ",
        "conclusions": "We have shown that our approach of making complete spectroscopic surveys has overcome the selection effects evident in previous cluster samples of dwarf galaxies. The ``new'' members we find are not difficult to detect optically, but in each case have been excluded from existing compilations because of their high surface brightness or large distance from the cluster centre. We have detected unresolved compact objects, high surface brightness dwarf galaxies and radio galaxies. When it comes to clusters, galaxies do not have to be optically faint to be ``hidden'', they just have to be unusual.  Apart from the compact objects, most of the galaxies we found show evidence of high rates of star formation. In the complete FLAIR-II sample our data show that star formation is about twice as common in the dwarf galaxies as would have been implied by the earlier morphological classifications. Star formation is still important in the Fornax cluster dwarfs, but not in the cluster core where the density of star-forming galaxies is reduced. This is consistent with the radio data which show that the neutral hydrogen in the cluster is much less centrally concentrated than the optical luminosity."
    },
    "0004/astro-ph0004100_arXiv.txt": {
        "abstract": "First detections of thermal water vapor absorption lines have been made toward Orion IRc2 using the {\\it Short Wavelength Spectrometer} (SWS) on board the {\\it Infrared Space Observatory} (ISO). Grating spectra covering wavelengths 25--45 $\\mu$m yield 19 pure rotational lines, originating from energy levels 200--750~K above ground. Fabry-Perot spectra of 5 transitions resolve the line profiles and reveal the H$_{2}$O gas kinematics. The fact that all lines are seen in absorption is in striking contrast with data from the ISO {\\it Long Wavelength Spectrometer} (LWS), where the H$_2$O lines appear in emission. At least one line displays a P-Cygni type profile, which suggests that the water is located in an expanding shell centered on or near IRc2. The expansion velocity is 18 km s$^{-1}$, in agreement with the value inferred from H$_{2}$O maser observations by Genzel et al.\\ (1981). Because the continuum is intense and likely formed in or near the water-containing gas, the excitation of the observed transitions is dominated by radiative processes. A simple, generalised curve-of-growth method is presented and used to analyze the data. A mean excitation temperature of 72 K and a total H$_2$O column density of $1.5\\times 10^{18}$ cm$^{-2}$ are inferred, each with an estimated maximum uncertainty of 20\\%. Combined with the H$_2$ column density derived from ISO observations of the pure rotational H$_2$ lines, and an assumed temperature of 200--350 K, the inferred H$_2$O abundance is 2--5$\\times 10^{-4}$ in the warm shocked gas. This abundance is similar to that found recently by Harwit et al. (1998) toward Orion using data from the LWS, but higher than that found for most other shocked regions by, for example, Liseau et al. (1996). ",
        "introduction": "Water is one of the prime species for probing the interaction between young stars and their surroundings, both in terms of its abundance and its peculiar excitation. In high temperature gas, appropriate to shocks and hot cores for instance, all of the oxygen not locked up in CO is predicted to be driven into H$_2$O at temperatures above 230 K by the O + H$_2$ $\\to$ OH + H and OH + H$_2$ $\\to$ H$_2$O + H reactions, resulting in very bright H$_2$O lines (e.g., Hollenbach \\& McKee 1979, Neufeld \\& Melnick 1987, Kaufman \\& Neufeld 1996, Charnley 1997). In contrast, the H$_2$O abundance may be at least two orders of magnitude lower in surrounding colder gas (e.g., Zmuidzinas et al.\\ 1995). In addition to collisions, the H$_2$O excitation and line profiles can be strongly affected by mid- and far-infrared radiation from warm dust (e.g., Phillips et al.\\ 1978, Takahashi et al. 1983, 1985), providing detailed information on the physical parameters of the gas and its location with respect to the radiation sources.  Interstellar water has been difficult to detect, apart from its presence as an ice mantle on dust grains or through its maser emission, owing to the severe telluric absorption encountered at Earth-based observatories. Nevertheless, over the last twenty years, many searches for lines of gas-phase H$_2$O and its isotopomers have been made from the ground and airborne altitudes, in particular toward Orion (e.g., Waters et al.\\ 1980, Phillips et al.\\ 1978, Jacq et al. 1990, Wannier et al.\\ 1991, Cernicharo et al.\\ 1994, Zmuidzinas et al.\\ 1995, Tauber et al.\\ 1996, Timmermann et al.\\ 1996, Gensheimer et al.\\ 1996, Cernicharo et al.\\ 1999a,b). However, due to the choice of line, wavelength and beam size, different components are often probed, whilst many ground-based observations refer to masing lines, for which sophisticated shock and maser models are required in order to extract physical parameters.  One of the major aims of the {\\it Infrared Space Observatory} ({\\it ISO}) mission has been the routine measurement of thermal gas-phase water lines, and their use as diagnostics of the chemical and physical conditions within molecular clouds. Far-infrared pure rotational H$_2$O emission lines in the 50--200 $\\mu$m wavelength range have been detected with the {\\it Long Wavelength Spectrometer} (LWS) of ISO in a number of star-forming regions (e.g., Liseau et al.\\ 1996, Ceccarelli et al.\\ 1998), including Sgr B2 (Cernicharo et al.\\ 1997a) and Orion (Cernicharo et al.\\ 1997b, 1999a; Harwit et al.\\ 1998). Additionally, van Dishoeck \\& Helmich (1996), van Dishoeck (1998) and Dartois et al.\\ (1998) have observed absorption around 6 $\\mu$m in the $\\nu_{2}$=1--0 band toward a number of deeply embedded, massive young stars. Typical H$_2$O abundances of $10^{-5}$ with respect to H$_2$ have been derived from these data. Similar observations have recently been reported for Orion BN/IRc2 by van Dishoeck et al.\\ (1998) and Gonz\\'{a}lez-Alfonso et al.\\ (1998), although in this case emission is also detected. Wright et al.\\ (1997) have however shown that the detection of the corresponding pure rotational lines at $\\sim$ 30--200 $\\mu$m in most sources observed at 6 $\\mu$m is still difficult. The observations of Orion-IRc2 presented here form a notable exception.  Many of the earlier searches for H$_2$O lines have been performed toward the BN/IRc2 complex of infrared sources in Orion, because of the extraordinary brightness of many atomic and molecular lines in this region compared with other clouds (e.g., Genzel \\& Stutzki 1989, Blake 1997). See van Dishoeck et al.\\ (1998) and references therein for a detailed description of the geometry and the diverse range of physical phenomena in this region. In this paper we present the first detection of numerous pure rotational water lines in absorption toward IRc2 in the 25--45 $\\mu$m interval with the {\\it Short Wavelength Spectrometer} (SWS) (de Graauw et al.\\ 1996) on board {\\it ISO}. Some of these lines have been velocity resolved with the Fabry-Perot, enabling direct information on the location of the absorbing gas to be inferred. These data complement the earlier ground-based data, as well as observations of pure rotational lines with the LWS obtained by Cernicharo et al.\\ (1997b, 1999a) and Harwit et al.\\ (1998) in a much larger beam.  Our interpretation of the ISO spectra of Orion IRc2 suggests that the 25--45 $\\mu$m H$_2$O spectrum originates in a region where the intrinsically strong lines couple efficiently to an intense continuum. In principle, the formation of such a spectrum should be described for a stratified atmosphere in which lines and continuum are treated self-consistently. We show here that the observed features of the spectrum can be described well by a simple ``generalised curve-of-growth'', which includes the effects of coupling to a strong continuum. In the following we describe our observations and present our results, followed by a discussion of the location of the absorbing water vapour, its excitation and finally its abundance. ",
        "conclusions": "Through use of the ISO--SWS in its grating mode, 19 pure rotational absorption lines of water have been detected toward Orion IRc2 for the first time. Fabry-Perot spectra of 5 lines reveal that the water is located in an outflow expanding at a velocity of 18 km s$^{-1}$. The strong mid-infrared continuum toward IRc2 plays a dominant role in the excitation of the molecule and the line formation, which can be modeled using a simple, generalised curve-of-growth technique. This yields a total water column density of order 1.5$\\times10^{18}$ cm$^{-2}$ and excitation temperature of 72 K, similar to the dust continuum colour temperature. Both derived quantities have a maximum uncertainty of about 20\\%. The data provide support for large abundances of H$_2$O in the outflows of massive stars. Simultaneous analysis of the complete ISO--SWS and LWS data set on H$_2$O, OH and CO may provide further information on the abundance and excitation of these molecules in the various physical components within the complex Orion environment."
    },
    "0004/astro-ph0004336_arXiv.txt": {
        "abstract": "MeV seed photons produced in shocks in a variable ultra-relativistic outflow gain energy by the Fermi mechanism, because the photons Compton scatter off relativistically colliding shells. The Fermi-modified high-energy photon spectrum has a non-universal slope and a universal cutoff. A significant increase in the total radiative efficiency is possible.  In some gamma ray bursts, most of the power might be emitted at the high-energy cutoff for this mechanism, which would be close to 100 MeV for outflows with a mean bulk Lorentz factor of 100. ",
        "introduction": "The most common model of gamma-ray burst (GRB) sources involves a relativistic outflow in which shocks occur and radiate away a fraction of the bulk kinetic energy. For typical model parameters the synchrotron emission peaks around 1 keV in the comoving frame, or $\\sim 100$ keV in the observer frame.  Generally, this is considered to be the primary spectrum observed.  However, in internal shocks in the neighborhood of the flow photosphere, and also in external shocks in some cases (e.g. Madau, Blandford, \\& Rees 2000), the shocks can have a non-negligible Thomson scattering depth, which results in upscattering of these primary photons. Single scattering on individual shock-accelerated electrons with Lorentz factors $\\gamma _e\\sim 300$ produces photons with energies $\\sim \\gamma _e^2$ keV in the comoving, or $\\sim \\Gamma \\gamma _e^2 \\sim 10\\Gamma _2$ GeV in the observer frame, where $\\Gamma =100\\Gamma _2$ is the bulk Lorentz factor.  Here we concentrate on a different, multiple scattering component. This is associated with mildly relativistic motions of different ejecta shells or turbulent cells resulting from multiple shock interactions, which contain more energy than the shock-accelerated highly relativistic electrons. Multiple interacting shells are naturally expected in internal shocks, and also in external shocks when a longer lasting modulated outflow runs into the first decelerated shell (e.g. Fenimore \\& Ramirez-Ruiz 2000; Kumar \\& Piran 2000).  Repeated scatterings using the energy of these bulk motions boost the photon energy through the equivalent of the Fermi acceleration mechanism of particles (Blandford \\& Payne 1981). This is related to Thompson's (1994) photon scattering off Alvf\\'en waves, but our mechanism relies instead on relative bulk motions. It differs also in using synchrotron photons instead of thermal photons as its source term, and hence leads to different characteristic energies. This bulk Comptonization results in a spectrum extending at least up to $\\sim$ MeV in the comoving frame and $\\sim 100\\Gamma_2$ MeV in the observer frame. The spectral power or luminosity per decade can increase as steeply as linearly in the photon energy. This provides a natural explanation for those GRB spectra (e.g. Preece et al 1999) which show a positive $\\nu F_\\nu$ slope above the MeV range (generally $\\leq +1$), which cannot be explained by direct synchrotron radiation from Fermi shock-accelerated electrons.  The component made up of bulk-scattered photons can extend up to a maximum observed energy $\\sim 100 \\Gamma_2$ MeV. Beyond this energy, Klein-Nishina and electron recoil effects set in, and the spectrum reverts to being dominated by the seed spectrum (with negative or flat power law slope) of the unscattered photons above the synchrotron peak.  In this work we adopt a test photon model, which assumes no back reaction on the plasma. Several potentially important effects are neglected (upscattered photons can heat electrons in the colliding shells and produce pairs, light pressure ensures that the total comoving energy of the scattered photons cannot exceed the total kinetic energy of the relative shell motions). These effects will be investigated elswehere (Gruzinov, \\Mesz, \\& Rees 2000). The simplified approach that we use here retains the essential properties of the bulk motion comptonization phenomenon, and allows us to explore the main qualitative features it introduces in the spectra.  In \\S 2 we specify the GRB model, in \\S 3 we give an analytical model of Fermi acceleration of photons by colliding shell, and in \\S 4 we describe our Monte Carlo simulations. The results are discussed and related to current and future observations in \\S 5. ",
        "conclusions": "The Fermi acceleration of photons in the standard GRB fireball model has several consequences of theoretical and observational interest. The first is that it can naturally produce $\\nu F_\\nu$ spectra which are harder than the input synchrotron spectrum, including the possibility of an increasing $\\nu F_{\\nu }$ above the usual break found in the Band parameterization of spectra (Band et al 1999, Preece et al 1999). The latter is a property that is hard to obtain with a synchrotron model. Such rising $\\nu F_\\nu \\propto \\nu^\\beta$ are expected, in this model, to have $\\beta \\siml 1$, and this implies that in some GRB most of the energy is at energies $h\\nu \\sim 100 (\\Gamma /10^2)$ MeV, well above the BATSE instrument band on the Compton Gamma Ray Observatory (CGRO). The current observational situation is that BATSE finds approximately 16 \\% of the spectra to be rising at $\\simg 1$ MeV, and the observed rise is not faster than  $\\beta =1$ (Preece et al 1999). These fits generally cut off above 1.8 MeV, and the break is usually not much lower than this, so there is some uncertainty. The COMPTEL instrument on CGRO, sensitive up to 30 MeV, has analyzed $\\sim 30$ bursts (Schoenfelder et al 2000), and an analysis of the slopes indicates in several cases $\\nu F_\\nu$ slopes $\\beta\\sim 0$, with one burst of $\\beta\\sim 0.5$ (Kippen et al 1999). The EGRET experiment on CGRO has detected $\\sim 30$ bursts with the scintillation counters in the 1-200 MeV range, and $\\sim$ 7 bursts with the spark chambers in the 100 MeV-30 GeV range. In this range the spectra are largely noise-dominated (Schaefer et al 1998, Bromm \\& Schaefer 1999). However, the scintillation spectral slopes (Catelli, Dingus \\& Schneid, 1997) are compatible with $\\beta\\siml 0$, although some could be positive and others negative, which is compatible with the analysis of spark chamber data (e.g. Sommers, 1994; Hurley et al, 1994). Large area detectors such as GLAST should be able to obtain more definite answers in the 20 MeV-300 GeV range.   Another implication of the photon acceleration described here is that it provides a natural mechanism to increase the efficiency of conversion of baryon bulk motion into photon energy. This is of interest since in general the internal shock synchrotron efficiency for radiating in the BATSE band is limited to 1-10\\% (Kumar 2000; Spada, Panaitescu \\& \\Mesz 2000; see however Fenimore \\& Ramirez-Ruiz 2000, and observation-based estimates by Freedman \\& Waxman 2000). The increase in the radiative efficiency is simply given by the increase in the value of $\\nu F_\\nu$ at different energies, or by the integral $\\int F_\\nu d\\nu$ in the range of interest. In the generic examples  shown, this increase is substantial. Of course, the spectra shown in Figures \\ref{fig:2} and \\ref{fig:3} are test photon spectra, which do not take into account the back-reaction of radiation. The latter can become important when a substantial fraction of the bulk energy has been converted to radiation through Fermi acceleration, and the natural limit for this effect can be estimated as $\\sim 50\\%$ radiative efficiency.  The calculations presented here are meant to illustrate the consequences of photon acceleration by bulk motions. A self-consistent calculation is needed in order to explore the back-reaction of the radiation pressure and pair formation on the shell dynamics.  Pair formation has an angle averaged cross section $\\sigma _{\\gamma \\gamma}\\sim \\sigma_T/8$ at a threshold which is similar to that for Klein-Nishina effects. It is not very important in low compactness situations (shocks at large radii), while in high compactness cases the exponential tail of the photon distribution would lead to a pair cascade, which lowers the cut off in the comoving spectra of Figures \\ref{fig:2},\\ref{fig:3} to $\\siml$ 0.5 MeV, possibly with some pile-up of photons at this energy (Gruzinov, \\Mesz \\& Rees, 2000)."
    },
    "0004/astro-ph0004385_arXiv.txt": {
        "abstract": "\\noindent \\rightskip=0pt Experimental study of the anisotropy in the cosmic microwave background (CMB) is gathering momentum.  The eagerly awaited {\\sl Boomerang\\/} results have lived up to expectations.  They provide convincing evidence in favor of the standard paradigm: the Universe is close to flat and with primordial fluctuations which are redolent of inflation.  Further scrutiny reveals something even more exciting however -- two hints that there may be some unforeseen physical effects. Firstly the primary acoustic peak appears at slightly larger scales than expected.  Although this may be explicable through a combination of mundane effects, we suggest it is also prudent to consider the possibility that the Universe might be marginally closed.  The other hint is provided by a second peak which appears less prominent than expected.  This may indicate one of a number of possibilities, including increased damping length or tilted initial conditions, but also breaking of coherence or features in the initial power spectrum.  Further data should test whether the current concordance model needs only to be tweaked, or to be enhanced in some fundamental way. ",
        "introduction": "The study of the Cosmic Microwave Background (CMB) anisotropy holds the promise of answering many of our fundamental questions about the Universe and the origin of the large-scale structure (see e.g.~Bond~\\cite{Bond}; Bennett, Turner \\& White~\\cite{BenTurWhi}; Lawrence, Scott \\& White~\\cite{Romans}). The development of CMB research can be split into 5 main phases. Firstly, the mere existence of the CMB showed that the early Universe was hot and dense.  Secondly, the blackbody nature of the CMB spectrum and its isotropic distribution implied that the Universe is approximately homogeneous on large scales. The third step came with the detection of anisotropies, confirming that structure grew through gravitational instability. Now we are entering the fourth stage, where the basic cosmological paradigm is defined. The recently released {\\sl Boomerang\\/} data (de Bernardis et al.~\\cite{BOOM}) provide support for a model with adiabatic initial conditions and a Universe with approximately flat geometry. The fact that our theories are holding up so well gives us further reason to believe that the CMB can be used as a precision cosmological tool. With the imminent launch of {\\sl MAP\\/}, we are on the verge of the fifth phase, which involves determining the precise values of the fundamental cosmological parameters to figure out exactly what kind of Universe we live in.  Most of the unmined cosmological information available from the CMB anisotropy is encoded in the acoustic signatures, the series of peaks and troughs in the spectrum at subdegree scales, which we are only now beginning to probe experimentally. Because the properties of the photon-baryon oscillations are determined by the background, while the driving force is described by the model for the perturbations, the acoustic signatures provide a unique opportunity to probe both the background cosmology and the model for structure formation. For example the position of the first peak, or indeed any other feature, provides a measure of the angular diameter distance to last scattering. The relative heights of the peaks provide information about the baryon `drag' on the photons and thus the baryon-to-photon ratio. The relative peak locations provide information on the perturbations as they crossed the horizon and thus indirectly on the mechanism for their production (see e.g.~Hu, Sugiyama \\& Silk~\\cite{HuSugSil}).  In the last year or so there have been several new CMB data sets which have begun to reveal the structure contained in the acoustic peaks (see e.g.~Lineweaver~\\cite{Charley}; Dodelson \\& Knox~\\cite{DodKno}; Melchiorri et al.~\\cite{Mel}; Pierpaoli, Scott \\& White~\\cite{Science}; Efstathiou~\\cite{Efs}; Tegmark \\& Zaldarriaga~\\cite{TegZal}; Lahav et al.~\\cite{LahBriHobLasSod}; Le Dour et al.~\\cite{LeDour}~for analyses of these data). Now with the first estimate of the power spectrum from a sub-set of the Antarctic flight of the {\\sl Boomerang\\/} experiment we are entering a whole new regime of precision.  There are 3 striking things about this new power spectrum estimate.  Firstly, and most importantly, it corroborates the basic picture of cosmological structure formation -- the shape is a confirmation of flat models of the sort inspired by inflation, dominated by a cosmological constant, as has become the standard paradigm.  Secondly, however, the position of the first peak appears at slightly larger angular scales than might have been expected.  And lastly, another possibility for something unexpected comes through a hint that the second peak may not be as pronounced as most models would predict.  We will make some general comments about the existence of the first acoustic peak, and then in the rest of this paper we focus on these latter two surprising features of the new data (see also Hu~\\cite{Hu}). ",
        "conclusions": "The {\\sl Boomerang\\/} data provide a stunning confirmation of the reality of acoustic oscillations in the photon-baryon fluid at last scattering.  The fact that the peak is at $\\ell\\,{\\sim}\\,200$ argues that the Universe is close to spatially flat.  The fact that the second peak appears to be smaller than naively expected, while explicable within standard models, could be a clue to something novel in our model of structure formation.  We have argued that the high first peak relative to the second is suggestive of tilt in the primordial power spectrum, a late epoch of matter-radiation equality and a low redshift of reionization.  The slightly leftwards position (relative to the precisely flat expectation) of the first peak argues for a short distance to last scattering, and in combination with the former this argues that the Universe may be (marginally) spatially closed. Whether the best-fitting models are significantly closed will require more high precision data.  But in any case, it is now clear that closed models need to be considered on at least an equal footing with open models when searching the cosmological parameter space.  The key to making further progress will be the detection of a third peak. Models with a high baryon content will have a high third peak, tilted models will have a lower third peak.  Lack of coherence in the oscillations would be more exciting still, since this would be harder to explain. The detection of a second feature in the power spectrum would pin down the fundamental mode of the baryon-photon fluid at last scattering and put us well on our way towards reconstructing the model of structure formation.  Further measurement of the second peak should come with analysis of the full {\\sl Boomerang\\/} 98 data-set, together with data from the {\\sl MAXIMA}\\footnote{{\\tt http://cfpa.berkeley.edu/group/cmb/}\\quad After this paper was submitted the results of the {\\sl MAXIMA}-1 flight were released (Hanany~\\cite{Hanany}). They show impressive confirmation of the basic picture presented by the {\\sl Boomerang\\/} data, including the weakness of the second peak.  Although the {\\sl MAXIMA}-1 data show no strong preference for closed models, we wish to emphasize that there is still a large region of closed model parameter space which needs to be explored.} {\\sl VSA}\\footnote{{\\tt http://www.mrao.cam.ac.uk/telescopes/vsa/index.html}}, {\\sl DASI}\\footnote{{\\tt http://astro.uchicago.edu/dasi/}} and {\\sl CBI}\\footnote{{\\tt http://astro.caltech.edu/$\\sim$tjp/CBI/}} experiments. In addition long-duration CMB balloon flights in the next couple of seasons, as well as the imminent launch of {\\sl MAP\\/}\\footnote{{\\tt http://map.gsfc.nasa.gov/}}, should produce much more precise measurements of the relevant $\\ell$ range.  The new {\\sl Boomerang\\/} results have shown a remarkable confirmation of the conventional picture for structure formation.  On top of that, it is exciting that the data show some hints of a couple of surprises.  To make it easier to fit the first peak position, it may be worth bearing in mind the possibility that the Universe may be spatially closed. And, for consistency with the structure of the subsidiary peaks, it is worth keeping an open mind to the possibility that there may yet be some important physical effects which are not contained within the simplest versions of the current standard paradigm.  \\bigskip"
    },
    "0004/astro-ph0004251_arXiv.txt": {
        "abstract": "A high resolution ($\\lambda/\\Delta\\lambda \\simeq 110\\,000$), very high $S/N$ ($\\ga 600$) spectrum of the metal-poor turnoff star \\object{G\\,271-162} has been obtained in connection with the commissioning of UVES at VLT/Kueyen. Using both 1D hydrostatic and 3D hydrodynamical model atmospheres, the lithium isotope ratio has been estimated from the \\Lione\\,670.8\\,nm line by means of spectral synthesis. The necessary stellar line broadening (1D: macroturbulence + rotation, 3D: rotation) has been determined from unblended \\Kone, \\Caone\\ and \\Feone\\ lines. The 3D line profiles agree very well with the observed profiles, including the characteristic line asymmetries. Both the 1D and 3D analyses reveal a possible detection of \\Lisix\\ in \\object{G\\,271-162}, $\\sixseven = 0.02\\pm0.01$ ($1\\sigma$). It is  discussed if the smaller amount of \\Lisix\\ in \\object{G\\,271-162} than in the similar halo star \\object{HD\\,84937} could be due to differences in stellar mass and/or metallicity or whether it may reflect an intrinsic scatter of \\sixseven\\ in the ISM at a given metallicity. ",
        "introduction": "Due to the special status of \\Lisix\\ for astrophysics and cosmology, much work has been devoted to the search for this isotope in metal-poor stars ever since the first detection in \\object{HD\\,84937} by Smith et al. (\\cite{smith93}). The reason for this interest is threefold: {\\em i)} The presence of \\Lisix\\ in the envelope of halo stars severely limits the possible depletion of \\Liseven, and thus allows a more accurate determination of the primordial \\Liseven\\ abundance (Ryan et al. \\cite{ryan99}; Asplund \\& Carlsson \\cite{asplund00}); {\\em ii)} \\Lisix\\ abundances provide an additional test of theories for the production of the light elements (Li, Be and B) by cosmic ray processes; {\\em iii)} Since \\Lisix\\ is an even more fragile nuclei than \\Liseven\\, it is a sensitive probe of possible mixing events during the stellar life.  To date \\Lisix\\ is claimed to have been detected in two metal-poor halo stars, \\object{HD\\,84937} and \\object{BD\\,+26 3578}, and two metal-poor disk stars, \\object{HD\\,68284} and \\object{HD\\,130551} (Smith et al. \\cite{smith98}; Hobbs \\& Thorburn \\cite{hobbs97}; Cayrel et al. \\cite{cayrel99}; Nissen et al. \\cite{nissen99}). Clearly, an observational test of models for the formation and evolution of \\Lisix\\ and for the depletion of \\Lisix\\ in stellar envelopes requires a much larger set of \\Lisix\\ data spanning a large metallicity range. With the advent of high-resolution spectrographs on 8m-class telescopes this is now becoming feasible.  The determinations of \\Lisix\\ abundances are based on the increased width and asymmetry of the \\Lione\\ 670.8\\,nm doublet introduced by the isotope shift (0.16\\,\\AA\\ = 7.1\\,\\kmprs) of \\Lisix. Since the line is not resolved, the derived \\Lisix\\ abundance depends on the adopted stellar line broadening as estimated from other spectral lines. Until now all such analyses have relied on 1D hydrostatic model atmospheres which cannot predict the inherent line asymmetries introduced by the convective motions in the atmosphere and thus is a potential source of uncertainty. An attractive alternative is provided by the new generation of 3D hydrodynamical model atmospheres (e.g. Stein \\& Nordlund \\cite{stein98}; Asplund et al. \\cite{asplund99}, \\cite{asp00}) which self-consistently compute the time-dependent convective velocity fields and thus are able to predict the convective line asymmetries.  In the present {\\em Letter} we analyze a high resolution, very high $S/N$ spectrum of the metal-poor halo star \\object{G271-162} obtained during the commissioning of UVES on VLT/Kueyen using both 1D and 3D model atmospheres. Thereby we also investigate the ability of UVES to provide high quality spectra as needed in many astrophysical studies. ",
        "conclusions": "Both the 1D and 3D analyses suggest a possible detection of \\Lisix\\ in \\object{G\\,271-162} at the level of $\\sixseven\\ = 0.02\\pm0.01$. Given the remaining uncertainties in the determination of the stellar line broadening (e.g. NLTE effects) the detection should, however, be considered preliminary. The similarity between the 1D and 3D results and the excellent agreement between the observed and predicted profiles in 3D, could be interpreted as support for a positive detection though. It also lends added confidence to previous detections of \\Lisix\\ in metal-poor stars based on 1D investigations (e.g. Smith et al. \\cite{smith98}; Nissen et al. \\cite{nissen99}). Evidently, the convective line asymmetries are rather negligible compared to the isotopic shift of the \\Lisix\\ doublet.  Regardless of whether \\Lisix\\ has been detected in \\object{G\\,271-162} or not, the \\sixseven\\ ratio appears in any case to be smaller than in \\object{HD\\,84937} for which a weighted average of the results of Hobbs \\& Thorburn (\\cite{hobbs97}), Smith et al. (\\cite{smith98}) and Cayrel et al. (\\cite{cayrel99}) is $\\sixseven = 0.059 \\pm 0.016 \\,\\, (1\\sigma)$. This is quite puzzling, because the two stars have almost identical parameters according to the $uvby$-$\\beta$ photometry of Schuster \\& Nissen (\\cite{schuster88}). Table 1 lists the measured indices for the three halo stars with \\Lisix\\ measurements after correction for a small reddening of G\\,271-162, $E(b-y) = 0.020$, as derived from the the $\\beta$ - $(b-y)_0$ calibration of Schuster \\& Nissen (\\cite{schuster89}). Table 2 shows the derived parameters using the IRFM calibration of \\teff\\ vs. $(b-y)_0$ by Alonso et al. (\\cite{alonso96a}), the \\feh\\ calibration of Schuster \\& Nissen (\\cite{schuster89}) and the $M_V$ calibration of Nissen \\& Schuster (\\cite{nissen91}). The directly measured IRFM temperatures of \\object{HD\\,84937} and \\object{BD\\,+26\\,3578} are 6330 and 6310 (Alonso et al. \\cite{alonso96b}) in excellent agreement with the values in Table 2, whereas \\object{G\\,271-162} has not been measured. The absolute magnitude of \\object{HD\\,84937} from the Hipparcos parallax (ESA, \\cite{esa97}) is $M_V = 3.82 \\pm 0.19$ in good agreement with the photometric value within the quoted errors. For the two other stars the relative error of the Hipparcos parallax is far too large to estimate $M_V$ with any reasonable accuracy. Finally, a preliminary abundance analysis shows that \\object{HD\\,84937} and \\object{G\\,271-162} have the same metal abundances within $\\pm 0.15$~dex.  \\begin{table} \\caption[]{Str\\\"{o}mgren photometry from Schuster \\& Nissen (\\cite{schuster88})} \\begin{tabular}{lrcccc} \\hline\\noalign{\\smallskip} ID & $V$ &  $(b-y)_0$ & $m_0$ & $c_0$ & $\\beta$  \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} HD\\,84937   & 8.33 & 0.303 & 0.056 & 0.354 & 2.613      \\\\ G\\,271-162  &10.35 & 0.306 & 0.055 & 0.355 & 2.602      \\\\ BD\\,+26\\,3578& 9.37 & 0.308 & 0.045 & 0.366 & 2.600      \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption[]{Stellar parameters and Li isotope ratios} \\begin{tabular}{lcccc} \\hline\\noalign{\\smallskip} ID & $\\teff$ & \\feh &  $M_V$ & \\sixseven    \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} HD\\,84937   & 6315\\,K  & $-2.14$ &  3.58  & $0.059 \\pm 0.016$ \\\\ G\\,271-162  & 6295\\,K  & $-2.15$ &  3.53  & $0.020 \\pm 0.010$ \\\\ BD\\,+26\\,3578& 6280\\,K  & $-2.60$ &  3.08  & $0.050 \\pm 0.030$ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}   \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{cc161.fig6.ps}} \\caption{The position of \\object{HD\\,84937}, \\object{BD+26\\,3578} and \\object{G\\,271-162} in the \\mvlgt\\ diagram compared to evolutionary tracks from VandenBerg et al. (\\cite{vandenberg00}) with masses and metallicities indicated} \\label{f:HR} \\end{figure}  Fig. \\ref{f:HR} shows the location of the three stars in the \\mvlgt\\ diagram. Clearly, \\object{HD\\,84937} and \\object{G\\,271-162} have nearly the same mass. This makes it difficult to explain the lower \\Lisix\\ abundance in \\object{G\\,271-162} as a depletion effect and raises the interesting question of a possible intrinsic scatter of \\sixseven\\ in the ISM at a given metallicity.  In this connection we note that some models for the formation of the light elements by cosmic ray processes in the early Galaxy predict a scatter of one order of magnitude in the abundance of \\Lisix , Be and B relative to Fe, e.g. the bimodal superbubble model of Parizot \\& Drury (\\cite{parizot99}) and the supernovae-driven chemical evolution model for the Galactic halo by Suzuki et al. (\\cite{suzuki99}).  It is clear that the ability of UVES on VLT/Kueyen to gather high quality, high resolution and very high $S/N$ spectra of even $V \\simeq 10\\fm5$ stars as evident from the present study, has opened up the possibility for a large survey of \\sixseven\\ for stars with a wide range of metallicities. An observing program with this exact aim using UVES/VLT is currently in progress, which is expected to be of great importance both for an improved understanding of Big Bang nucleosynthesis, cosmic chemical evolution of the light elements and stellar mixing events."
    },
    "0004/astro-ph0004298_arXiv.txt": {
        "abstract": "We present images and spectra of the Cepheus E (Cep E) region at both optical and infrared wavelengths. Only the brightest region of the southern lobe of the Cep E outflow reveals optical emission, suggesting that the extinction close to the outflow source plays an important r\\^ole in the observed difference between the optical and IR morphologies. Cep E is a unique object since it provides a link between the spectroscopic properties of the optical Herbig-Haro (HH) objects and those of deeply embedded outflows.  The observed H$_2$ infrared lines allow us to determine an excitation temperature of $\\sim 2300$~K, an Ortho-to-Para ratio of $\\sim 3$, and an H$_2$ (1,0)/(2,1) S(1) line ratio of $\\sim 9$. These results are consistent with the values observed for HH objects with detected NIR emission lines, with shock excitation as the main mechanism for their formation, and also with the values observed for embedded, NIR flows.  The optical spectroscopic characteristics of Cep E (HH~377) appear to be similar to the ones of low excitation HH objects. However, the electron density determined from the [SII]6731/6717 line ratio for this object ($n_e=$ 4100 cm$^{-3})$, and the [OI]6300/H$\\alpha$, [SII](6717+6731)/H$\\alpha$ ratios  are higher than the values of all of the previously studied low excitation HH objects. This result is likely to be the consequence of an anomalously high environmental density in the HH~377 outflow.  The ionization fraction obtained for HH~377 is $x_e \\sim 1\\%$ From this result, together  with the observed [OI]6300/H$\\alpha$ line ratio, we conclude that the observed H$\\alpha$ line emission is collisionally  excited. From a comparison with shock models, we also conclude that the extinction towards  HH~377 is very low. Comparing the observed H$\\beta$ and H$\\alpha$ fluxes of HH~377 with model predictions, we determine a shock speed  between 15 and 20 km s$^{-1}$, although somewhat higher velocities also produce spectra with line ratios that qualitatively agree with the observations of HH~377. ",
        "introduction": "The Cepheus E (Cep E) outflow was first detected in the $^{12}$CO J=1-0 transition in some of the early radio studies of star formation in molecular clouds \\markcite{sar77}. In his catalog of molecular outflows, Fukui (1989) first pointed out the presence of a bipolar, high velocity outflow in this region. The bipolar nature of Cep E became clear with the K$^\\prime$ image of the molecular outflow obtained by Hodapp (1994), which included the stronger NIR H$_2$ lines, and revealed a relatively compact system (with a size of $\\sim$ 1.5\\arcmin). Subsequent studies in the near/mid infrared have shown that the outflow is quite bright in the H$_2$ (1,0) S(1) 2.12 $\\mu$m~line, consistent with models of shock excited H$_2$ gas \\markcite{eis96,lad97,nor98}. Suttner et al. (1997) have used three-dimensional hydrodynamic simulations of highly collimated  molecular outflows in order to model the morphology of Cep E. These simulations, however, required a very high density {\\it in} the jet ($10^5$ cm$^{-3}$), which are not consistent with the hot and dense CO bullets recently found in the flow \\markcite{hat99}, with densities of $\\sim 10^4$ cm$^{-3}$.  The IRAS 23011+6126 source was originally identified as the main candidate for the outflow source. However, the presence of multiple outflows in near infrared and radio wavelengths \\markcite{eis96,lad97} indicates the existence of at least two sources, which have recently been confirmed by OVRO observations at 1.3 and 2.6mm \\markcite{tes98}. The sources are embedded and invisible at optical and near infrared wavelengths, and are likely to be Class I or Class 0 protostellar objects \\markcite{lef96,and93}. In addition, Noriega-Crespo et al. (1998)  detect one source at 6.9 $\\mu$m using ISOCAM, which is well detected in all IRAS bands.  The present study has been motivated by the detection of emission at optical wavelengths in a small section of the southern lobe of the Cepheus E outflow by Noriega-Crespo (1997) and Devine et al. (1997). Noriega-Crespo (1997) mentions that H$\\alpha$ and [SII] 6717/31 images reveal a compact knot which is the optical counterpart of the southern bowshock observed at 2 $\\mu$m by Eisl\\\"offel et al. (1996). This optical knot has been named HH~377 \\markcite{dev97}.  In this study, we explore the link between the physical properties of the outflow as determined from optical imaging and spectroscopy, and compare these results with those obtained from observations in the near infrared. Our goal is to understand the development of very young stellar outflows (we notice that Cep E has a dynamical age of $\\sim 3\\times 10^3$ years (Noriega-Crespo et al. 1998)) and the relationship between the mechanisms that produce the infrared and optical emission.   The paper is organized as follows. In Section 2 we describe the different observations obtained for this work, and comment on the reduction and calibration techniques. In Section 3, we present the results obtained from our infrared and optical observations. Finally, in Section 4 we compare the physical properties of Cepheus E deduced from the optical and  the NIR observations with other Herbig Haro objects. ",
        "conclusions": "We have carried out a spectroscopic and imaging study of the molecular hydrogen and optical atomic/ionized  emission in the  Cepheus E outflow. Our main results are:   \\noindent - For deriving the excitation state of molecular gas in Cep E we use line ratios from our H-band and K-band spectra. We find  that T$_{exc}$= 2260$\\pm$110 and 2340$\\pm$100 K for the northern and southern lobes, respectively, which are consistent with the  T$_{exc}$ measured  in other HH objects. The (1,0)/(2,1) S(1) ratios (8.50 and 9.40), and the Ortho/Para ratios  (values $\\sim 3$)  in  both lobes are also consistent with the values observed in collisionally excited objects.  \\noindent - Contrasting with the complex structure of the H$_{2}$ outflow, the optical emission is a compact, well resolved  knot (HH~377), that nearly coincides with the southern NIR lobe. The [S~II] emission of HH~377 is clearly brighter and more compact than  the H$\\alpha$ emission and its angular size is about 0.02 pc (at a distance of 0.75 kpc). The [S~II] and H$\\alpha$ peak emission spatially coincide and appear offset a few arcseconds upstream from the H$_2$ peak emission.  \\noindent - Our spectroscopic optical analysis reveals that  HH~377 has characteristics typical of low excitation Herbig-Haro objects. This is confirmed when comparing the relative fluxes of HH~377 with those of  other HH objects using  line ratio diagrams. However, HH~377 presents anomalous [SII](6717+6731)/H$\\alpha$ line ratio, larger than those obtained for objects classified as low excitation HH objects.  The electron density, $n_e$= 4100 cm$^{-3}$, determined for this object from [SII] lines would be the highest  density measured in low excitation HH objects. This value is similar to the electronic densities measured in high excitation HH objects.  \\noindent -  We estimate an ionization fraction $x_e \\sim 1\\%$ for HH~377. Together with the observed [OI]6300/H$\\alpha$ ratio, this result implies that the observed H$\\alpha$ line has to be collisionally excited. This result supports the low reddening obtained for HH~377 obtained through the shock models. Using this ionization fraction, a post-shock electron density $N_e \\sim 10^4$ cm $^{-3}$ and a compression of $\\sim 10$,  we obtain a pre-shock density ~$\\sim$10$^5$ cm$^{-3}$~for HH~377. This exceptionally high pre-shock density is very unusual for HH objects.  \\noindent - From the shock model predictions and the H$\\beta$ and H$\\alpha$ observed fluxes, we find that a shock speed between 15 and 20 km s$^{-1}$ gives the correct  total fluxes for HH~377. This velocity appears to be somewhat lower than the one deduced from the observed line ratios.  \\noindent - From a comparison between optical and infrared luminosities in HH~377 we find the possibility that the Mach disk produces the optical emission and the bow shock produces the IR emission.  \\noindent -  We have determined a visual extinction  $A_V = 2.72$ ($E(B-V)$=0.88) assuming a recombination cascade H$\\alpha$/H$\\beta$=3 Balmer decrement. If we use the H$\\alpha$/H$\\beta \\sim 6$ decrement predicted by the preferred shock wave models we obtain an $A_V \\sim 0.77$ ($E(B-V) \\sim 0.24$) extinction. Interestingly,  Lefloch et al. (1996) have obtained an $A_V = 3.4$ from the mm continuum of the southern lobe of the Cep E outflow, in qualitatively good agreement with the extinction obtained assuming H$\\alpha$/H$\\beta$=3 (see above). This confusing situation involving the optical extinction towards HH~377 and the Balmer decrement predicted from low velocity shock models could be clarified with future observations of the blue and IR [SII] lines (or, alternatively, the IR [Fe II] lines) of this object, in order to have a model-independent determination of the extinction. The overall extinction along the Cep E flow varies drastically. One could speculate different reasons (e.g. a large inclination with respect to the  plane of the sky, a dense molecular gas core surrounding the source with a rapidly decreasing density profile and/or a non-homogeneous ISM). We can not answer this question with our present data.  \\noindent - Comparing the {\\it optical} line ratios observed in south lobe of Cep E to atomic/ionic plane-parallel shock models (J-type) presented in the literature (Hartigan et al. 1994), we find that low velocity (v$_s=20$-30) shocks are the closest to reproduce the observations. These models have a high pre-shock density (n$_0=10^4$ cm$^{-3}$), and magnetic fields between 30 and 300 $\\mu$G. These conditions (high density, low ionization fraction and strong magnetic field) are quite appropriate for the development of {\\t molecular} C-type shock as well. Previous comparisons of the H$_2$ near-ir spectra with molecular shocks (Ladd \\& Hodapp, 1997), indeed indicated a preference for C-type shocks with $\\sim 35$ \\kms. Preliminary results from ISO Long Wavelength Spectrometer (50 - 200\\mum) support also this view, given Cep E rich H$_2$O spectra (Noriega-Crespo 2000), a characteristic signature for C-type shocks (Kaufman \\& Neufeld 1996ab; Noriega-Crespo et al. 2000)  \\noindent  - Finally, our analysis seems to confirm that Cep E corresponds to an outflow in its earliest developing phases. Its short dynamical age of few 10$^3$ yrs, the high gas density estimated (at least 10$^5$ \\cc) and the inhomogeneous nature of its extinction, suggest that the outflow is breaking through its placental molecular core."
    },
    "0004/astro-ph0004067_arXiv.txt": {
        "abstract": "We propose a unified picture of high magnetic field radio pulsars and magnetars by arguing that they are all rotating high-field neutron stars, but have different orientations of their magnetic axes with respective to their rotation axes. In strong magnetic fields where photon splitting suppresses pair creation near the surface, the high-field pulsars can have active inner accelerators while the anomalous X-ray pulsars cannot. This can account for the very different observed emission characteristics of the anomalous X-ray pulsar 1E 2259+586 and the high field radio pulsar PSR J1814-1744. A predicted consequence of this picture is that radio pulsars having surface magnetic field greater than about $2\\times 10^{14}$ G should not exist. ",
        "introduction": "There is growing evidence that two sub-groups of objects, namely soft $\\gamma$-ray repeaters (hereafter SGRs) and anomalous X-ray pulsars (hereafter AXPs) are magnetars (e.g. Kouveliotou et al. 1998, 1999; Hurley et al. 1999; Mereghetti \\& Stella 1995; Wilson et al. 1999; Kaspi, Chakrabarty \\& Steinberger 1999), a type of objects with dipolar magnetic fields much stronger than the critical magnetic field (Duncan \\& Thompson 1992; Paczynski 1992; Usov 1992; Thompson \\& Duncan 1995, 1996). These objects occupy a unique phase space in their combination of long, monotonically increasing periods and high period derivatives, and are believed to be a distinct species from the normal radio pulsars in that most of them are radio quiet, except for a possible detection of radio emission from SGR 1900+14 (Shitov 1999; Shitov, Pugachev \\& Kutuzov 2000). However, the recent Parkes multi-beam radio pulsar survey (e.g. Manchester et al. 2000; Camilo et al. 2000b) discovered three pulsars with dipolar field strength higher than the critical value (i.e. high magnetic field pulsars, hereafter HBPs); and one of them, PSR J1814-1744, has spin parameters quite similar to the AXP 1E 2259+586 (Camilo et al. 2000; Kaspi et al. 2000). Furthermore, a search of the archival X-ray data from the HBP PSR J1814-1744 indicates that the upper limit of the X-ray luminosity of this pulsar is approximately 1/10 that of 1E 2259+586; this led to the suggestion that HBPs and AXPs may have distinct evolutionary paths, despite their proximity in period-period derivative phase space (Pivoraroff, Kaspi \\& Camilo 2000). Here we propose a possible interpretation of the distinct emission properties of the HBPs and AXPs (especially PSR J1814-1744 and 1E 2259+586) using a simple geometric effect. ",
        "conclusions": "In this letter we discuss possible formation of the inner accelerators in a magnetar environment for the first time and come to a unified picture for AXPs and HBPs by arguing that they are all rotating high-field neutron stars, but have different orientations of the magnetic axes with respective to the rotation axes. If photon splitting suppresses pair creation near the surface, the HBPs can have active inner accelerators while the AXPs can not.  This suggestion may also have implications for another type of magnetar, i.e., the SGRs. These objects react much differently from the AXPs by exhibiting irregular short bursts and occasional giant flares, which are interpreted as crust cracking and large-scale magnetic field reconnection, respectively, within the framework of the magnetar model (Thompson \\& Duncan 1995). It remains unclear whether they are experiencing a different evolutionary stage than that of AXPs or whether they are intrinsically different objects. Only two of them (SGR 1806-20 and SGR 1900+14) have $\\dot P$ measurements, but determination of their dipolar magnetic fields is complicated by the contribution of the relativistic winds to the spin-down (Harding et al. 1999). It is notable that the constraints of both the SNR age and the magnetar energy requirements lead to a polar field $B_p({\\rm SGR})\\sim 10^{14}$ G (Harding et al. 1999), which may lie below the SCLF death line for PRs. Thus they may also have active inner accelerators if they are actually PRs. Here we suggest a possibility that SGRs might also have active accelerators while AXPs do not, and the active behaviors of SGRs may have some connections with their inner accelerators. For example, the constant extraction of electrons from the pole may somehow more frequently trigger instability within the crust. One expectation of this scenario is pulsed radio emission from the SGRs, which may account for the pulsed radio emission from SGR 1900+14 (Shitov 1999; Shitov et al. 2000).  Theoretically, the question of whether photon splitting occurs in all three modes permitted by QED or only in one mode in superstrong magnetic fields is difficult to tackle. If radio pulsar surveys discover any pulsar above the SCLF death line for PRs (the dashed line in Fig.1), these pulsars must have active V gaps with pair breakdown near the surface. This would strongly imply that only one mode of photon splitting occurs in fields above a few times $10^{14}$ G, and thus sheds important light on a fundamental physics process. It is worth noting that the location of the SCLF death line for PRs depends on the surface temperature of the neutron star, so that the location of the dashed line in Fig.1 may rise or drop. Thus detections in both radio and X-ray bands are desirable.  We thank Matthew Baring, Alex Muslimov, and Zaven Arzoumanian for interesting discussions and helpful comments."
    },
    "0004/astro-ph0004317_arXiv.txt": {
        "abstract": "White dwarfs and neutron stars are stellar objects with masses comparable to that of our sun. However, as the endpoint stages of stellar evolution, these objects do not sustain any thermonuclear burning and therefore can no longer support the gravitational load of their own mass by generating thermal pressure. Rather, matter in their interiors is compressed to much higher densities than commonly found in normal stars, and pressure is created by degenerate fermion kinetic energy and particle interactions. As a result, white dwarfs and neutron stars offer unique cosmic laboratories for studying matter at very high densities. In this review we discuss the basic properties of condensed matter at extreme densities and summarize the extent to which these properties can be examined by observations of compact objects. ",
        "introduction": "Astronomical phenomena provide many examples where matter exists in extreme conditions not found in terrestrial environments. One example is the high density of degenerate matter in ``compact objects'' - the relics of stars that have ceased burning thermonuclear fuel, and thereby no longer generate thermal pressure to support themselves against gravitational collapse. By contracting appreciably from their original sizes, the interiors of compact objects reach sufficiently high densities to produce nonthermal pressure via degenerate fermion pressure and particle interactions. Compact objects provide cosmic laboratories for studying the properties of matter at high densities.  Firm observational evidence and well-founded theoretical understanding both exist for two classes of compact objects which support themselves against collapse by cold, degenerate fermion pressure: {\\bf white dwarfs}, whose interiors resemble a very dense solid, with an ion lattice surrounded by degenerate electrons, and {\\bf neutron stars}, whose cores resemble a giant atomic nucleus - a mixture of interacting nucleons and electrons, and possibly other elementary particles and condensates. White dwarfs are supported by the pressure of degenerate electrons, while neutron stars are supported by pressure due to a combination of nucleon degeneracy and nuclear interactions. These unique states of matter are achieved by significant compression of stellar material. Table~\\ref{tab:sizes} compares the principal physical quantities of a typical white dwarf and neutron star with those of the sun\\footnote{Throughout this review we will be using units which are the standard in astrophysical research: cgs for microscopic properties of matter and solar units (denoted by $\\odot$) for macroscopic properties of astronomical objects.}.  \\begin{table}[h] \\caption{Parameters for the sun and a typical white dwarf and neutron star. \\label{tab:sizes}} \\begin{center} \\begin{tabular}{c c c c c c} object & mass                & radius & mean density & mean pressure & $GM/Rc^2\\;\\;^{b}$ \\\\ & ($M_\\odot\\;\\;^{a}$) &  (km)  & (\\gmcmc)  & (\\dncms) &     \\\\ \\hline Sun & 1          & $\\sim 7\\times 10^5$ & $\\sim 1$     & $\\sim 10^{12}$ & $\\sim 10^{-6}$    \\\\ White Dwarf & $\\leq 1.4$ & $\\sim 5\\times 10^3$ & $\\sim10^7$   & $\\sim 10^{24}$ & $\\sim 10^{-4}$\\\\ Neutron Star &   1-3   & $\\sim 10$           & $\\sim10^{14}$ & $\\sim 10^{34}$ & $\\sim 10^{-1}$ \\\\ \\end{tabular} \\end{center} $^{a}$ - $M_\\odot\\equiv 1.989\\times10^{33}\\;$gm = one solar mass.\\\\ $^{b}$ - This ratio measures the importance of relativistic gravitation, i.e., general relativity.  \\end{table}  Condensed matter in compact objects spans an enormous range of densities, which we loosely refer to as ``high densities''. These extend from about $7\\:$\\gmcmc (e.g., the density of terrestrial $^{56}_{26}$Fe), at the surface of a cold neutron star or white dwarf, to as much as $\\rho\\approx 10^{15}\\;$\\gmcmc, several times the density in atomic nuclei, in the cores of neutron stars. Matter at the various densities found in compact objects exhibits a variety of novel properties. Electromagnetic, strong, and weak interactions all play an important role in determining the character of compact objects. Since these objects are bound by gravity, they are a meeting point of all four of the fundamental forces of nature. Correspondingly, the astrophysics of white dwarfs and neutron stars incorporates a wide variety of physics including nuclear, particle, solid state and gravitation physics, to name a few areas.  In this review we briefly survey the theory of condensed matter at high densities in compact objects and illustrate how the basic theory is tested through astronomical observations. Since we must cover fifteen orders of magnitude in density, our presentation is at most introductory in nature, and we encourage the interested reader to pursue the cited references. A detailed introduction to the physics of high density matter and compact objects can be found in the textbook {\\it Black Holes, White Dwarfs and Neutron Stars: the Physics of Compact Objects}, by Shapiro and Teukolsky \\cite{ShaTeu83}.  We begin by considering the fundamental nature of cold ($T=0$) high density matter in  \\Sec~\\ref{sect:EoS}, and distinguish between different regimes of high density. In \\Sec~\\ref{sect:WDNSconf} we connect these microscopic properties with the fundamental macroscopic parameters of a compact object through the hydrostatic equilibrium dependence of mass and radius on central density. We summarize the fundamental predictions regarding the structure of white dwarfs and neutron stars. In \\Sec~\\ref{sect:WDs}-\\ref{sect:NSs} we examine how observations of white dwarfs and neutron stars can be used to probe the properties high density matter. We briefly discuss the perturbative effects of a finite temperature in \\Sec~\\ref{sect:Therm}. ",
        "conclusions": "\\label{sect:CONC}  The theoretical and observational study of compact objects remains one of the most exciting fields in modern astronomy. In essence, this research is also an exploration of the properties of condensed matter at extreme densities. Predictions regarding the properties of white dwarfs and neutron stars serve to test our understanding of matter at these high densities, while theories of high density matter serve as a basis for interpreting observational results regarding these objects. Most exciting, these objects bring together all four of the fundamental forces of nature and probe regimes not accessible in the terrestrial laboratory. They provide the most numerous and accessible sample of objects where relativistic gravitation - general relativity - plays a role in determining their physical properties.  In this review we have described the tight interconnection between the microscopic (local) properties of condensed matter at high densities and the macroscopic (global) properties of white dwarfs and neutron stars. While the fundamental principles of cold, high density matter are believed to be well understood, and are generally consistent with observations, key questions still remain, and new observations may give rise to new puzzles. The current boom in capabilities of Earth-bound telescopes and satellite instrumentation promises that many more puzzles - and hopefully, answers - are in store regarding the nature of cosmic matter at high densities.   \\def\\vol#1{{\\bf #1}} \\def\\aap{{\\it Astron.~Astrophys.}} \\def\\aj{{\\it Astron.~J.}} \\def\\apj{{\\it Astrophys.~J.}} \\def\\apjl{{\\it Astrophys.~J.~Lett.}} \\def\\apjs{{\\it Astrophys.~J.~Supp.}} \\def\\araa{{\\it Ann.~Rev~Astron.~Astrophys.}} \\def\\arnp{{\\it Ann.~Rev~Nuc.~Pat.~Sci.}} \\def\\nat{{\\it Nature}} \\def\\mnras{{\\it Mon.~Not.~Roy.~Astron.~Soc.}} \\def\\nphysa{{\\it Nucl.~Phys.~A}} \\def\\physletb{{\\it Phys.~Lett.~B}} \\def\\physrep{{\\it Phys.~Rep.}} \\def\\physrev{{\\it Phys.~Rev.}} \\def\\prc{{\\it Phys.~Rev.~{\\bf C}}} \\def\\prl{{\\it Phys.~Rev.~Lett.}} \\def\\rmphys{{\\it Rev.~Mod.~Phys}}"
    },
    "0004/astro-ph0004121_arXiv.txt": {
        "abstract": " ",
        "introduction": "The ultimate motivation for this article is the problem of explaining one of the salient observational features of isolated  (non-binary) pulsars, which is that comparatively long periods of continuous ``spin down'' of the observed frequency $\\Omega$ are occasionally interrupted by small ``glitches''. Such a glitch consists of a sudden small increase, $\\delta\\Omega$ say, that partially cancels the continuous negative variation  $\\Delta \\Omega$ that has been accumulated since the preceding glitch.  Since very soon after its  discovery in 1968, it has been generally agreed that the pulsar phenomenon  is attributable to a strong magnetic field anchored in the outer crust layers of a central neutron star. The observed frequency $\\Omega$ is to be interpreted as the rotation frequency of the outer crust layer, whose continuous spin down is evidently due to the continuous decrease of the angular momentum  $J$ due to radiation from the external magnetosphere. After thirty years of work, two basic problems remain.  The first is to account for  the spectrum (from radio to X-ray and beyond) and the detailed pulse structure of the radiation, which are presumed to depend on  the still very poorly understood workings of the magnetosphere.  The second problem -- the one with which the present article is concerned -- is to account for the frequency ``glitches''. It is generally recognised that the glitches must be explained in terms of what goes on in the interior of the neutron star, and it is also generally believed that the glitch phenomenon is essentially related to  the  property of  solidity that is predicted (on the basis of simple, generally accepted theoretical considerations) to characterise the crust of the neutron star after it has fallen below the relevant extremely high melting temperature, which occurs very soon after its formation.   The purpose of this article is to draw attention to the potential importance, as a mechanism for glitches, of the stresses induced in the crust just by the effective force arising from the deficit of centrifugal buoyancy that will be present whenever there is differential rotation.  It is to be noticed that centrifugal buoyancy is a phenomenom that  has been previously considered in the context of neutron stars, at least with reference to  one of its possible consequences, namely Ekman pumping. This is a mechanism that can considerably shorten the timescale needed for the redistribution of angular momentum (in  comparison with viscous diffusion characterized by the timescale  given by $\\tau_{\\rm visc}\\approx R_\\ast^{\\,2}/\\nu_\\ast$ where $\\nu_\\ast$ is the typical  kinetic viscosity coefficient and $R_\\ast$ is the relevant stellar radial length scale) and thus the damping of differential rotation in cases for which  (as will be the case in a typical pulsar) the star is rotating fast enough for the corresponding rotation timescale $\\tau_{\\rm rot}=2\\pi/\\Omega$ to be short compared with $\\tau_{\\rm visc}$. In such circumstances, ``Ekman pumping''  will supplement the very slow diffusive transport by more rapid convective transport propelled by centrifugal buoyancy forces. The ensuing ``Ekman timescale'' $\\tau_{_{\\rm E}}$ for the effective damping of differential rotation in such cases will be given roughly by the geometric mean of the pure diffusion and rotation timescales, i.e. $\\tau_{_{\\rm E}}\\approx \\sqrt{\\tau_{\\rm rot}\\tau_{\\rm visc}}$.  While it has been recognized that either Ekman pumping or magnetic coupling is in general efficient to bring into corotation the {\\it core plasma} with the crust \\cite{Easson79}, it is expected that Ekman pumping is quite inefficient (see e.g. \\cite{Epstein95}) for the uncharged {\\it crust neutron superfluid}  that is believed (see e.g. \\cite{sauls}) to permeate the lower layers of the crust in the density range from $10^{11}$ to about $10^{14}$ gm/cm$^3$. This means that the convectively accelerated Ekman timescale, $\\tau_{_{\\rm E}}\\approx R_\\ast \\sqrt{2\\pi/\\nu_\\ast\\Omega}$, is too long to prevent the development of significant differential rotation. The negligibility, in such cases, of Ekman pumping is attributable to the effective negligibility of viscosity, but should not be construed as implying the negligibility of centrifugal buoyancy forces. In previous discussions of such scenarios -- and in particular of the simplified strictly stationary limit in which the effective viscosity is neglected, so that no possibility of Ekman pumping can arise at all -- the role of centrifugal buoyancy forces has been rather generally overlooked. The upshot of the present investigation of stationary differentially rotating configurations is to show that in such cases the general neglect of the centrifugal buoyancy effect is quite unjustified, and that on the contrary this effect is potentially capable by itself of providing the dominant contribution to the crust stresses that are ultimately released in ``glitches''. ",
        "conclusions": "In the lower crust region that seems most likely to be relevant for the explanation of the large glitches observed in the Vela pulsar one would expect  the corotating constituent to be characterised by a density $\\rho_{\\rm c}$ (attributable mainly to protons and bound neutrons in the atomic type ions forming a solid lattice) having a range of values that is roughly comparable with that of the corresponding neutron superfluid density $\\rho_{\\rm n}$ (quantitatively round about $10^{13}$ g/cm$^3$).  Thus although they are of opposite sign (tending to push the crust material outward in the case (\\ref{65}) of vortex pinning, but to push it inwards in the case (\\ref{67}) for which pinning is absent) the alternative formulae (\\ref{72}) and (\\ref{72bis}) both predict the same rough order of magnitude for the stress induced on the crust by the existence of a difference between the angular velocity $\\Omega_{\\rm n}$ of the neutron superfluid constituent and the (externally observable) angular velocity $\\Omega$ characterising the crust.  The implication is that, as a candidate for explaining the large magnitude of the discontinuous changes $\\delta\\Omega$ that are commonly observed in a pulsar such as Vela, the previously overlooked buoyancy deficit mechanism characterised by the formula (\\ref{72}), i.e. \\be f_{\\rm s}^{\\, i}\\simeq \\rho_{\\rm c}\\big(\\Omega_{\\rm n}-\\Omega\\big)\\, \\Omega_{\\rm n}\\nabla^i\\varpi^2\\, , \\eqn{73}\\fe (pushing outward along the cylindrical radial direction) seems at first sight to be just as promising as the more thoroughly investigated vortex pinning mechanism, which, if the chemical contribution $f_{\\rm x}^{\\, i}$ were unimportant, would be given according to (\\ref{72bis}) by \\be f_{\\rm s}^{\\, i}\\simeq -\\rho_{\\rm n}\\big(\\Omega_{\\rm n}-\\Omega\\big)\\, \\Omega_{\\rm n}\\nabla^i\\varpi^2\\ ,\\eqn{74} \\fe (pushing inward along the cylindrical radial direction). In order to obtain definitive conclusions it is clear however that much more work on both kinds of mechanism will be needed. In particular it will be necessary to pay more attention than hitherto to the role of the chemical excess force (\\ref{55}).  The present situation can be summarised by the statement that the large magnitude of the observed glitches in Vela provides  strong evidence for the existence of angular velocity differences -- and hence for the existence of superfluidity -- in the pulsar interior, but that it is premature to claim  it also provides strong evidence for vortex pinning because stresses of comparable magnitude could be produced in the absence of pinning by the centrifugal buoyancy deficit mechanism.   \\bigskip \\noindent {\\bf Acknowledgments} \\medskip  We wish to thank R. Prix and P. Hansel for very valuable discussions. One of us (D.S.) would like to acknowledge ``Jumelage France-Arm\\'enie\" exchange programme for financial support."
    },
    "0004/astro-ph0004192_arXiv.txt": {
        "abstract": "Stellar fluxes from the 2MASS catalog are used to remove the contribution due to Galactic stars from the intensity measured by DIRBE in four regions in the North and South Galactic polar caps. After subtracting the interplanetary and Galactic foregrounds, a consistent residual intensity of \\KkJy\\ or \\KnW\\ at 2.2 $\\mu$m is found. At 1.25~$\\mu$m the residuals show more scatter and are a much smaller fraction of the foreground, leading to a weak limit on the CIRB of \\JkJy\\ or \\JnW\\ (1 $\\sigma$). ",
        "introduction": "The Diffuse InfraRed Background Experiment (DIRBE) on the COsmic Background Explorer ({\\sl COBE}, see \\citet{Bo92}) observed the entire sky in 10 infrared wavelengths from 1.25 to 240 \\um. \\citet{HAKDO98} discuss the determination of the Cosmic InfraRed Background (CIRB) by removing foreground emission from the DIRBE data. This paper detected the CIRB at 140 and 240 \\um, but only gives upper limits at shorter wavelengths.  From 5 to 100 \\um, the zodiacal light foreground due to thermal emission from interplanetary dust grains is so large that no reliable estimates of the CIRB can be made from a position 1 AU from the Sun \\citep{KWFRA98}.  In the shorter wavelengths from 1.25 to 3.5 \\um, the zodiacal light is fainter, but uncertainties in modeling the foreground due to Galactic stars are too large to allow a determination of the CIRB \\citep{AOWSH98}.  Recently, \\citet{GWC00} removed the Galactic star foreground by directly measuring the stars in a $2\\deg\\times2\\deg$ box using ground-based telescopes and then subtracting the stellar contribution from the DIRBE intensity on a pixel-by-pixel basis.  This field, a DIRBE dark spot, was selected using DIRBE data to have a minimal number of bright Galactic stars.  In addition, \\citet{WR00} used a histogram fitting method to remove the stellar foreground from the DIRBE data in a less model-dependent way than that used by \\citet{AOWSH98}. \\citet{GWC00} and \\citet{WR00} obtained consistent estimates of the CIRB at 2.2 and 3.5 \\um.  With the recent 2$^{nd}$ incremental release of 2MASS data, it is now possible to apply the direct subtraction method of \\citet{GWC00} to four additional DIRBE dark spots scattered around the North and South Galactic polar caps.  \\citet{KO00} have claimed a detection of the fluctuations of the CIRB. \\citet{KO00} also give the range 0.05 to 0.1 for the ratio of the fluctuations in the DIRBE beam to the mean intensity for the CIRB. But this ratio and fluctuation combine to give a range of CIRB values that is incompatible with the \\citet{HAKDO98} upper limits on the CIRB, especially at 1.25 \\um. Furthermore, the claimed cosmic fluctuations are larger than the residuals in the DIRBE$-$2MASS fits presented in \\S\\ref{sec:analysis}. In this paper, \\citet{KO00} is treated as an upper limit on the CIRB which is compatible with previous limits and the results found here. \\citet{Wr01} will discuss the possible cosmic fluctuation signal in the DIRBE$-$2MASS residuals in more detail. ",
        "conclusions": "\\begin{figure}[tbp] \\plotone{f7.eps} \\caption{% Comparison of CIRB values to previous determinations and upper limits. Lower limits from source counts from \\protect\\citet{SIBK99} at 850 \\um, \\protect\\citet{1997fisu.conf..159F} at 15 \\& 6.7 \\um, and \\protect\\citet{MP00} at 2.2 to 0.3 \\um. Solid upper limits from \\protect\\citet{HAKDO98}, open upper limit symbols using $\\gamma$-rays from \\citet{FMMRW98} and \\citet{SF98}. Open squares at 240 \\& 140 \\um\\ from \\protect\\citet{HAKDO98}, open circles at 100 \\& 60 \\um\\ from \\citet{FDS00}, while the filled circle far IR data points are the \\citet{HAKDO98} results modified by using this paper's zodiacal model. Dashed curve is from \\protect\\citet{FDMBS98}. Filled circles from 3.5 to 1.25 \\um\\ are an average of \\protect\\citet{GWC00}, \\protect\\citet{WR00} and this paper. Open circles from 0.8 to 0.3 \\um\\ are from \\protect\\citet{1999hrug.conf..487B}. Open squares are from \\protect\\citet{DWW77} and \\protect\\citet{T83}. The open diamond at 0.15 \\um\\ is from \\protect\\citet{HBM91}. \\label{fig:COIBR-prev}} \\end{figure}  Subtracting the 2MASS catalog from the zodi-subtracted DIRBE data yields a statistically significant, isotropic background at 2.2 \\um\\ of \\KkJy\\ which is consistent with the earlier results from \\citet{GWC00} $(16.4 \\pm 4.4\\;\\kJysr)$ and \\citet{WR00} $(16.9 \\pm 4.4\\;\\kJysr)$ within the systematic error associated with the modeling the zodiacal dust cloud. Averaging the results of \\citet{GWC00}, \\citet{WR00} and this paper gives a CIRB at 2.2 \\um\\ of $16 \\pm 4\\;\\kJysr$.  This averaging has not reduced the estimated error which is dominated by systematic effects that affect all three results equally. The foreground due to interplanetary dust at 1.25 \\um\\ is too large to allow a CIRB detection, but an improved upper limit is found. Note that the Zodi-Subtracted Mission Average maps which used the \\citet{KWFRA98} zodiacal light model give a CIRB that is 13.75 \\kJysr\\ larger at 1.25 \\um\\ and 6.08 \\kJysr\\ larger at 2.2 \\um\\ than results obtained here using the zodiacal light model described in \\citet{Wr98} and \\citet{GWC00} based on the very strong no-zodi principle of \\citet{Wr97}. Figure \\ref{fig:CBR-JK} shows a plot of the \\Jband\\ intensity \\vs\\ \\Kband\\ intensity averaged over the four dark spots analyzed in this paper.  Three values are plotted: the average total intensity $\\langle \\mbox{D} \\rangle$, the average zodi-subtracted intensity $\\langle \\mbox{DZ} \\rangle$, and the CIRB estimates. The \\citet{HAKDO98} upper limits on the CIRB, the \\cite{DA98} correlation and the $1 \\sigma$ error bars from this paper are shown as well.  This figure emphasizes the large subtractions that are involved in determining the CIRB from data taken 1~AU from the Sun: the zodiacal light is about 16 times larger than the CIRB at 1.25 \\um\\ and 8 times larger than the CIRB at 2.2 \\um. Galactic stars are a problem in the large DIRBE beam, but in the selected dark spots the effect of stars is 4 times less than that of the zodiacal light.  \\citet{1999hrug.conf..487B} has measured the optical extragalactic background light and obtained results at $\\lambda = 0.8,\\;0.55,\\; \\mbox{\\&}\\;0.3\\;\\um$ which are consistent with a reasonable extrapolation through the uncertain \\Jband\\ result found here, as shown in Figure \\ref{fig:COIBR-prev}. Both \\citet{1999hrug.conf..487B} and this work face challenging and uncertain corrections for the zodiacal light, but the two papers use very different techniques and should not have systematic errors in common.  Thus the lack of a discontinuity in the spectrum between 0.8 and 1.25 \\um\\ is an indication in favor of the background level reported here. The model shown in Figure \\ref{fig:COIBR-prev} is the $\\Lambda$CDM-Salpeter model from \\citet{PBSN99} which appears to fit the observed far IR to near IR to optical ratios. But the model was multiplied by 1.84 to match the level of the observed background."
    },
    "0004/nucl-th0004017_arXiv.txt": {
        "abstract": "\\parbox{14cm}{\\rm We present shell model calculations of both the structure of $^{17}$F and the reactions $^{16}$O(p,$\\gamma$)$^{17}$F, $^{16}$O(p,p)$^{16}$O. We use the ZBM interaction which provides a fair description of the properties of $^{16}\\mbox{O}$ and neighbouring nuclei and, in particular it takes account for the complicated correlations in coexisting low-lying states of $^{16}\\mbox{O}$.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004009_arXiv.txt": {
        "abstract": "We investigate, on the basis of CCD Str\\\"omgren photometry, the ages and metallicities of six LMC clusters together with their surrounding field population. The clusters and metallicities are: NGC~1651 (in the range $[Fe/H]=-0.65$ dex to $-0.41$ dex), NGC~1711 ($-0.57\\pm0.17$ dex), NGC~1806 ($-0.71\\pm0.23$ dex), NGC~2031 ($-0.52\\pm0.21$ dex) and NGC~2136/37 ($-0.55\\pm0.23$ dex) and NGC~2257 ($-1.63\\pm0.21$ dex).  The metallicities for NGC~1651, NGC~1711, NGC~1806 and NGC~2031 have been determined for the first time (NGC~2031 and NGC~2136/37 are interesting for the Cepheid distance scale).  In the cluster surroundings, we found about 650 field stars that were suitable to be used for a determination of an age-metallicity relation (AMR).  Our method is to estimate ages for individual stars on the basis of Str\\\"omgren isochrones with individually measured metallicities.  With this method we are able to sample the AMR of the field population up to 8 Gyr.  Our metallicity data are incompatible with models predicting many metal-poor stars (G-dwarf problem). The metallicity of the field population increased by a factor of six, starting around 2 Gyr ago. The proposed AMR is consistent with the AMR of the LMC cluster system (including ESO 121 SC03 and three clusters with an age of 4 Gyr).  The proposed AMR is incompatible with the recently proposed AMR by Pagel \\& Tautvai\\u{s}vien\\.e (\\cite{pagel98}). ",
        "introduction": "\\label{chap:intro}  In spite of its enormous importance for understanding galaxy evolution in adequate detail, the chemical enrichment process in galaxies is still poorly known, which is especially true for the field star component.  The Large Magellanic Cloud (LMC) is a natural target to study the chemical evolution because of its proximity. Also, its structure seems to be less complex than that of the Milky Way which might imply that the chemical enrichment history can be described by a simple global age-metallicity relationship (AMR).  First efforts to determine the AMR of LMC clusters have been made with integrated broad band photometry of clusters (Westerlund \\cite{westerlund97}).  Recent work continuing these studies is, for example, Bica et al. (\\cite{bica98}) and Girardi et al. (\\cite{girardi95}). Another major step towards an understanding of the LMC cluster AMR has been undertaken by Olszewski et al. (\\cite{olszewski91}), who used medium resolution spectroscopy of individual giants to measure the metallicity for around 70 clusters, with a quoted uncertainty of $\\pm 0.2$ dex.  In addition many photometric studies of stars in LMC clusters (e.g. with the Washington system by Bica et al. \\cite{bica98}) contributed to the unveiling of the cluster AMR.  The current wisdom on the cluster AMR that has been established by these studies is, that the mean metallicity of younger clusters is distinctly higher than that of old clusters by more than 1.2~dex. However, it is difficult to trace the AMR over the entire LMC history with this cluster sample, since for a long time, only one cluster (ESO121-SC03) with an age between 3~Gyr and 11~Gyr had been found (Mateo et al. \\cite{mateo86}, Bica et al. \\cite{bica98}).  Recently, Sarajedini (\\cite{sarajedini98}) found three more clusters with an age of about 4~Gyr (NGC~2121, NGC~2155 and SL~663).  The AMR as derived from LMC clusters shows a very large scatter (Olszewski et al.  ~\\cite{olszewski91}), which, if intrinsic and not due to measurement uncertainties, would argue for a more complex chemical enrichment history.  In addition there are hints that at least some clusters have smaller mean metallicities than the surrounding field population (e.g. Bica et al.~\\cite{bica98}, Richtler et al. \\cite{richtler89}). Thus possibly the chemical evolution of the cluster and field stars is to some degree decoupled. However, this is not without contradiction (e.g. Korn et al. \\cite{korn00}).  Santos Jr. et al. (\\cite{santos99}) claimed that the metallicity dispersion of the field seems to be smaller than that of the cluster system of similar age.  For the field population the metallicity distribution is known primarily for the young stars since mainly F \\& G supergiants have been spectroscopically investigated (e.g. Hill et al. \\cite{hillV95}, Luck \\& Lambert \\cite{luck92}, Russell \\& Bessell \\cite{russell89}). A compilation of young LMC field stars abundances which have been derived with high resolution spectroscopy can be found in the appendix (Table~\\ref{tab:fieldme}).  Th\\'evenin \\& Jasniewicz (\\cite{thevenin92}) study 9 field stars in the LMC with medium resolution spectroscopy (5 $\\AA$) and found an average abundance of $[Fe/H]=-0.25\\pm0.08$ which is higher than the mean value of field stars that has been derived with high resolution spectroscopy ($-0.38\\pm0.11$ dex). Dopita et al. (\\cite{dopita97}) measured element abundances of planetary nebulae (PNs) in the LMC and derived their age by modelling the hot, central star. They found four PNs that are older than $4$ Gyr. Their AMR shows only little enrichment from $15$ to $5$ Gyr ago, while the metallicity doubled in the last $2-3$ Gyr.  The study of the older stellar field component has been limited to studies using broad band photometry (e.g. Holtzman et al. \\cite{holtzman99} and Elson et al. \\cite{elson97}).  We used a different approach and measured the metallicity of individual stars by using the medium wide Str\\\"omgren filter system, that gives a good metallicity discrimination for giants and supergiants red-wards of $b-y = 0.4$~mag. This method has already been used by Grebel \\& Richtler (\\cite{grebel92}), Hilker et al. (\\cite{hilker95b}) and Hilker et al.  (\\cite{hilker95a}) to determine age and metallicity of NGC~330, NGC~1866 and NGC~2136/37. Ardeberg et al. (\\cite{ardeberg97}) used HST observations transformed into the Str\\\"omgren system to derive the SFH and the metallicity of LMC bar stars. Their investigation differs from our approach by the calibration they employed which is based on bluer stars and includes the gravity dependent $c1$ Str\\\"omgren colour index.  In the current work we investigate mainly young LMC clusters and their surrounding fields, namely NGC~1651, NGC~1711, NGC~1806, NGC~2031, and NGC~2257, an old cluster.  We have also re-analysed NGC~2136/37 because of the availability of Str\\\"omgren isochrones and a new calibration for photometric metallicities, which improves the calibration for more metal poor stars. This ensures the homogeneity of the sample and also tests if systematic shifts are present between the older investigations and the new one.  An important aspect of this new work is exactly this homogeneity of the metallicities allowing one to assess the real magnitude of the intrinsic dispersion among metallicities of clusters of similar age.  Two of the clusters (NGC~2136 and NGC~2031) are particularly interesting because they contain Cepheid variables, whose metallicities are important to know for distance scale problems. NGC~1866 might serve as example.  Its metallicity has been determined by Hilker et al. (\\cite{hilker95b}) via Str\\\"omgren photometry which was used for the distance determination using its Cepheid members by Gieren et al. (\\cite{gieren94}). ",
        "conclusions": "\\subsection{The Age-Metallicity Relation}  \\begin{table*}[t] \\begin{center} \\caption{ Results for the 5 clusters investigated in this work. The error includes the calibration uncertainty.} \\label{table:result} \\begin{tabular}{ccccc} \\hline Cluster & $E_{B-V}$ & Metallicity [dex] & log(Age [y]) & Remarks \\\\[0.5ex] \\hline NGC~1651 & $0.01$ to $0.05$ & $-0.65$ to $-0.45 $  & $9.4$ to $9.1$ & reddening problematic\\\\ NGC~1711 & $0.09 \\pm 0.03$ & $-0.57\\pm0.17$ & $ 7.7\\pm0.05 $ & reddening of the field is larger\\\\ NGC~1806 & $0.16 \\pm 0.06$ & $-0.71\\pm0.24$ & $8.7\\pm0.1$ & \\\\ NGC~2031 & $0.09 \\pm 0.05$ & $-0.52\\pm0.21$ & $8.2\\pm0.1$ & \\\\ NGC~2136/37 & $ 0.09 \\pm 0.05$ & $-0.55\\pm0.23$ & $8.0\\pm0.1$ & no differences between the two clusters\\\\ NGC~2257 & $0.04 \\pm 0.04$ & $-1.63\\pm0.21$ & $10.2\\pm0.1$& \\\\[0.5ex]\\hline \\end{tabular} \\end{center} \\end{table*}  In Table\\ref{table:result} we summarise the resulting ages and metallicities of the investigated clusters.  In order to compare our results with the literature, we compiled a list of clusters with ages and metallicities according to various sources (all published after 1989).  The data is tabulated in Table\\ref{tab:LitData}. In addition, we used the compilation of Sagar \\& Pandey (\\cite{sagar89}), from which only clusters have been selected with a limiting magnitude below $V\\!=\\!21$. This limit shall serve as a rough quality criterion that is comparable to the more recent data and explains why most (photographic) papers cited by Sagar \\& Pandey are excluded.  These clusters are plotted together with the newly investigated clusters and our field AMR in Fig.~\\ref{fig:LitData}. The solid line is the field AMR accompanied by two dotted lines which mark 1$\\sigma$ borders. If the older clusters are excluded, a weak correlation appears for the clusters: clusters younger $1$ Gyr have a mean metallicity of $[Fe/H]=-0.34$ with a standard deviation of 0.14 and in the age range $1-2.5$ Gyr the mean metallicity is $[Fe/H]=-0.71$ with a standard deviation of $0.17$ (11 cluster).   \\begin{figure*}[t] \\resizebox{12cm}{!}{\\includegraphics{fig32.ps}} \\hfill \\parbox[b]{55mm}{ \\caption{Published ages and metallicities for LMC clusters, given in Table\\ref{tab:LitData} (open triangles) and in the compilation by Sagar \\& Pandey (\\cite{sagar89}) with a limiting magnitude of fainter than $V=20$ (open squares). The cluster data from this investigation are marked with solid circles. The solid line connects the points in our derived AMR for the field population. The dotted line surrounding the solid line marks the standard deviation of the metallicity around a given age. In addition three models for the AMR are plotted: Pagel \\& Tautvai\\u{s}vien\\.e (\\cite{pagel98}) as short dashed line and two models calculated by Geha et al. (\\cite{geha98}) as long dashed line.} \\label{fig:LitData} } \\end{figure*}  The mean metallicity of our young clusters ($<10^{9.0}$ yr) is $-0.57 \\pm 0.04$ dex and thus lower than what we found using the newer cluster sample from the literature. The field of the same age has a mean metallicity of $-0.4\\pm0.2$ dex and is in good agreement with spectroscopic measurements ($-0.38\\pm0.11$ dex) of young field stars. The latter comparison is reasonable since {\\bf a)} most of the stars for which high resolution spectroscopy has been obtained are located at a similar radial distance and {\\bf b)} no radial gradient can be seen in the spectroscopic sample. We compiled a list of high resolution spectroscopic measurements of LMC stars in Table~\\ref{tab:fieldme}.  Bica et al. (\\cite{bica98}) obtained ages and metallicities of 13 outer clusters in the LMC using Washington photometry.  The mean metallicity of all the surrounding field stars is $\\langle[Fe/H]\\rangle\\simeq-0.6\\pm0.1$. The mean metallicities of our field populations seem to be systematically more metal poor than this value thus indicating a possible zero point difference of the order of 0.2 dex and comparable to the probable shift between our cluster metallicities and the ones taken from the literature.  However, such a difference between cluster and field stars has not been seen in a study by Korn et al. (\\cite{korn00}) who employed high resolution spectroscopy of supergiants.  The AMR for stars older than $3$ Gyr is consistent with little or even no enrichment until 8 Gyr ago.  The AMR in this age range agrees well with the $4$ Gyr old clusters studied by Sarajedini (\\cite{sarajedini98}) and also ESO 121 SC03 is in agreement with the derived field star AMR, especially when taking the systematic underestimation of the metallicity for older stars on the order of $0.05-0.1$ dex into account (see Sect.~3). Therefore, we do not see the necessity that ESO 121 SC03 belongs to a dwarf galaxy that is in the process of merging with the LMC as proposed by Bica et al. {\\cite{bica98}).  The field population of NGC~1651 and NGC~2257 is considerably different from that around the other clusters: around NGC~1651 we find two distinct field populations, around NGC~2257 only one, thus these fields cannot be compared to the other fields, where a mixture of populations have been detected. These fields contain a significantly larger fraction of old stars than the other fields, what is expected from their location in the LMC (e.g. Santos Jr. et al. \\cite{santos99}).  If cluster and field are compared it becomes apparent that our AMR does not argue for an extremely decoupled enrichment history between cluster and field stars, only hints can be seen that the younger clusters are slightly more metal poor than the surrounding field population of the same age. Bica et al. (\\cite{bica98}) found the same behaviour for several of their (young) clusters and the surrounding field population.  One has to consider the possibility, that these low mean cluster abundances are a result of the statistically larger effect of blending towards the cluster. This has been proposed by Bessell (\\cite{bessell93}) to explain the low Str\\\"omgren metallicity of NGC~330 measured by Grebel \\& Richtler (\\cite{grebel92}). In this work of Grebel \\& Richtler (\\cite{grebel92}) the mean metallicity found for the surrounding field population ($-0.74$ dex) agreed well with later on performed spectroscopic measurements ($-0.69$ dex, Hill \\cite{hillV99}) ( for a more comprehensive discussion on NGC~330 the reader is referred to the work by Gonzalez \\& Wallerstein \\cite{gonzalez99}).  Having this agreement in mind, one can estimate from the difference of mean field and cluster metallicity that the contamination has a minor effect on our derived metallicities, accounting possibly for a systematic deviation of less than $<0.15$ dex.  Our AMR is inconsistent with a recent calculation presented by Pagel \\& Tautvai\\u{s}vien\\.e (\\cite{pagel98}) based on LMC clusters and on planetary nebulae observed by Dopita et al. (\\cite{dopita97}), that predicts a steeper increase of the metallicity in earlier time, thus older stars should have a higher metallicity than what we observe (see Fig.~\\ref{fig:LitData}). The AMR is more consistent with closed box model calculations performed by Geha et al. (\\cite{geha98}). They present theoretical enrichment models for the two SFHs put forward by Holtzman et al. (1997) and by Vallenari et al. (1996a,b). These SFHs agree in the sense, that a long period of low star formation activity was followed by a sudden increase about 2 Gyr ago.  With the Vallenari et al.- SFH, the metallicity increased by a factor of five during the last 2 Gyr, while a modest increase of a factor of three resulted from the Holtzman et al.-SFH.  Dopita et al. (\\cite{dopita97}) published an AMR for the LMC based on planetary nebulae and found that the metallicity only doubled in the last 2-3 Gyr which is seen in our AMR as well. Another common feature is that a distinct enrichment (if any) between 4 and 9 Gyr cannot be seen. A comparison of the Dopita et al.-values with ours is made difficult by the fact that they measured $\\alpha$-element abundances instead of $[Fe/H]$, but they stated that {\\it \"there is no evidence in this sample of any \"halo\" abundance object\"}. If we would apply a constant shift of $-0.35$ on the [O/H]-abundance, to correct approximately the [O/Fe] overabundance in the LMC in comparison to the Milky Way the metallicity of the PNs with an age of $10^{8.8}-10^{9.8}$ yr would nicely be in coincidence with our measurements.  However, the $[O/Fe]$ variation in dependence of $[Fe/O]$ is still under discussion (see e.g. Russell \\& Dopita \\cite{russell92} or Pagel \\& Tautvai\\u{s}vien\\.e \\cite{pagel98}).  Judging from the field around NGC~2257 we find that no radial metallicity gradient can be seen, since the field stars are consistent with the AMR derived from the inner fields. Taking also NGC~1651 into account we find that in fields where no recent star formation happened the stellar population is dominated by a population with an age between $2$ and $4$ Gyr. Thus deriving a global SFH on a limited sample is quite uncertain.   \\subsubsection{The Star Formation History}  The manner in which we derived the field star SFH contains several points that may induce biases. One reason is that the isochrones have only a crude spacing in the parameters age and metallicity. Therefore, simulations are helpful for a discussion of the SFH of the field population as described above (Sect.~12).  To derive a SFH from our data is more difficult than to derive an AMR, since one has to know not only the age of a star with a given metallicity, but the amount of stars with a given age has to be quite precise. As a result the AMR is quite robust against for example reddening variations compared to the AND.  Two SFHs are shown that illustrate how to interpret the AND (Fig.~\\ref{fig:bestSFH}). One SFH has a constant SFR during the whole LMC evolution and thus serves to give an impression how the selection effects behave (left two panels in Fig.~\\ref{fig:bestSFH}).The SFR of the second SFH was constant until $10^{9.7}$ yr ago, then it increased by a factor of $5$ until $10^{8.4}$ yr ago, before the SFR dropped to its old low level. The AND \\& AMR resulting from this SFH is plotted in the right panel of Fig.~\\ref{fig:bestSFH}.  The constant SFR is marginally inconsistent with our data, which holds for a different reddening correction of $E_{b-y} = \\pm 0.02$. This is not true for exact behaviour of the SFH: for example a decrease in the reddening of $E_{B-V}=0.02$ results in a SFH in which a much larger increase (around a factor of $10$) is necessary to describe the observations.  However, the general trend, namely, the increase of the SFR around $10^{9.5\\pm0.2}$ yr ($2-5$~Gyr) ago and the necessary declining SFR some $10^{8.5}$~yr ago in these fields is more robust. Since stars in the LMC should be mixed (at least azimuthally) after $\\approx 1$ Gyr (Gallager et al. \\cite{gallagher96}) the SFH of the older stars should be a measure for the average SFH of the LMC in the radial distance of the investigated clusters.  As a rule of thumb an increase in the applied reddening correction of $E_{b-y}=0.1$ results for a single age population in a decrease in age by a factor of $0.7$.   \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{fig33.ps}} \\caption{ Simulation with a quasi continuous SFR. The left panel shows the AMR and AND for a constant SFR over the last $10$ Gyr.  In the right graph the symbols and solid line correspond to a SFH which was constant until $10^{9.7}$ yr, at $10^{9.6}$ yr it increased by a factor of $4$ until $10^{8.6}$ yr ago after that the SFR dropped to the same level as in the beginning.  The dashed line shows the result of our composite field.} \\label{fig:bestSFH} \\end{figure}   Vallenari et al. (1996a,b) proposed, on the basis of ground based observations, a SFH in which the SFR increased about a factor of ten 2 Gyr ago, thus only around 5 \\% of the stars should be older than 4 Gyr. This has recently also been found by Elson et al. (\\cite{elson97}) with HST observations. A different SFH was advanced by Holtzman et al. (\\cite{holtzman97}), Geha et al. (\\cite{geha98}) and Holtzman et al.  (\\cite{holtzman99}) also based on HST observations. In their model, approximately half of the stars are older than 4 Gyr.  In our data the fraction of stars older than 4 Gyr is $40 \\pm 20$\\%, but we note that already a small additional reddening of of $E_{b-y}=0.015$ leaves only $\\approx 15$\\% of the stars older than 4 Gyr. Olsen (\\cite{olsen99}) used Washington photometry of the LMC field population and derived a SFH which is compatible with the one proposed by Holtzman et al. (\\cite{holtzman99}).  Summarizing, despite the uncertainty in the amount of the increase, the SFH is consistent with an increased SFR that started roughly $3\\pm 1$ Gyr ago.  Interestingly the sparsely populated outer fields are tentatively populated by mainly a population with ages between $2$ and $4$ Gyr. If this result will hold for a larger sample of outlying fields this could mean that the stellar body of the inner part of the LMC contains more younger {\\it and} older stars compared to these intermediate age stars than the more remote parts of this galaxy. However, spectroscopic studies of several of these candidate stars are needed, especially because the tilted red clump of NGC~1651 could be due to a He overabundance and thus isochrones might be misleading.  The observed SFH is inconsistent with a starformation history in which no star has been born between $4$ and $8$~Gyr. The simulations showed that virtually no star should have been recovered with an age of more than $10^{9.3}$~yr, even if the differential reddening is as large as $E_{B-V}=0.07$ and the binary fraction is $70\\%$. However, a large amount of old ($>10$~Gyr) CN anomal stars could mimic starformation between $4$ and $8$ Gyr.   A last remark on the cluster formation rate: it has been noted several times that there apparently was a long period in the LMC where no clusters (or a few) have been formed.  Recently, Larsen \\& Richtler (1999) performed a search for bright star clusters in 21 face-on galaxies. They found a correlation of the specific cluster frequency with parameters indicating the SFR. The age gap of the LMC cluster thus could reflect the low SFR during this period, where the condition for cluster formation where not present."
    },
    "0004/astro-ph0004186_arXiv.txt": {
        "abstract": "We present $BVRI$ light curves of the afterglow of GRB000301C, one of the brightest ever detected at a day time scale interval after GRB trigger. The monitoring started 1.5 days after the GRB and ended one month later. Inspection of the extremely well sampled $R$ band light curve and comparison with $BVI$ data has revealed complex behavior, with a long term flux decrease and various short time scale features superimposed. These features are uncommon among other observed afterglows, and might trace either intrinsic variability within the relativistic shock (re-acceleration and re-energization) or inhomogeneities in the medium in which the shock propagates. ",
        "introduction": "Fundamental progress on the knowledge of Gamma-Ray Bursts (GRBs) has been made possible by detection of their optical counterparts.  Of nearly 40 GRBs accurately and rapidly localized so far by BSAX, BATSE/RXTE, IPN, and promptly followed up in the optical, only about 50\\% exhibited optical afterglows\\footnote{http://www.aip.de/$\\sim$jcg/grbgen.html}, suggesting that these sources are rapidly fading, or heavily obscured. The best monitored afterglows (GRBs 970228, 970508, 980326, 980519, 990123, 990510) exhibit a variety of behaviors, indicating that the shape of the optical decay must be determined not only by the intrinsic physics, but also by the nature, structure and composition of the surrounding medium. Therefore, optical light curves of GRB counterparts need to be frequently sampled for long time intervals, to follow the evolution of the afterglow and to allow mapping the characteristics of the medium.  GRB000301C was detected by the IPN and by the RXTE ASM on 2000 March 1.4 UT with an error box of 50 arcmin$^2$ (Smith et al. 2000). Its field was acquired starting $\\sim$1.5 days later by various optical, infrared and radio telescopes.  The optical afterglow was independently detected by Fynbo et al. (2000) and by us (Bernabei et al. 2000a), and is among the brightest ever observed. Near-infrared detection and monitoring of the afterglow are reported in Rhoads \\& Fruchter (2000). Observations of the counterpart at radio and millimetric wavelengths have been reported by Berger \\& Frail (2000) and Bertoldi (2000), respectively.  Ultraviolet spectroscopy with the STIS instrument onboard HST allowed the determination of the redshift (Smette et al. 2000), then refined by optical ground-based spectroscopy ($z$ = 2.03, Castro et al. 2000). The good sampling and the brightness of the GRB000301C afterglow have allowed a detailed study of its evolution up to 15 days after the explosion. In this paper we present the results of the optical monitoring conducted at Loiano, Calar Alto, Sierra Nevada, Nainital and Canary Islands.  \\begin{table*} \\caption[]{Journal of the optical observations of the GRB000301C afterglow} \\begin{center} \\begin{tabular}{rccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Exposure start & Telescope & Filter & Exp. time & Seeing & Magnitude$^1$ \\\\ \\multicolumn{1}{c}{(UT)} & & & (minutes) & (arcsecs) &  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  2000 Mar 2.906 & UPSO     & R & 70 & 1.4 & 20.42 $\\pm$ 0.04$^2$ \\\\ 3.144 & CAHA   & R & 5    & 1.1 & 20.25 $\\pm$ 0.05 \\\\ 3.179 & CAHA   & B & 15   & 1.1 & 21.07 $\\pm$ 0.05 \\\\ 3.185 & Loiano & R & 16.7 & 2   & 20.16 $\\pm$ 0.05 \\\\ 3.205 & CAHA   & R & 5    & 1.1 & 20.25 $\\pm$ 0.05 \\\\ 3.210 & CAHA   & I & 10   & 1.1 & 19.94 $\\pm$ 0.07 \\\\ 3.219 & CAHA   & V & 15   & 1.1 & 20.57 $\\pm$ 0.05 \\\\ 3.232 & CAHA   & B & 15   & 1.1 & 21.10 $\\pm$ 0.12 \\\\ 3.913 & UPSO   & R & 50   & 1.2 & 20.51 $\\pm$ 0.04 \\\\ 4.038 & CAHA   & R & 15   & 1.6 & 20.53 $\\pm$ 0.06 \\\\ 4.149 & Loiano & R & 36.7 & 3   & $>$ 20.25$^3$ \\\\ 4.165 & Loiano & B & 20   & 3   & $>$ 21.0  \\\\ 5.135 & SNO    & R & 20   & 2   & 20.47 $\\pm$ 0.07 \\\\ 5.152 & SNO    & B & 20   & 2   & 21.60 $\\pm$ 0.20 \\\\ 5.172 & SNO    & V & 20   & 2   & 21.04 $\\pm$ 0.20 \\\\ 5.930 & UPSO   & R & 85   & 1.3 & 21.14 $\\pm$ 0.06 \\\\ 6.135 & Loiano & R & 30   & 1.7 & 21.65 $\\pm$ 0.20 \\\\ 6.163 & Loiano & B & 30   & 1.7 & 22.45 $\\pm$ 0.15 \\\\ 6.185 & Loiano & I & 16.7 & 1.7 & 20.82 $\\pm$ 0.15 \\\\ 6.968 & UPSO   & R & 35   & 1.6 & $>$ 21.6 \\\\ 7.125 & Loiano & R & 30   & 1.7 & 21.68 $\\pm$ 0.15 \\\\ 7.149 & Loiano & B & 35   & 1.7 & 22.43 $\\pm$ 0.10 \\\\ 7.177 & Loiano & I & 20   & 1.7 & 21.20 $\\pm$ 0.15 \\\\ 7.894 & UPSO   & R & 105  & 1.6 & 22.00 $\\pm$ 0.15 \\\\ 8.146 & Loiano & R & 30   & 1.6 & 21.68 $\\pm$ 0.10 \\\\ 8.170 & Loiano & I & 30   & 1.6 & 21.61 $\\pm$ 0.10 \\\\ 8.924 & UPSO   & R & 75   & 1.3 & 22.04 $\\pm$ 0.20 \\\\ Apr 5.213 & TNG & B & 20  & 0.5 & $>$ 25.5 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{6}{l}{$^1$Magnitudes of the GRB counterpart, not corrected for interstellar absorption}\\\\ \\multicolumn{6}{l}{$^2$Uncertainties of the magnitudes are at 1$\\sigma$ confidence level; lower limits at 3$\\sigma$}\\\\ \\multicolumn{6}{l}{$^3$Note that this measurement is reported as a detection in Bernabei et al. (2000b)}\\\\  \\end{tabular} \\end{center} \\end{table*}   \\begin{figure*} \\vspace{-1.5cm} \\begin{center} \\epsfig{figure=cd141f1.ps,width=14cm} \\end{center} \\vspace{-1.8cm} \\caption[]{$BVRI$ light curves of GRB000301C afterglow, based on the data presented in this paper and in the literature (see text). Filled symbols represent data presented in this work, while open symbols refer to measurements published by other authors. We have consistently referred all magnitudes to the calibration zero point of Henden (2000). To the statistical uncertainties a 5\\% systematic error has been added in quadrature (see text). No Galactic extinction correction, nor host galaxy flux subtraction has been applied. The GRB start time, indicated with $t_0$, corresponds to 2000 March 1.410845 UT} \\end{figure*}  \\begin{figure} \\vspace{-1.2cm} \\begin{center} \\epsfig{figure=cd141f2.ps,width=11.5cm} \\end{center} \\vspace{-1.8cm} \\caption[]{Colors of GRB000301C afterglow (data are from this paper and from the literature, see text).  These are reported as filled circles when computed between pairs of measurements spaced apart in time by no more than 0.5 hr, and as stars when the temporal separation is larger than 0.5 hr, but smaller than 9 hr. As in Fig. 1, calibration by Henden (2000) has been adopted and a 5\\% systematic error has been added in quadrature (see text).  The GRB start time, indicated with $t_0$, corresponds to 2000 March 1.410845 UT}  \\end{figure} ",
        "conclusions": "Our optical monitoring of the bright GRB000301C afterglow has provided one of the best sampled afterglow datasets, especially in the $R$ filter. The long term behavior of this optical afterglow is better described by a continuous steepening, rather than by a single power-law, as expected in afterglows developing in laterally spreading jets (Sari et al. 1999; Rhoads 1999) or decelerating to non-relativistic regimes (Dai \\& Lu 1999), and seen in few other cases.  Among equally well monitored GRB afterglows, GR000301C appears peculiar in that several shorter time scale variations are superimposed on the long term decrease.  The reality of two of these (3.1-3.7 days and 5-7 days after the event) is supported by their appearance in more than one band. The first two points of the $R$ and $I$ band light curves might suggest a rise and could be reminiscent of the early (1-2 days after the GRB trigger) light curve of GRB970228 and GRB970508 (Guarnieri et al. 1997; Pedersen et al. 1998), although in the latter the initial increase was more structured. In the present case we cannot exclude that the flux is declining since the start of the monitoring, and hour time scale flares modulate this decrease.  Some isolated short term variability events are seen in GRB980703 (Vreeswijk et al. 1999) and GRB990123 (Castro-Tirado et al. 1999) and are almost totally absent in GRB990510 (e.g., Stanek et al. 1999).  Recently, various scenarios have been developed in which intrinsic re-energization of the blast wave, or irregularities of the dense interstellar medium in which the blast is expanding can account for the observed behavior (Panaitescu et al. 1998;  M\\'esz\\'aros et al. 1998; Sari \\& M\\'esz\\'aros 2000; Wang \\& Loeb 2000; Dai \\& Lu 2000).  In particular, a flattening of the afterglow light curve, similar to that exhibited by GRB000301C in the $R$ band on days 3.1-3.7 and 5-7 days after the GRB, is predicted by Kumar \\& Piran (2000) as a consequence of the collision of a slow shell ejected at a late time after the GRB with an outer shell decelerated by its propagation in the circumburst medium (see their Fig. 5). We note that the temporal occurrence of the observed flattenings could be consistent with a ``colliding shells\" interpretation, while it is incompatible with the time scale implied by an hypernova scenario (see Rhoads \\& Fruchter 2000).  The lack of a clear correlation between the $R$ band light curve and the $I$ and $K$ band light curves (see Rhoads \\& Fruchter 2000 for the latter) might be due to non strict simultaneity of the data points.  In fact, the $R-I$ and $R-K$ colors as a function of time show only marginally significant deviations from constancy (Fig. 2c and 2d), and these are mainly exhibited by color values derived from pairs of measurements separated in time by more than 0.5 hours, the shortest variability time scale observed in this afterglow. Our findings underline the critical importance of intensive multiwavelength observations of afterglow sources."
    },
    "0004/astro-ph0004135_arXiv.txt": {
        "abstract": "We report the discovery of an \\ion{O}{6} absorption system at \\zabs\\ = 0.14232 in a high resolution FUV spectrum of PG 0953+415 obtained with the Space Telescope Imaging Spectrograph (STIS). Both lines of the \\ion{O}{6} $\\lambda \\lambda$ 1032, 1038 doublet and multicomponent \\ion{H}{1} \\lya absorption are detected, but the \\ion{N}{5} doublet and the strong lines of \\ion{C}{2} and \\ion{Si}{3} are not apparent. We examine the ionization mechanism of the \\ion{O}{6} absorber and find that while theoretical considerations favor collisional ionization, it is difficult to observationally rule out photoionization. If the absorber is collisionally ionized, it may not be in equilibrium due to the rapid cooling of gas in the appropriate temperature range. Non-equilibrium collisionally ionized models are shown to be consistent with the observations. A WIYN survey of galaxy redshifts near the sight line has revealed a galaxy at a projected distance of 395 kpc separated by $\\sim$130 \\kms\\ from this absorber, and three additional galaxies are found within $\\lesssim$ 130 \\kms\\ of this redshift with projected separations ranging from 1.0 Mpc to 3.0 Mpc. All of these galaxies are luminous (0.6 -- 4.0 $L*$), and two of them show the [\\ion{O}{2}] $\\lambda$3727 emission line indicative of active star formation. The galaxies with [\\ion{O}{2}] emission are probably normal spirals. Combining the STIS observations of PG0953+415 with previous high signal-to-noise observations of H1821+643 with the Goddard High Resolution Spectrograph (GHRS), we find two \\ion{O}{6} systems with $W_{\\rm r} >$ 60 m\\AA\\ and $z <$ 0.3 over a total redshift path $\\Delta z$ of only 0.10. Both of these QSOs were originally observed to study the low $z$ \\lya lines and should not be biased in favor of \\ion{O}{6} detection.  If these sight lines are representative, they imply a large number of \\ion{O}{6} absorbers per unit redshift, $dN/dz \\sim$ 20. The corresponding value of $dN/dz$ for low $z$ \\lya lines with $W_{\\rm r} >$ 50 m\\AA\\ is $102 \\pm 16$.  We use this sample to obtain a first estimate of the cosmological mass density of the \\ion{O}{6} systems at $z \\approx$ 0.  If further observations confirm the large $dN/dz$ derived for the \\ion{O}{6} systems, then these absorbers trace a significant reservoir of baryonic matter at low redshift. ",
        "introduction": "Hydrodynamic simulations of cosmological structure growth predict that when the initial density perturbations collapse, gas should be shock-heated to temperatures of $10^{5}$ -- $10^{7} \\ ^{\\circ}$K (Ostriker \\& Cen\\markcite{oc96} 1996; Dav\\'{e} et al.\\markcite{dave99} 1999; Cen \\& Ostriker\\markcite{co99a} 1999a). The fraction of the gas which has been heated to these temperatures increases with decreasing redshift, and at the present epoch, the model of Cen \\& Ostriker\\markcite{co99a} (1999a) predicts that 47 \\% of the baryons (by mass) are in this shock-heated phase, hereafter referred to as warm/hot gas (to distinguish it from the hotter gas in rich galaxy clusters which are readily detected X-ray sources). This warm/hot gas prediction has not been adequately tested by observations because the soft X-rays emitted by gas at these temperatures are difficult to detect with current instrumentation, especially at lower temperatures where corrections for foreground absorption and emission are complicated. However, it may be possible to detect gas in the lower half of this temperature range via absorption lines of species such as \\ion{O}{6}, \\ion{Ne}{8}, or \\ion{Mg}{10} in the spectrum of a background QSO (Verner, Tytler, \\& Barthel\\markcite{vern94} 1994). It is important to search for this warm/hot gas as part of the census of matter in the universe and because it could affect the formation and evolution of galaxies and galaxy groups and clusters (e.g., Blanton et al.\\markcite{blan99} 2000).  There are some indications that warm/hot gas is present in some galaxy groups. For example, Mulchaey et al.\\markcite{mul96} (1996) have noted that ROSAT observations show that poor galaxy groups which are rich in elliptical galaxies tend to exhibit X-ray emission (E $>$ 0.5 keV) while spiral-rich groups do not. They suggest that spiral-rich groups contain cooler ($\\lesssim 4\\times 10^{6}$ K) intragroup gas which is not easily detected in X-rays, and they predict that the intragroup medium of spiral-rich groups will produce absorption lines of \\ion{O}{6}, but not \\ion{C}{4} or \\ion{N}{5} because their column densities are too low at $T > 5 \\times 10^{5} \\ ^{\\circ}$K. Interestingly, Savage, Tripp, \\& Lu\\markcite{stl98} (1998) have recently identified a QSO absorber in the spectrum of H1821+643 ($z_{\\rm QSO}$ = 0.297) which fits this description: \\ion{O}{6} absorption lines with two nearby spiral galaxies but no accompanying \\ion{C}{4} or \\ion{N}{5} lines. However, the absorption could be due to the halo of the closer spiral galaxy (which is at a projected distance of 105 $h_{75}^{-1}$ kpc) rather than the intragroup medium, and Savage et al.\\markcite{stl98} show that the \\ion{O}{6} absorption could plausibly arise in very low density photoionized gas. The \\ion{O}{6} doublet has been identified in several other intervening absorption systems\\footnote{Strong O VI absorption lines have also been detected in ``associated'' absorption line systems with \\zabs\\ $\\approx \\ z_{\\rm QSO}$ (e.g., Papovich et al.\\markcite{pap} 2000 and references therein). These are a rather different class of absorber which are often known to be very close to the QSO (Hamann \\& Ferland\\markcite{hf99} 1999). In this paper, we have focused our analysis and discussion on the intervening systems.} both at moderate redshifts (e.g., Bergeron et al.\\markcite{berg94} 1994; Burles \\& Tytler\\markcite{bt96} 1996; Jannuzi et al.\\markcite{jan98} 1998; Lopez et al.\\markcite{lop99} 1999; Churchill \\& Charlton\\markcite{cc99} 1999) and at high redshifts (Kirkman \\& Tytler\\markcite{kt97} \\markcite{kt99} 1997,1999). Composite spectra and statistical techniques have also been used to show that \\ion{O}{6} absorption is present at high $z$ (Lu \\& Savage\\markcite{ls93} 1993; Dav\\'{e} et al.\\markcite{dave98} 1998). In most cases it has been difficult to pin down the ionization mechanism definitively, partly due to the low resolution of the observations made with first-generation {\\it Hubble Space Telescope (HST)} spectrographs, but in several cases there is evidence that the \\ion{O}{6} systems occur in multiphase absorbing media. An important step in this approach to the search for warm/hot gas is to determine whether the \\ion{O}{6} absorbers trace collisionally ionized gas or photoionized gas.  As part of the program described by Tripp, Lu, \\& Savage\\markcite{tls98} (1998) to study low $z$ \\lya absorption line systems, we have recently observed the low redshift QSO PG0953+415 ($z_{\\rm QSO}$ = 0.239) with the E140M echelle mode of the Space Telescope Imaging Spectrograph (STIS). This high resolution FUV spectrum has revealed another highly ionized \\ion{O}{6} absorber associated with a group of spiral galaxies, and in this paper we present our analysis of this particular absorbing system. In \\S 2 we review the observations and data reductions including measurements of galaxy redshifts with the WIYN telescope. We present in \\S 3 the absorption line measurements. We constrain the temperature of the gas and examine its ionization in \\S 4, and we discuss the implications of the observations in \\S 5. Throughout this paper we assume $H_{0} = 75 h_{75}$ \\kms\\ Mpc$^{-1}$ and $q_{0}$ = 0.0. Also, all wavelengths and redshifts reported here are heliocentric, but in this direction heliocentric and LSR velocities are nearly identical.\\footnote{Assuming the Sun is moving in the direction $l = 53^{\\circ}, b = 25^{\\circ}$ at 16.5 \\kms\\ (Mihalas \\& Binney\\markcite{mb81} 1981), $v_{\\rm LSR} = v_{\\rm helio} - 0.1$ \\kms .} ",
        "conclusions": "\\subsection{Association with a Galaxy Group}  The fact that the \\zabs\\ = 0.14232 \\ion{O}{6} absorption line system is associated with a group of galaxies strongly indicates that the absorber is an {\\it intervening} system rather than an {\\it intrinsic} absorber which was ejected or somehow accelerated to high displacement velocity by the QSO. While this may seem a trivial conclusion since the \\ion{O}{6} is displaced from the QSO redshift by $\\sim$23,000 \\kms , during the last few years observations of absorption variability have established that some highly ionized intrinsic QSO absorbers are separated from $z_{\\rm QSO}$ by such large velocities (e.g. Hamann, Barlow, \\& Junkkarinen\\markcite{hbj97} 1997) and yet are relatively narrow (i.e., not traditional broad absorption line outflows). Therefore, it is important to find evidence that a given \\ion{O}{6} absorber is indeed intervening even if $\\Delta v$ is large. Of course, this association with a galaxy group does not necessarily indicate that the absorption arises in the intragroup medium; it could be due to gas within one of the galaxies in the group. We briefly discuss some possibilities below.  The close proximity of the \\ion{O}{6} system to a galaxy group provides a theoretical prejudice in favor of collisional ionization. The presence of galaxies requires that intergalactic gas has collapsed in this region of space, and simple arguments (e.g., \\S 3 in Cen \\& Ostriker \\markcite{co99a} 1999a) suggest that substantial shock-heating probably occurred as a result.  Therefore collisionally ionized gas is expected in the vicinity (if not along the pencil-beam probed by the QSO). It is interesting that this absorption system and associated galaxies fit the prediction of Mulchaey et al.\\markcite{mul96} (1996) discussed in \\S 1: a group which is possibly spiral-rich and has associated \\ion{O}{6} absorption, as expected based on their postulated collisionally ionized intragroup medium. However, we really are not sure that this is a spiral-rich group (or even a bound group) with only four known galaxies, and it is also possible that the \\ion{O}{6} absorption arises in the gaseous halo of a single galaxy. However, models of galaxy gaseous halos (e.g., Mo \\& Miralda-Escud\\'{e} \\markcite{mo} 1996) also usually produce \\ion{O}{6} absorption in a collisionally ionized hot phase. Similarly, models which produce QSO absorption line gas in supernova-driven winds from dwarf galaxies (e.g. Wang\\markcite{wang} 1995) also involve substantial shock-heating and collisionally ionized gas.  The kinematics of the absorption also provide useful information about the possible origins of the system at \\zabs\\ = 0.14232. A fundamental question is whether the absorption is due to the intragroup medium or to the ISM of a single galaxy that happens to intercept the line-of-sight. The \\ion{H}{1} \\lya\\ profile has the ``leading edge asymmetry'' that various authors (e.g., Lanzetta \\& Bowen\\markcite{lanz92} 1992) have discussed as the signature of a moderately edge-on rotating disk. Of course, this is not a unique interpretation of such profile asymmetry, but it provides some evidence in favor of the single-galaxy interpretation.  The simplicity of the \\ion{O}{6} profile also favors this interpretation.  The velocity dispersion of poor galaxy groups is typically one to a few hundred \\kms\\ (Zabludoff \\& Mulchaey\\markcite{zab} 1998), so one might expect the \\ion{O}{6} to be spread over a larger velocity range if the absorption is due to the intragroup medium.  Here, though, we must recognize that the observed \\ion{O}{6} profiles are noisy and we only detect \\ion{O}{6} at the velocity of the {\\it strongest} component of the \\ion{H}{1} profile.  If the hot gas is concentrated at the center of the group, then this may be the only location along the line-of-sight where we have sufficient sensitivity to detect it, and there may be \\ion{O}{6} absorption at other velocities which has fallen below our detection threshold.  To sort out the various possibilities, it would be very helpful to obtain additional STIS observations to improve the signal-to-noise and search for \\ion{O}{6} at other velocities.  \\subsection{Number Density and Cosmological Mass Density}  One means to test the warm/hot gas prediction of cosmological simulations (\\S 1) is to compare the number of \\ion{O}{6} absorbers observed per unit redshift, $dN/dz$, to the number statistically predicted from many random pencil-beams through the cosmological simulations.  Is the observed $dN/dz$ consistent with the number predicted by the cosmological models? We can also estimate the mean cosmological mass density traced by the \\ion{O}{6} systems at low $z$.  For these purposes, we combine the STIS observations of PG0953+415 with the GHRS observations of H1821+643 from Tripp et al.\\markcite{tls98} (1998).\\footnote{After this paper was completed, Tripp, Savage, \\& Jenkins\\markcite{tsj2000} (2000) carried out an analysis of new STIS echelle observations of H1821+643 with the E140M mode, and we refer the reader to that paper for the O VI number density and cosmological mass density derived from an independent data set.}  We emphasize that the manner in which these QSOs were selected should not bias the sample to enhance the number of \\ion{O}{6} systems detected compared to sampling many random directions. Both QSOs were originally observed to study the relationship between weak \\lya clouds and galaxies\\footnote{$HST$ program IDs 6155 and 7747, see Tripp et al.\\markcite{tls98} (1998) for details.} at $z <$ 0.3. These two objects were selected for this project simply because they are among the brightest known QSOs with $z >$ 0.2, criteria which were required to substantially increase the sample of {\\it weak} \\lya absorbers with a minimal amount of {\\it HST} time.  No consideration was given to factors which might indicate an enhanced likelihood of detecting the \\ion{O}{6} doublet when the targets were selected.  To estimate $dN/dz$ of the \\ion{O}{6} systems, we define a sample of \\ion{O}{6} lines with $W_{\\rm r} >$ 60 m\\AA\\ for both lines of the doublet,\\footnote{The 3$\\sigma$ detection limit throughout the region where O VI absorbers can be detected is 60 m\\AA\\ or better in the spectrum of PG0953+415 and 50 m\\AA\\ or better in the spectrum of H1821+643.} and set the maximum absorber redshift, $z_{\\rm max}$, for each sight line to exclude any absorbers within $\\mid \\Delta v \\mid \\ \\leq$ 5000 \\kms\\ of the QSO redshift and thereby avoid contaminating the sample with associated/intrinsic absorbers which are close to the QSO. One might argue that 5000 \\kms\\ is insufficient (see above). However, we find that the two \\ion{O}{6} systems in the final sample are associated with galaxy groups, which suggests that these are indeed intervening absorbers. The lower redshift cutoff for each sight line, $z_{\\rm min}$, was determined by the lowest wavelength in the observed spectrum (with a small buffer to ensure that a line would be recognized if at that $\\lambda$) in the case of H1821+643 and by the wavelength at which the S/N is unacceptably low in the case of PG0953+415. These criteria resulted in a sample of two \\ion{O}{6} absorbers within a total redshift path of $\\Delta z$ = 0.100 (after a significant correction of 0.067 for regions of the spectra in which we cannot detect either of the \\ion{O}{6} lines\\footnote{Both of the O VI lines were required to fall in unblocked regions of the spectra so that the doublet could be securely identified.} because they would be blocked by strong ISM or extragalactic lines from other redshift systems). Therefore, $dN/dz \\sim$ 20 for \\ion{O}{6} systems with $W_{\\rm r} \\geq$ 60 m\\AA\\ at $z <$ 0.3. Using the confidence limits from Gehrels\\markcite{geh} (1986) for a sample of two absorbers, we derive 4 $< dN/dz <$ 63 for these \\ion{O}{6} absorbers at the 90\\% confidence level.  For comparison, Tripp et al.\\markcite{tls98} (1998) find $dN/dz = 102 \\pm 16$ for \\ion{H}{1} \\lya lines with $W_{\\rm r} \\geq$ 50 m\\AA\\ at $z <$ 0.3, and Weymann et al.\\markcite{wey} (1998) report $dN/dz = 33 \\pm 4$ for \\lya absorbers at $z$ = 0 with $W_{\\rm r} \\geq$ 240 m\\AA\\ based on a large sample at $z <$ 1.5. In the case of low to moderate redshift \\ion{Mg}{2} absorbers, $dN/dz = 0.97 \\pm 0.10$ for $W_{\\rm r} \\geq$ 300 m\\AA\\ (Steidel \\& Sargent \\markcite{ss92} 1992) and $dN/dz = 2.65 \\pm 0.15$ for $W_{\\rm r} \\geq$ 20 m\\AA\\ (Churchill et al.\\markcite{crcv99} 1999; see also \\S 6 in Tripp, Lu, \\& Savage\\markcite{tls97} 1997). Evidently, these weak low $z$ \\ion{O}{6} systems have a substantially larger cross section and/or covering factor than the \\ion{Mg}{2} absorbers. Similarly, the {\\it stronger} \\ion{O}{6} absorbers at higher redshifts are less common: Burles \\& Tytler\\markcite{bt96} (1996) report $dN/dz = 1.0\\pm 0.6$ for \\ion{O}{6} systems with $W_{\\rm r} \\geq$ 210 m\\AA\\ at $<z_{\\rm abs}>$ = 0.9, and similar results (with smaller uncertainties) are derived from the larger sample provided by the FOS Quasar Absorption Line Key Project (B. Jannuzi \\& R. Weymann 2000, private communication). This comparison with the \\ion{O}{6} $dN/dz$ derived from FOS data must be interpreted carefully, however, because the FOS samples are dominated by lines which are substantially stronger and at substantially higher redshifts than the \\ion{O}{6} absorbers discussed in this paper.  With a larger sample of low $z$ weak \\ion{O}{6} absorbers and smaller uncertainties in their number density, comparison of $dN/dz$ to the space density of objects such as dwarf galaxies may provide insight on the nature of the \\ion{O}{6} systems.  Next we estimate the baryonic content of the \\ion{O}{6} absorbers, expressed as the cosmological mass density $\\Omega _{b}$(\\ion{O}{6}), following previous analogous calculations for damped \\lya\\ systems (e.g., Lanzetta et al.\\markcite{lanz91} 1991) as well as \\ion{O}{6} absorbers (Burles \\& Tytler\\markcite{bt96} 1996).  To estimate the density of baryons in the gaseous component of the universe traced by \\ion{O}{6}, we require information about the metallicity of the gas and the \\ion{O}{6} ionization fraction, $f$(\\ion{O}{6}) = \\ion{O}{6}/O$_{\\rm total}$. In collisional ionization equilibrium, $f$(\\ion{O}{6}) peaks at $\\sim$0.2 (Sutherland \\& Dopita\\markcite{sd93} 1993), and similarly low peak fractions are predicted by non-equilibrium collisional models (e.g., Shapiro \\& Moore\\markcite{shape76} 1976; Benjamin \\& Shapiro\\markcite{benj} 2000). The \\ion{O}{6} ion fraction is not much larger at peak value in photoionized gas (see appendix). Therefore we will adopt $f$(\\ion{O}{6}) $\\sim$ 0.2 to set a lower limit on the baryonic content of the \\ion{O}{6} absorbers, and the following calculation is relatively independent of how the gas is ionized or whether or not it is due to the intragroup medium. For the mean cosmic metallicity of the \\ion{O}{6} absorbers, it is less clear what value to adopt, but to set a lower limit on $\\Omega _{b}$(\\ion{O}{6}) we should set (O/H) to a high but plausible value.  As noted in \\S 4.2.1, Davis et al.\\markcite{dmm99} (1999) and Hwang et al.\\markcite{hwa} (1999) have derived metallicities of 1/3--1/2 solar for intragroup gas in several X-ray bright groups.  Therefore we will initially use 1/2 solar metallicity for the calculation of $\\Omega _{b}$(\\ion{O}{6}) and then discuss how it scales with (O/H).  The mean cosmological mass density in the \\ion{O}{6} absorbers, in units of the current critical density $\\rho _{c}$, can be estimated as \\begin{equation} \\Omega _{b}({\\rm O \\ VI}) = \\frac{\\mu m_{\\rm H} H_{0}}{\\rho _{c} c f({\\rm O \\ VI})} \\left( \\frac{\\rm O}{\\rm H} \\right)^{-1}_{\\rm O \\ VI} \\frac{\\sum_{i} N_{i}({\\rm O \\ VI})}{\\sum_{i} \\Delta X_{i}} \\end{equation} where $\\mu$ is the mean atomic weight (taken to be 1.3), (O/H)$_{\\rm O \\ VI}$ is the assumed mean oxygen abundance by number in the \\ion{O}{6} absorption systems, $m_{\\rm H}$ is the mass of hydrogen, $N_{i}$(\\ion{O}{6}) is the total \\ion{O}{6} column density and $\\Delta X_{i}$ is the absorption distance interval (Bahcall \\& Peebles\\markcite{bah69} 1969) probed to the $i$th QSO, \\begin{equation} \\Delta X_{i} = \\case{1}{2} \\{ [(1 + z_{\\rm max})^{2} - 1] - [(1 + z_{\\rm min})^{2} - 1]\\} \\end{equation} assuming $q_{0}$ = 0.\\footnote{Over the redshift range probed by the sight lines to PG0953+415 and H1821+643, results are insensitive to the value assumed for $q_{0}$.} As in the calculation of $dN/dz$, we correct $\\Delta X_{i}$ for spectral regions blocked by strong lines. Combining the PG0953+415 STIS data and the GHRS observations of H1821+643 from Tripp et al.\\markcite{tls98} (1998) with $z_{\\rm min}$ and $z_{\\rm max}$ set to the same values used for the derivation of $dN/dz$, we obtain $\\Omega _{b}({\\rm O \\ VI}) \\gtrsim 0.0006 h_{75} ^{-1}$ assuming the mean O abundance is 1/2 solar.  This is a lower limit not only because $f$(\\ion{O}{6}) and (O/H) were set to their approximate upper limits, but also because we have applied an equivalent width cutoff to define the sample; if \\ion{O}{6} absorbers with $W_{\\rm r} \\leq$ 60 m\\AA\\ significantly increase $\\sum_{i} N_{i}$(\\ion{O}{6}), then the true $\\Omega _{b}$(\\ion{O}{6}) will be higher.  Note that $\\Omega _{b}({\\rm O \\ VI})$ is inversely proportional to (O/H).  Decreasing the metallicity to 1/10 solar, for example, increases the lower limit on the \\ion{O}{6} absorber baryon content to $\\Omega _{b}({\\rm O \\ VI}) \\gtrsim 0.003 h_{75} ^{-1}$.  To demonstrate the level of uncertainty in the cosmological mass density estimate due to small number statistics, we can recalculate $\\Omega _{b}$(\\ion{O}{6}) using an alternative expression analogous to equation (9) from Rao \\& Turnshek\\markcite{rt2000} (2000), \\begin{equation} \\Omega _{b}({\\rm O \\ VI}) = \\frac{\\mu m_{\\rm H} H_{0}}{\\rho _{c} c f({\\rm O \\ VI})} \\left( \\frac{\\rm O}{\\rm H} \\right)^{-1}_{\\rm O \\ VI} \\left(\\frac{dN}{dz}\\right) \\frac{<N_{\\rm O \\ VI}>}{(1+z)} \\end{equation} where $<N_{\\rm O \\ VI}>$ is the mean \\ion{O}{6} column density of the absorption systems in the sample and again we have assumed $q_{0}$ = 0. The advantage of this alternative expression for $\\Omega _{b}$(\\ion{O}{6}) is that we can employ the Gehrels\\markcite{geh} (1986) small sample statistics to estimate the uncertainty in $dN/dz$ and then propagate the uncertainty into the estimate of $\\Omega _{b}$(\\ion{O}{6}). With $dN/dz$ = $20^{+26}_{-13}$ and the other parameters set to the values assumed above, we find for the 1/10 solar metallicity case $\\Omega _{b}({\\rm O \\ VI}) \\gtrsim 0.003^{+0.004}_{-0.002} \\ h_{75} ^{-1}$ (error bars are $1\\sigma$ uncertainties).  Bearing in mind that there is still considerable uncertainty in the lower limit on $\\Omega _{b}$(\\ion{O}{6}) due to the small sample, small redshift path, and uncertain mean metallicity, this preliminary estimate suggests that the \\ion{O}{6} absorbers may indeed harbor a significant fraction of the baryons in the universe at low $z$.  The lower limit assuming (O/H) = 1/10 solar is comparable to the cosmological mass density of stars, \\ion{H}{1}, and X-ray emitting galaxy group and cluster gas at low redshift (Fukugita, Hogan, \\& Peebles\\markcite{fhp98} 1998), for example."
    },
    "0004/astro-ph0004303_arXiv.txt": {
        "abstract": "The collapse of marginally Jeans unstable primordial gas clouds in the presence of UV radiation field is discussed. Assuming that the dynamical collapse proceeds approximately in an isothermal self-similar fashion, we investigate the thermal evolution of collapsing central core until H$_2$ cooling dominates photoheating and the temperature drops to below $10^4$K.  Consequently, the mass of the cooled core is evaluated as $M_{\\rm cool}=3.6\\times 10^6 M_\\odot \\left(I_{21}/1\\right)^{-0.32}$.  This scale depends only on the incident UV intensity, and provides a lower limit to the mass of collapsed objects in the UV radiation field. ",
        "introduction": "\\label{INTRO} Formation of primordial objects, such as young galaxies and globular clusters, in the early universe is a fundamental problem in modern cosmology.  Rapid progress of observations in the last decade has pressed theoreticians to construct the physically correct and proper theory to solve this problem. Recently, formation of primordial objects has been investigated mainly in the following two contexts; one is the formation of ``first stars'' or ``first luminous objects'' (e.g. \\cite{Tegmark97,Abel98,ON98,NS99}), and the other is that of ``second generation objects'' (e.g. \\cite{HRL97,KBS97,ON99}). These two populations are expected to arise in quite different physical environments. For instance, the former is likely to be born from the primordial gas under little influences of the external radiation field, except for that of the cosmic microwave background radiation (CMB).  The latter, on the other hand, is largely affected by the external UV radiation produced by the former. The UV photons not only ionize and heat up the pregalactic gas clouds, but also dissociate H$_2$ in the clouds, which is an important coolant in metal-free gas clouds. In this paper, we pay particular attention to the formation of the latter population.  At high redshifts, say $z\\gg 5$, ionized bubbles around photon sources, such as young galaxies or AGN's, are still so small that most of the gas in the universe is not ionized and photoheated. In this case, the external radiation affects the formation of primordial objects only via H$_2$ photodissociation \\cite{HRL97}. On the other hand, gas clouds near the ionizing sources, or those at lower redshifts ($z\\lsim 5$) are ionized and photoheated, before they collapse and cool. Therefore, the formation of population III objects close to the ionizing sources and the formation of galaxies at low redshifts are likely to start from hot ionized media. Kepner, Babul \\& Spergel (1997) investigated this problem in the context of dwarf galaxy formation under the assumption that pregalactic clouds are in hydrostatic equilibrium. They found that the gas transit from H$^+$ phase to H and H$_2$ ones, as radiative cooling proceeds.  Corbelli, Galli \\& Palla (1997) also found similar phenomena in the rotationally supported hydrostatic gaseous disc. However, gravitational collapse of ionized gas proceeds almost isothermally \\cite{UI84,TW96,KI99,SU00}, and such a collapse becomes inevitably dynamical \\cite{Larson69}.  In this paper, we investigate H$_2$ cooling in the dynamically collapsing core in marginally Jeans unstable clouds. If these clouds are self-gravitating, they are likely to collapse almost spherically with the temperature kept nearly at $\\sim 10^4$K, until the external radiation field is shielded by the clouds themselves. We thus employ the isothermal Larson-Penston similarity solution \\cite{Larson69} and solve explicitly non-equilibrium chemical reactions and energy equation to trace the thermal history of a collapsing central core. Consequently, the mass and the size of the cooled core are assessed, and their astrophysical implications are discussed.  \\newpage ",
        "conclusions": ""
    },
    "0004/astro-ph0004073_arXiv.txt": {
        "abstract": "This paper presents an analysis of the UV spectrum of some nebulae with clearly identified illuminating stars, all observed by the IUE satellite.  The data show remarkable properties of the UV spectrum of the nebulae. Each spectrum is the product of the star spectrum and a linear function of $1/\\lambda$. There is no peculiar behaviour in the spectrums at 2200~\\AA{}: no bump created in the spectrum of a nebula and no excess of scattering. When moving away from the star, the surface brightness of a nebula decreases as the inverse of the square of the angular distance to the star.  These results can logically be interpreted in terms of scattering of starlight. They imply constant properties of the interstellar grains in the UV and in the directions of space sampled by the nebulae, and probably a strong forward scattering phase function. There is no evidence for any particular type of grain which would specifically extinguish starlight at 2200~\\AA{}.  Concerning the UV spectrum of a star, this may imply a revisal of the traditional interpretation of the 2200~\\AA{} bump. ",
        "introduction": "Introduction} In the UV, as in the optical, the light which we receive from a nebula should in great part be the light of the illuminating star scattered by dust grains in the nebula. If no other process than scattering of starlight by the average population of interstellar grains is involved, the ratio of the nebula to the star spectrum depends on the optical depth of the nebula, on the angle $\\varphi$ of scattering, and on the albedo and phase function of interstellar grains.  The original idea of this work was to study the behavior of the spectrum of nebulae in the 2200~\\AA{} region. In this spectral region, the spectrum of the stars is known to have a characteristic depression (the 2200~\\AA{} bump) when there is interstellar matter between the star and the observer. The possible causes of the bump are reviewed by Bless~\\& Savage \\cite{bless72}. Admittedly, the bump is an extinction feature (eventually a purely absorbing one, Witt et al. \\cite{witt92}) which arises due to a particular type of interstellar grain.  Since a nebula placed in front of a star produces a bump in the spectrum of the star, what will the bump spectral region of a nebula illuminated by a star at close angular distance be? If the particles which extinguish starlight at 2200~\\AA{} are present in the nebula, its spectrum should have an evident feature at 2200~\\AA{}. This feature will be an absorption one if the bump carriers are purely absorbing particles. It will be an excess of emission if the carriers scatter starlight.  Section \\ref{data} presents the IUE observations used in this work. It consists of the spectrums of a few well known nebulae, all illuminated by close bright stars.  Section~\\ref{analyse} to Section~\\ref{analyse2}, is an analysis of the data. It is a remarkable fact, and the most significant result of this analysis, that all the nebulae have a spectrum which is exactly the product of the illuminating star spectrum and a linear function of $1/\\lambda$ (Section~\\ref{analyse}). Some nebulae do not have a bump, and when they do have one it is always proportional to the bump of the illuminating star. No bump is created in the nebulae.  In Section~\\ref{int}, the relation between the spectrum of a nebula and the spectrum of the illuminating star is interpreted as scattering of starlight by a medium with an optical depth which varies as $1/\\lambda$.  A second result, the power law dependence of the surface brightness on the angular distance to the star, is arrived at and discussed in Sections~\\ref{analyse2} and \\ref{rdis}. The large variations of the surface brightness level of a nebula with the distance to the star may be an important argument in favor of a strong forward scattering phase function in the UV (Section~\\ref{rdis}). It implies that a considerable amount of starlight is scattered in directions close to the star.  Different authors (e.g. A.N. Witt and collaborators) have found a strong forward scattering phase function of the grains in the UV. These papers are associated with models which also predict changes of the albedo and/or the phase function with the wavelength. From analysis of the present data, I cannot agree with these conclusions. Contrary to these studies, I find constant albedo and phase function of the grains in all the UV, and perhaps up to the near infrared, in all directions of space. In Section~\\ref{grains}, a few of these papers are briefly discussed along with the consequences the present paper have on the properties of interstellar grains. It is suggested that there is no excess of extinction in the spectrum of the stars at 2200~\\AA{}.  An Appendix will review elementary properties of scattering and the notations used in the paper. ",
        "conclusions": ""
    },
    "0004/astro-ph0004245_arXiv.txt": {
        "abstract": " ",
        "introduction": "Soft X-ray transients (SXTs), or X-ray novae (e.g. Chen et al. 1997) are binary systems which exhibit luminous X-ray and optical outburst (Tanaka, Shibazaki 1996).  They uniquely provide the most compelling evidence for the existence of steller mass black holes using radial velocity studies, giving mass functions exceeding the maximum mass of a stable neutron star ($\\sim 3\\,M_\\odot$), and we know eight such black hole candidates (van Paradijs, McClintock 1995; Bailyn et al. 1998; Orosz et al. 1998).  Their outburst light curves often have an common feature, that is, after a rise of a few days, the X-ray intensity comes to the maximum, which was typically followed by the exponential decay with an e-folding time of $\\sim 40\\,{\\rm d}$ (Chen et al. 1997). At the maximum, the X-ray luminosity reaches $10^{38-39}\\;{\\rm erg\\,s^{-1}}$ and the optical to X-ray flux ratio of $\\sim 500$ is typical in SXTs (Tanaka, Shibazaki 1996).  Many astronomers have recently believed that the compact object of SXTs in quiescence is surrounded by the accretion disk whose inner part is the geometrically thick and optically thin advection dominated accretion flow (ADAF) and the outer part is the geometrically thin and optically thick disk which becomes thermally unstable when the disk becomes too hot for hydrogen to remain neutral (Narayan, Yi 1995; Shakura, Sunyaev 1973; Osaki 1974).  This model of the outburst mechanism is called the disk instability model and satisfactorily explains the outburst cycle of a few tens of years and the outburst duration by the viscous diffusion time scale of the accretion disk in the cool and hot states, respectively (Mineshige 1996).  Almost all observed SXTs are distributed on the galactic disk and few had been discovered in the galactic halo (Chen et al. 1997; Bradt et al. 2000).  White, van Paradijs (1996) have studied the galactic distribution of BHCs low-mass X-ray binaries and found an rms value for the distance from the galactic disk to be $\\sim 0.4$ kpc.  We can easily understand these biased distribution because massive stars, which are responsible for producing neutron stars and black holes, have been generated in the disk rather than the halo.  A new SXT, XTE J1118+480 whose galactic latitude is high ($\\sim 62^\\circ$), has been discovered at an intensity of 39 mCrab with the All-Sky Monitor (ASM; Levine et al. 1996) on the {\\it Rossi X-Ray Timing Explorer} (RXTE) in 2000 March 29 (Remillard et al. 2000).  The X-ray spectrum just after the discovery is similar to Cyg X-1 in its hard state (Remillard et al. 2000).  The hard X-ray spectrum is well characterized by a power law with a photon index of 2.1 and the source is visible up to 120 keV which implies that it is a possible black hole X-ray transient (Wilson, McCollough 2000). In this paper, we report the discovery of the optical counterpart of XTE J1118+480 and the optical short-time variability, and discuss the binary nature and a distance estimation. ",
        "conclusions": "XTE J1118+480 shows some features different from those of typical X-ray transients, that is, the multiple peaks within three months observed both in X-ray and optical, an extremely low optical to X-ray flux ratio, and the high galactic latitude of $62^\\circ$.  For the first outburst shown in figure 3, the light curve profile around the peak is relatively sharp compared with the second one and no clear time lag is seen between the X-ray and optical outburst.  We can interpret this correlation through the occurrence of an ``inside-out'' type outburst in accretion disk which will generate the rapid rise in X-ray and simultaneous optical brightening, however this picture cannot predict the fast optical flare at the end of outburst (Smak 1984).  On the other hand, we suggest the ``outside-in'' type outburst for the second outburst to explain the optical precursor.  To observe these different types of outburst indicates that they may be caused by the different instability faculty such as normal and super outbursts of SU\\,UMa-type dwarf novae (Warner 1995).  From the X-ray flux of 39 mCrab reported in Remillard et al. (2000) and the simultaneous magnitude of optical counterpart of $\\sim$ 13 mag, the optical to X-ray flux ratio is calculated as $\\sim 5$ which is by two orders of magnitude smaller than the typical value of $\\sim 500$.  The X-ray intensity particularly shows a low value rather than optical, and this suggests an idea that it may be a high inclination system (Garcia et al. 2000).  The HST observation revealed the continuum slope somewhat flatter, which suggests the intrinsic X-ray flux is relatively low and hence removes the suggestion of the high inclination system (Haswell et al. 2000).  A high inclination system generally provides a high probability of occurrence of eclipses by secondary, whereas we have not detected them.  It is almost certain that the periodicity appeared in figure 3 reflects the orbital motion of the underlying binary, through the reprocessing the X-rays by the secondary, or possibly superhump modulations as seen in SU\\,UMa-type dwarf novae (O'Donoghue, Charles 1996).  We therefore suggest the orbital period of $0.17078$ d, which is relatively short compared with other low mass X-ray binaries (Ritter, Kolb 1998).  Since the quiescent magnitude is 18.8 mag, if we assume a K dwarf secondary star and no other optical source in quiescence, the distance is at least 1.5 kpc, that is, 1.3 kpc above the galactic plane.  The color of the quiescent counterpart in the USNO catalog ($b-r=+0.6$) is much bluer than that of a K or M dwarf, implying the substatial contribution from the accretion disk. The above distance estimate should be thus regarded as a lower limit, implying that XTE J1118+480 is a galactic halo X-ray transient. We can generally expect that such a galactic halo object is older than that in disk, and this is consistent with the short orbital period of 0.17078 d which implies a more evolved binary system.  It should be noted, however, if we assume a short orbital period of 0.17078 d, the secondary may be a smaller star than in other typical SXTs.  This case may be comparable to the system GRO J0422+32, whose orbital parameter and secondary type are 0.212 d and M2V, respectively (Orosz, Bailyn 1995; Filippenko et al. 1995).  An M2V star secondary leads to the lower limit of the distance of 0.55 kpc.  Optical or infrared spectroscopic observations are definitely needed to unambiguously determine the orbital period and the type of the secondary.  On the point of the high galactic latitude SXTs, 3U 0042+32 is similar to XTE J1118+480 (Ricketts, Cooke 1977).  The outburst in 1977 of 3U 0042+32 was detected by {\\it Uhuru} which classified it as a high galactic latitude ($\\sim -30^\\circ$) X-ray transient.  During the active phase in 1977 Feburuary, 3U 0042+32 experienced at least four distinct outbursts whose interval is estimated as 11.6 d (Watson, Ricketts 1978).  Compared with XTE J1118+480, the outburst duration seems to be quite shorter (decay e-folding time $= 2.8$ d).  Another noteworthy, possibly related, high galactic latitude X-ray binary is MS 1603+2600 = UW CrB (Morris et al. 1990).  This object is a persistent source whose orbital period is 111 min.  Although Hakala et al. (1998) suggested this X-ray binary may correspond to a quiescent state of SXTs, the discovery of a short-period transient system XTE J1118+480 may suggest that these two systems comprise a new variety of halo X-ray binaries."
    },
    "0004/astro-ph0004163_arXiv.txt": {
        "abstract": "The recently published precise spectrum of cosmic ray protons from the Alpha Magnetic Spectrometer has been examined in some detail from the standpoint of a search for deviations from a smooth, simple, power law. We find a significant excess (~$\\simeq$ 10\\% over $\\sim$ 0.3 interval in logE~) centered on 50 GeV in the published data. It is possible that the 'unfolding technique' adopted by the experimenters causes an overestimate of the excess but it is difficult to reduce it much below about (~5$\\pm$2~)\\%.  \\noindent We have examined other recent data, too. There is also evidence here, for an excess in the same energy region, although of only (~1.6$\\pm$0.9~)\\%. A value of (~3$\\pm$1.5~)\\% would be consistent with all the data. There are hints of similar excesses for heavier nuclei and electrons.  \\noindent Possible explanations are put forward for an excess, should it prove to be genuine. ",
        "introduction": "It is commonly asserted that the energy spectrum of cosmic rays is almost featureless, being characterised by a power law spectrum which slowly steepens near 3$\\cdot$10$^{15}$ eV (~the knee~) and continues with a new exponent to near 10$^{19}$ eV, where it flattens somewhat (~the ankle~).  In fact, we (~Erlykin and Wolfendale, to be referred to as EW~) have claimed that near the knee the situation is more complex (~\\cite{EW1} and references therein~). The complexity is attributed to the effect of a single, local and recent supernova, the remnant of which has accelerated particles, particularly oxygen and iron nuclei, with a spectrum of the form E$^{-2}$ up to a rather sharp maximum at $\\simeq 4 \\cdot 10^{14}Z$ eV. It is not unreasonable, statistically \\cite{Fatem}, that such a single 'source' should show itself in a rather narrow band of energy. ",
        "conclusions": "Inspection of the very recent, precise, AMS measurements of the primary cosmic ray proton spectrum shows an interesting excess in the region of 50 GeV. Although the quoted values may overestimate the excess - perhaps by a factor 2 - it seems unlikely that the whole effect will go away.  Including results from other proton experiments, too, an overall excess averaged over $\\Delta$ logE = 0.15 centred on 50 GeV of (~3$\\pm$1.5)\\% would fit all the data.  It should be possible to confirm or deny a value of this magnitude by a more sophisticated study of the AMS results.  More importantly, perhaps, is the general point that modern experiments of apparently high precision, should focus on examinations of the detailed shape of the energy spectrum of protons, of heavier nuclei and of electrons, too. Fine structure should be visible at some level; perhaps it is already starting to be seen ?\\\\ Concerning the explanation, if the heavier nuclei do, in fact, show similar effects to those for protons (~at present the results are clearly very marginal~) - at the same rigidity - it will be possible to rule out exotic processes as being responsible. Weak, local supernova remnants, would then be the preferred mechanism. However, it is premature to rule out exotic processes, yet.\\\\ {\\bf Acknowledgements}\\\\ We are grateful to Professor Yu.Galaktionov for useful discussions and comments. The Royal Society is thanked for the provision of financial support by way of a Joint Project. Mr. P.Kiraly kindly provided helpful comments, as did the referee, and we are grateful to them."
    },
    "0004/astro-ph0004355_arXiv.txt": {
        "abstract": "The TeV and X-ray data obtained by the imaging Cherenkov telescope CAT and X-ray satellite BeppoSAX during the remarkable flare of Mkn~501 in April~16,~1997 are used to constrain  the flux of the Cosmic Infrared Background (CIB) using  different CIB  models. We show that a non-negligible absorption of $\\gamma$-rays due to the CIB could take place already in the low-energy (sub-TeV) domain of the spectrum of Mkn~501. This implies that the data of the low-energy threshold  CAT telescope contain very important information about the CIB at short wavelengths, 0.4~$\\mu$m$\\leq \\lambda \\leq$~3.~\\mic. The analysis of almost simultaneous  spectroscopic measurements of Mkn~501 in a high state by CAT and BeppoSAX in the framework  of  the standard homogeneous Synchrotron-Self-Compton (\\mbox{SSC}) framework model leads to the conclusion that the density of the near-infrared background with typical ``starlight spectrum'' around 1~$\\mu$m should be between 5 and $35 \\, \\rm nW \\, m^{-2} \\, sr^{-1}$ (99$\\%$~CL), with most likely value around 20~$\\rm nW \\, m^{-2} \\, sr^{-1}$. Also we argue that  the CAT \\gr~data alone allow rather robust upper limits on the CIB, $\\lambda F_\\lambda \\leq 60 \\ \\rm nW \\, m^{-2} \\, sr^{-1}$ at 1~\\mic, taking into account  that for any reasonable scenario of \\gr~production the differential intrinsic spectrum of \\grs~hardly could be flatter than ${\\rm d} N/{\\rm d} E \\propto E^{-1}$. This estimate agrees, within statistical and systematic uncertainties,  with recent reports about tentative detections of the CIB  at 2.2 and 3.5 $\\mu$m by the Diffuse Infrared Background Experiment (DIRBE), as well as  with the measurements of the background radiation at optical wavelengths from absolute photometry. The high flux of CIB at $\\leq$~few~\\mic~ wavelengths implies a significant distortion of the shape of the initial (source) spectrum of $\\gamma$-rays from Mkn~501 at sub-TeV energies. The ``reconstructed'' intrinsic $\\gamma$-ray spectrum shows a distinct  peak in the Spectral Energy Distribution (SED) around 2~TeV with a flux by a factor of 3 higher than the measured flux.  The energy spectrum of gamma radiation from both sides of the peak has power-law behavior with photon index $\\alpha \\simeq 1.5$ below 2~TeV, and $\\alpha \\simeq 2.5$ above 2~TeV. This agrees with predictions of \\mbox{SSC} model. We also discuss the impact of the intergalactic absorption  effect in derivation of the \\mbox{SSC} parameters for the jet in Mkn~501. ",
        "introduction": "The Cosmic Infrared Background (CIB) is contributed mostly by the red-shifted `stellar' and `dust'  radiation components, and therefore carries vital cosmological information about the epoch of galaxy formation and their evolution in time.  The derivation of information about CIB from direct measurements is a hard task which requires an effective removal  of  heavy contamination caused by foregrounds of different origin, in particular  by the  zodiacal light, stellar and interstellar emission  of our Galaxy, etc (Arendt et~al. 1998;   Kelsall et~al. 1998). The conclusions of this  approach in the near-infrared domain are to a large extent   model-dependent because they are generally  based  on comprehensive  modeling of the foregrounds; the far-infrared background is determined with a better accuracy as it dominates  the foreground emission (Hauser et~al. 1998).  Very High Energy (VHE) \\gr~astronomy provides an independent and complementary approach for the study of the CIB. The idea is simple, and based on the detection of absorption features in the \\gr~spectra of distant extragalactic objects  caused by interactions  of VHE \\grs~with the CIB photons in their way from a  source to the observer (Nikishov 1962; Gould and Schreder  1967, Stecker et~al. 1992). The recent  detections of  \\grs~from two BL~Lac objects, Mkn~421 and Mkn~501, with  spectra  extending  up to 10~TeV and beyond, open an  interesting path for the realization of  this exciting cosmological aspect of VHE gamma-ray  astronomy. Obviously,  the success of  the `$\\gamma$-astronomical'  approach  essentially depends on two crucial conditions: (i) accurate \\gr~spectrometry, and (ii) good understanding of the  intrinsic (source) spectra of TeV  $\\gamma$-rays.  The observations of  Mkn~501, the  second  closest X-ray selected BL~Lac object  with a redshift $z \\simeq 0.034$, during its extraordinary outburst in 1997 yielded  unique data,   which initiated important theoretical studies of the  physical  conditions in the relativistic  jets of BL Lac objects (see {\\it e.g.} Tavecchio ~et~al. 1998 (hereafter TMG);  Kirk \\& Mastichiadis 1999, Hillas 1999; Coppi \\& Aharonian 1999a, Bednarek \\& Protheroe 1999, Krawczynski ~et~al. 1999), as well as interesting efforts to set meaningful   upper limits on the CIB flux (see   Biller ~et~al. 1998,  Stanev \\& Franceschini 1998, Barrau 1998, Stecker \\& De Jager 1998,  Stecker 1999, Aharonian ~et~al. 1999,  Coppi \\& Aharonian 1999b; Konopelko ~et~al. 1999).  During the 1997 outburst, lasted  several months,   Mkn~501 showed  dramatic variations of fluxes both  in X-rays (BeppoSAX:  Pian ~et~al. 1998,  RXTE: Lamer \\& Wagner 1998)    and  TeV \\grs~ (Whipple: Catanese ~et~al. 1997,   Samuelson ~et~al.  1998; HEGRA: Aharonian ~et~al.  1997, 1999; Telescope Array: Hayashida ~et~al. 1998; CAT:  Djannati-Ata\\\"{\\i} ~et~al.  1999). More importantly, the high TeV fluxes allowed monitoring of the energy spectrum of Mkn~501 on a diurnal basis, especially during very strong flares. On several occasions truly simultaneous observations of Mkn~501 were available in  TeV and X-ray band  (Krawczynski ~et~al. 1999). A special interest presents the famous April~16,~1997 flare which was observed by  BeppoSAX (Pian ~et~al. 1998) and low-energy threshold  ($\\sim 300 \\, \\rm GeV$) Whipple and   CAT atmospheric   Cherenkov telescopes (Catanese ~et~al. 1997; Djannati-Ata\\\"{\\i} ~et~al. 1999). Because of very large fluxes, the energy spectrum of the flare was obtained   with good accuracy in broad dynamical ranges in both X-ray (0.1-100~keV) and TeV \\gr~(0.3-10~TeV)  regimes. Within today's  most popular, the  so-called synchrotron-self-Compton (\\mbox{SSC}) model of non-thermal high energy radiation of Mkn~501,  the data of April~16,~1997 flare allow to compute the intrinsic source spectrum of $\\gamma$-rays. This information,  coupled with almost simultaneous spectral measurements of $\\gamma$-radiation by CAT,  provides a good opportunity to analyze  the intergalactic absorption signature in the observed VHE \\gr~spectrum.  Since the optical depth  $\\tau_{\\gamma \\gamma}$  of \\grs~in the intergalactic medium increases with energy (for any reasonable spectral shape of the CIB), most stringent constraints on the CIB come from the very energetic tail of the  \\gr~ spectrum  of Mkn~501. For $E \\geq$~10~TeV the $\\gamma - \\gamma$ absorption is dominated  by the `dust' component of CIB at far infrared (FIR)  wavelengths, $\\lambda \\geq$ 10~\\mic. Therefore, it is generally believed that a deep probe of CIB at  optical and near-infrared (NIR) wavelengths  is contingent only upon the discovery of more distant ({\\it e.g.} with $z \\geq 0.1$)  VHE \\gr~sources.  In this paper, however,  we show that the accurate measurements of the spectrum of Mkn~501 by CAT at sub-TeV energies already provide meaningful upper limits on the CIB flux at wavelengths between 0.4 and a few \\mic. Moreover, the analysis of the X-ray and \\gr~spectra of the April~16,~1997 flare within the homogeneous \\mbox{SSC} model  allows rather conclusive estimates of the CIB flux at such  short wavelengths. The CIB flux $\\lambda F_{\\lambda} \\sim 5-35 \\, \\rm nW\\ m^{-2} sr^{-1}$ at $\\lambda \\sim 1$~\\mic~gives a reasonable slope, $\\nu F_\\nu \\propto E^{0.5}$, in the `reconstructed'  Spectral Energy  Distribution (SED) of \\grs~at low  energies -~as expected within the framework of the \\mbox{SSC} model. The significant intergalactic absorption of sub-TeV \\grs~ leads to both the shift of the position ($E \\simeq 2 \\, \\rm TeV$ ) and increase of the flux ($\\nu F_\\nu \\simeq 10^{-9} \\, \\rm erg\\ cm^{-2}\\ s^{-1}$) of the so-called Compton peak in the \\gr~spectrum. An analytical approach has been recently proposed by (TMG) for derivation of constraints on the  jet parameters of TeV blazars. With these revised spectral parameters, new results in  a well defined  self-consistent \\mbox{SSC} parameter-space for the jet of Mkn~501 in the high state are obtained. ",
        "conclusions": "The {accurate spectrometric observations of the $\\gamma$-ray spectrum of Mkn~501 during the remarkable April~16,~1997 flare by the low-threshold  imaging  atmospheric Cherenkov telescope CAT in the energy region from 300~GeV to 10~TeV is used for extraction of important information about the CIB at wavelengths $\\geq$ 0.4~\\mic. The interpretation of the spectrum of sub-TeV \\grs, together with simultaneously obtained X-ray data of BeppoSAX,  requires, within the  one-zone \\mbox{SSC} model,  rather high NIR background at a level close to  $20 \\, \\rm nW \\, m^{-2} \\, sr^{-1}$ at 1~\\mic. Such high flux of CIB implies an essential distortion of the shape of the initial (source) spectrum of $\\gamma$-rays from Mkn~501 not only at multi-TeV, but also at sub-TeV energies. The `reconstructed' intrinsic $\\gamma$-ray spectrum shows a distinct  peak in the Spectral Energy Distribution around 2~TeV with a flux by a factor of 3 higher than the measured flux.  Moreover,  the energy spectrum of gamma radiation from both sides of the peak has power-law behavior with spectral indices $\\alpha \\simeq 0.5$ below 2~TeV, and $\\alpha \\simeq 1.5$ above 2~TeV, which  perfectly  agrees with predictions of the \\mbox{SSC} model.  We have shown that the intergalactic absorption has  non-negligible impact on  the construction of self-consistent \\mbox{SSC} parameters.  And finally, we argue that  the CAT \\gr~data alone allow rather robust upper limits on the CIB, $\\lambda F_\\lambda \\leq 60 \\ \\rm nW \\, m^{-2} \\, sr^{-1}$ at 1~\\mic, taking into account  that for any reasonable scenario of \\gr~production the differential intrinsic spectrum of \\grs~hardly could be flatter than ${\\rm d} N/{\\rm d} E \\propto E^{-1}$."
    },
    "0004/astro-ph0004025_arXiv.txt": {
        "abstract": "Terrestrial planets may be detected using the gravitational microlensing technique. This was demonstrated in the high magnification event MACHO-98-BLG-35. Observing strategies aimed at measuring the abundance of terrestrial planets are discussed, using both existing telescopes and planned telescopes. ",
        "introduction": "Significant advances have been made in recent years in the study of extra-solar planets using the radial velocity technique and, more recently, the transit technique (Mayor, these proceedings; Fischer, these proceedings). Approximately thirty Jupiter-mass planets have been detected by the radial velocity technique. Initial indications from transit measurements indicate they are gas giants similar to Jupiter. However, they have been detected in eccentric orbits at $>$ 0.2 AU or circular orbits at $<$ 0.2 AU, quite dissimilar to Jupiter's orbit. These surprising findings will assist our understanding of planetary systems and planetary formation.  The gravitational microlensing technique complements the above studies because it is sensitive to giant and terrestrial planets at orbital radii of a few AU. In this technique, perturbations of standard microlensing caused by planets are searched for. Two versions of the technique have been utilised to date. These involve large perturbations of low-magnification events, and small perturbations of high-magnification events, respectively. Most searches have been made in the galactic bulge where the observed rate of microlensing is maximal.  In the following sections strategies for optimising the detection rate of planets by microlensing, especially low-mass terrestrial planets, are discussed. The ultimate goal is to determine their abundances. To illustrate the technique, a high magnification event is described in the following section. Other examples of both high and low magnification events are described by Gaudi in these proceedings. ",
        "conclusions": ""
    },
    "0004/astro-ph0004213_arXiv.txt": {
        "abstract": "{We report on low-resolution multi-object spectroscopy of 30 faint targets ($R \\simeq 24-25$) in the HDF-S and AXAF deep field obtained with the VLT Focal Reducer/low dispersion Spectrograph (FORS1). Eight high-redshift galaxies with $2.75< z < 4$ have been identified. The spectroscopic redshifts are in good agreement with the photometric ones with a dispersion $\\sigma_z = 0.07$ at $z<2$ and $\\sigma_z = 0.16$ at $z>2$. The inferred star formation rates of the individual objects are moderate, ranging from a few to a few tens $M_{\\odot}~ {\\rm yr}^{-1}$. Five out of the eight high-z objects do not show prominent emission lines. One object has a spectrum typical of an AGN. In the AXAF field two relatively close pairs of galaxies have been identified, with separations of 8.7 and 3.1 proper Mpc and mean redshifts of 3.11 and 3.93, respectively. } ",
        "introduction": "Observations of galaxies, now extending up to a redshift $z \\sim 6$ (\\cite{hu99}), are starting to provide quantitative information on basic properties - number densities, luminosities, colors, sizes, morphologies, star formation rates (SFR), chemical abundances, dynamics and clustering - over a large span of cosmic time.  These data are beginning to sketch out a direct picture not only of the physical processes taking place in the assembly of the first galaxies, but also of the formation and evolution of large scale structure (LSS) from the primordial density fluctuations.  Quantitative information is now available about the evolution of the neutral hydrogen and metal content of the universe since $z\\simeq 4$, the galaxy luminosity function since $z\\simeq 1$, the morphology of field and cluster galaxies since $z\\simeq 0.8$. A recent dramatic addition to the general picture has been the discovery of a large population of actively star-forming galaxies at $z\\simeq 3$ (\\cite{SGDA96}, hereafter SGDA96). The ``Lyman break'' color-selection technique (\\cite{SH92}) has proved a reliable and highly efficient method to select galaxies in large numbers at $z\\ga 2.5$, providing the first opportunity for statistical studies of evolutionary processes in galaxies beyond $z=1$.  Follow-up spectroscopy of the UV drop-out candidates on the Keck telescopes shows most to lie in the expected redshift range, $2.5 \\leq z < 3.5$, with successful redshift measurement for more than $70 \\%$. The Lyman-break galaxies have spectra resembling those of nearby starburst galaxies, are strongly clustered, with a co-moving correlation length similar to present-day galaxies. SGDA96 inferred typical SFRs of $1-6 h^{-2} {\\rm M}_{\\odot} {\\rm yr}^{-1}$ for their galaxies, assuming a critical density universe. Dust corrections based on the UV continuum slope and on near-IR spectroscopy of a few objects suggest values larger by a mean factor of about 7 (\\cite{pettini98}). From the width of saturated interstellar absorption lines, SGDA96 inferred tentative 1D velocity dispersions in the range $\\sigma_{1D} = 180 - 320 {\\rm km \\, s}^{-1}$, but Pettini et al. (1998) measure in the IR significantly narrower line-widths $\\sigma_{1D} = 55 - 190 {\\rm km \\, s}^{-1}$ for the Balmer and [OIII] emission lines, albeit in a sample of only five objects.  A programme has been started with the ESO VLT to study systematically galaxies at $z \\sim 4$ with the aim to clarify the earliest phases of the processes leading to the formation of galaxies and LSS, reaching a redshift domain where observations are more cosmologically discriminant (\\cite{arnouts99}) and taking advantage of a wide photometric coverage (in particular in the IR) to obtain mass estimates of the detected objects.  We report here the results of pilot observations carried out during the commissioning and the science verification of the FORS1 instrument at the VLT-UT1. ",
        "conclusions": "At present, the spectroscopy of candidate Lyman-break galaxies has been restricted to an area of 13.5 sq.arcmin in the HDF-S in which $UBVRIJHK$ imaging is available and 25 sq.arcmin in the AXAF deep field ($22$ sq.arcmin covered in $UBVRIJK$ and 3 sq.arcmin in $UBVRI$ only). {\\it After} the FORS1 spectroscopic observations, which were based on a preliminary list of Lyman-break candidates produced shortly after the EIS imaging observations, we carried out a more refined selection of galaxies with $z>2.75$ on the basis of a photometric redshift code (described in Arnouts et al. 1999). In the HDF-S $25$ candidates have been found down to a limiting mag of $I_{AB} = 24.5$, while in the AXAF deep field $36$ candidates have been selected down to $R_{AB} = 25$. Of the total $61$ candidates $10$ turned out to have been observed during Commissioning and Science Verification: 8 of them have been confirmed to be at high redshift, 2 resulted in inconclusive spectra. \\begin{figure} \\psfig{figure=cristiani.fig2.ps,width=90mm} \\caption{Comparison of photometric and spectroscopic redshifts in the HDF-S and AXAF deep field.} \\label{fig:DZ} \\end{figure} Fig.~\\ref{fig:DZ} shows the comparison between photometric and spectroscopic redshifts for the 8 galaxies with $z>2.75$ observed so far in the HDF-S and AXAF field. The resulting dispersion is $\\sigma_z (z>2.75) = 0.16$.  At lower redshift, including some preliminary results in the HDF-S of Dennefeld et al. (2000, in preparation, see {\\tt http://www.iap.fr/hst/tmrresults.html}, the dispersion turns out to be $\\sigma_z (z<2)= 0.07$ (47 objects).  The properties of the high-z galaxies confirmed with the present observations are summarized in Table~\\ref{tab:properties}. They have been inferred from the GISSEL models best-fitting the photometric data (\\cite{arnouts99}) imposing the redshift at the spectroscopic value. The Star formation rate (Column~2) is estimated with the UV continuum flux at 150 nm derived from the best fitting model. For a Salpeter IMF ($0.1 M_{\\odot} < M < 125_{\\odot}$) with constant SFR, a galaxy with SFR $=1~ M_{\\odot}$ yr $^{-1}$ produces $L(150$ nm $)=10^{40.15}$ erg ~ s$^{-1}$ \\AA$^{-1}$ (\\cite{madau96}). Column~4 lists the SFR computed with the correction of the intrinsic extinction (reported in Column~3), as obtained form the best-fit procedure. The Calzetti extinction law (\\cite{calzetti97}) has been adopted. Column~5 and 6 show the estimated age \\footnote{In this paper $H_o =50$ and $q_o=0.5$ are assumed throughout.} and stellar mass. As typically found in surveys based on the ``Lyman-break'' technique, the inferred star formation rates are moderate, ranging from a few to a few tens $M_{\\odot}$ yr$^{-1}$. Five out of eight high-z objects do not show prominent emission lines. AXAF EIS-U21 has a spectrum typical of an AGN, showing Ly-$\\alpha$, CIV and possibly SiIV in emission with a P-Cyg profile.  \\begin{table} \\caption{Properties of the Galaxies with $z>2.5$ in the HDF-S and AXAF fields.} \\begin{tabular}{lccccc} \\hline \\hline Identifier&SFR&$E_{BV}$&SFR&Age&Mass\\\\ &Uncorr.& &Corr.& & \\\\ &$M_{\\odot}/$yr&mag&$M_{\\odot}/$yr&Gyr&$\\log(M_{\\odot})$\\\\ \\hline HDF-EIS43& ~8 & 0.0 & ~8 & 0.7 & 10.6 \\\\ HDF-EIS47& 15 & 0.1 & 38 & 0.1 & 11.2 \\\\ HDF-EIS53& 12 & 0.1 & 19 & 0.1 & 10.8 \\\\ AX-EISU28& ~6 & 0.2 & 13 & 1.4 & 10.9 \\\\ AX-EISU12& ~6 & 0.0 & ~6 & 0.1 & 10.0 \\\\ AX-EISU21& AGN \\\\ AX-EISB07& 12 & 0.0 & 12 & 1.0 & 10.5 \\\\ AX-EISB02& 12 & 0.0 & 12 & 0.7 & 10.5 \\\\ \\hline \\label{tab:properties} \\end{tabular} \\end{table}  It is interesting to note that two relatively close pairs of galaxies are observed in the AXAF field.  EISU28/EISU12 and EISB07/EISB02 are separated of only 8.7 and 3.1 proper Mpc, respectively. Given the small number of objects in the present sample any statistical conclusion is obviously impossible, but it appears natural to link  the occurrence of the two pairs to the redshift ``spikes'' observed by Steidel et al. (1998) at $z\\sim3$. Future observations of the remaining high-z galaxy candidates and the extension of the surveyed area (see {\\tt http://www.eso.org/science/eis/}) will make it possible to address also this issue on a more quantitative basis."
    },
    "0004/astro-ph0004096_arXiv.txt": {
        "abstract": "s{In a model recently proposed by Albrecht and Skordis\\cite{Albrecht:99} it was suggested that the observed accelerated expansion of the universe could be caused by a scalar field which is trapped in a local minimum of an exponential potential modified by a polynomial prefactor. We show that scalar field cosmologies with this kind of local minimum in the potential are stable and do not decay to the true vacuum if they fulfill the observational constraints from the Type Ia Supernovae experiments. Further we briefly sketch how this potential could be related to a potential of interacting D-branes.} ",
        "introduction": "Recent observations by the Supernovae Cosmology Project (SCP)\\cite{Perlmutter:98} and the High-Z Supernovae Search Team\\cite{Riess:98} have revealed that there is an energy component in the universe which is dark and has negative pressure. This is confirmed if one combines observations of the anisotropy in the cosmic microwave background (CMB) radiation and clusters\\cite{Knox:00}. The simplest way to achieve accelerated expansion of the universe is by introducing an {\\em ad hoc} cosmological constant. The findings of the SCP are that the probability for a non-vanishing cosmological constant is 99\\%\\cite{Perlmutter:98}. If one assumes a flat universe, which seems to be confirmed by recent CMB observations, the universe consists of 30\\% matter and 70\\% cosmological constant or dark energy component\\cite{Knox:00}. A more general way to obtain accelerated expansion of the universe is by introducing a scalar field $\\phi$ which either slowly rolls down a potential or is trapped in a local minimum. In this way the universe eventually becomes vacuum dominated and therefore the expansion accelerates. During recent years dark energy or quintessence models have been proposed\\cite{models:88-00}, with some of these models needing tuning of the initial conditions to fulfill the observational constraints, but the majority requiring only a tuning of the parameters of the model. The novel feature of the model by Albrecht and Skordis\\cite{Albrecht:99} is that the parameters involved are roughly of order one in Planck units ($M_{\\rm Pl} \\approx 2.44 \\times 10^{18} {\\rm GeV}$) and only moderate tuning is required to be within the observational constraints. Throughout this paper we use natural units ($\\hbar = c = 1$) and set the Planck mass $M_{\\rm Pl} = \\left(\\hbar c^5/ 8\\pi G\\right)^{1/2} = 1$. ",
        "conclusions": "The considerations above show that the Albrecht and Skordis model is a stable solution and the probability that the field decays to the true vacuum state is negligible. We also observed that if we move the feature in the potential to smaller field values $\\phi \\leq 1$ the tunneling rate is significant and $B \\sim {\\cal O}(1)$ in Planck units. However a feature in the potential at such small field values could not fulfill the observational constraints and resolve the dark energy problem, since the field gets trapped in the local minimum too early. Since the false vacuum state does not decay for the relevant range of parameters, a classical description is sufficient and the field stays trapped in the local minimum. The dark energy model with a polynomial exponential is therefore a possible solution to the problem of the missing dark energy. One might wonder if it is possible to connect such a potential to fundamental physics. A potential with similar features to the one discussed above is an inverse polynomial exponential of the form $V(\\phi) = \\left[(\\phi-\\beta)^2+\\delta\\right]^{-1}\\exp(-\\lambda\\phi)$. It may be possible  to relate this potential to the massive bulk modes of interacting D3-branes, which involves similar contributions\\cite{Dvali:98}. In this case the false vacuum state is also very stable with $B \\approx 10^{120}$. The question of how realistic a rational exponential potential which resolves the dark energy problem is remains open to this point but is under current investigation.\\\\ {\\bf Acknowledgment}\\\\ We would like to thank A.~Albrecht for fruitful collaboration\\cite{Weller:00} and acknowledge financial support from the German Academic Exchange Service (DAAD)."
    },
    "0004/astro-ph0004269_arXiv.txt": {
        "abstract": "Microlensing is the only technique likely, within the next 5 years, to constrain the frequency of Jupiter-analogs.  The PLANET collaboration has monitored nearly 100 microlensing events of which more than 20 have sensitivity to the perturbations that would be caused by a Jovian-mass companion to the primary lens. No clear signatures of such planets have been detected.  These null results indicate that Jupiter mass planets with separations of 1.5-3 AU occur in less than 1/3 of systems.  A similar limit applies to planets of 3 Jupiter masses between 1-4 AU. ",
        "introduction": "A Galactic microlensing event occurs when a massive, compact object (the lens) passes near to our line-of-sight to a more distant star (the source).  If the lens, observer, and source are perfectly aligned, then the lens images the source into a ring, called the Einstein ring, which has angular radius of\\footnote{For the scaling relation on the far right of equations (1), (2), and (3), we have assumed $\\dos=8~\\kpc$ and $\\dol=6.5~\\kpc$, typical distances to the lens and source for microlensing events toward the bulge.} \\begin{equation} \\thetae \\equiv\\left[{4GM \\over c^2} {\\dls \\over \\dol \\, \\dos}\\right]^{1/2} \\sim 480 \\, \\mu{\\rm as} \\left({M \\over M_{\\odot}}\\right)^{1/2}, \\label{eqn:thetae} \\end{equation} where $M$ is the mass of the lens, and $\\dls$, $\\dos$, $\\dol$ are the lens-source, observer-source, and observer-lens distances, respectively.  This corresponds to a physical distance at the lens plane of \\begin{equation} \\re = \\thetae \\dol \\sim 3~\\au \\left({M \\over M_{\\odot}}\\right)^{1/2}. \\label{eqn:re} \\end{equation}   If the lens is not perfectly aligned with the line-of-sight, then the lens splits the source into two images.  The separation of these images is ${\\cal O}(\\thetae)$ and hence unresolvable.  However, the source is also magnified by the lens, by an amount that depends on the angular separation between the lens and source in units of $\\thetae$.  Since the lens, observer, and source are all in relative motion, this magnification is a function of time: a `microlensing event.'  The time scale for such an event is \\begin{equation} \\te \\equiv { \\thetae  \\over \\murel} \\sim 40~\\days \\left({M\\over M_{\\odot}}\\right)^{1/2}, \\end{equation} where $\\murel$ is the relative lens-source proper motion.  If the primary lens has a planetary companion, and the position of this companion happens to be in the path of one of the two images created during the primary event, then the planet will perturb the light from this image, creating a deviation from the primary light curve (see Figure 1).  The duration of this perturbation is $\\sim \\sqrt{q} \\te$, where $q$ is the mass ratio between the planet and primary.  Hence, for a Jupiter/Sun mass ratio ($q\\simeq10^{-3}$), the perturbation time scale is ${\\cal O}({\\rm day})$.  These short-duration deviations are the signatures of planets orbiting the primary lenses. Note that since the perturbation time scale is considerably less than $\\te$, the majority of the light curve will be indistinguishable from a single lens.  Three parameters determine the magnitude of the perturbation, and hence define the observables.  These are mass ratio $q$, the instantaneous angular separation $d$ between the planet and primary in units of $\\re$, and the angle $\\alpha$ between the projected planet/star axis and the path of the source.  As $q$ decreases, the perturbation time scale decreases, although the magnitude of the deviation does not necessarily decrease.  Thus very small mass ratio planets ($q \\la 10^{-5}$) can be detected using microlensing, although the detection probability is small. The lower limit to the detectable $q$ is set practically by the sampling of the primary event, and ultimately by the finite size of the source stars (Bennett \\& Rhie 1996). A microlensing event is generally alerted only if the minimum angular impact parameter in units of $\\thetae$ satisfies $\\umin \\leq 1$, which corresponds to image positions between $(0.6-1.6)\\thetae$. Since the planet must be near one of these images to create a perturbation, microlensing is most sensitive to planets with separations $0.6 \\la d \\la 1.6$, the `lensing zone.'  The angle $\\alpha$, which is of no physical interest, is uniformly distributed. Only certain values of $\\alpha$ will create detectable deviations. Thus integration over $\\alpha$ defines a geometric detection probability.  \\begin{figure}[t] \\plotfiddle{fig1.eps}{0.5cm}{0}{70}{70}{-230}{-300} \\vskip6.0cm \\caption{ Left: The images (dotted ovals) are shown for several different positions of the source (solid circles), along with the primary lens (dot) and Einstein ring (long dashed circle).  If the primary lens has a planet near the path of one of the images, i.e. within the short-dashed lines, then the planet will perturb the light from the source, creating a deviation to the single lens light cure. Right: The magnification as a function of time is shown for the case of a single lens (solid) and accompanying planet (dotted) located at the position of the X in the top panel. If the planet was located at the + instead, then there would be no detectable perturbation, and the resulting light curve would be identical to the solid curve. } \\end{figure}   Microlensing as a method to detect extrasolar planets was first suggested by Mao \\& \\Packy (1991), and was expanded upon by Gould \\& Loeb (1992) who demonstrated that if all lenses had a Jupiter analog, than $\\sim 20\\%$ of all light curves should exhibit $\\ga 5\\%$ deviations.  Since these two seminal papers, many authors have explored the use of microlensing to detect planets.  It is not our intention to provide a comprehensive review of this field.  However, of particular relevance is the paper by Griest \\& Safizadeh (1998, GS98) who demonstrated that, for high-magnification events (those with maximum magnification $\\amax > 10$), the detection probability for planets in the lensing zone is $\\sim 100\\%$.  Thus high-magnification events are an extremely efficient means of detecting extrasolar planets. The results of GS98 also imply that multiple planets in the lensing zone should betray their presence in high-magnification events (Gaudi, Naber \\& Sackett 1998). ",
        "conclusions": "Microlensing offers a unique and complementary method of detecting extrasolar planets.  Although many light curves have been monitored in the hopes of detecting the short-duration signature of planetary companions to the primary lenses, no convincing planetary detections have yet been made, despite the fact that data of sufficient quality are being acquired to detect such companions.  These null results indicate that Jupiter-mass companions with separations in the `lensing zone,' $1.5-3~\\au$, occur is less than 1/3 of systems.  The potential for this field is enormous.  Current microlensing searches for planets will continue to monitor events alerted toward the bulge, and either push these limits down to levels probed by radial velocity surveys ($\\sim 5\\%$), or finally detect planets, and measure the frequency of companions at separations more relevant to our solar system.  Next generation microlensing planet searches have the promise of obtaining a robust statistical estimate of the fraction of stars with planets of mass as low as that of the Earth."
    },
    "0004/astro-ph0004119_arXiv.txt": {
        "abstract": "We present a detailed study of the cluster pair SL\\,353 \\& SL\\,349. This candidate binary cluster is located at the northwestern rim of the LMC bar. Based on photometric data we find that both clusters are coeval with an age of 550$\\pm$100 Myr. We use the Ca\\,II triplet in the spectra of individual red giants to derive radial velocities. Both components of the binary cluster candidate show very similar mean velocities ($\\approx274\\pm10$ km\\,$\\mbox{s}^{-1}$ for SL\\,349 and $\\approx279\\pm4$ km\\,$\\mbox{s}^{-1}$ for SL\\,353) while the field stars show lower velocities ($\\approx240\\pm19$ km\\,$\\mbox{s}^{-1}$). These findings suggest a common origin of the two clusters from the same GMC. In this sense the cluster pair may constitute a true binary cluster. We furthermore investigate the stellar densities in and around the star clusters and compare them with isopleths created from artificial, interacting as well as non-interacting star clusters. Gravitational interaction leads to a distortion which can also be found in the observed pair. ",
        "introduction": "The Magellanic Clouds offer the unique possibility to study star clusters in general and binary clusters in particular. These two companion galaxies are close enough to resolve single stars, but distant enough to make the detection of close pairs of star clusters an easy task. About a decade ago, Bhatia \\& Hatzidimitriou (\\cite{bh}), Hatzidimitriou \\& Bhatia (\\cite{hb}), and Bhatia et al. (\\cite{brht}) surveyed the Magellanic Clouds in order to catalogue binary cluster candidates. To qualify as a binary cluster candidate, the maximum projected centre-to-centre separation of the components of a pair was chosen to be $\\approx 1\\farcm3$, which corresponds to $\\approx$ 19 pc in the LMC if a distance modulus of 18.5 mag is adopted. A binary cluster with larger separation may become detached by the external tidal forces while shorter separations may lead to mergers (Sugimoto \\& Makino \\cite{sm}, Bhatia \\cite{bhatia}). The number of chance-pairs of objects uniformly distributed in space can be estimated adopting a formula presented by Page (1975): Roughly half of the pairs found may be explained by mere chance line-up. This suggests that at least a subset of them must be true binary clusters, i.e., clusters that are formed together and/or may interact or even be gravitationally bound.  Depending on their masses and separations binary clusters will eventually merge or become detached. At one stage during the merger process the former binary cluster could have one single but elliptical core (Bhatia \\& McGillivray \\cite{bm}). Recently de\\,Oliveira et al. (\\cite{obd}) performed numerical simulations of star cluster encounters which could represent a possible scenario to explain the ellipticities found in several star clusters in the Magellanic Clouds.  Star clusters form in giant molecular clouds (GMCs), but the details of cluster formation are not yet understood (Elme\\-green et al. \\cite{eepz}). Fujimoto \\& Kumai (\\cite{fk}) suggest that binary or multiple star clusters form through strong, oblique collisions between massive gas clouds in high-velocity random motion, resulting in compressed sub-clouds revolving around each other. Since the components of a cluster pair formed together, they should be coeval or at least have a small age difference compatible with cluster formation time scales. Binary clusters are expected to form more easily in galaxies like the Magellanic Clouds, whereas in the Milky Way the required large-scale high-velocity random motions are lacking. Indeed only few binary open cluster candidates and no globular cluster pairs are known in our Galaxy (see for example Subramaniam \\& Sagar \\cite{ss}). However, binary globular clusters are not expected to survive the gravitational forces of the Milky Way (Surdin \\cite{surdin}).  In case of tidal capture two clusters would be gravitationally bound, but age differences are likely. Encounters of clusters can be traced using isodensity maps (de\\,Oliveira et al. \\cite{dodb}, Leon et al. \\cite{lbv}). Though van\\,den\\,Bergh (\\cite{vdbergh}) suggests that tidal capture becomes more probable in dwarf galaxies like the Magellanic Clouds with small velocity dispersion of the cluster system, Vallenari et al. (\\cite{vbc}) estimated a cluster encounter rate of $dN/dt\\sim1\\cdot(10^{9} \\mbox{yr})^{-1}$. This makes tidal capture of young clusters very unlikely.  The formation of low-mass star clusters tends to proceed hierarchically in large molecular complexes over several $10^7$ years (e.g., Efremov \\& Elmegreen \\cite{ee}). Leon et al. (\\cite{lbv}) suggest that in these groups the cluster encounter rate is higher and thus tidal capture is more likely: Binary clusters are not born together as a pair with similar ages but are formed later through gravitational capture. An observational test of this scenario would require the detection of evidence of tidal interactions between clusters, whose age differences need to be compatible with the survival times of GMCs.  Another binary cluster formation scenario is introduced by Theis (\\cite{theis}) and Ehlerov\\'a et al. (\\cite{epth}): Exploding supernovae close to the centre of a GMC sweep up the outer cloud material within a few Myrs and accumulate it in the shell. The large amount of matter makes the shell prone to gravitational fragmentation and may eventually lead to the formation of many open cluster-like associations (Theis et al. \\cite{teph}, Ehlerov\\'a et al. \\cite{epth}). In case of a dense ambient medium outside the cloud or a very massive original molecular cloud the shell can be strongly decelerated resulting in a gravitationally bound system of fragments or stars which can recollapse. A galactic tidal field acting on this recollapsing shell can split it into two or more large clusters, thus forming coeval twin globulars (Theis \\cite{theis}). These clusters may stay together for a long time, though they are gravitationally unbound. The evolution of their spatial separation mainly depends on the shape of the shock front at the time of fragmentation.  We are studying binary cluster candidates in the Magellanic Clouds to investigate whether the cluster pairs may be of common origin and if they show evidence for interaction. While it is so far impossible to measure true, deprojected distances between apparent binary clusters, an analysis of their properties can give clues to a possible common origin.  The binary cluster candidate SL\\,353 (or BRHT\\,33b, see Bhatia et al. \\cite{brht}) and SL\\,349 (BRHT\\,33a) is located in the outer western part of the LMC bar. Based on integrated colours, Bica et al. (\\cite{bcdsp}) suggest that SL\\,353 \\& SL\\,349 are coeval clusters of  SWB type V (Searle et al. \\cite{swb}) which is in agreement with the findings of Vallenari et al. (\\cite{vbc}). Leon et al. (\\cite{lbv}) suggest that the two clusters may even be physically connected. Based on ages for a large sample of star clusters, derived on the base of integrated colours, Bica et al. (\\cite{bcdsp}) propose an age gradient of the LMC bar. Younger clusters are predominantly found in the eastern part, while older clusters of SWB type III and higher are concentrated to its western end.  This paper is organized as follows. In Sect. \\ref{phot} we describe the photometric data in general. Stellar density maps are presented in Sect. \\ref{stardens}. The following section describes the colour magnitude diagrams (CMD) for the components of the cluster pair. Ages for each cluster are derived and compared with previous studies. The spectroscopic data are described in Sect. \\ref{spectroscopy}, and radial velocities are derived in Sect. \\ref{radvel}. A comparison of the observed isopleths with simulated ones for interacting and non-interacting star clusters is given in Sect. \\ref{nbody}. In the last Sect. \\ref{summary} we summarize and discuss the results. ",
        "conclusions": "\\label{summary}  Fitting isochrones based on the Geneva models (Schaerer et al. \\cite{smms}) to the CMDs, we find that the components of the double cluster SL\\,353 \\& SL\\.349 are coeval within the accuracy of our data, and we derive an age of 550$\\pm$100 Myr for both clusters.  The clusters are sufficiently old that the Ca\\,II triplet visible in the spectra of red giants could be used to derive radial velocities. 22 stars in and around the two star clusters were investigated. Most stars located inside the cluster areas show similar velocities of $\\approx277\\pm7$ km\\,$\\mbox{s}^{-1}$ whereas almost all field stars show lower velocities of $\\approx240\\pm19$ km\\,$\\mbox{s}^{-1}$. Two foreground stars were identified through their velocities which are too low to belong to the LMC population (14 km\\,$\\mbox{s}^{-1}$ and 84 km\\,$\\mbox{s}^{-1}$, respectively).  Both components of the binary cluster candidate are of the same age and, furthermore, have very similar mean radial velocities of $\\approx274\\pm10$ km\\,$\\mbox{s}^{-1}$ for SL\\,349 and $\\approx279\\pm4$ km\\,$\\mbox{s}^{-1}$ for SL\\,353. These findings support that both clusters may have formed at similar times and from the same GMC. In this sense they may constitute a true binary cluster.  We investigate the stellar densities in and around the binary cluster SL\\,353 \\& SL\\,349 and find an enhanced density between the two clusters. The smaller cluster SL\\,349 shows a distortion towards SL\\,353, and similar distortions can be seen around the more massive cluster.  Vallenari et al. (\\cite{vbc}) found a distortion of the isophotal contours, similar to our Fig. \\ref{sl353dens}, and a twisting of the isopleths which they regard as a sign of interaction and physical connection between the two clusters. It is remarkable that the age of this binary cluster is higher than the theoretical survival time of few $10^{7}$ yrs for physically connected cluster pairs suggested by Bhatia (\\cite{bhatia}). Leon et al. (\\cite{lbv}) suggest that SL\\,349 \\& SL\\,353 are part of a larger star cluster group in which interacting binary clusters are formed later via tidal capture, thus explaining an age larger than the theoretical survival time. However, another explanation could be that the survival time might be larger in the LMC bar where the tidal field might be weaker, as proposed by Elson et al. (\\cite{elson}, their Fig.\\,13).  Whether the distortion of the clusters is indeed a sign of possible interaction cannot be decided on the basis of our imaging data alone.  Based on observational parameters we created two artificial star clusters. The isopleths of the non-interacting pair as well as of a simulated interacting pair of star clusters were compared with the observed density plot. The density plot of the artificial, interacting pair shows a distortion of the more massive cluster which can also be found in the observed isopleth, however, the observed distortion of the smaller cluster SL\\,349 is not seen in the artificial density plot. It seems likely that this cluster pair shows signs of interaction, however, this does not necessarily imply that both clusters are gravitationally bound."
    },
    "0004/physics0004062_arXiv.txt": {
        "abstract": "One of the topical problems of contemporary physics is a possible variability of the fundamental constants. Here we consider possible variability of two dimensionless constants which are most important for calculation of atomic and molecular spectra (in particular, the X-ray ones): the fine-structure constant $\\alpha=e^2/\\hbar c$ and the proton-to-electron mass ratio $\\mu=m_p/m_e$. Values of the physical constants in the early epochs are estimated directly from observations of quasars -- the most powerful sources of radiation, whose spectra were formed when the Universe was several times younger than now. A critical analysis of the available results leads to the conclusion that present-day data do not reveal any statistically significant evidence for variations of the fundamental constants under study. The most reliable upper limits to possible variation rates at the 95\\% confidence level, obtained in our work, read: $$ |\\dot\\alpha/\\alpha| < 1.4\\times 10^{-14}{\\rm~yr}^{-1}, \\quad |\\dot\\mu/\\mu| < 1.5\\times10^{-14}{\\rm~yr}^{-1} $$ on the average over the last $10^{10}$ yr. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004082_arXiv.txt": {
        "abstract": "We analyze the temporal behavior of Network Bright Points (NBPs) using a set of data acquired during coordinated observations between ground-based observatories (mainly at the NSO/Sacramento Peak) and the Michelson Doppler Interferometer onboard SOHO.  We find that, at any time during the observational sequence, all the NBPs visible in the NaD$_2$ images are co-spatial within 1$^{\\prime\\prime}$ with locations of enhanced magnetic field. The ``excess'' of NaD$_2$ intensity in NBPs, i.e. the emission over the average value of quiet regions, is directly related to the magnetic flux density. This property implies that, in analogy with the Ca II K line, the NaD$_2$ line center emission can be used as a proxy for magnetic structures.  We also compare the oscillation properties of NBPs and internetwork areas. At photospheric levels no differences between the two structures are found in power spectra, but analysis of phase and coherence spectra suggests the presence of downward propagating waves in the internetwork. At chromospheric levels some differences are evident in the power spectrum between NBPs and internetwork. At levels contributing to the \\Na~emission the NBPs show a strongly reduced amplitude of oscillations at the $p$ - mode frequencies. At levels contributing to the \\Ha \\ core emission, the amplitude of network oscillations is higher than the internetwork ones. The power spectrum of NBPs at this wavelength shows an important peak at 2.2 mHz (7 minutes), not present in the internetwork areas. Its coherence spectrum with \\Ha \\ wings shows very low coherence at this frequency, implying that the oscillations at these chromospheric levels are not directly coupled with those present in lower layers. ",
        "introduction": "\\label{s_int}  The chromospheric bright network has long been observed in narrowband spectroheliograms taken in the H and K cores of the Ca II resonance lines. The Ca II network typically shows H and K profiles with high double peaks and enhanced line wings that persist for extended periods of time (longer than 10 minutes, see, e.g., Rutten \\& Uitenbroek 1991 ). The chromospheric network emission pattern is cospatial with small-scale magnetic field concentrations, and defines the supergranular network boundaries. It is this atmospheric component that produces the correlation between H and K excess line-core flux and magnetic activity of cool stars (\\cite{schri89}).  The dynamics of the network elements, compared with the internetwork or quiet chromosphere, has been extensively studied (especially from the observational point of view) since these small-scale structures can be important in channeling the energy from photospheric layers to the transition region and corona (\\cite{kne86}, 1993, \\cite{deu90}, \\cite{kul92}, % Al et al. 1998).  An assessment of the spectral characteristic properties of Network Bright Points (NBPs) at different layers in the atmosphere has been provided by Lites et al. (1993) using spectral observations in the range of the Ca II H line. In their work, these authors analyzed spectrographic observations of a single network bright patch and of several internetwork points. The wavelength shifts of photospheric and chromospheric lines allowed them to perform a compared analysis between the dynamics of the two atmospheric components. One of the relevant characteristics they describe is that at chromospheric levels (Ca II H$_3$) the NBPs show long period oscillations ($\\nu <$ 3 mHz) not correlated with oscillations in the lower atmosphere, while they lack power at higher temporal frequencies. The internetwork regions display instead enhanced power at higher frequencies, well correlated with photospheric oscillations. The presence of these low frequency oscillations in the network has been confirmed by Lites (1994) also for the chromospheric He I 10830 line, in contrast to Bocchialini et al. (1994) which observe, for the same line, oscillations only in the 5 minutes range.  An enhanced power in the low frequency range for network points with respect to the internetwork has also been observed by Kneer \\& von Uexk\\\"{u}ll (1986) in the center of the chromospheric \\Ha~ line. These authors however interpret this feature as not due to oscillations, but of mainly stochastic origin, and attribute it to erratic motions of the corresponding photospheric footpoints.  The problem is still open, and further observations to better address this issue are required (\\cite{lit94}). In particular one would need observations: 1- on a larger number of NBPs, to improve on the statistics; 2- at different heights in the atmosphere, since the analysis of the coherence between fluctuations at different levels can help exploring the nature of oscillations. To this end, a reliable method for the identification of the same physical structure at different atmospheric levels is mandatory, since the inclination of the magnetic field could displace the chromospheric network points with respect to the corresponding photospheric ones.  In this paper, we address some of these issues, and present observational results on the NBPs and internetwork characteristics as derived from a multiwavelength analysis. The observations were obtained in August 1996, during  a coordinated observing program between ground-based observatories and the Solar and Heliospheric Observatory (SOHO). For the ground-based observations we used the cluster of instruments at the NSO/Sa\\-cra\\-men\\-to Peak R.B. Dunn Solar Telescope (NSO/SP-DST), that could provide a complete coverage at lower atmospheric levels. The dataset used is described in Sect. \\ref{s_obs}. General properties of a sample of NBPs, followed from the photosphere up to the chromosphere and including their relationship with the magnetic structures, are given in Sect. \\ref{s_char}. The temporal development of the NBPs is described in Sect. \\ref{s_lcurves}. Sect. \\ref{s_power} and \\ref{s_phase} provide an analysis of the power, phase difference and coherence spectra for the fluctuations observed separately within the NBPs and the surrounding internetwork. Finally, discussion and conclusions are given in Sect. \\ref{s_concl}.  \\begin{table*}[ht] \\caption{Summary of the observations.  }\\label{tbl-1} \\begin{center} \\scriptsize \\begin{tabular*}{18.2cm}{llllll} Instrument & FOV & Spat. resol. & Observing $\\lambda$ (\\AA) & FWHM (\\AA) & $\\Delta$t (s) \\\\ & & & & & \\\\ UBF & 2$^\\prime \\times 2^\\prime$ & $0.5^{\\prime\\prime} \\times 0.5^{\\prime\\prime}$ & 5889.9 (NaD$_2$) & 0.2 & 12 \\\\ & & & 6562.8 (H$\\alpha$) & 0.25 & \\\\ & & & 6561.3 (H$\\alpha -1.5$ \\AA) & 0.25 & \\\\ Zeiss & 2$^\\prime \\times 2^\\prime$ & $0.5^{\\prime\\prime} \\times 0.5^{\\prime\\prime}$ & 6564.3 (H$\\alpha +1.5$ \\AA) & 0.25 & 3 \\\\ White Light & 2$^\\prime \\times 2^\\prime$ & $0.5^{\\prime\\prime} \\times 0.5^{\\prime\\prime}$ & 5500 & 100 & 3 \\\\ HSG & $0.75^{\\prime\\prime}\\times 2^\\prime$ & $0.75^{\\prime\\prime}\\times 0.36^{\\prime\\prime}$ & 3904$-$3941 (CaII K) & 0.035 &   \\\\ MDI &10$^\\prime \\times 6^\\prime$ &$0.6^{\\prime\\prime}\\times 0.6^{\\prime\\prime}$& 6768 (Ni I) & 0.1 &60 \\\\ \\end{tabular*} \\end{center} \\end{table*} ",
        "conclusions": "\\label{s_concl}  The observations presented in this paper allowed us to define the characteristics of network bright points at different atmospheric heights, and to compare them with those of the surrounding internetwork areas.  We improved on the existing statistics using a good-sized sample of NBPs, and the same number of ``test'' internetwork areas, defined in a comparable way. The method we adopted to study the temporal evolution of NBPs insures that each bright structure is properly followed in time and position at each height. In fact, the evaluation of the light curves and their properties after a spatial averaging over a well defined area guarantees that we are studying the same NBP at all heights, and avoids the problem (first pointed out by \\cite{lit94}) of a possible structure displacement  due to the magnetic field inclination. Given the characteristic horizontal size of the NBPs, the analysis and the comparison of power spectra and phase differences concern the propagation of waves pertaining to a horizontal wavenumber of about 3 Mm$^{-1}$.  The quasi-simultaneous series of \\Na~images and of MDI maps allowed us to establish for the first time a correspondence between \\Na~bright network and magnetic network at high spatial and temporal resolution. A correspondence between bright chromospheric structures (Ca II, Ly$\\alpha$, Mg I and UV continuum) and magnetic structures had been observed before, but not at this high temporal resolution. We also established for the NBPs a quantitative relationship between the Na excess and the corresponding absolute value of magnetic flux density. This relationship is best expressed by a power law with an exponent very close to the one found by Schrijver et al. (1989, 1996) for the Ca II - K excess, and indicates that the emission in \\Na~may be used as a proxy for the magnetic flux density.  The NBPs considered in this work have the following properties: - are bright in the Ca II wings and in the Ca II K$_2$ peaks; - are visible in the \\Na~ images for about 1 hr; - coincide spatially with the magnetic structures; - are nearby or within a lower activity region. The general characteristics found for these NBPs do not differ from the ones derived in absolutely quiet regions (\\cite{deu90}, \\cite{lit93}).  Our results referring to photospheric and chromospheric properties are so summarized:  {\\it At photospheric levels:} No difference is detected between network and internetwork power spectra, either in intensity or in velocity,   within the limits of sensitivity and accuracy of the instruments used for this work. The phase difference spectra between photospheric signatures  in general do not show different characteristics for network or internetwork. However, when analyzing the phase difference between \\Ha \\ red wing and white light images($\\Delta h \\leq 100$ km),  we find $ 5\\degr \\leq \\Delta\\Phi$ $\\leq 10\\degr$ \\  in the frequency window 1.5 -  2.5  mHz and in the internetwork. (The $\\Delta\\Phi$ value is more uncertain in the network, due to a lower coherence value). A phase lag of this amplitude and sign is usually considered a signature of gravity waves directed radially inward. A possible explanation for their origin might be sought in recent  models of convection, described as a non-local process driven by {\\it cooling} at the solar surface rather than by heating from the lower layers (Spruit, 1997). One can imagine that the downward flowing cooled plasma can trigger some inward directed waves, and hence justify the fact that the external layers ``lead'' the deeper layers. The general inhibition of convection in magnetic structures might be the reason for the lack of this signature in network points.  The power spectrum of the magnetic flux variations in NBPs shows a small but significant peak around 3 mHz, that could be related to a ``transformation'' of acoustic waves into MHD waves. However, the phase difference and coherence spectra between magnetic  flux and velocity (B$-$V) for the NBPs indicate a very low correlation between the two signals so we cannot conclude anything on the presence of MHD waves within the network points.  {\\it At chromospheric levels:} Network and internetwork areas have a rather different behaviour in the power spectra. We do not see any evidence for the typical chromospheric period of 3 minutes (but it must be reminded that they are best seen in velocity variations rather than intensity). In the low chromospheric levels, where \\Na~originates, the NBPs power spectrum is compressed at all frequencies if compared to the internetwork, while in the high chromosphere, where \\Ha \\ originates, the power of NBPs is higher than the one of internetwork. This opposite effect may be an indication that the magnetic field disturbs and reduces the amplitude of oscillations already present in the low chromosphere while it assumes a leading r\\^{o}le in the high chromosphere.  In the layers contributing to the \\Na~emission it seems that the oscillations present in network points change regime with respect to both the photosphere and the high chromosphere and we think that it would be important to perform observations of NBPs in tha Na line, with high spectral resolution. Unfortunately we cannot analyze the phase difference spectrum for \\Na~intensity fluctuations with respect to  others formed at different layers, since the \\Na~intensity fluctuations, measured with the UBF filter (FWHM$=0.2$\\AA), are more related to velocity than to temperature fluctuations (see Sect \\ref{s_power_chrom}).  The power spectrum of \\Ha \\ intensity in NBPs has the more relevant peak at 2.2 mHz, but this signal is not correlated with the photospheric fluctuations, as indicated by the very low coherence measured at all frequencies between the \\Ha \\ core and the blue and red wings. We can then confirm, using a larger sample of NBPs, the presence of the peak found by Lites et al. (1993) around 2 mHz in the power spectrum of K3 velocity fluctuations for one network point. Kalkofen (1997) and Hasan \\& Kalkofen (1999) proposed  an explanation for this peak in terms of transverse  magneto-acoustic waves in magnetic flux tubes, excited by granular buffeting in the solar photosphere. In their model the low coherence between photospheric and chromospheric signatures could be explained by a partial conversion of the transverse waves to longitudinal modes in the higher chromosphere.  A general result of our analysis, valid from the low photosphere to the high chromosphere, is that the NBPs always show  a coherence lower than the internetwork, pointing out that the presence of the magnetic field changes the propagation regime of waves  with respect to the non-magnetic regions."
    },
    "0004/astro-ph0004207_arXiv.txt": {
        "abstract": "We present a new method for fitting peculiar velocity models to complete flux limited magnitude-redshifts catalogues, using the luminosity function of the sources as a distance indicator. The method is characterised by its robustness. In particular, no assumptions are made concerning the spatial distribution of sources and their luminosity function. Moreover, selection effects in redshift are allowed. Furthermore the inclusion of additional observables correlated with the absolute magnitude -- such as for example rotation velocity information as described by the Tully-Fisher relation -- is straightforward.  As an illustration of the method, the predicted IRAS peculiar velocity model characterised by the density parameter $\\beta$ is tested on two samples. The application of our method to the Tully-Fisher MarkIII MAT sample leads to a value of $\\beta=0.6 \\pm 0.125$, fully consistent with the results obtained previously by the VELMOD and ITF methods on similar datasets. Unlike these methods however, we make a very conservative use of the Tully-Fisher information. Specifically, we require to assume neither the linearity of the Tully-Fisher relation nor a gaussian distribution of its residuals. Moreover, the robustness of the method implies that no Malmquist corrections are required.  A second application is carried out, using the fluxes of the IRAS 1.2 Jy sample as the distance indicator. In this case the effective depth of the volume in which the velocity model is compared to the data is almost twice the effective depth of the MarkIII MAT sample. The results suggest that the predicted IRAS velocity model, while successful in reproducing locally the cosmic flow, fails to describe the kinematics on larger scales. ",
        "introduction": "\\label{Introduction}  The study of the large-scale motions of galaxies in the Universe may provide valuable information concerning the dynamics of large-scale structures and the nature of the underlying dark matter. According to the gravitational instability scenario, the peculiar velocity field (i.e. the deviation from the smooth Hubble flow) may be used to infer the power spectrum of the mass fluctuations on intermediate scales and to constrain the cosmological density parameter $\\Omega$ (see for example Dekel 1994).  Since the discovery of the Great Attractor (Lynden-Bell et al. 1988), the field has proven to be particularly active.  Various observationnal programs have been completed, providing large and accurate datasets: e.g. the W91CL and W91PP samples (Willick 1990); MAT sample (Mathewson et al. 1992); HM sample (Han \\& Mould 1992); CF sample (Courteau et al. 1993); Abell BCG sample (Lauer \\& Postman 1994); SCI sample (Giovanelli et al. 1997a); KLUN sample (Theureau et al. 1997); nearby SNIa sample (Riess et al. 1997); MarkIII dataset (Willick et al. 1997b); SBF survey (Tonry et al. 1997); SFI sample (Giovanelli et al. 1998); SMAC sample (Hudson et al. 1999); EFAR and ENEAR samples (Colless et al. 1999 and Wegner et al. 1999); SCII sample (Dale et al. 1999); Shellflow survey (Courteau et al. 1999); LP10k survey (Willick 1999a).  Another significant advance during the past decade has been the improved understanding of the statistical formalism underlying the use of galaxy distance indicators -- and in particular the principles and practical methods of correcting for Malmquist bias. (See for example Hendry \\& Simmons 1990, Teerikorpi 1990, Bicknell 1992, Landy \\& Szalay 1992, Triay et al. 1994, Willick 1994, Hendry \\& Simmons 1994, Sandage 1994, Willick et al. 1995, Freudling et al. 1995, Willick et al. 1996, Rauzy \\& Triay 1996, Ekholm 1996, Triay et al. 1996, Rauzy 1997, Willick et al. 1997b, Giovanelli et al. 1997b, Theureau et al. 1998, Teerikorpi et al. 1999).  Several methods for extracting dynamical and kinematical information from distance indicator datasets have been proposed, i.e. the POTENT method (Bertschinger \\& Dekel 1989, Dekel et al. 1990, Bertschinger et al. 1990, Dekel et al. 1999 and references therein) and its variants (Rauzy et al. 1993 and 1995, Newsam et al. 1995); the ITF method (Nusser \\& Davis 1995, Davis et al. 1996, Da Costa et al. 1998); the VELMOD method (Willick et al. 1997a, Willick \\& Strauss 1998).  The comparison between the peculiar velocity or density fields inferred from distance indicator data with their corresponding fields derived from whole-sky redshift surveys has been one of the major issues addressed throughout the last decade. The question here is whether the spatial distribution of luminous matter e.g. the galaxies, traces the underlying mass fluctuations and if not, what are the properties of the ``biasing'' between the two fields?  Up to now, the point has not received any consensual answer. Indeed, the application of POTENT to various distance indicator datasets (Sigad et al. 1998 and references therein) favours a value of $\\beta_I = \\Omega_0^{0.6}/b_I \\simeq 1$ for the linear ``biasing'' density parameter, while the VELMOD and ITF fitting methods lead to a value of $\\beta_I \\simeq 0.5$ (Davis et al. 1996, Willick et al. 1997a, Riess et al. 1997, Da Costa et al. 1998, Willick \\& Strauss 1998). The origins of this significant discrepancy have not yet been elucidated (see for example Strauss 1999, Willick 1999b). At least one of these methods is suffering from some systematic effects, i.e. some statistical biases plaguing the estimate of the parameter $\\beta$ and not accounted for in the error analysis.  This remark leads us to the object of the present paper. The POTENT, VELMOD and ITF methods all require, at some stage of the analysis, to assume some a priori working hypotheses concerning the characteristics of the distance indicator dataset. For Tully-Fisher data for example, it will be assumed that the Tully-Fisher law is well described by a linear relation. These methods apply moreover under the hypothesis that the observational selection effects obey particular conditions. How are the results affected if one or more of the working assumptions fails to be satisfied by the dataset is generally a question not addressed in the error analysis.  The philosophy of the method we present herein is to reduce as far as possible the number of a priori hypotheses concerning the distance indicator sample. A direct consequence is that the range of application of the method will be considerably broadened.  The statistical background of the method is presented section 2. Its potential is illustrated by testing the predicted IRAS peculiar velocity model on two samples. In section 3, we perform the analysis using the fluxes of the IRAS 1.2 Jy survey as the distance indicator. We have deliberately chosen this sample looking to demonstrate the wide range of application of the method. Where the POTENT, VELMOD or ITF methods would not have been successful in extracting kinematical information from this dataset, our method does. We also treat a more classical case, the Tully-Fisher MarkIII MAT sample, in section 4. Finally, in section 5 we summarise our conclusions. ",
        "conclusions": "We presented a method for fitting peculiar velocity models to complete flux limited magnitude-redshift catalogues, using the luminosity function of the sources as a distance indicator, i.e. assuming that the distribution function of the absolute magnitudes of the galaxies does not depend on the spatial position.  Our method is based on a null-correlation approach. For a given peculiar velocity field model parametrised by a parameter $\\beta$, we defined a random variable $\\zeta_\\beta$, computable from the observed redshifts and apparent magnitudes of the sampled galaxies, which has the property of being statistically independent on the position in space (and thus on the modelled radial peculiar velocities themselves) if and only if the parameter $\\beta$ matches its true value $\\beta^\\star$. Therefore any test of independence between the random variable $\\zeta_\\beta$ and the modelled velocities or similar quantities provides us with an unbiased estimate of the value of $\\beta^\\star$. The method can be easily generalised to velocity models parametrised by an $N$-dimensional vector ${\\bf \\beta}=(\\beta_1,\\beta_2,...,\\beta_N)$.  The method is characterised by its robustness. No assumptions are made concerning the spatial distribution of sources and their luminosity function and selection effects in redshifts are also allowed. The required strict completeness in apparent magnitude can moreover be checked independently (Rauzy, in preparation). Furthermore the inclusion of additional observables correlated with the absolute magnitude is straightforward.  The predicted IRAS peculiar velocity model characterised by the density parameter $\\beta$ has been tested on two samples, the Tully-Fisher MarkIII MAT sample and the 60 $\\mu$m IRAS 1.2 Jy sample using the fluxes as the distance indicator.  The application of our method to the MarkIII MAT sample gives a value of $\\beta=0.6 \\pm 0.125$, in excellent agreement with the results obtained previously by the VELMOD and ITF methods on similar datasets. Our method is however more robust than these two fitting methods. In particular, we make a very conservative use of the Tully-Fisher information. We do not require to assume the linearity of the Tully-Fisher relation nor a gaussian distribution of its residuals.  We showed that our method allows to extract some valuable informations on the peculiar velocity field from the fluxes  of the IRAS 1.2 Jy sample. The poor accuracy of the distance indicator (due to the broad spread of the luminosity function) is balanced in this case thanks to the large number of galaxies contained in the sample. The IRAS sample permits to probe the cosmic flow at larger scales. Indeed, the mean effective depth of the volume in which the velocity model is compared to the data is almost twice the mean effective depth of the MarkIII MAT sample.  The application of our method to an IRAS subsample truncated in distance, of an effective depth similar to the MarkIII MAT sample, gives a value of $\\beta$ in accord with the values obtained using Tully-Fisher data. On the other hand when the application is performed on the whole sample, we found that the predicted IRAS velocity models with $\\beta \\ge 0.5$ can be rejected with a confidence level of $95\\%$. These results suggest that the predicted IRAS velocity model, while successful in reproducing locally the cosmic flow, fails to describe the kinematics on larger scales.  Note that these results do not lead to dismiss the linear ``biasing'' paradigm. As the errors on the predicted IRAS velocity field increase with distances, it could be that the predictions at the scales considered herein, i.e. beyond $5000$ km s$^{-1}$, drastically differ from the true cosmic flow (see for example Davis et al. 1995)."
    },
    "0004/astro-ph0004031_arXiv.txt": {
        "abstract": "We present three-dimensional numerical simulations of the rise and fragmentation of twisted, initially horizontal magnetic flux tubes which evolve into emerging $\\Omega$-loops. The flux tubes rise buoyantly through an adiabatically stratified plasma that represents the solar convection zone.  The MHD equations are solved in the anelastic approximation, and the results are compared with studies of flux tube fragmentation in two dimensions.  We find that if the initial amount of field line twist is below a critical value, the degree of fragmentation at the apex of a rising $\\Omega$-loop depends on its three-dimensional geometry: the greater the apex curvature of a given $\\Omega$-loop, the lesser the degree of fragmentation of the loop as it approaches the photosphere.  Thus, the amount of initial twist necessary for the loop to retain its cohesion can be reduced substantially from the two-dimensional limit.  The simulations also suggest that as a fragmented flux tube emerges through a relatively quiet portion of the solar disk, extended crescent-shaped magnetic features of opposite polarity should form and steadily recede from one another.  These features eventually coalesce after the fragmented portion of the $\\Omega$-loop emerges through the photosphere. ",
        "introduction": "The largest concentrations of magnetic flux on the Sun occur in active regions.  Great progress has been made over the past decade in understanding the connections between the magnetic field in active regions, observed at the surface of the Sun, to the magnetic field deep in the solar interior. Active regions have a bipolar structure, suggesting that they are the tops of magnetic flux loops which have risen from deep in the solar interior. On average, active regions are oriented in the E-W direction (Hale's Polarity Law) suggesting that the underlying field geometry is toroidal. The persistence of Hale's law for periods of several years during a given solar cycle suggests that magnetic flux must be stored in a relatively stable region of the solar interior.  Several stability arguments \\citep{svb82,vb82,fms93,fms95} show that the only place where such fields can be confined stably for periods of several years is below the solar convection zone.  On the other hand, if magnetic fields are placed any significant distance below the top of the radiative zone, they are so stable they could not emerge on the time scale of a solar cycle.  We are thus led to the conclusion that the most likely origin of active region magnetic fields is from a toroidally oriented field layer residing in the ``convective overshoot region'', a thin, slightly convectively stable layer just beneath the convection zone.  This layer also seems to coincide with the ``tachocline'' \\citep{k96,c99}, where the solar rotation rate transitions from solid body behavior in the radiative zone, to the observed latitudinally dependent rotation rate we see at the Sun's surface. This suggests that not only are solar magnetic fields stored in the convective overshoot layer, the overshoot layer is also the most likely site for the solar cycle dynamo \\citep{gmd89,dg91,p93,mc97a,mc97b,d97,dc99}.  Over the past decade, most efforts to study the emergence of active region magnetic fields have employed the ``thin flux tube'' approximation.  This model assumes that magnetic flux tubes behave as distinct tube-like entities, surrounded by field-free plasma.  The approximation further assumes that the tube diameter is small compared to all other length scales in the problem, and that pressure balance exists across the tube at all times.  After adopting these assumptions, it is straightforward to derive an equation of motion for the dynamics of the tube from the momentum equation in MHD. Thin flux tube models of emerging active regions have proven very successful in explaining many properties of active regions in terms of flux tube dynamics in the solar interior.  For example, they have successfully explained the variation of active region tilt with respect to the E-W direction as a function of solar latitude \\citep{dh93,dc93,ffh95}, the asymmetric orientation of the magnetic field after emergence \\citep{vp90,msc94,cmv96}, and the observed scatter in tilts as a function of active region size \\citep{lf96}.  In spite of these successes, recent two-dimensional MHD simulations of flux tube emergence have shown results which seem to invalidate many assumptions that are adopted in the thin flux tube approximation. \\citet{s79} and \\citet{lfa96} find that an initially buoyant, untwisted flux tube will fragment into two counter-rotating tube elements which then separate from one another, essentially destroying the tube's initial identity.  \\citet{mie96} and \\citet{fzl98} have demonstrated via two-dimensional MHD simulations that in order to prevent a flux tube from fragmenting, enough twist must be introduced into the tube to provide a cohesive force to balance the hydrodynamic forces acting to rip it apart.  That critical twist is defined, roughly, by that necessary to make the Alfv\\'en speed from the azimuthal component of the field at least as great as the relative velocity between the tube and the field free plasma surrounding it \\citep{llf96,em98,fzl98}.  But when global levels of twist in active regions \\citep{pcm95,lfp98} are determined from vector magnetograms, the amplitude of the observed twist is typically far smaller than this critical value \\citep{llp99}.  When plotted as a function of active region latitude, the twists exhibit large scatter, but superimposed on this apparently random behavior there is a slight, but clearly discernible trend for active regions in the northern hemisphere to be negatively twisted, while those in the south are positively twisted.  \\citet{lfp98} have developed a theoretical model which not only explains the latitudinal variation of twist, but also can account for the large fluctuations in twist from active region to active region.  In this model, an initially untwisted flux tube rises through the convection zone in accordance with the thin flux tube approximation.  Coriolis forces acting on convective eddies produce a non-zero average kinetic helicity, which is proportional to latitude.  The kinetic helicity acts to ``writhe'' the flux tube, which is then twisted in the opposite direction to preserve its magnetic helicity.  \\citet{lfp98} showed that this model can explain the observed data.  Yet the model assumes from the beginning that the thin flux tube approximation can be used, even for an initially untwisted tube, while the two-dimensional  MHD simulations suggest that this is invalid. Is there some way out of this quandary, which we dub ``Longcope's Paradox''?  In this paper, we describe MHD simulations of flux tube fragmentation in three dimensions.  The result of these simulations is that the critical degree of twist necessary to prevent fragmentation is reduced dramatically by the presence of flux tube curvature, as will be present in an emerging $\\Omega$-loop.  We find that for a fixed amount of twist, the degree of fragmentation is a function of the tube's curvature, and transitions asymptotically to the two-dimensional limit as the curvature approaches zero.  Even for flux tubes with little to no initial twist, the fragmenting magnetic morphology at the apex of an emerging $\\Omega$-loop is considerably less dispersed than the two-dimensional simulations would indicate --- a finding consistent with the simulation results of \\citet{dn98}.  It is not clear at this time whether our results resolve Longcope's paradox or not, but they certainly ameliorate the problem a great deal.  The remainder of this paper is organized as follows: In Section~\\ref{method} we briefly discuss the formalism of the anelastic approximation employed in our models, together with the numerical methods used to solve the system of equations.  We also describe the range of initial configurations that we use to explore the relationship between flux tube fragmentation, tube geometry, and the initial twist of the field lines.  At the beginning of Section~\\ref{results}, we define what is meant by a flux tube in the context of our three-dimensional MHD simulations, and further define a quantitative measure of the degree of fragmentation of such a tube.  We then present the results of our numerical simulations and discuss the implications of the models.  Finally, in Section~\\ref{conclusions} we summarize our conclusions.  \\clearpage ",
        "conclusions": "We have performed detailed numerical simulations of the rise of twisted magnetic flux tubes embedded in a non-magnetic, stratified plasma using a code which solves the three-dimensional MHD equations in the anelastic approximation. The evolving magnetic field is described in terms of its volumetric flux distribution as a magnetic flux tube which may fragment into separate, distinct tubes during its rise toward the photospheric boundary.  We find that the degree of fragmentation of the evolving magnetic flux tube depends not only on the initial ratio of the azimuthal to axial components of the field along the tube (as was the case in two dimensions), but also on the three-dimensional geometry of the tube as it rises through the convection zone.  The principal results of our analysis are the following: \\begin{enumerate} \\item If the ratio of azimuthal to axial components of the field along a magnetic flux tube exceeds a critical limit, then it will retain its cohesion and not fragment as it rises through the plasma.  The critical limit occurs when the rise speed is approximately equal to the maximum Alfv\\'en speed of the azimuthal component of the field.  This reinforces the conclusions of \\citet{em98}, \\citet{mie96}, \\citet{lfa96}, and \\citet{fzl98}. \\item If the initial amount of field line twist is less than the critical value, and if the magnetic flux tube rises toward the photosphere as an $\\Omega$-loop, then the degree of apex fragmentation depends on the curvature of the loop --- the greater the apex curvature, the lesser the degree of fragmentation for a fixed amount of initial twist. \\item  In two dimensions, counter-rotating vortices are effectively infinite in extent and generate long-range flows which eventually prevent the continued vertical rise of tube fragments. This artificial geometric constraint is relaxed in three dimensions, and although the fragments continue to experience forces due to the interaction of the vortex pairs, those forces are only important along a short section of the tube.  Thus, the fragments are able to rise to the surface.  The forces due to vortex interaction depend upon the geometry of the $\\Omega$-loop --- the greater the apex curvature, the lesser their magnitude.  Thus, highly curved loops exhibit less fragmentation during their rise. \\item Differential circulation between the apex and footpoint of an $\\Omega$-loop leads to the introduction of new magnetic twist of opposite sign in each leg of a loop fragment. This twist reduces the circulation about the apex of each fragment, and further reduces the forces acting to separate the fragments of the flux tube. This result is consistent with the predictions of \\citet{em98} and \\citet{m97}. \\item  Though these models do not admit to a direct, detailed comparison with observations, it is possible to infer certain general observational characteristics of emerging magnetic flux if the field configuration is that of a fragmented $\\Omega$-loop rising through the solar surface.  If the magnetic field erupts through a relatively quiet portion of the Sun, then we expect that one should observe concentrations of vertical flux which resemble expanding oval shapes. This type of feature should quickly evolve into longer-lived crescent-shaped regions of opposite polarity that steadily move away from one another.  These extended regions should then coalesce once the fragmented portions of the $\\Omega$-loop have emerged through the solar surface.  \\end{enumerate}  \\clearpage"
    },
    "0004/astro-ph0004341_arXiv.txt": {
        "abstract": "Following the discovery of the cosmic microwave background, the hot big-bang model has become the standard cosmological model.  In this theory, small primordial fluctuations are subsequently amplified by gravity to form the large-scale structure seen today. Different theories for unified models of particle physics, lead to different predictions for the statistical properties of the primordial fluctuations, that can be divided in two classes: gaussian and non-gaussian.  Convincing evidence against or for gaussian initial conditions would rule out many scenarios and point us towards a physical theory for the origin of structures.  The statistical distribution of cosmological perturbations, as we observe them, can deviate from the gaussian distribution in several different ways. Even if perturbations start off gaussian, non-linear gravitational evolution can introduce non-gaussian features. Additionally, our knowledge of the Universe comes principally from the study of luminous material such as galaxies, but galaxies might not be faithful tracers of the underlying mass distribution. The relationship between fluctuations in the mass and  in the galaxies distribution ({\\it bias}), is often assumed to be local, but could well be non-linear. Moreover, galaxy catalogues use the redshift as third spatial coordinate: the resulting redshift-space map of the galaxy distribution is non-linearly distorted by peculiar velocities. Non-linear gravitational evolution, biasing, and redshift-space distortion introduce non-gaussianity, even in an initially gaussian fluctuation field.  I will investigate the statistical tools that allow us, in principle, to disentangle the  above different effects, and the observational datasets we require to do so in practice. ",
        "introduction": "Until recently in cosmology, non-gaussianity has been a synonymous of non-linearity; but, in the last 5 years or so, more and more objects like the galaxy of [1] at redshift 5.6 have been found. For the first time a galaxy has been found at higher redshift than the most distant known quasar. More recently, a galaxy at redshift almost 7 has been found [2]. The standard ``inflationary'' cosmological model with gaussian initial conditions predicts that these objects should be very rare. It is becoming increasingly difficult to accommodate the existence of so many high-redshift galaxies under the assumption that non-gaussianity is equivalent to non-linearity, that is postulating gaussian initial conditions. Non-gaussianity does not necessarily imply non-linearity: there might be some primordial non-gaussianity and it is necessary to ``find a way'' to distinguish the two effects. ",
        "conclusions": "We have shown that non-gaussianity does not necessarily mean non-linearity, but it is possible to distinguish different kinds of non-linearity: e.g. bias, gravitational evolution, redshift space distortions. For the physically motivated  non-gaussian models we considered, it turns out that CMB bispectrum is better that LSS bispectrum to detect primordial non-gaussianity: if the future CMB missions will produce maps that are consistent with the gaussian hypothesis, any non-gaussianity seen in the LSS bispectrum can be unambiguously attributed to the effects of non-linearities. Thus, if this is the case, from on-going LSS surveys such as SDSS and 2dF we will be able to know the bias with few \\% accuracy. We have also shown that, to measure the bias parameter, ongoing 3D surveys are much better that 2D ones, even with full sky coverage, but the method developed has applications in different areas such as CMB and gravitational lensing studies. To conclude, we have seen different ways to disentangle primordial non-gaussianity from effects of non-linearity: CMB bispectrum, LSS trispectrum and the abundance of high-redshift objects such as galaxies and clusters. These methods probe the Universe a different scales and at different times and in addition to that they are sensitive to different moments of the distribution. We should therefore conclude that these methods are complementary rather than mutually exclusive.  \\vspace*{1cm}  \\begin{small} {\\bf Acknowledgments}  I would like to thank my collaborators in this work Alan Heavens, Sabino Matarrese, Marc Kamionkowski, Limin Wang and Raul Jimenez.  I also would like to thank the organizers for a very enjoyable workshop. \\end{small}"
    },
    "0004/astro-ph0004406_arXiv.txt": {
        "abstract": "General Relativity predicts that binary systems of stars produce gravitational waves of significant intensity. Here we are particularly interested in the cataclysmic variable binaries (CVs). These systems emit low frequency gravitational waves, $ f < 10^{-3} Hz$. We present here a catalog of CVs and argue that part of them are capable of being detected by the Laser Interferometer Space Antenna (LISA). ",
        "introduction": "Detection of gravitational radiation from astrophysical sources will mark a breakthrough in the history of astronomy (see, e.g., Thorne \\cite{thor87} and Schutz \\cite{schu96}). Experimental efforts to search for these space-time wrinkles have been under development for the past twenty years (Thorne \\cite{thor95,thor96}). With the advent of technological improvements in several crucial aspects of the detection process we will soon be ready to turn them a physical reality (Schutz \\cite{schu96}, Thorne \\cite{thor95}, Finn \\&  Chernoff \\cite{finn93}).  In particular, the Laser Interferometric Space Antenna (LISA) is designed to detect low frequency gravitational waves in the frequency range $10^{-4} - 1$ Hz, which are not possible to detect on the Earth because of seismic noise. There is a lot of very interesting astrophysical phenomena which are believed to generate GWs in the frequency band detectable by LISA, namely: formation of supermassive  black holes (SMBHs), SMBH-SMBH binary coalescence, compact stars orbiting around SMBHs (in, e.g., galactic nuclei), a wide variety of binaries, such as pairs of close white dwarfs (WDs), pairs of neutron stars, neutron star and black hole binaries, pairs of contacting normal stars, normal stars and white dwarfs (cataclysmic) binaries, and pairs of stellar black holes.  Due to the fact the GWs are produced by a large variety of astrophysical sources and cosmological phenomena it is quite probable that the Universe is pervaded by a background of such waves. Binary stars of a variety of stars (ordinary, compact or combinations of them), Population III stars, phase transitions in the early Universe, cosmic strings are examples of sources able to produce a background of GWs.  As the GWs possess a very weak interaction with matter passing through it unharmed, relic radiation (spectral properties for example) once detected can provide information on the physical conditions from the era in which they were produced. In principle it will be possible, for example, to get information from the epoch when the galaxies and stars started to form and evolve.  Concerning our galaxy, it presents a large number of binary systems, which produce a GW background named binary confusion noise (see Hils, Bender \\& Webbink \\cite{hils90}, Bender \\& Hils \\cite{bend97}). Some of the galactic binary sources are: close white dwarfs binaries (CWDBs), neutron star binaries (NSBs), unevolved binaries, WUMs binaries and cataclysmic binaries.  The binary systems are the most understood of all sources of GWs (see, e.g., Thorne 1987). Knowing  the masses of the stars, the orbital parameters and their estimated distances, one can calculate the details of the GW produced.  The LISA's sensitivity as well as the binary confusion noise will determine in the end if one is able to discriminate the signal of a particular astrophysical source.  The first papers concerning the gravitational radiation from binaries systems was written by Mironowskii (1966), who studied in particular the W UMa stars, and by Forward \\& Berman (1967), approximately 30 years ago. After that many other studies concerning the evaluation of GWs background produced by various types of binary stars in the Galaxy followed (see, e.g., Douglass \\& Braginsky 1979, Lipunov \\& Postnov 1987, Lipunov, Postnov \\& Prokhorov 1987, Evans, Iben \\& Smarr 1987, Hils, Bender \\& Webbink 1990, Bender \\& Hils 1997, Webbink \\& Han 1998, Hils 1998)  Here we are particularly interested in the cataclysmic variable binaries as sources of GWs, such a system is formed by a white dwarf and a low mass secondary star. The total number of such a kind of binary is estimated to amount $10^6$ in the Galaxy (see, e.g. Hils, Bender \\& Webbink 1990). These systems produce low frequency GWs, namely, $f_{gw} < 10^{-3}$, which could be detected by LISA.  We are not concerned here with the calculation of a confusion noise produced by such binaries, our aim is similar to the study by Douglass \\& Braginsky (1979) who evaluate the dimensionless amplitude h for a series of specific low frequency GW binaries. Based mainly on the 6th edition of the catalogue of cataclysmic binaries, low mass X-ray binaries and related objects (Ritter \\& Kolb 1998) we have catalogued almost 160 CV systems for which it is possible to evaluate the GW amplitude. We have catalogued firstly those CVs with known distances, orbital period and masses, quantities necessary to evaluate the GW amplitude produced by such objects; secondly we have catalogued those systems for which the distances and the orbital periods are known, the masses being obtained from a mass-period relationships.  The remainder of the paper is as follows: Section 2 deals with the cataclysmic variables. Section 3 addresses the gravitational waves from cataclysmic variables. The discussion and conclusions are summarized in Section 4. ",
        "conclusions": "The CVs produce GWs which could in principle be detected by the LISA antenna, since CVs produce low frequency GWs in the frequency band where LISA is sensitive. Due to the fact that a positive detection of a CV by the LISA antenna might be improved once we know the sources beforehand we compile in the present study a catalogue of CVs, for which we know at least their orbital periods and distances.  We argue that the present study is of interest since in the literature one has not found a systematic identification of possible detectable GW CVs, since an early study made by Douglass \\& Braginsky (1979) twenty years ago, and also a preliminary study by Aguiar et al. (1998). We have been able to catalogue approximately 160 CVs, from which a reasonable part of them could be detected once the LISA antenna become operative.   We argue that it would be of interest whether other groups performed a similar study for the other binary systems which produce low frequency GWs in the frequency band where the LISA antenna is sensitive.  It is worth mentioning that a positive detection of a binary system through its gravitational emission, with some help of electromagnetic data observations, could lead one to know all the parameters related to the binary system, namely, the masses of the stars, their distances to the earth, the period of the system and their orientation angles."
    },
    "0004/astro-ph0004295_arXiv.txt": {
        "abstract": "We have observed the nearly face-on spiral galaxy NGC 5668 with the TAURUS II Fabry-Perot interferometer at the William Herschel Telescope using the $H\\alpha$ line to study the  kinematics of the ionized gas. From the extracted data cube we construct intensity, velocity and velocity dispersion maps. We calculate the rotation curve in the innermost 2 arcmin and we use the residual velocity field to look for regions with important vertical  motions. By comparing the geometry of these regions in the residual velocity field with the geometry in the intensity and velocity dispersion maps we are able to select some regions which are very likely to be shells or chimneys in the disk. The geometry and size of these regions are very similar to the shells or chimneys detected in other galaxies by different means. Moreover, it is worth noting than this galaxy has been reported to have a population of neutral hydrogen high velocity clouds (Schulman et al. 1996) which, according to these observations, could have been originated by chimneys similar to those reported in this paper. ",
        "introduction": "There is now great observational evidence that  disk-halo interactions in galaxies as well as the structure of the interstellar medium (ISM) is closely related to star formation processes in the disks of spiral galaxies (see the review by Dahlem 1997). Big shells develop around the brightest star forming regions, induced by the energy input of supernovae and the strong stellar winds produced by high-mass stars. These shells can grow enough to be able to break the disk, allowing large amounts of gas to  blow out from the disk along these big chimneys (Norman \\& Ikeuchi 1989). The very hot gas going out through these chimneys cools as it rises until it eventually recombines and condenses to form clouds of neutral gas that fall back to the plane (Shapiro \\& Field 1976, Bregman 1980). This {\\em fountain} model then provides an explanation for the origin of high velocity clouds (HVC's) that have been observed in our galaxy and a in few external galaxies (with the galaxy studied in this paper being one of those few (Schulman et al. 1996, hereafter S96)), although alternative explanations have been proposed as well (Blitz et al. 1999). Good reviews on this topic can be found in Wakker \\& van Woerden (1997) and van der Hulst (1996, 1997). On the other side, the expanding shells can induce new star formation (sequential star formation (SSF)) at their edges. This effect has already been observed in some galaxies (see for example Thilker et al. 1998, hereafter T98) and it is also observed to take place in the case of NGC 5668 in our observations.  According to the chimney model, the structure of the ISM, and in particular whether the chimney phenomenon takes place  or not, is controlled by the amount of star formation. The study of the properties of these phenomena in a sample of nearby galaxies can greatly help to understand the structure of the ISM and the nature of disk-halo interactions. Observations of the neutral gas and narrow band imaging of the ionized gas have already been extensively used to study these phenomena (see the review by Dahlem 1997). Scanning long-slit $\\mathrm{H\\alpha}$ spectroscopy has also been used to study these phenomena (Saito et al., 1992; Tomita et al., 1993, 1994). In this work we make a first attempt to use optical Fabry-Perot spectroscopy in a nearly face-on spiral galaxy to directly study these vertical motions. Therefore we have chosen the spiral galaxy NGC 5668 which is already known to have HVC's and an important rate of star formation. These facts make it a perfect candidate for us to detect important vertical motions in its disk.  In Sect. 2 we describe the general properties of the galaxy NGC 5668, with particular emphasis on the observation of the HVC's. Sect. 3 deals with the observations and data reduction (including calculation of intensity, velocity and velocity dispersion maps). Sect. 4 is devoted to the calculation of the rotation model for the galaxy from the observed velocity field. Sect. 5 describes how the residual velocity field is used to look for systematic deviations of circular rotation and how comparison of the geometry of the residual velocity field with that of the intensity and velocity dispersion in some regions can be used to detect real shells and/or chimneys. Finally, in Sect. 6 we report the shell candidates found in NGC 5668 and some of their properties. ",
        "conclusions": "We have carefully analyzed the data obtained for NGC 5668 with the Fabry-Perot interferometer TAURUS II at the WHT in order seek a connection between the star formation processes and vertical motions in spiral galaxies. We have found that there is a clear correlation between the morphology of the regions with a high residual velocity (HRVR's) and the intensity of the $\\mathrm{H\\alpha}$ emission, showing that the HRVR's are indeed regions with important vertical motions associated with star formation processes. Although we are not able to calculate the ages and energetics of these features, comparison with observations in other galaxies strongly supports the hypothesis that the structures detected present a wide age range, from young expanding shells in a bright HII region, to evolved chimneys blowing out hot gas to the halo, surrounded by several bright HII regions. The formation of these regions was probably induced by the pressure exerted by the expanding shell/chimney on the ambient gas. An alternative explanation to these features is that they are produced by infall/collision of gas clouds with the disk (e.g. Saito et al., 1992) followed by induced active star formation. Although this scenario explains in a very natural and simple way the fact that the velocity offsets are one-sided, this is also quite normal in expanding shells. If they are formed not exactly in the equator but slightly off-plane, they will grow up mostly in the low density side, therefore showing only one-sided offsets in the velocity structure. In fact, with the present data there is no way to decide whether the HRVR's are moving into or out of the disk. On the other hand, the different structures found in the HRVR's (from compact HII regions to {\\em rings} of HII regions) are fully compatible with an evolutionary pattern in the chimney model. Therefore, although we can not rule out the infalling hipotheses with the present observations, we find the chimney scenario followed by sequential star formation in the shell borders to be a more likely explanation for what is happening in NGC 5668. These observations provide in any case a clear link between the star formation processes in the disk with other observed phenomena in NGC 5668 like high HVC's of neutral hydrogen. Observations of this kind have been made for other galaxies with a lower star formation rate and without HVC's, and the residual velocity field is not reported to show these features (see for example Jim\\'enez-Vicente et al. 1999, for a similar study of NGC 3938). This fact seems to suggest that these features and the existence of HVC's in the disk are closely related. High resolution HI observations of NGC 5668 would be desirable to confirm the connection between those phenomena."
    },
    "0004/astro-ph0004010_arXiv.txt": {
        "abstract": "Solutions to equilibrium sequences of irrotational binary polytropic stars in Newtonian gravity are expanded in a power of $\\e=a_0/R$, where $R$ and $a_0$ are the orbital separation of the binary system and the radius of each star for $R=\\infty$. For each order of $\\e$, we should solve ordinary differential equations for arbitrary polytropic indices $n$. We show solutions for polytropic indices $n= 0.5, 1, 1.5$ and $2$ up to $\\e^6$ orders. Our semi-analytic solutions can be used to check the validity of numerical solutions. ",
        "introduction": "Coalescing binary neutron stars (BNSs) are considered to be one of the most promising sources of gravitational waves for laser interferometers such as TAMA300\\cite{TAMA300}, GEO600\\cite{GEO600}, VIRGO\\cite{VIRGO} and LIGO\\cite{LIGO,Thorne94}. If we have accurate theoretical templates of inspiraling phase of BNSs, we can determine the mass and the spin of neutron stars from the gravitational wave signals in the inspiraling phase\\cite{Thorne95}. We may also extract the informations on the equation of state of a neutron star from the signals in the premerging phase\\cite{Lindblom92}. For this purpose it is important to complete theoretical templates of gravitational waves in the premerging phase as well as in the inspiraling phase. Moreover, the study of BNSs in the premerging phase is motivated by the need to provide the realistic initial condition for the merging phase simulations\\cite{Shibata99,SU99,ON99}.  In order to obtain accurate theoretical templates of gravitational waves from BNSs around the premerging phase\\cite{Rasio99}, a general relativistic hydrostatic problem with compressible equation of state must be numerically solved. In solving this problem, we suppose that BNSs reach quasi-equilibrium states because the timescale of the orbital decay driven by the radiation reaction is much longer than the orbital period of BNSs until their innermost stable circular orbit (ISCO). (The formalism for solving quasi-equilibrium figures of irrotational binary neutron stars in general relativity is given in the references \\cite{BGM97,Asada98,Shibata98,Teukolsky98}.) Furthermore, we regard the internal state of a neutron star as an irrotational or nearly irrotational one because the viscosity of a neutron star is not large enough for synchronization even near the ISCO\\cite{Kochanek92,BC92}.  Recently Bonazzola, Gourgoulhon and Marck have numerically studied the irrotational binary neutron stars in general relativity with the conformally flat condition\\cite{BGM99a,BGM99b,BGGM99}. More recently, Ury\\=u and Eriguchi have also numerically solved the same problem as that of Bonazzola, Gourgoulhon and Marck\\cite{UE99}. In order to check the validity of these numerical calculations, it is necessary to compare them with analytic or semi-analytic ones. However, we have not obtained such analytic solutions in general relativity yet. There are solutions of irrotational BNSs only in the first post-Newtonian (1PN) approximation of general relativity. For example, Lombardi, Rasio and Shapiro have semi-analytically studied irrotational BNSs by using the energy variational method\\cite{Lombardi97}, and one of the authors of the present paper (KT) has analytically calculated equilibrium sequences of irrotational BNSs by using the tensor virial method\\cite{Taniguchi99}. The former one restricts the internal motion within the plane orthogonal to the rotational axis assuming the shape of the star to an ellipsoidal one in order to treat compressible equation of state. While in the latter one the velocity component along the orbital axis is included at 1PN order and the shape of the star at 1PN order is not restricted to an ellipsoidal one although the fluid is incompressible.  The situation is also the same in Newtonian gravity, that is, even in Newtonian gravity useful analytic or semi-analytic solutions for equilibrium sequences of irrotational binary systems are not obtained. Let us consider numerically constructed stationary structures of irrotational binary stars computed by Ury\\=u and Eriguchi in Newtonian gravity\\cite{UE98a,UE98b}. One may think that semi-analytic solutions by Lai, Rasio and Shapiro\\cite{LRS94} may be used to check the validity of numerical solutions. However, in numerical solutions of Ury\\=u and Eriguchi the velocity component along the orbital axis exists while in those of Lai, Rasio and Shapiro such a component is assumed to be zero from the beginning. Therefore new analytic or semi-analytic solutions are needed to check numerical solutions even in Newtonian gravity. Such a check of numerical solutions is extremely important because in the numerical calculation, there is a possibility to obtain another solution although the binding energy of a binary neutron star is almost the same value, and to lead a different conclusion\\cite{MMW99,MW99}.  In the previous paper\\cite{TN99}, we showed such a new, almost analytic solution to an equilibrium of irrotational binary polytropic stars for the polytropic index $n=1$ in Newtonian gravity by expanding all physical quantities in a power of $\\e \\equiv a_0/R$, where $R$ and $a_0$ are the orbital separation of the binary system and the radius of each star for $R=\\infty$. In that paper, we have extended the method developed by Chandrasekhar more than 65 years ago for corotating fluids\\cite{Ch33,Kovetz68} to the one for irrotational fluids. In this paper, we show semi-analytic solutions for arbitrary polytropic indices by numerically solving ordinary differential equations.  This paper is organized as follows. In \\S 2, we formulate the method to solve the irrotational binary polytropic stars. In \\S 3, the physical values, i.e., the central density of a star, the orbital angular velocity, the total energy and total angular momentum of the binary system are calculated. The numerical results are presented in \\S 4. Section 5 is devoted to summary and discussions.  Although a binary system consists of two stars, we pay particular attention to one of two stars. We call it star 1 whose mass is $M_1$ and the companion one star 2 whose mass is $M_2$. In this paper, we adopt two corotating coordinate systems. First one is a Cartesian coordinate system $\\bX$ whose origin is located at the center of mass of the binary system. For calculational convenience, we choose the orbital axis as $X_3$, and we take the direction of $X_1$ from the center of mass of star 2 to that of star 1. The second coordinate system is the spherical one $\\br =(r,\\th,\\vp)$ whose origin is located at the center of mass of star 1. We use units of $G=1$. ",
        "conclusions": "\\subsection{Summary}  In this paper, we have calculated the equilibrium solutions of irrotational binary polytropic stars in Newtonian gravity by expanding all physical quantities in a power of $\\e$. We have presented the results of the cases of several polytropic indices ($n=0.5, 1, 1.5$, and $2$). In particular, we have shown the velocity fields by solving the equation of continuity. It is found that there exists the small velocity component along the orbital axis (see \\S \\ref{sec:results}). It agrees with the numerical calculations performed by Ury\\=u and Eriguchi\\cite{UE98b} and Bonazzola, Gourgoulhon and Marck\\cite{BGM99a,BGM99b,BGM99c}.  Furthermore, we have given the figures and tables of the total energy, total angular momentum and orbital angular velocity along the equilibrium sequences for each polytropic indices. We can see from these figures that our results agree with those of Lai, Rasio and Shapiro\\cite{LRS94} and Ury\\=u and Eriguchi\\cite{UE98b} for $R/a_0 >3$.  Since our solutions are correct if $\\e \\ll 1$, they can be used to check the validity of numerical solutions. For any numerical codes, one can ask to solve an equilibrium for large $R$, and compare numerically derived velocity distribution and so on with our semi-analytic solutions.  However, since we expanded physical quantities up to $O(\\e^6)$, it may not be enough to discuss about the behavior of solutions for small $R$. In order to apply our solutions in the case of small $R$ and to check the validity of numerical codes in this case, further higher order calculations are needed.   \\subsection{Discussions}  It is important to compare the velocity field which we obtain by solving the continuity equation with that given by Lai, Rasio and Shapiro. We can see the velocity field they give in their papers\\cite{LRS94,LRS93} or in the Chandrasekhar's textbook\\cite{Ch69}. In the irrotational case, it becomes \\beqa (u_{LRS})_1 &=&{2a_1^2 \\over a_1^2 +a_2^2} \\O x_2, \\\\ (u_{LRS})_2 &=&-{2a_2^2 \\over a_1^2 +a_2^2} \\O x_1, \\\\ (u_{LRS})_3 &=&0, \\eeqa in the corotating frame. Here $a_i$ denotes the length of the principal axis parallel to the $x_i$-axis. We can rewrite these component of the velocity field as \\beqa (u_{LRS})_1 &=&\\O x_2 +\\Bigl( {a_1^2 -a_2^2 \\over a_1^2 +a_2^2} \\Bigr) \\O x_2, \\\\ (u_{LRS})_2 &=&-\\O x_1 +\\Bigl( {a_1^2 -a_2^2 \\over a_1^2 +a_2^2} \\Bigr) \\O x_1. \\eeqa The second terms in the above equations are order $\\e^3$ because the deviation between $a_1$ and $a_2$ is produced by the tidal force and the effect of the tidal force is order $\\e^3$.  On the other hand, if we restrict the form of the function in our velocity field as $\\four \\phi_2 =(b_2/\\xi_1) \\xi^2$, $\\five \\phi_3 =(\\a b_3/\\xi_1) \\xi^3$ and $\\six \\phi_4 =(\\a^2 b_4/\\xi_1) \\xi^4$ as in the case of the incompressible fluid, we can express our velocity field by using Eqs. (\\ref{Eq:velo_corot_1}), (\\ref{Eq:velo_corot_2}) and (\\ref{Eq:velo_corot_3}) as \\beqa (u_{present})_1 &=&\\O x_2 +\\O \\bigl[ 6\\e^3 b_2 x_2 -48\\e^4 b_3 x_1 x_2 +90 \\e^5 b_4 x_2 (r^2 -5x_1^2) \\bigr], \\\\ (u_{present})_2 &=&-\\O x_1 +\\O \\bigl[ 6\\e^3 b_2 x_1 +6\\e^4 b_3 (r^2 -5x_1^2 +2x_2^2) +30\\e^5 b_4 r(3r^2 +7x_1^2 -6x_2^2) \\bigr], \\\\ (u_{present})_3 &=&\\O \\bigl[ 12\\e^4 b_3 x_2 x_3 +180\\e^5 b_4 x_1 x_2 x_3 \\bigr], \\eeqa where $r^2=x_1^2 +x_2^2 +x_3^2$. Therefore, we find that the form of the velocity field we obtain in the case of the incompressible fluid coincides with that given by Lai, Rasio and Shapiro up to $O(\\e^3)$ including the value of $b_2$. This means that the velocity field of Lai, Rasio and Shapiro is correct only in the case of the ``incompressible'' equation of state and ``ellipsoidal'' figures, because the velocity field at order $\\e^3$ is produced by the ellipsoidal deformation of star 1 (see Appendix \\ref{order:4th}).  Finally, we discuss about the configuration of each star. When we pay attention to star 1, the equation for the stellar surface is written as \\beqa \\Xi (\\th, \\vp) &=&\\xi_1 +\\e^3 S_3(\\th, \\vp) +\\e^4 S_4(\\th, \\vp) +\\e^5 S_5(\\th, \\vp) +\\e^6 S_6(\\th, \\vp), \\\\ &=&\\xi_1 +\\e^3 {\\three \\psi_2(\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} P_2 (\\sin \\th \\cos \\vp) +\\e^4 {\\four \\psi_3 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} P_3 (\\sin \\th \\cos \\vp) +\\e^5 {\\five \\psi_4(\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} P_4 (\\sin \\th \\cos \\vp) \\nonumber \\\\ &&+\\e^6 {1 \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) \\Bigl\\{ {18 \\over 35} P_4 (\\sin \\th \\cos \\vp) +{2 \\over 7} P_2 (\\sin \\th \\cos \\vp) +{1 \\over 5} \\Bigr\\} \\nonumber \\\\ &&\\hspace{50pt}+\\six \\psi_0 (\\xi_1) +\\six \\psi_2 (\\xi_1) P_2 (\\sin \\th \\cos \\vp) +\\six \\psi_{22} (\\xi_1) P_2^2 (\\cos \\th) \\cos 2\\vp \\nonumber \\\\ &&\\hspace{50pt}+\\six \\psi_4 (\\xi_1) P_4 (\\sin \\th \\cos \\vp) +\\six \\psi_5 (\\xi_1) P_5 (\\sin \\th \\cos \\vp) \\Bigr], \\label{Eq:stellar_surface} \\eeqa where $S_i$ are defined in Appendix \\ref{Ap:1} and we have used the relation (\\ref{Eq:ap_relation}). Accordingly, we can express the length of the principal axis. Although the real length of the axis is written as $\\a \\Xi$, we show the results divided by $\\a$.  \\vspace{0.3cm} \\noindent (1) The opposite direction to star 2: \\beqa \\Xi \\Bigl( {\\pi \\over 2}, 0 \\Bigr) &=&\\xi_1 +\\e^3 {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} +\\e^4 {\\four \\psi_3 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} +\\e^5 {\\five \\psi_4 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\nonumber \\\\ &&+\\e^6 {1 \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +\\six \\psi_0 (\\xi_1) +\\six \\psi_2 (\\xi_1) +3 \\six \\psi_{22} (\\xi_1) \\nonumber \\\\ &&\\hspace{40pt}+\\six \\psi_4 (\\xi_1) +\\six \\psi_5 (\\xi_1) \\Bigr], \\eeqa  \\noindent (2) The direction to star 2: \\beqa \\Xi \\Bigl( {\\pi \\over 2}, \\pi \\Bigr) &=&\\xi_1 +\\e^3 {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} -\\e^4 {\\four \\psi_3 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} +\\e^5 {\\five \\psi_4 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\nonumber \\\\ &&+\\e^6 {1 \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +\\six \\psi_0 (\\xi_1) +\\six \\psi_2 (\\xi_1) +3 \\six \\psi_{22} (\\xi_1) \\nonumber \\\\ &&\\hspace{40pt}+\\six \\psi_4 (\\xi_1) -\\six \\psi_5 (\\xi_1) \\Bigr], \\eeqa  \\noindent (3) The (opposite) direction to the orbital motion: \\beqa \\Xi \\Bigl( {\\pi \\over 2}, {\\pi \\over 2} \\Bigr) &=&\\Xi \\Bigl( {\\pi \\over 2}, {3\\pi \\over 2} \\Bigr), \\nonumber \\\\ &=&\\xi_1 -\\e^3 {\\three \\psi_2 (\\xi_1) \\over 2|\\Th_{0,\\xi} (\\xi_1)|} +\\e^5 {3\\five \\psi_4 (\\xi_1) \\over 8|\\Th_{0,\\xi} (\\xi_1)|} \\nonumber \\\\ &&+\\e^6 {1 \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ {\\three \\psi_2 (\\xi_1) \\over 4|\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +\\six \\psi_0 (\\xi_1) -{1 \\over 2}\\six \\psi_2 (\\xi_1) -3 \\six \\psi_{22} (\\xi_1) \\nonumber \\\\ &&\\hspace{40pt}+{3 \\over 8}\\six \\psi_4 (\\xi_1) \\Bigr], \\eeqa  \\noindent (4) The direction parallel to the rotational axis: \\beqa \\Xi(0, 0) &=&\\Xi(-\\pi, 0), \\nonumber \\\\ &=&\\xi_1 -\\e^3 {\\three \\psi_2 (\\xi_1) \\over 2|\\Th_{0,\\xi} (\\xi_1)|} +\\e^5 {3\\five \\psi_4 (\\xi_1) \\over 8|\\Th_{0,\\xi} (\\xi_1)|} \\nonumber \\\\ &&+\\e^6 {1 \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ {\\three \\psi_2 (\\xi_1) \\over 4|\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +\\six \\psi_0 (\\xi_1) -{1 \\over 2}\\six \\psi_2 (\\xi_1) +{3 \\over 8}\\six \\psi_4 (\\xi_1) \\Bigr]. \\eeqa We can see from these equations and Tables \\ref{table1} -- \\ref{table8} that the axis to star 2 is the longest, and the deviation between the axis to star 2 and that opposite to star 2 appears at order $\\e^4$. On the contrary, the deviation between the axis to the orbital motion and that parallel to the rotational axis appears at order $\\e^6$, and the difference is the effect of the deformation induced by the spin of the figure ($\\six \\psi_{22}$).  When we see the quadrupole moments in Eq. (\\ref{Eq:stellar_surface}), we find that the coeffient of the higher order term is not large as \\beq \\e^3 {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ 1 +\\e^3 \\Bigl\\{ {2 \\over 7|\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +{\\six \\psi_2 (\\xi_1) \\over \\three \\psi_2 (\\xi_1)} \\Bigr\\} \\Bigr] P_2 (\\sin \\th \\cos \\vp). \\eeq Therefore we can expect the convergence of these terms. However, there is another quadrupole moment induced by the spin of the figure. The term $\\e^6 (\\six \\psi_{22} (\\xi_1)/|\\Th_{0,\\xi} (\\xi_1)|) P_2^2 (\\cos \\th) \\cos 2\\vp$ seems to be as effective as the leading quadrupole term for the smaller orbital separation. This means that the terms concerned with the spin of the figure which appear at order $\\e^9$ in the total energy may change the behavior of the total energy. Furthermore, the hexadecapole moments in Eq. (\\ref{Eq:stellar_surface}), \\beq \\e^5 {\\five \\psi_4 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl[ 1 +\\e \\Bigl\\{ {18 \\over 35\\five \\psi_4 (\\xi_1)} {\\three \\psi_2 (\\xi_1) \\over |\\Th_{0,\\xi} (\\xi_1)|} \\Bigl( {\\three \\psi_2 (\\xi_1) \\over \\xi_1} +{d\\three \\psi_2 \\over d\\xi} (\\xi_1) \\Bigr) +{\\six \\psi_4 (\\xi_1) \\over \\five \\psi_4 (\\xi_1)} \\Bigr\\} \\Bigr] P_4 (\\sin \\th \\cos \\vp), \\eeq does not seem to converge at order $\\e^6$. However, since these terms will appear at order $\\e^{10}$ in the total energy, they do not have so much effect. Anyway, if we discuss about the behavior of the total energy, the total angular momentum and so on for the smaller orbital separation such as $R/a_0 <3$, we must calculate at least up to order $\\e^9$.    \\label{Ap:1}  In the following appendices, we derive equations order by order of $\\e$, and describe in detail. In this appendix, we summarize the equations which should be solved numerically with their boundary conditions in each order.  First of all, we give the total equations which include all terms up to $O(\\e^6)$. The equations for determination of the velocity potential and stellar configuration are written as \\beqa \\t{\\D} \\t{\\Phi} &=& -n \\Bigl[ \\t{\\n} \\t{\\Phi} -(\\t{\\bf \\O} \\times \\bxi)_{orb} -{\\e \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\Bigr] \\cdot {\\t{\\n} \\Th \\over \\Th}, \\\\ \\t{\\D} \\Th &=&-\\Th^n -{\\O^2 \\xi_1^2 \\over 8\\pi \\r_c \\e^2} \\t{\\D} \\Bigl[ (\\t{\\n} \\t{\\Phi})^2 -2(\\t{\\n} \\t{\\Phi}) \\cdot \\Bigl\\{ (\\t{\\bf \\O} \\times \\bxi)_{orb} +{\\e \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\Bigr\\} \\Bigr], \\eeqa where \\beqa \\t{\\D} &=&\\a^2 \\D ={1 \\over \\xi^2} {\\p \\over \\p \\xi} \\Bigl( \\xi^2 {\\p \\over \\p \\xi} \\Bigr) +{1 \\over \\xi^2 \\sin \\th} {\\p \\over \\p \\th} \\Bigl( \\sin \\th {\\p \\over \\p \\th} \\Bigr) +{1 \\over \\xi^2 \\sin^2 \\th} {\\p^2 \\over \\p \\vp^2}, \\\\ \\t{\\n} &=&\\a \\n ={\\p \\over \\p \\xi} \\hat{\\bxi} +{1 \\over \\xi} {\\p \\over \\p \\th} \\hat{\\bth} +{1 \\over \\xi \\sin \\th} {\\p \\over \\p \\vp} \\hat{\\bvp}. \\label{Eq:nomnabla} \\eeqa  The boundary condition for the velocity field is \\beq \\Bigl[ \\t{\\n} \\t{\\Phi} -(\\t{\\bf \\O} \\times \\bxi)_{orb} -{\\e \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\Bigr] \\cdot (\\t{\\n} \\Th) \\Bigl|_{surf} =0. \\eeq  The internal and external gravitational potentials which should be matched at the stellar surface are \\beqa \\t{U}_{int} &\\equiv& {U^{1 \\ra 1} \\over 4\\pi \\r_c \\a^2} \\nonumber \\\\ &=&\\Th +{\\O^2 \\xi_1^2 \\over 8\\pi \\r_c \\e^2} \\Bigl[ (\\t{\\n} \\t{\\Phi})^2 -2(\\t{\\n} \\t{\\Phi}) \\cdot \\Bigl\\{ (\\t{\\bf \\O} \\times \\bxi)_{orb} +{\\e \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\Bigr\\} \\Bigr] +\\t{U}_0 \\nonumber \\\\ &&-{\\mu \\over 1+p} \\e \\sum_{l=0}^{\\infty} (-1)^l \\e^l \\Bigl( {\\xi \\over \\xi_1} \\Bigr)^l P_l (\\sin \\th \\cos \\vp) -{3\\mu \\over 2M_{tot}} \\Bigl( {\\bI_{11}' \\over a_0^2} \\Bigr) \\e^3 \\Bigl[ 1 -3\\e {\\xi \\over \\xi_1} P_1 (\\sin \\th \\cos \\vp)+ \\cdots \\Bigr], \\\\ \\t{U}_{ext} &=& {\\k_0 \\over \\xi} +\\e \\sum_{l,m} {\\k_{1:lm} \\over \\xi^{l+1}} Y_l^m +\\e^2 \\sum_{l,m} {\\k_{2:lm} \\over \\xi^{l+1}} Y_l^m +\\e^3 \\sum_{l,m} {\\k_{3:lm} \\over \\xi^{l+1}} Y_l^m +\\e^4 \\sum_{l,m} {\\k_{4:lm} \\over \\xi^{l+1}} Y_l^m +\\e^5 \\sum_{l,m} {\\k_{5:lm} \\over \\xi^{l+1}} Y_l^m \\nonumber \\\\ &&+\\e^6 \\sum_{l,m} {\\k_{6:lm} \\over \\xi^{l+1}} Y_l^m + \\cdots, \\eeqa where $\\k_0$ and $\\k_{i:lm}$ are multipole moments. $\\t{U}_0$ is constant and expanded as \\beq \\t{U}_0 =c_0 +\\e c_1 +\\e^2 c_2 +\\e^3 c_3 +\\e^4 c_4 +\\e^5 c_5 +\\e^6 c_6. \\eeq The velocity potential and configuration function are expanded up to $O(\\e^6)$ as \\beqa \\t{\\Phi} &=&\\t{\\Phi}_0 +\\e \\t{\\Phi}_1 +\\e^2 \\t{\\Phi}_2 +\\e^3 \\t{\\Phi}_3 +\\e^4 \\t{\\Phi}_4 +\\e^5 \\t{\\Phi}_5 +\\e^6 \\t{\\Phi}_6, \\\\ \\Th &=& \\Th_0 +\\e \\Th_1 +\\e^2 \\Th_2 +\\e^3 \\Th_3 +\\e^4 \\Th_4 +\\e^5 \\Th_5 +\\e^6 \\Th_6. \\eeqa  The boundary conditions for $\\Th$ are \\beqa \\t{U}_{int}(\\Xi) &=&\\t{U}_{ext} (\\Xi), \\label{Eq:bc1} \\\\ {\\p \\t{U}_{int} \\over \\p \\xi} (\\Xi) &=&{\\p \\t{U}_{ext} \\over \\p \\xi} (\\Xi), \\label{Eq:bc2} \\eeqa where $\\Xi$ denotes the first zero point of the function $\\Th$, i.e., the stellar surface and is formally expressed as \\beq \\Xi(\\th,\\vp) =\\xi_1 +\\e S_1(\\th,\\vp) +\\e^2 S_2(\\th,\\vp) +\\e^3 S_3(\\th,\\vp) +\\e^4 S_4(\\th,\\vp) +\\e^5 S_5(\\th,\\vp) +\\e^6 S_6(\\th,\\vp). \\eeq The regularity conditions at the center of each star is \\beq {\\p \\Th \\over \\p \\xi} (\\xi=0) =0, \\eeq and also we normalize the central value of $\\Th$ as \\beq \\Th(\\xi=0) =1. \\eeq  In the following, we show the equation for the velocity potential, its boundary condition, the equation for determination of the figure and its boundary conditions (\\ref{Eq:bc1}) and (\\ref{Eq:bc2}).  \\subsection{0th order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_0 =-n \\Bigl[ \\t{\\n} \\t{\\Phi}_0 -(\\t{\\bf \\O} \\times \\bxi)_{orb} \\Bigr] \\cdot {\\t{\\n} \\Th_0 \\over \\Th_0}, \\eeq and its boundary condition is \\beq \\Bigl[ \\t{\\n} \\t{\\Phi}_0 -(\\t{\\bf \\O} \\times \\bxi)_{orb} \\Bigr] \\cdot (\\t{\\n} \\Th_0) \\Bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq {1 \\over \\xi^2} {d \\over d\\xi} \\Bigl( \\xi^2 {d\\Th_0 \\over d\\xi} \\Bigr) =-\\Th_0^n, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_0 ={\\k_0 \\over \\xi_1}, \\\\ &&{d \\Th_0 \\over d\\xi} (\\xi_1) =-{\\k_0 \\over \\xi_1^2}. \\eeqa The regularity condition at the center of the star is \\beq {d \\Th_0 \\over d\\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_0$ at the center of the star is \\beq \\Th_0 (\\xi=0) =1. \\eeq   \\subsection{1st order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_1 =-n \\Bigl[ \\t{\\n} \\t{\\Phi}_1 -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\Bigr] \\cdot {\\t{\\n} \\Th_0 \\over \\Th_0}, \\\\ \\eeq and its boundary condition is \\beq (\\t{\\n} \\t{\\Phi}_1) \\cdot (\\t{\\n} \\Th_0) \\bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_1 =-n\\Th_0^{n-1} \\Th_1, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_1 -{\\mu (3+2p) \\over 2(1+p)^2} =-{\\k_0 \\over \\xi_1^2} S_1 +\\sum_{l,m} {\\k_{1:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_1 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{\\p \\Th_1 \\over \\p \\xi}(\\xi_1) ={2\\k_0 \\over \\xi_1^3} S_1 -\\sum_{l,m} (l+1) {\\k_{1:lm} \\over \\xi_1^{l+2}} Y_l^m. \\eeqa The regularity condition at the center of the star is \\beq {\\p \\Th_1 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_1$ at the center of the star is \\beq \\Th_1 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$.   \\subsection{2nd order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_2 =-n (\\t{\\n} \\t{\\Phi}_2) \\cdot {\\t{\\n} \\Th_0 \\over \\Th_0}, \\eeq and its boundary condition is \\beq (\\t{\\n} \\t{\\Phi}_2) \\cdot (\\t{\\n} \\Th_0) \\bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_2 =-n \\Th_0^{n-1} \\Th_2, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_2 =-{\\k_0 \\over \\xi_1^2} S_2 +\\sum_{l,m} {\\k_{2:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_2 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{\\p \\Th_2 \\over \\p \\xi}(\\xi_1) ={2\\k_0 \\over \\xi_1^3} S_2 -\\sum_{l,m} (l+1) {\\k_{2:lm} \\over \\xi_1^{l+2}} Y_l^m. \\eeqa The regularity condition at the center of the star is \\beq {\\p \\Th_2 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_2$ at the center of the star is \\beq \\Th_2 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$.   \\subsection{3rd order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_3 =-n (\\t{\\n} \\t{\\Phi}_3) \\cdot {\\t{\\n} \\Th_0 \\over \\Th_0}, \\eeq and its boundary condition is \\beq (\\t{\\n} \\t{\\Phi}_3) \\cdot (\\t{\\n} \\Th_0) \\bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_3 =-n\\Th_0^{n-1} \\Th_3, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_3 -{\\mu \\over 1+p} P_2(\\sin \\th \\cos \\vp) =-{\\k_0 \\over \\xi_1^2} S_3 +\\sum_{l,m} {\\k_{3:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_3 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{\\p \\Th_3 \\over \\p \\xi}(\\xi_1) -{2\\mu \\over \\xi_1 (1+p)} P_2(\\sin \\th \\cos \\vp) ={2\\k_0 \\over \\xi_1^3} S_3 -\\sum_{l,m} (l+1) {\\k_{3:lm} \\over \\xi_1^{l+2}} Y_l^m. \\eeqa The regularity condition at the center of the star is \\beq {\\p \\Th_3 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_3$ at the center of the star is \\beq \\Th_3 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$.   \\subsection{4th order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_4 =-{n \\over \\Th_0} \\Bigl[ (\\t{\\n} \\t{\\Phi}_4) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_3) \\Bigr], \\eeq and its boundary condition is \\beq \\Bigl[ (\\t{\\n} \\t{\\Phi}_4) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_3) \\Bigr] \\Bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_4 =-n\\Th_0^{n-1} \\Th_4, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_4 +{\\mu \\over 1+p} P_3(\\sin \\th \\cos \\vp) =-{\\k_0 \\over \\xi_1^2} S_4 +\\sum_{l,m} {\\k_{4:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_4 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{\\p \\Th_4 \\over \\p \\xi}(\\xi_1) +{3\\mu \\over \\xi_1(1+p)} P_3(\\sin \\th \\cos \\vp) ={2\\k_0 \\over \\xi_1^3} S_4 -\\sum_{l,m} (l+1) {\\k_{4:lm} \\over \\xi_1^{l+2}} Y_l^m. \\eeqa The regularity condition at the center of the star is \\beq {\\p \\Th_4 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_4$ at the center of the star is \\beq \\Th_4 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$.   \\subsection{5th order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_5 =-{n \\over \\Th_0} \\Bigl[ (\\t{\\n} \\t{\\Phi}_5) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_4) \\Bigr], \\eeq and its boundary condition is \\beq \\Bigl[ (\\t{\\n} \\t{\\Phi}_5) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_4) \\Bigr] \\Bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_5 =-n\\Th_0^{n-1} \\Th_5, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_5 -{\\mu \\over 1+p} P_4(\\sin \\th \\cos \\vp) =-{\\k_0 \\over \\xi_1^2} S_5 +\\sum_{l,m} {\\k_{5:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_5 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{\\p \\Th_5 \\over \\p \\xi}(\\xi_1) -{4\\mu \\over \\xi_1(1+p)} P_4(\\sin \\th \\cos \\vp) ={2\\k_0 \\over \\xi_1^3} S_5 -\\sum_{l,m} (l+1) {\\k_{5:lm} \\over \\xi_1^{l+2}} Y_l^m. \\eeqa The regularity condition at the center of the star is \\beq {\\p \\Th_5 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_5$ at the center of the star is \\beq \\Th_5 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$.   \\subsection{6th order} The equation for the velocity potential is \\beq \\t{\\D} \\t{\\Phi}_6 =-{n \\over \\Th_0} \\Bigl[ (\\t{\\n} \\t{\\Phi}_6) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_5) \\Bigr], \\eeq and its boundary condition is \\beq \\Bigl[ (\\t{\\n} \\t{\\Phi}_6) \\cdot (\\t{\\n} \\Th_0) -{1 \\over \\xi_1} (\\t{\\bf \\O} \\times \\bxi)_{fig} \\cdot (\\t{\\n} \\Th_5) \\Bigr] \\Bigl|_{\\xi_1} =0. \\eeq The equation for determination of the figure is \\beq \\t{\\D} \\Th_6 =-n\\Th_0^{n-1} \\Th_6 -{1 \\over 2} n(n-1) \\Th_0^{n-2} \\Th_3^2 +{\\mu \\over \\xi_1} \\t{\\D} {\\p \\t{\\Phi}_4 \\over \\p \\vp}, \\eeq and its boundary conditions at the stellar surface are \\beqa &&c_6 -{3\\mu \\b{\\bI}_{11}' \\over 2M_{tot} a_0^2} -{\\mu \\d \\over 2(1+p)^2} -{2\\mu \\over \\xi_1 (1+p)} S_3 P_2(\\sin \\th \\cos \\vp) +{\\mu \\over 1+p} P_5(\\sin \\th \\cos \\vp) -{\\mu \\over \\xi_1} {\\p \\t{\\Phi}_4 \\over \\p \\vp}(\\xi_1) \\nonumber \\\\ &&\\hspace{30pt}={\\k_0 \\over \\xi_1^2} \\Bigl( {S_3^2 \\over \\xi_1} -S_6 \\Bigr) -\\sum_{l,m} (l+1) {\\k_{3:lm} \\over \\xi_1^{l+2}} S_3 Y_l^m +\\sum_{l,m} {\\k_{6:lm} \\over \\xi_1^{l+1}} Y_l^m, \\\\ &&S_6 {d^2 \\Th_0 \\over d\\xi^2}(\\xi_1) +{1 \\over 2} S_3^2 {d^3 \\Th_0 \\over d\\xi^3}(\\xi_1) +S_3 {\\p^2 \\Th_3 \\over \\p \\xi^2} (\\xi_1) +{\\p \\Th_6 \\over \\p \\xi}(\\xi_1) -{2\\mu \\over \\xi_1^2 (1+p)} S_3 P_2(\\sin \\th \\cos \\vp) +{5\\mu \\over \\xi_1 (1+p)} P_5(\\sin \\th \\cos \\vp) \\nonumber \\\\ &&\\hspace{30pt}-{\\mu \\over \\xi_1} {\\p^2 \\t{\\Phi}_4 \\over \\p \\xi \\p \\vp} (\\xi_1) =-{\\k_0 \\over \\xi_1^3} \\Bigl( {3S_3^2 \\over \\xi_1} -2S_6 \\Bigr) +\\sum_{l=0}^{\\infty} (l+1)(l+2) {\\k_{3:lm} \\over \\xi_1^{l+3}} S_3 Y_l^m -\\sum_{l=0}^{\\infty} (l+1) {\\k_{6:lm} \\over \\xi_1^{l+2}} Y_l^m, \\eeqa where \\beq \\d \\equiv {9 \\over 2a_0^2} \\Bigl[ {\\b{\\bI}_{11} \\over M_1} +\\Bigl( {a_0' \\over a_0} \\Bigr)^3 {\\b{\\bI}_{11}' \\over M_2} \\Bigr]. \\eeq The regularity condition at the center of the star is \\beq {\\p \\Th_6 \\over \\p \\xi} (\\xi=0) =0, \\eeq and the normalization of $\\Th_6$ at the center of the star is \\beq \\Th_6 (\\xi=0) =0, \\eeq because we take $\\Th_0 (\\xi=0)=1$."
    },
    "0004/astro-ph0004226_arXiv.txt": {
        "abstract": "A number of authors have reported filaments connecting bright structures in high-resolution X-ray images, and in some cases these have been taken as evidence for a physical connection between the structures, which might be thought to provide support for a model with non-cosmological redshifts. In this paper I point out two problems which are inherent in the interpretation of smoothed photon-limited data of this kind, and develop some simple techniques for the assessment of the reality of X-ray filaments, which can be applied to either simply smoothed or adaptively smoothed data. To illustrate the usefulness of these techniques, I apply them to archival {\\it ROSAT} observations of galaxies and quasars previously analysed by others. I show that several reported filamentary structures connecting X-ray sources are not in fact significantly detected. ",
        "introduction": "With the advent of high-resolution X-ray telescopes it is now routine to see structure in X-ray images. Assessing the level at which one should believe this structure presents more of a challenge in X-ray observations than in optical or radio images of comparable resolution, because X-ray data are very often photon-noise limited. The problem is made worse by the common (and necessary) practice of smoothing the data with a Gaussian. This is carried out in order to make the images conform to our expectations of what the `true' sky image with infinite exposure time would look like; but it is important to remember that in general it does {\\it not} in fact take away the photon-limited nature of the underlying data.  A number of smoothed X-ray images of extragalactic objects appear, on published contour maps, to show quite clear `filamentary' connections between X-ray sources in the field. (The term `filament' does not have a single definition in the literature; I shall use it to mean any extended, apparently linear connection between two sources.) In this paper I discuss the techniques necessary to determine whether such filamentary connections are real, and apply them to some sample observations of galaxies and quasars, previously analysed by other authors (Arp 1996, Dahlem \\etal\\ 1996, Arp 1997), which I have taken from the {\\it ROSAT} archives. ",
        "conclusions": "Using Monte Carlo techniques to find the distribution of pixel values in smoothed noise fields and to assess the reality of connections between adjacent sources, I have found that the filamentary X-ray connections between low- and high-redshift Seyfert galaxies and quasars reported by Arp (1996) are not statistically significant. I have confirmed that the X-ray extension of two low-redshift galaxies located near high-redshift radio-loud quasars, reported in the same paper, is significant. The filament extending south from NGC 3628, discussed by Dahlem \\etal\\ (1996), was shown to be marginally significant. The connection between NGC 4151 and a BL Lac object reported by Arp (1997) is also apparently not statistically significant.  These results suggest that it is necessary to be cautious in the interpretation of smoothed X-ray images, and I urge authors to apply the techniques I have described to any situation where the detection or non-detection of faint extended X-ray emission has important scientific implications."
    },
    "0004/astro-ph0004156_arXiv.txt": {
        "abstract": "We investigate some aspects of quintessence models with a non-minimally coupled scalar field and in particular we show that it can behave as a component of matter with $-3 \\lesssim P/\\rho \\lesssim 0$. We study the properties of gravitational waves in this class of models and discuss their energy spectrum and the cosmic microwave background anisotropies they induce. We also show that gravitational waves are damped by the anisotropic stress of the radiation and that their energy spectrum may help to distinguish between inverse power law potential and supergravity motivated potential. We finish by a discussion on the constraints arising from their density parameter $\\Omega_\\GW$. ",
        "introduction": "\\label{par1}  Recent astrophysical and cosmological observations such as the luminosity distance-redshift relation for the supernovae type~Ia~\\cite{sn1a_1,sn1a_2,sn1a_3}, the recent observations of the cosmic microwave background temperature anisotropies~\\cite{boomerang}, gravitional lensing~\\cite{lens} and velocity fields~\\cite{roman} tend to indicate that a large fraction of the matter of the universe today is composed of matter with negative pressure (see \\EG~\\cite{indication} for a comparison of the different observations). Recent analyses~\\cite{efsta,perlmutter99} seem to indicate that the energy density $\\rho$ and the pressure $P$ of this fluid satisfies \\begin{equation} -1\\leq P/\\rho \\leq -0.6, \\end{equation} which is compatible with a cosmological constant $\\Lambda$ for which $P/\\rho=-1$ (see also~\\cite{phantom} for arguments in favour of $P/\\rho<-1$). A typical value of $\\Omega_\\Lambda\\simeq 0.7$ for its energy density in units of the critical density of the universe corresponds to an energy scale of order $5\\times10^{-47}\\,\\hbox{GeV}^4$ which is very far from what is expected from high energy physics; this is the well known cosmological constant problem~\\cite{weinberg}. To circumvent this problem different solutions have been proposed starting from the idea of a dynamical cosmological constant~\\cite{coble} to lead to the class of models known as {\\it quintessence}~\\cite{caldwell}, where a spatially homogeneous scalar field $\\phi$ is rolling down a potential decreasing when $\\phi$ tends to infinity. An example of such a potential which has been widely studied is the inverse power law potential. It can be obtained from some high energy physics models, \\EG{} where supersymmetry is broken through fermion condensates~\\cite{binetruy}. Recently, it has been argued~\\cite{brax} that supergravity has to be taken into account since today one expects the scalar field to be of order of the Planck mass $M_\\PL$ and corrections to the potential appear at this energy. This leads to a better agreement with observations~\\cite{brax2}.  An important point about this family of models is the existence of scaling solutions~\\cite{scaling1,scaling2} (refered to as {\\it tracking solutions}), \\IE{} such that $\\phi$ evolves as the scale factor of the universe at a given power. These solutions are attractors of the dynamical system describing the evolution of the scale factor and of the scalar field. This implies that the present time behaviour of the field is almost independent of its initial conditions~\\cite{steinhard,steinhard2}. This property allows to address~\\cite{zlatev} (i) the {\\it coincidence problem}, \\IE{} the fact that $\\phi$ starts to dominate today and (ii) the {\\it fine tuning problem}, \\IE{} the fact that one does not have to fine tune the initial condition of the field $\\phi$.  One of us extended these models to include a non-minimal coupling $\\xi\\bar Rf(\\phi)$ between the scalar field and the scalar curvature $\\bar R$~\\cite{uzan99}. Such a coupling term appears \\EG~when quantising fields in curved spacetime~\\cite{birrel,ford87} and in multi-dimensional theories~\\cite{maeda,acceta,mata}. It was shown that when $f(\\phi)=\\phi^2/2$ tracking solutions still exist~\\cite{uzan99} and this result was generalised~\\cite{amendola} to any coupling function $f$ and potential $V$ satisfying $V(\\phi)\\propto f^n(\\phi)$. However, such a coupling is constrained by the variations of the constants of nature~\\cite{caroll} which fix bounds on $\\xi$~\\cite{chiba99}. A way to circumvent this problem is to consider quintessence models in the framework of scalar-tensor theories~\\cite{bartolo,bertolami} where a double attractor mechanism can occur, \\IE{} of the scalar-tensor theory towards general relativity and of the scalar field $\\phi$ towards its tracking solution.   Among all the possible observations of cosmology, gravitational waves give an insight on epochs where there was a variation of the background dynamics since every such variation affects the shape of the stochastic graviton background spectrum~\\cite{gw1,gw2}. We can then view our universe as containing a sea of stochastic gravitational waves from primordial origin, as predicted by most models of structure formations such as inflation~\\cite{gw1,gw2} (see also~\\cite{allen94} for a review) and topological defects scenarios~\\cite{vs}. Their spectrum extends typically from $10^{-18}$~Hz (for wavelengths of order of the size of the Hubble radius today) to $\\simeq 10^{10}$~Hz (the smallest mode that has been inflated out of the Hubble radius) and they could be detected or constrained by coming experiments such as LIGO~\\cite{ligo}, VIRGO~\\cite{virgo} (at $\\simeq 10^{2}$~Hz) and LISA~\\cite{lisa} (at $\\simeq 10^{-4}$~Hz). Gravitational waves, which are perturbations in the metric of the universe have also an effect on the cosmic microwave (CMB) temperature anisotropy~\\cite{soft1,soft2,soft3,white92,turner93} and polarisation~\\cite{pola} allowing to extract information on their amplitude from the measure of the CMB anisotropies. For instance, bounds on the energy density spectrum of these cosmological gravitational waves in units of the critical density, $\\Omega_\\GW$, have been obtained from the CMB~\\cite{soft1,soft2,soft3} $$ \\left.\\frac{\\ddd \\Omega_\\GW}{\\ddd \\ln \\omega}\\right|_{10^{-18}\\,\\mathrm{Hz}} \\lesssim 10^{-10}. $$  Gravitational waves are also a very good probe of the conditions in the early universe since they decouple early in its history and can help \\EG{} testing the initial conditions of $\\phi$. An example was put forward by Giovannini~\\cite{giovannini99,giovannini99b} who showed that in a class of quintessential inflation models~\\cite{vilenkin} there was an era dominated by the scalar field $\\phi$ before the radiation dominated era which implies that a large part of the gravitational wave energy of order $\\Omega_\\GW\\simeq 10^{-6}$ (about eight orders of magnitude higher than for standard inflation) was in the GHz region. This may happen in any scenario where the inflation ends with a kinetic phase~\\cite{ford87,spoko93} or when the dominant energy condition is violated \\cite{gio3}. On the other hand, the CMB temperature fluctuations give information on the history of the gravitational waves in between the last scattering surface and today through the integrated Sachs-Wolfe effect, whereas the polarisation of the CMB radiation gives mainly information on the gravitational waves at decoupling. These three observables (energy spectrum, CMB temperature and polarisation anisotropies) are thus complementary and we aim to present here a global study of the cosmological properties of the gravitational waves.  \\vspace{0.5cm}  The goals of this article are (i) to study in more details the cosmology with a non-minimal quintessence field and (ii) to study gravitational waves in this class of models. In \\S\\ref{par2} we set up the general framework and describe the two potentials we shall consider. In \\S\\ref{par3} we introduce and define the observable quantities associated with the gravitational waves: their energy density spectrum and their imprint on the CMB radiation anisotropies and its polarisation. In \\S\\ref{sec_damping}, we point out the general mechanism of damping by the anisotropic stress of the radiation. In \\S\\ref{par4} we discuss the parameters of the problem and investigate the tuning of the potential parameters; we also describe the evolution of the background spacetime and show that a non-minimally coupled quintessence field is a candidate for a ($\\omega<-1$)-matter. In \\S\\ref{par5} we describe the main properties of the gravitational waves. We finish in \\S\\ref{par6} by presenting numerical results and we underline the complementarity of the different observational quantities.  This work gives a detailed study of the observational effects of gravitational waves in the framework of quintessence, including some recent developments, and allowing for non-minimal coupling. This extends the work on quintessential inflation~\\cite{giovannini99} by including the effects on the CMB. It also extends the studies on $\\Lambda$CDM~\\cite{melchiorri99} to quintessence and is, as far as we know, a more complete study of the effect of gravitational waves on the CMB polarisation. We hope to show that a joint study of the gravitational wave detection experiments~\\cite{ligo,virgo,lisa} of the CMB experiments~\\cite{boomerang,map,planck} and of the polarisation experiments~\\cite{planck} can lead to a better determination of their properties. ",
        "conclusions": "\\label{par6}  In this article, we have studied some properties of quintessence models with a non-minimally coupled scalar field among which the spectrum of gravitational waves.  We have shown that such a quintessence field can behave as a fluid with $\\omega<-1$ and our models lead to $-3\\lesssim \\omega \\lesssim 0$ when the field dominates. We related the energy scale $\\Lambda$ of the potential to its slope $\\alpha$ and to the scalar field energy density today $\\Omega_\\phi^0$. In particular, we showed that $\\Lambda$ is almost independent of $\\Omega_\\phi^0$. The {\\it coincidence problem}, \\IE{} the fact that $\\Omega_\\phi^0\\sim1$ implies a tuning of $\\Lambda$ (roughly the precision on $\\Lambda$ has to one order of magnitude higher than the one on $\\Omega_\\phi^0$) which is however less severe than the fine tuning needed for a cosmological constant. This being fixed, the tracking mechanism allows to span a very wide range of initial conditions for the scalar field and there is no fine tuning in that respect.  We then showed that the combined study of the gravitational waves energy spectrum and of their imprint on the CMB radiation temperature and polarisation enables to extract many complementary informations on the models: \\begin{itemize} \\item the CMB mainly gives results on $\\xi$, $\\Omega_\\phi^0$ and $n_T$, \\item the energy spectrum gives results on the initial conditions of the scalar field. \\end{itemize} As pointed out in~\\cite{giovannini99,vilenkin}, there is an excess of gravitational waves today if inflation ends by a kinetic phase. In that case, one has to check that both $\\Omega_\\phi$ and $\\Omega_{\\rm GW}$ are negligible at the time of nucleosynthesis and we relate the amount of gravitational waves today to the reheating temperature and the time of equality between the kinetic scalar era and the radiation era.  We also pointed out that gravitational waves are damped by the anisotropic stress of radiation, which implies that the CMB anisotropy and polarisation spectra are lowered roughly by 10\\% for high multipoles. It was also shown that the amplitude of the gravitational waves spectrum for inverse power law potentials is $\\sim 30$\\% higher than for SUGRA-like potentials at high frequency. Indeed this is probably not detectable by coming experiments but it could ultimately lead to a signature of supergravity."
    },
    "0004/astro-ph0004374_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-ray bursts (GRBs) are a transient astrophysical phenomenon in which the emission is confined exclusively to high energies: they are detected at gamma-ray bands and last shortly (from a few milliseconds to several hundred seconds). Since the discovery of the events about thirty years ago (Klebesadel et al. 1973), more and more data have been obtained (e.g., Fishman et al. 1994; Meegan et al. 1994, 1996, 1998; Paciesas et al. 1997, 1999). Based on the data available so far, statistical studies have recently unveiled many basic properties of the objects. For example, two separate classes of GRBs with a division at the duration of 2 s were detected (Dezalay et al. 1992; Hurley 1992; Kouveliotou et al. 1993); a correlation between the peak spectral hardness and the peak intensity was discovered (Atteia et al. 1991, 1994, 1997; Belli 1992; Mitrofanov et al. 1992, 1996; Paciesas et al 1992; Nemiroff et al. 1994; Mallozzi et al. 1995; Dezalay et al. 1997).  More recently, the hardness-duration correlation was confirmed by Fishman (1999) with a large number of bursts observed with BATSE. Generally, the hardness ratio is defined as the fluence in channel 3 ($\\sim 100$ to $\\sim 300$ keV) divided by the fluence in channel 2 ($\\sim 50$ to $\\sim 100$ keV) (see, e.g., Fishman et al. 1994). This quantity is more likely to present an intrinsic property of the objects rather than an observational one. It is noticed that, besides the two channel fluences, there are two other channel fluences available in the BATSE burst catalogs, e.g., channels 1 ($\\sim 20$ to $\\sim 50$ keV) and 4 ($>300$ keV) (see, e.g., Paciesas et al. 1997). Therefore, one can define more than one hardness ratios. We wonder if there are any meaningful statistical relations between these different hardness ratios. We also want to make clear if there are any differences between the two classes of the objects in the distributions of the ratios.  In Section 2, correlations between various hardness ratios of GRBs will be studied, while in Section 3, differences between the two classes of the objects in the distributions of the ratios will be investigated. ",
        "conclusions": ""
    },
    "0004/astro-ph0004142_arXiv.txt": {
        "abstract": "The comparison of quasi-instantaneously measured HIPPARCOS proper motions with long-term averaged proper motions, derived from ground-based data, allows the identification of many stellar objects as $\\Delta\\mu$ binaries (Wielen et al., 1999, A\\&A 346, 675). We have used this method to find $\\Delta\\mu$ binaries among the fundamental stars, among the GC stars, and among the Tycho-2 stars. A finding list for about 4000 $\\Delta\\mu$ binaries is given under the URL http://www.ari.uni-heidelberg.de/dmubin in a machine-readable format. This information on the probable duplicity of the listed objects can be used for planning specific observing programs, for improving our knowledge of binary statistics, or for avoiding the use of these binaries in some astrophysical calibrations or as astrometric reference stars. ",
        "introduction": "In order to identify $\\Delta\\mu$ binaries (Wielen et al. 1999a) we compare the HIPPARCOS proper motions (ESA 1997), which are quasi-instantaneously measured within about three years, with long-term averaged, mean proper motions, which are derived by using ground-based data from a much longer observational period (often a century or more). If the difference $\\Delta\\mu$ between the short-term and the long-term proper motion is statistically significant, the object is probably a binary or a multiple system. The statistical significance of $\\Delta\\mu$ is measured by the test parameter $F_{\\Delta\\mu}$. If the errors of the proper-motion components were independent of the directions, $F_{\\Delta\\mu}$ would be equal to $\\Delta\\mu$/(mean error of $\\Delta\\mu$). We take into account the non-isotropy of the errors of $\\Delta\\mu$ and the correlations between the components of $\\Delta\\mu$. A value of $F_{\\Delta\\mu} >$ 3.44 corresponds to the same error probability as the familiar two-sided 3\\,$\\sigma$ criterion, and is therefore used by us for identifying $\\Delta\\mu$ binaries.  The long-term averaged proper motions are partially taken from available astrometric catalogues; partially they are derived from the mean positions given in the ground-based catalogues and the positions in the HIPPARCOS Catalogue (we call this proper motion $\\mu_0$). The ground-based catalogues are corrected for systematic errors by comparing them with the HIPPARCOS Catalogue.  For determining various values of $\\Delta\\mu$ with respect to HIPPARCOS, we have used the following astrometric catalogues: FK5 (Fricke et al. 1988, 1991), RSup (Schwan et al. 1993), GC (Boss et al. 1937), and TYCHO-2 (Hoeg et al. 2000). More details will be given in other papers. For the basic FK5 stars, Part I of the FK6 (Wielen et al. 1999b) gives already many additional informations (see also URL: http://www.ari.uni-heidelberg.de/fk6). ",
        "conclusions": ""
    },
    "0004/astro-ph0004004_arXiv.txt": {
        "abstract": "\\noindent Large-scale astrophysical systems are non-extensive due to their long-range force of gravity. Here we show an approach toward the statistical mechanics of such self-gravitating systems (SGS). This is a generalization of the standard statistical mechanics based on the new definition of entropy; Tsallis statistical mechanics. Developing the composition of entropy and the generalized Euler relation, we investigate the galaxy distributions in count-in-cell method. This is applied to the data of CfA II South redshift survey. ",
        "introduction": "% Astrophysical systems in the Universe are characterized by the gravitation. The structure formed through this long-range force is quite different from those formed through the other short-range forces. If the system does not strongly depend on the initial conditions of the Universe, we can apply statistical mechanics for describing such self-gravitating systems (SGS). However, we cannot directly apply the standard Boltzmann statistical mechanics for SGS since the long-range nature of gravity strongly violates the extensive property of the system which is the premise of statistical mechanics. Actually, the total energy increases much faster than the particle number $N$, the partition function $Z$ often becomes complex\\cite{iguchi99}, reflecting the fact that there is no absolute stable state in SGS.  In order to seek for workable statistical mechanics of SGS, we try an approach based on the new definition of entropy whose extensivity is violated from the beginning; Tsallis statistical mechanics.  We formulate the count-in-cell method for the large scale galaxy distributions in this new statistical mechanics. First we calculate the expression of the composite entropy and the generalized Euler relation in this new statistical mechanics. These are applied to the data of CfA II South redshift survey. The parameter q becomes negative, which represents the instability of gravity. ",
        "conclusions": "% We have constructed the non-extensive statistical mechanics based on the non-extensive entropy. Especially calculating the entropy of composite systems and deriving the generalized Euler relation in thermodynamics, we could evaluate the void probability function $f\\left( 0 \\right)$ and the probability to find $N$ galaxies $f\\left( N \\right)$. This result was applied to the CfAII South galaxy observations and we have obtained negative parameter $q$. This is thought to be another representation of the intrinsic instability of SGS. It will be also interesting to notice the fact that the multi-fractal scaling is observed in this CfAII South data within the scale-region from $500[{\\rm km/sec}]$ to $3000[{\\rm km/sec}]$\\cite{kurokawa99}. We would like to clarify possible connection between the non-extensive distributions and the multi-fractal nature in the context of gravity.  On the way we derive $f\\left( N \\right)$, we encountered ``scale ($\\alpha $) dependent temperature $T$''. If we put the values we obtained $q =  - 5.66847{\\rm  }$ and $s = 0.164142{\\rm  }$ to Eq.(\\ref{T-alpha}), we can explicitly plot the scale dependence of the temperature. It turns out to reduce with increasing scale and abruptly drops to zero at about $r \\approx 600[{\\rm km/sec}]$ or, assuming the cosmic expansion speed as $H = 72[{\\rm km/sec/Mpc]}$, at about $8.3[{\\rm Mpc}]$, which is almost the scale that the galaxy correlation function becomes unity. On the other hand, it is apparent that the galaxies do have peculiar velocity of order $1000[{\\rm km/sec}]$ at this scale. Therefore, at least, the quantity $T$ cannot be interpreted as the ordinary temperature as defined from the velocity dispersions. One of our next task will be to elucidate the meaning of $T$.  In relation with the astronomical velocity distributions, the authors\\cite{lavagno98} claim that the velocity distribution of the {\\it clusters of galaxies} can be well fitted by the Tsallis distribution with the parameter $q = 0.23_{ - 0.05}^{ + 0.07}$, which is apparently different from our negative value for {\\it galaxy} distributions. Further study on the velocity distributions in various scales (galaxies, clusters, super-clusters) would reveal the origin and evolution of the large-scale structure of Universe. We would like to report these results in the near future."
    },
    "0004/astro-ph0004281_arXiv.txt": {
        "abstract": " ",
        "introduction": "It was suggested in \\cite{bkPaczynski1} that in the convective region cooling of matter may be enhanced in presence URCA shells appearing when matter contains an isotope with a threshold Fermi energy for an electron capture, coresponding to a density less then the central one. Presence of such isotope leads to existence of a jump in the composition at a density, corresponding to a threshold energy. During convective motion the matter in eddies around this density crosses periodically the boundary. That implies continious beta capture and beta decay in the matter of these eddies.  Because of heating of a degenerate matter due to nonequilibrium beta processes \\cite{bkSeidov1}, with account of convective URCA shell different conclusions had been done with respect to stabilizing or destabilizing the carbon burning in the convective degenerate core \\cite{bkBruenn, bkErgma, bkPaczynski2, bkCouch, bkLazareff, bkIben1, bkIben2, bkBarkat, bkMochkovitch, bkStein}. Here we calculate damping of stellar oscillations in presence of URCA shell (see \\cite{bkIsern}), analyze physical processes in the convective URCA shells and formulate approximate quantitative approach to the solution of this problem. ",
        "conclusions": ""
    },
    "0004/astro-ph0004248_arXiv.txt": {
        "abstract": "We present the 2-D photometric decomposition of the Virgo galaxy \\object{IC3328}. The analysis of the global light distribution of this morphologically classified nucleated dwarf elliptical galaxy (dE1,N) reveals a tightly wound, bi-symmetric spiral structure with a diameter of 4.5\\,kpc, precisely centered on the nucleus of the dwarf. The amplitude of the spiral is only three percent of the dwarf's surface brightness making it the faintest and smallest spiral ever found in a galaxy. In terms of pitch angle and arm winding the spiral is similar to the intermediate-type galaxy M51, but it lacks the dust and prominent \\ion{H}{ii} regions which signal the presence of gas. The visual evidence of a spiral pattern in an early-type dwarf galaxy reopens the question on whether these dwarfs are genuine rotationally supported or anisotropic stellar systems. In the case of IC3328, we argue for a nearly face-on disk (dS0) galaxy with an estimated maximum rotation velocity of $v_{\\mathrm{c,max}}\\approx 55$\\kms. The faintness of the spiral and the small motions within it, suggests that we could be seeing swing-amplified noise. The other possibility is a tidal origin, caused by the near passage of a small companion. ",
        "introduction": "\\begin{figure*} \\centering\\leavevmode \\epsfxsize=16cm \\epsfbox[43 435 536 683]{jerjen.fig1.ps} \\caption[jerjen.fig1.ps]{The deep R-band CCD image of IC3328 (left panel) illustrates the overall morphology of this as dE1,N classified galaxy: a smooth radially decreasing light distribution with a centrally located nucleus. After the subtraction of the axis symmetric component, the residual image (right panel) reveals a prominent 2-armed spiral structure with a possible central bar. \\label{fig1}} \\end{figure*}  In this paper we report on the serendipitous discovery of a spiral structure in the Virgo cluster dwarf elliptical IC3328. The presence of spiral structure provides compelling evidence for the disk nature of that particular dwarf galaxy. The observations are described and the light distribution is analysed in Sect.2. In Sect.3 we estimate the kinematical properties of the dwarf galaxy from the observed light distribution, assuming a likely value for the mass-to-light ratio and distance. We speculate on the origin of the spiral pattern in Sect.4. The concluding section deals with the ramifications for the dwarf elliptical taxonomy arising from IC3328. ",
        "conclusions": ""
    },
    "0004/astro-ph0004412_arXiv.txt": {
        "abstract": "The very recent Boomerang results give an estimate of unprecedented precision of the Cosmic Microwave Background anisotropies on sub--degree scales. A puzzling feature for theoretical cosmology is the low amplitude of the second acoustic peak.  Through a qualitative discussion, we argue that a scarcely considered category of flat models, with a leptonic asymmetry, a high baryon density and a low cosmological constant seems to be in very good agreement with the data, while still being compatible with big bang nucleosynthesis and some other observational constraints.  Although this is certainly not the only way to explain the data, we believe that these models deserve to be included in forthcoming likelihood analyses. ",
        "introduction": "The measurement of Cosmic Microwave Background (CMB) anisotropies has been driving the attention of cosmologists over the past decade.  Very recently, De Bernardis et al. \\cite{boomerang} published the first results of the Boomerang balloon Antarctic flight. With these data, the (recent) story of precision cosmology climbs a new step. For the first time, the anisotropy power spectrum has been measured by a single experiment in a wide range of angular scales, from multipoles $l \\sim 50$ up to $l \\sim 600$, with many independent points, and error bars of order $\\pm$ 20 percent. The observation of a narrow peak, centered around multipoles $l \\simeq 200$, confirms the inflationary picture of an approximately flat Universe with adiabatic fluctuations.  This beautiful result was already suggested by previous CMB experiments (see \\cite{ddbb} for a recent review). On the other hand, the Boomerang anisotropy spectrum exhibits a puzzling behavior on small scales (high multipoles): in the range in which a pronounced secondary peak was expected, the data points are rather low, with an almost flat shape. It seems that the cosmological model most favored during the past year, which is a flat Cold Dark Matter (CDM) model with a large cosmological constant and ``standard'' parameter values (see below), can hardly account for this feature, unless some new ingredient is added.  After this paper was submitted, a detailed analysis of the data by the Boomerang team was released in ref.\\cite{Lange}, followed by another work \\cite{Balbi} including also the new MAXIMA data \\cite{Hanany}. We refer the reader to these works for a more exhaustive interpretation of this puzzling small--scale behavior. In the following discussion, we just intend to point out that a particular category of models, which are scarcely taken into account, seem to be in remarkable agreement with the new published data (as can be seen from Fig.1). We therefore believe that they deserve some attention, and should be included in future data analyses.  While this discussion was being completed, a nice paper by White et al. \\cite{Elena} was put on the preprint database, suggesting many possible explanations of the data, including a large baryon density $\\Omega_b$.\\footnote{In addition, the possibility of explaining the data with a high $\\Omega_b$ was confirmed contemporaneously \\cite{Tegmark} and soon after \\cite{Lange,Balbi} the release of this work.} The key point of this rapid communication is to recall that a large $\\Omega_b$ is still in agreement with the observed light element abundances, provided that a large neutrino asymmetry is also present \\cite{Kang}. With these two ingredients (large $\\Omega_b$ and neutrino asymmetry) the Boomerang data can be nicely fitted in a flat Universe context, especially with a low value of the cosmological constant. Interestingly, this class of models can satisfy some other cosmological requirements, such as the ones coming from the matter power spectrum and from the baryon fraction in clusters. ",
        "conclusions": "In conclusion, we believe that models with a large neutrino asymmetry deserve to be included in forthcoming precise comparisons with experimental data. In practice, this amounts in including simultaneously higher values of $N_{\\mathrm eff}$ and $h^2 \\Omega_b$ than the ones usually considered.  Since a high baryon density enhances odd peaks with respect to even ones, a natural outcome of these models is a large amplitude for the third peak. This will be probably the best way to test this scheme in a near future. If the third peak turns out to be also very low, the situation will be even more puzzling, and more complicated models (for instance, with a Broken--Scale--Invariant primordial spectrum \\cite{LPS} or with topological defects \\cite{Patrick}) may have to enter into the game."
    },
    "0004/astro-ph0004138_arXiv.txt": {
        "abstract": "We study observational constraints on cosmological models with a quintessence field in the form of a dynamical pseudo Nambu--Goldstone boson. After reviewing the properties of the solutions, from a dynamical systems phase space analysis, we consider the constraints on parameter values imposed by luminosity distances from the 60 type Ia supernovae published by Perlmutter \\etal, and also from gravitational lensing statistics of distant quasars. In the case of the type Ia supernovae we explicitly allow for the possibility of evolution of the peak luminosities of the supernovae sources, using simple empirical models which have been recently discussed in the literature. We find weak evidence to suggest that the models with supernovae evolution fit the data better in the context of the quintessence models in question. If source evolution is a reality then the greatest challenge facing these models is the tension between current value of the expansion age, $H_0t_0$, and the fraction of the critical energy density, $\\Omp$, corresponding to the scalar field. Nonetheless there are ranges of the free parameters which fit all available cosmological data. ",
        "introduction": "Scalar fields have played a central role in models of the very early universe for the past 20 years. In the past few years attention has turned to models in which a scalar field plays a dynamical role at late times, rather than simply being frozen in as a static relic vacuum energy. Such models, which have been dubbed ``quintessence'' models \\cite{CDS}, could in principle provide a dynamical solution to the cosmological constant problem -- namely the question of why the magnitude of the vacuum energy at the present epoch is so much smaller than one might na\\\"{\\i}vely expect from particle physics models such as various supergravity theories. A dynamical `solution' of the cosmological constant problem would amount to a demonstration that a particular dynamical evolution of the scalar quintessence field is a natural consequence of the cosmological field equations without fine-tuning of parameters, given some reasonable physical assumptions about the initial conditions.  The most notable recent observational evidence which has driven the theoretical interest is the measurement of the apparent magnitude-redshift relationship using type Ia supernovae (SNe Ia) \\cite{PS}. These results have been interpreted, in the context of a cosmological model containing pressureless dust and a cosmological constant, $\\LA$, as evidence that the universe is undergoing accelerated expansion at the present epoch (see \\cite{P98,R98} and references therein). The validity of this conclusion is currently open to some doubt, however. In particular, a recent analysis by Riess \\etal\\ \\cite{Rie1} indicates that the sample of type Ia supernovae shows a possible evolution in rise times from moderate ($z\\goesas 0.3$) to large ($z\\goesas1$) redshifts. Although the statistical significance of this result has been diminished -- from the $5.8\\si$ level \\cite{Rie1} to the $1.5\\si$ level \\cite{AKN} -- upon a more rigorous treatment of the uncertainties in the data \\cite{AKN}, it remains true that while a systematic evolution in the rise times of the supernovae is not conclusively ruled in, neither is it conclusively ruled out.  Given that an evolution in the shape of the light curves of the supernovae measured in their rest frame remains a real possibility, it would not be surprising if the peak luminosity -- which is the effective standard candle used -- were also to evolve. Riess \\etal\\ \\cite{Rie1} conclude that the type Ia supernovae data could conceivably be explained entirely within the context of an open Friedmann-Robertson-Walker universe together with a reasonable astrophysical evolution model, e.g., a consequence of a time variation of the abundances of relevant heavy elements in the environment of the white dwarf supernovae progenitors. Detailed astrophysical modelling -- see, e.g., \\cite{Arn} -- should hopefully eventually resolve the issue, although at this stage the difference between our theoretical understanding and the observations remains quite substantial \\cite{Rie2}.  In many recent papers it has been commonly assumed that the dynamical scalar field, $\\ph$, should obey an effective equation of state $P_\\ph\\simeq w\\rh_\\ph$ with $-1<w<0$, at the present epoch, in order to obtain a cosmological acceleration, i.e., a negative deceleration parameter $\\qn$. Indeed, the condition that $-1<w<0$ is often taken as a defining characteristic of ``quintessence'' \\cite{CDS}. The broad picture in this cosmological scenario is that the universe is currently in the early stage of an epoch of inflationary expansion. The motivation for this is that one could then hope to have a model cosmology in which observations such as the type Ia supernovae apparent magnitude--redshift relation could be explained by a cosmological acceleration in a similar fashion to models with a cosmological constant, but with the possibility of explaining why the magnitude of the vacuum energy density and the energy density in ordinary pressureless matter, $\\rh_m$, are comparable at the present epoch -- the so-called ``cosmic coincidence problem'' \\cite{ccp}.  One attractive feature of homogeneous isotropic cosmological models with dynamical scalar fields is that many of them possess ``cosmological scaling solutions'' \\cite{LS}, namely solutions which at late times have energy density components which depend on the cosmic scale factor according to $\\rh\\propto a^ {-m_1}$ and $\\rh_\\ph\\propto a^{-m_2}$ simultaneously, and which act as attractors in the phase space. If $m_2<m_1$, which is the case, for example, for simple power-law potentials with inverse powers \\cite{PR,ZWS} $V(\\ph)\\propto \\ph^{-\\al}$, or for certain power-law potentials with positive powers \\cite{LS} then the scalar field dominates at late times, producing a quintessence-dominated cosmology with accelerated expansion at late times. If $m_1=m_2$, which is the case for exponential potentials \\cite{expV1,expV2,WCL,FJ,VL,FW,CLW,MP}, then the scaling solutions are ``self-tuning'' \\cite{FJ} -- i.e., the dynamics of the scalar field follows that of the other dominant energy component, with a dependence $\\rh_\\ph\\propto a^{-4}$ in the radiation--dominated era and a dependence $\\rh_\\ph\\propto a^{-3}$ in the matter--dominated era.  Even if the ultimate late-time behaviour properties of the solutions is not precisely ``self-tuning'' in the above sense, models such as those with inverse power law potentials can still effectively act as ``tracking solutions'', since for a wide range of initial conditions, the solutions rapidly converge to a common, cosmic evolutionary track \\cite{ZWS}. Thus there are a number of ways in which one might hope to solve the ``cosmic coincidence problem'', though in practice a degree of tuning of the parameters has been necessary in all models studied to date.  In this paper, we consider a form of quintessence, an ultra-light pseudo Nambu-Goldstone-boson (PNGB) \\cite{HR} which is still relaxing to its vacuum state. From the viewpoint of quantum field theory PNGB models are the simplest way to have naturally ultra-low mass, spin-$0$ particles and hence perhaps the most natural candidate for a presently-existing minimally-coupled scalar field. The effective potential of a PNGB field $\\ph$ can be taken to be of the form \\cite{FHSW} \\be\\label{VPNGB}V(\\ph)=M^4[\\cos(\\ph/\\ff)+1]\\ ,\\ee where the constant term is to ensure that the vacuum energy vanishes at the minimum of the potential. This potential is characterized by two mass scales, a purely spontaneous symmetry breaking scale $\\ff$ and an explicit symmetry breaking scale $M$.  The effective PNGB mass is $m_\\ph\\goesas M^2/\\ff$. To obtain solutions with $\\Om_\\ph\\goesas1$, the energy scales are essentially fixed \\cite{HR} to values $M\\goesas10^{-3}\\w{eV}$, interestingly close to the neutrino mass scale for the MSW solution to the solar neutrino problem, and $\\ff\\goesas m_{PL}\\simeq10^{19}\\w{GeV}$, the Planck scale. Since these two energy scales have values which are reasonable from the viewpoint of particle physics, one might hope to explain the coincidence that the vacuum energy is dynamically important at the present epoch.  The cosmology of PNGB models has already been extensively studied in the literature \\cite{VL,FW,MP,FHSW,CDF,FW2}. In particular, a number of constraints have been placed on the parameters $M$ and $\\ff$ by various sets of observational data \\cite{VL,FW,CDF,FW2}. Most recently, Frieman and Waga \\cite{FW2} have set bounds based on the SNe Ia data of Riess \\etal \\cite{R98} (hereafter R98) on the one hand, and gravitational lensing surveys on the other. Comparing these bounds is of interest, since the SnIa data have been interpreted as favouring a cosmological constant, whereas gravitational lensing data has been used to place upper bounds on $\\Lambda$ \\cite{K95,Fal,Hel}. They therefore provide complementary tests of the parameter spaces of models with a non-trivial vacuum energy.  In this paper it is our intention to critically study these bounds. Firstly, we will consider how the bounds are affected by the initial value of the scalar field at the beginning of the matter--dominated epoch. Secondly, we wish to investigate how such bounds might be affected in the case of the PNGB model if the observed apparent faintness of type Ia supernovae is at least partly due to an intrinsic evolution of the sources over cosmological time scales, which in view of the results of \\cite{Rie1} would appear to be a very real possibility. The reason for focusing on the PNGB models in such an investigation is suggested by the fact that whereas many quintessence models have been singled out in the literature, perhaps somewhat artificially, simply because they have the property of yielding an accelerated expansion, many different possibilities arise in the PNGB case. Indeed, at very late times, the apparent magnitude--redshift relation for PNGB models ultimately coincides with that of the Einstein--de Sitter model, even though the density of ordinary matter can be low in the PNGB cosmologies. The requirement that the faintness of type Ia supernovae is entirely due to their cosmological distances places rather strong restrictions on the values of the parameters $M$ and $\\ff$ \\cite{FW2}, because it requires us to exist at an epoch of the PNGB cosmologies which is still quite far removed from our ultimate destiny. If these restrictions are relaxed because of evolutionary effects, then it is quite plausible that other regions of the parameter space of the PNGB models become viable alternatives. Since PNGB cosmologies could therefore still solve the ``missing energy problem'', even if the evidence for a cosmological acceleration proves to be ephemeral, we believe it is important to investigate this possibility quantitatively.  We will begin the paper with a qualitative analysis of the solutions, to provide some general insights which will help to guide our quantitative discussion. Although these properties are no doubt already known, to the best of our knowledge an analysis of the phase space of the solutions has never been presented in the literature. Having completed this analysis in Sec.\\ II we will go on to discuss a number of issues relating to numerical integration in Sec.\\ III, and relate the properties of the solutions found numerically to the exact analysis of Sec.\\ II. In Sec.\\ IV we present the main analysis of the constraints imposed on the $(M,\\ff)$ parameter space, allowing for the possibility of evolution of peak luminosities in the type Ia supernova sources. Bounds from gravitational lensing statistics are updated in Sec.\\ V, and the implications of our results are discussed at greater length in Sec. VI. ",
        "conclusions": "Let us now consider the overall implications of the constraints observed above.  Firstly, since empirical models with source evolution do appear to fit the data somewhat better, it would appear that we do have weak evidence for an underlying evolution of the peak luminosity of the type Ia supernovae sources, at least in the context of the PNGB quintessence models. It might be interesting to compare the case of other quintessence models, or the case of a cosmological constant. However, the PNGB model is qualitatively different from such models since its final state corresponds to one in which the ultimate destiny of the universe is to expand at the same rate as a spatially flat Friedmann--Robertson--Walker model, rather than to undergo an accelerated expansion. This is of course precisely why we chose the PNGB models as the basis of our investigation, rather than models in which a late--time accelerated expansion had been built in by hand. If we wish to test the hypothesis that the faintness of the type Ia supernovae is at least partly due to an intrinsic variation of their peak luminosities -- which is a very real possibility in view of the results of \\cite{Rie1} -- then a quintessence model which possesses a variety of possibilities for the present--day variation of the scale factor is probably the best type of model to investigate.  If only supernovae luminosity distances, (cf., Fig.\\ \\ref{SnIa-evolv1}, and gravitational lensing statistics, (cf., Fig.\\ \\ref{lenstat}), are compared then we see that there is a remarkable concordance between the two tests -- region II of Fig.\\ \\ref{SnIa-evolv1} coincides with a region included at even the $1\\si$ level in Fig.\\ \\ref{lenstat}. This is perhaps not surprising, since in view of Fig.\\ \\ref{qn} region II corresponds to parameter values for which the present day universe has already undergone almost one complete oscillation of the scalar field about the final critical point $\\Co_2$ of Fig.\\ \\ref{rhgm0}. It is thus already well on the way towards its asymptotic behaviour, which closely resembles that of a standard spatially flat Friedmann-Robertson-Walker model.  Due to the oscillatory behaviour, parameter values in region II correspond to models in which there has been a recent cosmological acceleration, (e.g., at $z\\goesas0.2$), but with a $q(z)$ which changes sign three times over the larger range of redshifts, $0<z<4$, in the quasar lensing sample, and therefore differing significantly from Friedmann--Lema\\^{\\i}tre models over this larger redshift range. Extending the SNe Ia sample to include objects at redshifts $0.1<z<0.4$ and $z>0.85$ in substantial numbers would greatly improve the ability to decide between models in regions I and II.  There is some cause for concern, however, if we consider the favoured values of $\\Omp$ and $H_0t_0$. In the case of the models with empirical evolution of the supernovae sources, we always found that the best-fit parameter values occured at the $\\Omp\\rarr1$ boundary of the ($M,\\ff$) parameter space. The overwhelming evidence of many astronomical observations over the past two decades \\cite{OS} would tend to indicate that $\\Om_{m0}\\simeq0.2\\pm0.1$, indicating that a vacuum energy fraction of $\\Omp\\goesas0.7$--$0.8$ is desirable, and $\\Omp\\lsim0.9$ in any case. Although parameter values with $\\Omp<0.8$ certainly fall within both the 2$\\si$ and 1$\\si$ portions of region II of Fig.\\ \\ref{SnIa-evolv1}, for all values of $b$, there are potentially serious problems if we wish to simultaneously obtain large values of $H_0t_0$. In view of recent estimates of the ages of globular clusters \\cite{Krauss}, a lower bound of 12 Gyr for the age of the Universe appears to currently indicated. With $h\\simeq0.65$ this would require $H_0t_0\\gsim0.8$. For $w_i=1.5$, parameter values with $H_0t_0>0.8$ coincide with values $\\Omp\\gsim 0.9$ in region II, which is phenomenologically problematic.  The tension between the values of $\\Omp$ and $H_0t_0$ is somewhat mitigated for lower values of $w_i$. For $w_i=0.2$, for example, we see from Figs.\\ \\ref{SnIa-evolv3} and \\ref{lenstat}(b) that the $\\Omp=0.7$ and $H_0t_0$ meet in Region II, and there is a small region of parameters there with $0.7\\lsim \\Omp\\lsim0.9$ and $H_0t_0\\gsim0.8$, which is also consistent with the other cosmological tests.  Even if the supernovae sources undergo evolution it is clear that parameter values in region I of \\ref{SnIabeta0} which are favoured in the absence of evolution of peak SnIa luminosities, are still included at the $2\\si$ level in the models with evolution, in view of Fig.\\ \\ref{SnIa-evolv1}.  Perhaps the most significant aspect of our results is the fact that the mere introduction of an additional dispersion, $b>0.17$, in the peak luminosities while leaving their mean value fixed (cf., Fig.\\ \\ref{bval}), gives rise to a change in the best-fit region of parameter space from region I to region II. (See \\cite{dlw} for further details.) One would imagine that an increased dispersion is likely to be a feature of many models of source evolution, even if evolutionary effects are of secondary importance. Thus even if the empirical models with non-zero $\\beta$ are somewhat artificial, more sophisticated scenarios could well lead to similar changes in regard to the fitting of cosmological parameters in the PNGB model.  Much tighter bounds on the parameter space of quintessence models, including the present model, will be obtained over the next decade as more supernovae data is collected. What we wish to emphasize, however, is that an effective vacuum energy which is cosmologically significant at the present epoch should not simply be thought of in terms of a ``cosmic acceleration''. A dynamical vacuum energy with a varying effective equation of state allows for many possibilities for the evolution of the universe, and overly restrictive assumptions, such as equating quintessence to models with a late period of continuous cosmological acceleration, should be avoided. If detailed astrophysical modeling of type Ia supernovae explosions ultimately shows that the dimness of distant supernova events is largely due to evolutionary effects, it does not spell the end for cosmologies with dynamical scalar fields."
    },
    "0004/astro-ph0004194_arXiv.txt": {
        "abstract": "We report on large-scale, regular morphological patterns found in the radio jet of the nearby radio galaxy NGC 6251. Investigating morphological properties of this radio jet {}from the nucleus to a radial distance of $\\sim$ 300 arcsec ($\\approx$ 140 kpc) mapped at 1662 MHz and 4885 MHz by Perley, Bridle, \\& Willis, we find three chains, each of which consists of five radio knots. We also find that eight radio knots in the first two chains consist of three small sub-knots (the triple-knotty substructures). We discuss the observational properties of these regular morphological patterns. ",
        "introduction": "Since the discovery of the powerful radio jet from quasar 3C273 (Hazard et al. 1963, see for a review Bridle \\& Perley 1984; Zensus 1997), the generation of radio jets has been one of the long-standing main problems in active galactic nuclei (e.g., Begelman, Blandford, \\& Rees 1984; Urry \\& Padovani 1995). The most probable energy sources have been considered to be either mass-accreting, supermassive, single black holes around which gravitational energy is transformed into huge kinetic and radiation energies with the help of gaseous accretion disks (e.g., Rees 1984) or the electromagnetic extraction of energy from spinning supermassive black holes (Blandford \\& Znajek 1977; see also Wilson \\& Colbert 1995). In order to understand the genesis of radio jets, it is important to find some reliable observational constraints on theoretical models.  It has been noticed that parsec-scale radio jets probed by VLBI techniques often show wiggles (e.g., Whitmore \\& Matese 1981; Roos 1988; Roos, Kaastra, \\& Hummel 1993). If a certain mechanism responsible for the formation of the wiggles (e.g., the precession of a supermassive black hole binary; Begelman, Blandford, \\& Rees 1980) has been working since the onset of radio jet activity, there could be some morphological evidence even in well-developed (i.e., $\\sim$ 100 kpc scale) radio jets. Furthermore, recent progress in three-dimensional numerical MHD simulations has enabled us to examine the detail of morphological properties of radio jets due to the so-called Kelvin-Holmheltz instability in the magnetic fluid (e.g., Rosen et al. 1998; Koide, Shibata, \\& Kudoh 1999).  Therefore, since morphological properties of actual radio jets provide important constraints on the physical process involved in radio jets, it is interesting to investigate overall morphological properties of some well-developed radio jets in detail.  For this purpose, in this paper, we investigate morphological properties of the radio jet of NGC 6251 in detail because the radio jet of NGC 6251 is one of the brightest known examples of a well-developed jet (Waggett, Warner, \\& Baldwin 1977; Cohen \\& Readhead 1979; Perley, Bridle, \\& Willis 1984; Jones et al. 1986; Jones \\& Wehrle 1994). We use a distance to NGC 6251, 94.4 Mpc, which is determined with the use of a recession velocity of NGC 6251 to the galactic standard of rest, $V_{\\rm GSR}$ = 7079 km s$^{-1}$ (de Vaucouleurs et al. 1991), and a Hubble constant, $H_{0}$ = 75 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have shown that there are two kinds of regular morphological structures in the radio jet of NGC 6251; 1) three Chains consisting of five radio knots , and 2) the triple-knotty substructures in the individual knots. Here we discuss some possible origins of these large-scale regular structures.  \\subsection{Chains}  The presence of three Chains suggests that a certain periodicity is involved in the radio jet activity of NGC 6251. As shown in Table 1, the separations of the $i$-th knots ($W$) are quite similar between two adjacent Chains; i.e., the average values are $\\overline{W}$(A-B) $\\simeq 48.0 \\pm 0.5$ kpc and $\\overline{W}$(B-C) $\\simeq 79.3 \\pm 2.9$ kpc. We estimate timescales corresponding to these separations. In order to perform this, both the viewing angle toward the jet ($\\theta_{\\rm jet}$) and the jet velocity ($v_{\\rm jet}$) are necessary for the kpc-scale radio jet.  Since the parsec-scale counterjet cannot be seen in the previous VLBI observations (Perley et al. 1984; Jones et al. 1986; Jones \\& Wehrle 1994; see however Sudou et al. 2000a, b), it is unlikely that the radio jet of NGC 6251 lies close to the celestial plane. Jones et al. (1986) estimate that the angle between the radio jet and our line of sight may be $\\theta_{\\rm jet} \\sim 45^\\circ$ based on the observed jet-to-counterjet intensity ratio.  We therefore adopt $\\theta_{\\rm jet} = 45^\\circ$ in later analysis. Another important quantity is the large-scale (i.e., kpc-scale) jet velocity $v_{\\rm jet}$ which is also difficult to be estimated (e.g., Perley et al. 1984). Based on several constraints (e.g., the energy flux required to power the radio jet, etc.), Perley et al. (1984) suggest that the large-scale jet velocity of NGC 6251 is subrelativistic; $v_{\\rm jet} \\leq 0.1c$. It is known that ram pressure confinement for the strongest double-lobed radio sources such as NGC 6251 requires $v_{\\rm jet} \\simeq 0.1c$ (e.g., Begelman et al. 1984). On the other hand, using the observed jet-to-counterjet brightness ratio, Jones et al. (1986) suggests  $v_{\\rm jet} \\cos \\theta_{\\rm jet} \\geq 0.6$ for brightness ratio $ \\geq 30$. Given $\\theta_{\\rm jet} = 45^\\circ$, they obtain $v_{\\rm jet} \\sim  0.84c$. Since this estimate seems more reliable, we adopt for simplicity $v_{\\rm jet} = 0.8c$ in later analysis. These assumptions (i.e., $\\theta_{\\rm jet} = 45^\\circ$ and $v_{\\rm jet} = 0.8c$) seem enough to estimate rough timescales related to the large-scale regular structures found in this study.  First we estimate timescales related to three Chains A, B, and C. Since the jet velocity is relativistic, we have to take account of the relativistic aberration effect. The true projected distance of the radio jet from the nucleus is estimated as $D_{\\rm jet}' = \\delta \\ D_{\\rm jet}$ where $\\delta$ is the Doppler factor defined as $\\delta = [\\gamma(1 - (v_{\\rm jet}/c) \\cos \\theta_{\\rm jet})]^{-1}$ where $\\gamma = [1 - (v_{\\rm jet}/c)^2]^{-1/2}$. Thus the true length of the radio jet is estimated as $L_{\\rm jet}^0  = D_{\\rm jet}' (\\sin \\theta_{\\rm jet})^{-1}$. Therefore, the related timescale $\\tau_{\\rm jet}$ is estimated as  \\begin{eqnarray} \\tau_{\\rm jet} & = & L_{\\rm jet}^0 / v_{\\rm jet} \\nonumber \\\\ & = & \\ D_{\\rm jet}'/ (0.8 c ~ {\\rm sin} \\theta_{\\rm jet}) \\nonumber \\\\ & \\simeq & 1.83 \\times 10^{11}\\  \\delta \\ D_{\\rm jet, 1}\\ v_{\\rm jet, 0.8}^{-1}\\ (\\sin \\theta_{\\rm jet, 45})^{-1}~ {\\rm s} \\nonumber \\\\ & \\simeq & 5.79  \\times 10^3\\   \\delta \\ D_{\\rm jet, 1}\\ v_{\\rm jet, 0.8}^{-1}\\ (\\sin \\theta_{\\rm jet, 45})^{-1}  ~ {\\rm y} \\end{eqnarray} where $D_{\\rm jet, 1}$ is the jet length projected onto the celestial plane in units of 1 kpc, $v_{\\rm jet, 0.8}$ is the jet velocity in units of $0.8c$, and $\\theta_{\\rm jet, 45}$ is the jet viewing angle in units of 45$^\\circ$. Given the velocity and viewing angle assumed above, we obtain $\\delta = 1.38$. The projected jet lengths of three Chains are $D_{\\rm jet}$(A) = $D_5$(A) $-$ $D_1$(A) = 32.4 $-$ 3.6 = 28.8 kpc, $D_{\\rm jet}$(B) = $D_5$(B) $-$ $D_1$(B) = 87.8 $-$ 45.5 = 42.3 kpc, and $D_{\\rm jet}$(C) = $D_5$(C) $-$ $D_1$(C) = 169.9 $-$ 120.7 = 49.2 kpc. Then we obtain the durations required to develop the radio jet for three Chains; $\\tau_{\\rm jet}$(A) $\\approx 2.3 \\times 10^5$ years, $\\tau_{\\rm jet}$(B) $\\approx 3.4 \\times 10^5$ years, and $\\tau_{\\rm jet}$(C) $\\approx 3.9 \\times 10^5$ years. These durations are shorten by one order of magnitude than the precession period estimated by Jones et al. (1986; see also Begelman et al. 1980), $\\tau_{\\rm prec} \\simeq 1.8 \\times 10^6$ y. This precession is proposed to explain the global wiggle pattern of the radio jet of NGC 6251.  As estimated above, the length of Chain becomes longer with increasing radial distance; i.e., $D_{\\rm jet}$(A) $<$ $D_{\\rm jet}$(B) $<$ $D_{\\rm jet}$(C). If this tendency is real, the radio jet of NGC 6251 must be accelerated even at several tens of kpc. An alternative idea may be that the radio jet is bending in a plane which encloses the radio jet and our line of sight to the jet. Since the position angle of the parsec-scale radio jet (PA = 302$^\\circ$.2 $\\pm$ 0$^\\circ$.8, for the epoch of 1950.0; Cohen \\& Readhead 1979) is slightly different from that of the kpc-scale jet (PA = 296$^\\circ$.5; Waggett et al. 1977), Cohen \\& Readhead (1979) suggest a possible bending of the radio jet of NGC 6251. Therefore, it is interesting to investigate this idea in more detail.  Here we assume that the true jet lengths of three Chains are nearly the same and the observed differences among them are attributed to the differences in the viewing angle toward them. If this is the case, we obtain the following relation;  \\begin{equation} {D_{\\rm jet}({\\rm A}) ~ \\delta({\\rm A}) \\over {\\rm sin} \\theta_{\\rm jet}({\\rm A})} = {D_{\\rm jet}({\\rm B}) ~ \\delta({\\rm B}) \\over {\\rm sin} \\theta_{\\rm jet}({\\rm B})} = {D_{\\rm jet}({\\rm C}) ~ \\delta({\\rm C}) \\over {\\rm sin} \\theta_{\\rm jet}({\\rm C})} \\end{equation} where $\\theta_{\\rm jet}({\\rm A})$, $\\theta_{\\rm jet}({\\rm B})$, and $\\theta_{\\rm jet}({\\rm C})$ are the average viewing angles toward Chain A, B, and C, and $\\delta(\\rm A)$,  $\\delta(\\rm B)$, and $\\delta(\\rm C)$ are the average Doppler factors toward Chain A, B, and C, respectively. If we adopt $\\theta_{\\rm jet}({\\rm C}) = 45.0^\\circ$ and $\\delta$(C) = 1.38, we obtain $\\theta_{\\rm jet}({\\rm A}) \\simeq 33.1^\\circ$ and $\\theta_{\\rm jet}({\\rm B}) \\simeq 41.3^\\circ$, with $\\delta({\\rm A}) \\simeq 1.82$ and $\\delta({\\rm B}) \\simeq 1.50$. It is therefore suggested that the direction of the radio jet is approaching the line of sight as time goes by, or the jet flow follows fixed but bent path. This result is schematically illustrated in Figure 8.  Recently, Sudou et al. (2000a, b) have found the counterjet at sub-parsec scale and estimated the viewing angle to the sub-pc scale jet $\\theta_{\\rm jet} \\simeq$ 17$^\\circ$ -- 31$^\\circ$. Since the viewing angle toward Chain A derived above is consistent with their new estimate, this bending jet model appears consistent with the observation. Therefore, it is not necessary to introduce the jet acceleration at kpc regions.  \\subsection{The Knots in Chains}  We investigate the separations of the knots in three Chains. An average separation of the knots in each Chain is; $\\overline{S} = \\sum_{i=1}^4 S_{[i+1]-1}/4 \\simeq 7.2 \\pm 2.9$ kpc for Chain A, $10.6 \\pm 4.5$ kpc for Chain B, and $12.3 \\pm 3.1$ kpc for Chain C. Therefore, the average separation appears to increase with increasing Chain number. It is noted that the fifth knot in each Chain is located at a larger distance than that expected from the separations for the remaining four knots (see column 5 of Table 1). If we omit the data of the fifth knot, we obtain average separations of $\\overline{S} = \\sum_{i=1}^3 S_{[i+1]-1}/3 \\simeq 5.6 \\pm 1.0$ kpc for Chain A, $8.1 \\pm 1.6$ kpc for Chain B, and $10.6 \\pm 1.2$ kpc for Chain C. Although the tendency  holds, the average separations are smaller than the former estimates respectively. If we adopt the bending jet model described in section 3.2 together with the relativistic aberration effect, the true separation, $S^0 = \\delta \\overline{S}/ \\sin \\theta_{\\rm jet}$, is; $S^0({\\rm A}) \\simeq 5.6 \\times 1.82 / \\sin 33.1^\\circ \\simeq 18.7 $ kpc, $S^0({\\rm B}) \\simeq 8.1 \\times 1.50 / \\sin 41.3^\\circ \\simeq 18.4$ kpc, and $S^0({\\rm C}) \\simeq 10.6 \\times 1.38 / \\sin 45^\\circ \\simeq 20.7$ kpc. These values are similar to each other.  We obtain an average separation of the knots in three Chains of $\\simeq$ 19.2 kpc, corresponding to a timescale of $\\simeq 7.8 \\times 10^4$ years given the jet velocity of $v_{\\rm jet} = 0.8c$.  \\subsection{The Triple-Knotty Substructure}  We discuss briefly observational properties of the triple-knotty substructures. As shown in Figure 7, the first two knots (A2 and  A3) show a clear triple-knotty substructures. However, the others show a range of irregular, complex morphologies although we give possible identifications of the triple knots for them by arrows. It is interesting to mention that such a triple-knotty substructure is also found in the radio jet of Centaurus A (Clarke et al. 1986).  \\subsection{Concluding Remarks}  In summary, the large-scale, regular morphological patterns involve the three kinds of structures; Chains, knots, and triple-knotty substructures in the knots. The concerned timescales for the first two structures are $\\sim 10^5$ years and $\\sim 10^4$ years, respectively. Although the longest timescale obtained for Chains may be related to the precession motion, the other two timescales are not understood easily.  Although some studies of the large-scale morphological properties of the radio jets have been carried out (e.g., Perley \\& Bridle 1984, Sparks et al. 1996, Perlman et al. 1999, Biretta et al. 1999, Bahcall et al. 1995, R$\\ddot{\\rm o}$ser et al. 1996, and Clarke et al. 1986), the analysis presented in this paper is the first trial to investigate large-scale morphological regularity of radio jets. Since this kind of analysis needs high-resolution and large-scale radio continuum mapping, it seems difficult to perform a systematic morphological study of radio jets. Therefore, at present, it is difficult to judge whether or not large-scale morphological patterns found in the radio jet of NGC 6251 are general properties. However, this kind of analysis will be important  to provide observational constraints on the theory for radio jets.   \\vspace{0.5cm}  We thank to Dr. D. Jones and anonymous referee for their useful comments and suggestions. This work was financially supported in part by Grant-in-Aids for the Scientific Research (Nos. 10044052, and 10304013) of the Japanese Ministry of Education, Culture, Sports, and Science.   \\newpage"
    },
    "0004/astro-ph0004311_arXiv.txt": {
        "abstract": "We present new multicolour CCD photometry of the central part of the globular cluster M3, mapping the precise position of $\\sim 120$ RR Lyrae stars (RRab, RRd, RRc) on the horizontal branch (HB). The location of the double-mode variables (RRd) is in perfect agreement with recent theoretical results. We find a significant internal spread of metallicity amongst the RRab variables. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004257_arXiv.txt": {
        "abstract": "The interaction of an expanding Pair-Electromagnetic pulse (PEM pulse) with a shell of baryonic matter surrounding a Black Hole with electromagnetic structure (EMBH) is analyzed for selected values of the baryonic mass at selected distances well outside the dyadosphere of an EMBH. The dyadosphere, the region in which a super critical field exists for the creation of $e^+e^-$ pairs, is here considered in the special case of a Reissner-Nordstrom geometry. The interaction of the PEM pulse with the baryonic matter is described using a simplified model of a slab of constant thickness in the laboratory frame (constant-thickness approximation) as well as performing the integration of the general relativistic hydrodynamical equations. The validation of the constant-thickness approximation, already presented in a previous paper Ruffini, et al.(1999) for a PEM pulse in vacuum, is here generalized to the presence of baryonic matter. It is found that for a baryonic shell of mass-energy less than 1\\% of the total energy of the dyadosphere, the constant-thickness approximation is in excellent agreement with full general relativistic computations. The approximation breaks down for larger values of the baryonic shell mass, however such cases are of less interest for observed Gamma Ray Bursts (GRBs). On the basis of numerical computations of the slab model for PEM pulses, we describe (i) the properties of relativistic evolution of a PEM pulse colliding with a baryonic shell; (ii) the details of the expected emission energy and observed temperature of the associated GRBs for a given value of the EMBH mass; $10^3M_\\odot$, and for baryonic mass-energies in the range $10^{-8}$ to $10^{-2}$ the total energy of the dyadosphere. ",
        "introduction": "That vacuum polarization process {\\it \\`a la} Heisenberg-Euler-Schwinger (\\cite{he}, \\cite{s}) can occur in the field of a Kerr Newmann EMBH and that they naturally lead to a model for gamma-ray bursts was pointed out in \\cite{dr}. The basic energetics of the process, governed by the Christodoulou-Ruffini mass formula, for an EMBH gives as shown in \\cite{rc}, \\begin{eqnarray} E^2&=&M^2c^4=\\left(M_{\\rm ir}c^2 + {Q^2\\over2\\rho_+}\\right)^2+{L^2c^2\\over \\rho_+^2},\\label{em}\\\\ S&=& 4\\pi \\rho_+^2=4\\pi(r_+^2+{L^2\\over c^2M^2})=16\\pi\\left({G^2\\over c^4}\\right) M^2_{\\rm ir}, \\label{sa} \\end{eqnarray} with \\begin{equation} {1\\over \\rho_+^4}\\left({G^2\\over c^8}\\right)\\left( Q^4+4L^2c^2\\right)\\leq 1, \\label{s1} \\end{equation} where $M, L$ and $Q$ are respectively the mass energy, the angular momentum and the charge of the EMBH and $M_{\\rm ir}$ is the irreducible mass, $r_{+}$ is the horizon radius, $\\rho_+$ is the quasi-spheroidal cylindrical coordinate of the horizon evaluated at the equatorial plane and $S$ is the horizon surface area. Extreme black holes satisfy the equality in Eq.(\\ref{s1}). The vacuum polarization process being ``reversible'' transformations in the sense of \\cite{rc} can extract an energy up to 29\\% of the mass-energy of an extremal rotating black hole and up to 50\\% of the mass-energy of an extremely magnetized and charged black hole.  Although in general such a process is endowed with axial symmetry, in order to clarify the pure interplay of the gravitational and electrodynamical phenomena and also for simplicity, we have examined (\\cite{rr} and \\cite{prxa}) the limiting cases of such a process in the field of a Reissner-Nordstrom geometry whose spherically symmetric metric is given by \\begin{equation} d^2s=g_{tt}(r)d^2t+g_{rr}(r)d^2r+r^2d^2\\theta +r^2\\sin^2\\theta d^2\\phi ~, \\label{s} \\end{equation} where $g_{tt}(r)= - \\left[1-{2GM\\over c^2r}+{Q^2G\\over c^4r^2}\\right] \\equiv - \\alpha^2(r)$ and $g_{rr}(r)= \\alpha^{-2}(r)$. The dyadosphere, defined (see Fig. 1 of \\cite{prxa}) as the region where the electric field exceeds the critical value, ${\\cal E}_{\\rm c}={m^2c^3\\over\\hbar e}$ (\\cite{he}, \\cite{s}), where $m$ and $e$ are the mass and charge of the electron, extends between the horizon radius \\begin{eqnarray} r_{+}&=&1.47 \\cdot 10^5\\mu (1+\\sqrt{1-\\xi^2})\\hskip0.1cm {\\rm cm}. \\label{r+} \\end{eqnarray} where we have introduced the dimensionless mass and charge parameters $\\mu={M\\over M_{\\odot}}$, $\\xi={Q\\over (M\\sqrt{G})}\\le 1$, and the outer radius \\begin{eqnarray} r_{\\rm ds} &=&1.12\\cdot 10^8\\sqrt{\\mu\\xi} \\hskip0.1cm {\\rm cm}. \\label{rc} \\end{eqnarray} Consequently the local number density of electron and positron pairs created in the dyadosphere as a function of radius \\begin{equation} n_{e^+e^-}(r) = {Q\\over 4\\pi r^2\\left({\\hbar\\over mc}\\right)e}\\left[1-\\left({r\\over r_{\\rm ds}}\\right)^2\\right] ~, \\label{nd} \\end{equation} and their energy density is given by \\begin{equation} \\epsilon(r) = {Q^2 \\over 8 \\pi r^4} \\biggl(1 - \\biggl({r \\over r_{\\rm ds}}\\biggr)^4\\biggr) ~, \\label{jayet} \\end{equation} (see Figs.[2] \\& [3] of \\cite{prxa}). The total energy of pairs converted from the static electric energy and deposited within the dyadosphere is then \\begin{equation} E_{\\rm dya}={1\\over2}{Q^2\\over r_+}(1-{r_+\\over r_{\\rm ds}})\\left[1- \\left({r_+\\over r_{\\rm ds}}\\right)^2\\right] ~, \\label{tee} \\end{equation} and the total number of electron and position pairs in the dyadosphere is \\begin{equation} N^\\circ_{e^+e^-}\\simeq {Q-Q_c\\over e}\\left[1+{ (r_{\\rm ds}-r_+)\\over {\\hbar\\over mc}}\\right], \\label{tn} \\end{equation} where $Q_c={\\cal E}_{\\rm c}r_+^2$ (see \\cite{prxa}). In the limit of ${r_+\\over r_{\\rm ds}}\\rightarrow 0$, Eq.(\\ref{tee}) leads to $E_{\\rm dya}\\rightarrow {1\\over2}{Q^2\\over r_+}$, which coincides with the energy extractable from EMBHs by reversible processes ($M_{\\rm ir}={\\rm const.}$), namely $E-M_{\\rm ir}={1\\over2}{Q^2\\over r_+}$(see Fig. 4 of \\cite{prxa}).  Due to the very large pair density given by Eq.(\\ref{nd}) and to the sizes of the cross-sections for the process $e^+e^-\\leftrightarrow \\gamma+\\gamma$, the system is expected to thermalize to a plasma configuration for which \\begin{equation} N_{e^+}=N_{e^-} \\sim N_{\\gamma} \\sim N^\\circ_{e^+e^-}, \\label{plasma} \\end{equation} where $N^\\circ_{e^+e^-}$ is the number of $e^+e^-$-pairs created in the dyadosphere(see \\cite{prxa}). These are the initial conditions for the evolution of the dyadosphere. In a previous paper (\\cite{rswx99}) we presented the temporal evolution of the dyadosphere in vacuum giving origin to an extremely sharply defined and extremely relativistic expanding pulse of pair and electromagnetical radiation of a constant length in the laboratory frame: the PEM pulse. In this paper we present results of the collision of the PEM pulse with a remnant of baryonic matter surrounding the just formed black hole. We assume the PEM pulse to collide with a shell of baryonic matter of constant density and at a radius of the order of 100 times the radius of the dyadosphere $r_{\\rm ds}$ Eq.(\\ref{rc}). The shells have this thickness of the order of 10 times $r_{\\rm ds}$. The mass-energies of the shells are taken to be $10^{-8}-10^{-2}$ of the total energy of the dyadosphere.  The work contains the following sections: Sect. \\ref{hydrodynamic} presents discussions of the hydrodynamical equations of a PEM pulse interacting with the baryonic shell; Sect. \\ref{baryon} defines the parameters of the baryonic shells, the behaviour of the PEM pulse colliding with a baryon shell as well as before and after the collision are presented; in Sect. \\ref{comparison} the results of a numerical calculation solving the hydrodynamical equations of Sect. \\ref{hydrodynamic} are compared to the results of the analytical model of slab approximation for selected parameters. ",
        "conclusions": "By the direct comparison with the numerical integration of the complete relativistic hydrodynamical equations, the use of the constant-thickness approximation has been validated for values of the parameter $B\\leq10^{-2}$. For $B\\geq 10^{-2}$ the amount of energy released at transparency in the burst decreases, and its energy drifts toward low energy values, which are of little interest for the astrophysics of GRB. We conclude that the constant-thickness approximation is valid for all astrophysically relevant situations for the analysis of the GRB at transparency. Based on this validation we have studied the evolution of a PEM pulse in the presence of selected amounts of baryonic matter. We have studied for a typical case of $M_{\\rm BH}=10^3M_\\odot, \\xi=0.1$ EMBH and $B= 10^{-2}$ baryonic shell, the thermal, the bulk Lorentz factor evolution of the PEM pulse as well as the kinetic energy left over in the baryonic matter after the decoupling of matter and radiation and the emission of the GRB. Additional results corresponding to a larger range of masses and charges of the EMBH and the correlations between the peak energy and the duration of GRBs to be expected in our model will be considered in a future paper, together with the results of analyzing the interaction of the kinetic energy left over in the baryonic matter, after the decoupling of matter and radiation, with the interstellar medium."
    },
    "0004/astro-ph0004061_arXiv.txt": {
        "abstract": "We report first observational results of multifrequency campaigns on the prominent Virgo blazars 3C~273 and 3C~279 which were carried out in January and February 1999. Both blazars are detected from radio to \\gray\\ energies. We present the measured X- to \\gray\\ spectra of both sources, and for 3C~279 we compare the 1999 broad-band (radio to \\gray) spectrum to measured previous ones. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004204.txt": {
        "abstract": "First, we review the main physical effects to be considered in the building of evolutionary models of rotating stars on the Upper Main-Sequence (MS). The internal rotation law evolves as a result of contraction and expansion, meridional circulation, diffusion processes and mass loss. In turn, differential rotation and mixing exert a feedback on circulation and diffusion, so that a consistent treatment is  necessary.  We review recent results on the evolution of internal rotation and the  surface rotational velocities for stars on the Upper MS, for red giants, supergiants and W--R stars. A fast rotation is enhancing the mass loss by stellar winds and reciprocally high mass loss is removing a lot of angular momentum. The problem of the ``break--up'' or $\\Omega$--limit is critically examined in connection with the origin of Be and LBV stars. The effects of rotation on the tracks in the HR diagram, the lifetimes, the isochrones, the blue to red supergiant ratios, the formation of W--R stars, the chemical abundances in massive stars as well as in red giants and AGB stars, are reviewed in relation to recent observations for stars in the Galaxy and  Magellanic Clouds. The effects of rotation on the final stages and on the chemical yields are examined, as well as the constraints placed by the periods of pulsars. On the whole, this review points out that stellar evolution is not only a function of mass M and metallicity Z, but of angular velocity $\\Omega$ as well. ",
        "introduction": "Stellar rotation is an example of an astronomical domain which has been studied for several centuries and where the developments are rather slow. A  short historical review since the discovery of the solar rotation by Galileo Galilei is given by Tassoul (1978). Few of the early works apply to real stars, since in general gaseous configurations were  not considered and  no account was given to radiative energy transport. The equations of rotating stars in radiative equilibrium were first considered by Milne (1923), von Zeipel (1924) and Eddington (1925); see also Tassoul (1990) for a more recent history. From the early days of stellar evolution, the studies of rotation and evolution have been closely associated. Soon after the first models showing that Main--Sequence (MS) stars move further into the giant and  supergiant region (Sandage \\& Schwarzschild 1952), rotation was used as a major test for the evolution. Oke \\& Greenstein (1954) and Sandage (1955) found that the observed rotational velocities were  consistent with the proposed evolutionary sequence.  Stellar evolution, as other fields of Science, proceeds using as a guideline the principle of Occam's razor, which says that the explanation relying on the smallest number of hypotheses is usually the one to be preferred. Thus, due to the many well known successes of the theory of stellar evolution, rotation was and is still generally considered as only a second order effect. However, over recent years a number of serious discrepancies between current models and observations have been noticed. They concern particularly the helium and nitrogen abundances in massive O-- and B--type stars and in giants and supergiants, as well as the distribution of stars in the HR diagram at various metallicities. The observations show that the role of rotation has been largely overlooked. All the model outputs (tracks in the HR diagram, lifetimes, actual masses, surface abundances, nucleosynthetic yields, supernova precursors, etc\\dots ) are greatly influenced  by rotation, thus  it turns out that stellar evolution is basically a function of mass M, metallicity Z and angular velocity $\\Omega$.  There are a number of reviews concerning stellar rotation, for example Strittmatter (1969), Fricke \\& Kippenhahn (1972), Tassoul (1978, 1990),  Zahn (1983, 1994) and Pinsonneault (1997). Here, we focus on rotation in Upper MS stars, where the effects are likely the largest ones. The consequences for blue, yellow and red supergiants, Wolf--Rayet (W--R) stars, as well as for red giants and Asymptotic Giant Branch (AGB) stars are also examined. The rotation of low mass stars, where spin--down due to magnetic coupling between the wind and the central body is important, has been treated in a recent review (Pinsonneault 1997); rotation and magnetic activity were also reviewed by Hartmann and Noyes (1987). The role of  rotation in  pre--Main Sequence evolution with accretion disks has been discussed by Bodenheimer (1995). ",
        "conclusions": ""
    },
    "0004/astro-ph0004075_arXiv.txt": {
        "abstract": "This is a review of the physics and cosmology of the cosmological constant.  Focusing on recent developments, I present a pedagogical overview of cosmology in the presence of a cosmological constant, observational constraints on its magnitude, and the physics of a small (and potentially nonzero) vacuum energy. ",
        "introduction": "\\label{section:introduction}  \\subsection{Truth and beauty} \\label{section:truth}  Science is rarely tidy.  We ultimately seek a unified explanatory framework characterized by elegance and simplicity; along the way, however, our aesthetic impulses must occasionally be sacrificed to the desire to encompass the largest possible range of phenomena ({\\it i.e.}, to fit the data). It is often the case that an otherwise compelling theory, in order to be brought into agreement with observation, requires some apparently unnatural modification. Some such modifications may eventually be discarded as unnecessary once the phenomena are better understood; at other times, advances in our theoretical understanding will reveal that a certain theoretical compromise is only superficially distasteful, when in fact it arises as the consequence of a beautiful underlying structure.  General relativity is a paradigmatic example of a scientific theory of impressive power and simplicity.  The cosmological constant, meanwhile, is a paradigmatic example of a modification, originally introduced \\cite{einstein} to help fit the data, which appears at least on the surface to be superfluous and unattractive.  Its original role, to allow static homogeneous solutions to Einstein's equations in the presence of matter, turned out to be unnecessary when the expansion of the universe was discovered \\cite{hubble}, and there have been a number of subsequent episodes in which a nonzero cosmological constant was put forward as an explanation for a set of observations and later withdrawn when the observational case evaporated.  Meanwhile, particle theorists have realized that the cosmological constant can be interpreted as a measure of the energy density of the vacuum.  This energy density is the sum of a number of apparently unrelated contributions, each of magnitude much larger than the upper limits on the cosmological constant today; the question of why the observed vacuum energy is so small in comparison to the scales of particle physics has become a celebrated puzzle, although it is usually thought to be easier to imagine an unknown mechanism which would set it precisely to zero than one which would suppress it by just the right amount to yield an observationally accessible cosmological constant.  This checkered history has led to a certain reluctance to consider further invocations of a nonzero cosmological constant; however, recent years have provided the best evidence yet that this elusive quantity does play an important dynamical role in the universe. This possibility, although still far from a certainty, makes it worthwhile to review the physics and astrophysics of the cosmological constant (and its modern equivalent, the energy of the vacuum).  There are a number of other reviews of various aspects of the cosmological constant; in the present article I will outline the most relevant issues, but not try to be completely comprehensive, focusing instead on providing a pedagogical introduction and explaining recent advances.  For astrophysical aspects, I did not try to duplicate much of the material in Carroll, Press and Turner \\cite{cpt}, which should be consulted for numerous useful formulae and a discussion of several kinds of observational tests not covered here.  Some earlier discussions include \\cite{felten,charlton,sandage88}, and subsequent reviews include \\cite{cohn,sahni,Turner:1998}.  The classic discussion of the physics of the cosmological constant is by Weinberg \\cite{weinberg}, with more recent work discussed by \\cite{cohn,sahni}.  For introductions to cosmology, see \\cite{kt,lindebook,peebles}.  \\subsection{Introducing the cosmological constant} \\label{section:introducing}  Einstein's original field equations are \\begin{equation} R_{\\mu\\nu} - {1\\over 2}Rg_{\\mu\\nu} = 8\\pi GT_{\\mu\\nu}\\ . \\label{einstein} \\end{equation} (I use conventions in which $c=1$, and will also set $\\hbar=1$ in most of the formulae to follow, but Newton's constant will be kept explicit.)  On very large scales the universe is spatially homogeneous and isotropic to an excellent approximation, which implies that its metric takes the Robertson-Walker form \\begin{equation} {\\rm d}s^2 = -{\\rm d}t^2 + a^2(t)R_0^2\\left[ {{{\\rm d}r^2}\\over{1-kr^2}} + r^2 {\\rm d}\\Omega^2\\right]\\ , \\label{rwmetric} \\end{equation} where ${\\rm d}\\Omega^2 = {\\rm d}\\theta^2 + \\sin^2\\theta {\\rm d} \\phi^2$ is the metric on a two-sphere.  The curvature parameter $k$ takes on values $+1$, $0$, or $-1$ for positively curved, flat, and negatively curved spatial sections, respectively. The scale factor characterizes the relative size of the spatial sections as a function of time; we have written it in a normalized form $a(t) = R(t)/R_0$, where the subscript $0$ will always refer to a quantity evaluated at the present time. The redshift $z$ undergone by radiation from a comoving object as it travels to us today is related to the scale factor at which it was emitted by \\begin{equation} a = {1\\over {(1+z)}}\\ . \\end{equation}  The energy-momentum sources may be modeled as a perfect fluid, specified by an energy density $\\rho$ and isotropic pressure $p$ in its rest frame.  The energy-momentum tensor of such a fluid is \\begin{equation} T_{\\mu\\nu} = (\\rho +p)U_\\mu U_\\nu + p g_{\\mu\\nu}\\ , \\label{tmunufluid} \\end{equation} where $U^\\mu$ is the fluid four-velocity. To obtain a Robertson-Walker solution to Einstein's equations, the rest frame of the fluid must be that of a comoving observer in the metric (\\ref{rwmetric}); in that case, Einstein's equations reduce to the two Friedmann equations \\begin{equation} H^2 \\equiv \\left({{\\dot a}\\over a}\\right)^2 = {{8\\pi G}\\over 3}\\rho - {k\\over{a^2R_0^2}}\\ , \\label{feq1} \\end{equation} where we have introduced the Hubble parameter $H\\equiv \\dot a /a$, and \\begin{equation} {{\\ddot a}\\over a} = -{4\\pi G \\over 3}(\\rho + 3p)\\ . \\label{feq2} \\end{equation}  Einstein was interested in finding static ($\\dot a = 0$) solutions, both due to his hope that general relativity would embody Mach's principle that matter determines inertia, and simply to account for the astronomical data as they were understood at the time.\\footnote{This account gives short shrift to the details of what actually happened; for historical background see \\cite{weinberg}.}  A static universe with a positive energy density is compatible with (\\ref{feq1}) if the spatial curvature is positive ($k=+1$) and the density is appropriately tuned; however, (\\ref{feq2}) implies that $\\ddot a$ will never vanish in such a spacetime if the pressure $p$ is also nonnegative (which is true for most forms of matter, and certainly for ordinary sources such as stars and gas).  Einstein therefore proposed a modification of his equations, to \\begin{equation} R_{\\mu\\nu} - {1\\over 2}Rg_{\\mu\\nu} + \\Lambda g_{\\mu\\nu} = 8\\pi GT_{\\mu\\nu}\\ , \\label{einsteinl} \\end{equation} where $\\Lambda$ is a new free parameter, the cosmological constant.  Indeed, the left-hand side of (\\ref{einsteinl}) is the most general local, coordinate-invariant, divergenceless, symmetric, two-index tensor we can construct solely from the metric and its first and second derivatives.  With this modification, the Friedmann equations become \\begin{equation} H^2 = {{8\\pi G}\\over 3}\\rho + {\\Lambda\\over 3} - {k\\over{a^2R_0^2}}\\ . \\label{feq1l} \\end{equation} and \\begin{equation} {{\\ddot a}\\over a} = -{4\\pi G \\over 3}(\\rho + 3p) + {\\Lambda \\over 3} \\ . \\label{feq2l} \\end{equation} These equations admit a static solution with positive spatial curvature and all the parameters $\\rho$, $p$, and $\\Lambda$ nonnegative.  This solution is called the ``Einstein static universe.\"  The discovery by Hubble that the universe is expanding eliminated the empirical need for a static world model (although the Einstein static universe continues to thrive in the toolboxes of theorists, as a crucial step in the construction of conformal diagrams).  It has also been criticized on the grounds that any small deviation from a perfect balance between the terms in (\\ref{feq2l}) will rapidly grow into a runaway departure from the static solution.  Pandora's box, however, is not so easily closed.  The disappearance of the original motivation for introducing the cosmological constant did not change its status as a legitimate addition to the gravitational field equations, or as a parameter to be constrained by observation.  The only way to purge $\\Lambda$ from cosmological discourse would be to measure all of the other terms in (\\ref{feq1l}) to sufficient precision to be able to conclude that the $\\Lambda/3$ term is negligibly small in comparison, a feat which has to date been out of reach.  As discussed below, there is better reason than ever before to believe that $\\Lambda$ is actually nonzero, and Einstein may not have blundered after all.  \\subsection{Vacuum energy} \\label{section:vacuumenergy}  The cosmological constant $\\Lambda$ is a dimensionful parameter with units of (length)$^{-2}$.  From the point of view of classical general relativity, there is no preferred choice for what the length scale defined by $\\Lambda$ might be.  Particle physics, however, brings a different perspective to the question.  The cosmological constant turns out to be a measure of the energy density of the vacuum --- the state of lowest energy --- and although we cannot calculate the vacuum energy with any confidence, this identification allows us to consider the scales of various contributions to the cosmological constant \\cite{zeld,bludman}.  Consider a single scalar field $\\phi$, with potential energy $V(\\phi)$.  The action can be written \\begin{equation} S = \\int d^4x\\, \\sqrt{-g}\\left[ {1\\over 2} g^{\\mu\\nu} \\partial_\\mu\\phi \\partial_\\nu\\phi - V(\\phi)\\right] \\end{equation} (where $g$ is the determinant of the metric tensor $g_{\\mu\\nu}$), and the corresponding energy-momentum tensor is \\begin{equation} T_{\\mu\\nu} = {1\\over 2}\\partial_\\mu\\phi \\partial_\\nu\\phi + {1\\over 2} (g^{\\rho\\sigma}\\partial_\\rho\\phi \\partial_\\sigma\\phi) g_{\\mu\\nu} - V(\\phi)g_{\\mu\\nu}\\ . \\end{equation} In this theory, the configuration with the lowest energy density (if it exists) will be one in which there is no contribution from kinetic or gradient energy, implying $\\partial_\\mu\\phi = 0$, for which $T_{\\mu\\nu}= - V(\\phi_0)g_{\\mu\\nu}$, where $\\phi_0$ is the value of $\\phi$ which minimizes $V(\\phi)$.  There is no reason in principle why $V(\\phi_0)$ should vanish. The vacuum energy-momentum tensor can thus be written \\begin{equation} T^{\\rm vac}_{\\mu\\nu} = -\\rho_{\\rm vac} g_{\\mu\\nu}\\ , \\label{tmunuvac} \\end{equation} with $\\rho_{\\rm vac}$ in this example given by $V(\\phi_0)$. (This form for the vacuum energy-momentum tensor can also be argued for on the more general grounds that it is the only Lorentz-invariant form for $T^{\\rm vac}_{\\mu\\nu}$.) The vacuum can therefore be thought of as a perfect fluid as in (\\ref{tmunufluid}), with \\begin{equation} p_{\\rm vac} = -\\rho_{\\rm vac}\\ . \\end{equation} The effect of an energy-momentum tensor of the form (\\ref{tmunuvac}) is equivalent to that of a cosmological constant, as can be seen by moving the $\\Lambda g_{\\mu\\nu}$ term in (\\ref{einsteinl}) to the right-hand side and setting \\begin{equation} \\rho_{\\rm vac} = \\rho_\\Lambda \\equiv {{\\Lambda}\\over{8\\pi G}}\\ . \\end{equation} This equivalence is the origin of the identification of the cosmological constant with the energy of the vacuum. In what follows, I will use the terms ``vacuum energy\" and ``cosmological constant\" essentially interchangeably.  It is not necessary to introduce scalar fields to obtain a nonzero vacuum energy.  The action for general relativity in the presence of a ``bare'' cosmological constant $\\Lambda_0$ is \\begin{equation} S = {1\\over 16\\pi G}\\int d^4x\\, \\sqrt{-g} (R - 2\\Lambda_0)\\ , \\label{action} \\end{equation} where $R$ is the Ricci scalar.  Extremizing this action (augmented by suitable matter terms) leads to the equations (\\ref{einsteinl}).  Thus, the cosmological constant can be thought of as simply a constant term in the Lagrange density of the theory.  Indeed, (\\ref{action}) is the most general covariant action we can construct out of the metric and its first and second derivatives, and is therefore a natural starting point for a theory of gravity.  Classically, then, the effective cosmological constant is the sum of a bare term $\\Lambda_0$ and the potential energy $V(\\phi)$, where the latter may change with time as the universe passes through different phases.  Quantum mechanics adds another contribution, from the zero-point energies associated with vacuum fluctuations. Consider a simple harmonic oscillator, {\\it i.e.}\\ a particle moving in a one-dimensional potential of the form $V(x)={1\\over 2}\\omega^2 x^2$. Classically, the ``vacuum'' for this system is the state in which the particle is motionless and at the minimum of the potential ($x=0$), for which the energy in this case vanishes. Quantum-mechanically, however, the uncertainty principle forbids us from isolating the particle both in position and momentum, and we find that the lowest energy state has an energy $E_0 = {1\\over 2} \\hbar\\omega$ (where I have temporarily re-introduced explicit factors of $\\hbar$ for clarity). Of course, in the absence of gravity either system actually has a vacuum energy which is completely arbitrary; we could add any constant to the potential (including, for example, $-{1\\over 2} \\hbar\\omega$) without changing the theory. It is important, however, that the zero-point energy depends on the system, in this case on the frequency $\\omega$.  A precisely analogous situation holds in field theory. A (free) quantum field can be thought of as a collection of an infinite number of harmonic oscillators in momentum space. Formally, the zero-point energy of such an infinite collection will be infinite.  (See \\cite{weinberg,cpt} for further details.)  If, however, we discard the very high-momentum modes on the grounds that we trust our theory only up to a certain ultraviolet momentum cutoff $k_{\\rm max}$, we find that the resulting energy density is of the form \\begin{equation} \\rho_{\\Lambda} \\sim \\hbar k_{\\rm max}^4\\ . \\end{equation} This answer could have been guessed by dimensional analysis; the numerical constants which have been neglected will depend on the precise theory under consideration. Again, in the absence of gravity this energy has no effect, and is traditionally discarded (by a process known as ``normal-ordering'').  However, gravity does exist, and the actual value of the vacuum energy has important consequences.  (And the vacuum fluctuations themselves are very real, as evidenced by the Casimir effect \\cite{casimir}.)  The net cosmological constant, from this point of view, is the sum of a number of apparently disparate contributions, including potential energies from scalar fields and zero-point fluctuations of each field theory degree of freedom, as well as a bare cosmological constant $\\Lambda_0$.  Unlike the last of these, in the first two cases we can at least make educated guesses at the magnitudes.  In the Weinberg-Salam electroweak model, the phases of broken and unbroken symmetry are distinguished by a potential energy difference of approximately $M_{\\rm EW} \\sim 200$~GeV (where 1~GeV $= 1.6\\times 10^{-3}$~erg); the universe is in the broken-symmetry phase during our current low-temperature epoch, and is believed to have been in the symmetric phase at sufficiently high temperatures in the early universe.  The effective cosmological constant is therefore different in the two epochs; absent some form of prearrangement, we would naturally expect a contribution to the vacuum energy today of order \\begin{equation} \\rho_\\Lambda^{\\rm EW} \\sim (200~{\\rm GeV})^4 \\sim 3\\times 10^{47} {\\rm ~erg/cm}^3\\ . \\end{equation} Similar contributions can arise even without invoking ``fundamental\" scalar fields.  In the strong interactions, chiral symmetry is believed to be broken by a nonzero expectation value of the quark bilinear $\\bar q q$ (which is itself a scalar, although constructed from fermions). In this case the energy difference between the symmetric and broken phases is of order the QCD scale $M_{\\rm QCD} \\sim 0.3$~GeV, and we would expect a corresponding contribution to the vacuum energy of order \\begin{equation} \\rho_\\Lambda^{\\rm QCD} \\sim (0.3~{\\rm GeV})^4 \\sim 1.6\\times 10^{36} {\\rm ~erg/cm}^3\\ . \\end{equation} These contributions are joined by those from any number of unknown phase transitions in the early universe, such as a possible contribution from grand unification of order $M_{\\rm GUT}\\sim 10^{16}$~GeV. In the case of vacuum fluctuations, we should choose our cutoff at the energy past which we no longer trust our field theory.  If we are confident that we can use ordinary quantum field theory all the way up to the Planck scale $M_{\\rm Pl} = (8\\pi G)^{-1/2} \\sim 10^{18}$~GeV, we expect a contribution of order \\begin{equation} \\rho_\\Lambda^{\\rm Pl} \\sim (10^{18}~{\\rm GeV})^4 \\sim 2\\times 10^{110} {\\rm ~erg/cm}^3\\ . \\label{planckrho} \\end{equation} Field theory may fail earlier, although quantum gravity is the only reason we have to believe it will fail at any specific scale.  As we will discuss later, cosmological observations imply \\begin{equation} |\\rho_\\Lambda^{\\rm (obs)}| \\leq (10^{-12}~{\\rm GeV})^4 \\sim 2\\times 10^{-10} {\\rm ~erg/cm}^3\\ , \\label{obsrho} \\end{equation} much smaller than any of the individual effects listed above.  The ratio of (\\ref{planckrho}) to (\\ref{obsrho}) is the origin of the famous discrepancy of 120 orders of magnitude between the theoretical and observational values of the cosmological constant. There is no obstacle to imagining that all of the large and apparently unrelated contributions listed add together, with different signs, to produce a net cosmological constant consistent with the limit (\\ref{obsrho}), other than the fact that it seems ridiculous.   We know of no special symmetry which could enforce a vanishing vacuum energy while remaining consistent with the known laws of physics; this conundrum is the ``cosmological constant problem''. In section \\ref{section:physics} we will discuss a number of issues related to this puzzle, which at this point remains one of the most significant unsolved problems in fundamental physics. ",
        "conclusions": "Observational evidence from a variety of sources currently points to a universe which is (at least approximately) spatially flat, with $(\\Omega_{\\rm M}, \\Omega_\\Lambda) \\approx (0.3, 0.7)$. The nucleosynthesis constraint implies that $\\Omega_{\\rm B}\\sim 0.04$, so the majority of the matter content must be in an unknown non-baryonic form.  \\begin{figure}[t] \\epsfxsize = 8cm \\centerline{\\epsfbox{omegalvsa.ps}} \\caption{$\\Omega_\\Lambda$ as a function of the scale factor $a$, for a universe in which $\\Omega_{{\\rm M}0}=0.3$, $\\Omega_{\\Lambda 0}=0.7$.  Indicated are the scale factors corresponding to the Planck era, the electroweak phase transition, and Big Bang Nucleosynthesis.} \\label{omegalvsa} \\end{figure}  Nobody would have guessed that we live in such a universe. Figure (\\ref{omegalvsa}) is a plot of $\\Omega_\\Lambda$ as a function of the scale factor $a$ for this cosmology.  At early times, the cosmological constant would have been negligible, while at later times the density of matter will be essentially zero and the universe will be empty.  We happen to live in that brief era, cosmologically speaking, when both matter and vacuum are of comparable magnitude.  Within the matter component, there are apparently contributions from baryons and from a non-baryonic source, both of which are also comparable (although at least their ratio is independent of time).  This scenario staggers under the burden of its unnaturalness, but nevertheless crosses the finish line well ahead of any competitors by agreeing so well with the data.  Apart from confirming (or disproving) this picture, a major challenge to cosmologists and physicists in the years to come will be to understand whether these apparently distasteful aspects of our universe are simply surprising coincidences, or actually reflect a beautiful underlying structure we do not as yet comprehend.  If we are fortunate, what appears unnatural at present will serve as a clue to a deeper understanding of fundamental physics."
    },
    "0004/astro-ph0004243_arXiv.txt": {
        "abstract": "We present PLANET observations of OGLE-1999-BUL-23, a binary-lens microlensing event towards the Galactic bulge. PLANET observations in the I and V bands cover the event from just before the first caustic crossing until the end of the event. In particular, a densely-sampled second caustic crossing enables us to derive the linear limb-darkening coefficients of the source star; $\\cv=0.786^{+0.080}_{-0.078}$ and $\\ci=0.632^{+0.047}_{-0.037}$. Combined analysis of the light curve and the color-magnitude diagram suggests that the source star is a G/K subgiant in the Galactic bulge ($\\Teff\\simeq 4800$ K). The resulting linear limb-darkening coefficient of the source is consistent with theoretical predictions, although it is likely that non-linearity of the stellar surface brightness profile complicates the interpretation, especially for the I band. The global light curve fit to the data indicates that the event is due to a binary lens of a mass ratio $q\\simeq 0.39$ and a projected separation $d\\simeq 2.42$. The lens/source relative proper motion is $(22.8\\pm 1.5)\\ \\kms\\ \\kpc^{-1}$, typical of bulge/bulge or bulge/disk events. ",
        "introduction": "In point-source-point-lens (PSPL) microlensing events, the light curve yields only one physically interesting parameter, the characteristic time scale of the event, $\\te$, which is a combination of the mass of the lens and the source-lens relative parallax and proper motion. However, more varieties than PSPL events have been observed in reality, and using deviations from the standard light curve, one can deduce more information about the lens and the source. The Probing Lensing Anomalies NETwork (PLANET) is an international collaboration that monitors events in search of such anomalous light curves using a network of telescopes in the southern hemisphere \\citep{al98}.  One example of information that can be extracted from anomalous events is the surface brightness profile of the source star \\citep{wi95}. In a binary or multiple lens system, the caustic is an extended structure. If the source passes near or across the caustic, drastic changes in magnification near the caustics can reveal the finite size of the source \\citep{go94, ne94, wi94, al97}, and one can even extract its surface-brightness profile \\citep{bc96, go96, sa97, va98}.  The fall-off of the surface brightness near the edge of the stellar disk with respect to its center, known as limb darkening, has been extensively observed in the Sun. Theories of stellar atmospheres predict limb darkening as a general phenomenon and give models for different types of stars. Therefore, measurement of limb darkening in distant stars other than the Sun would provide important observational constraints on the study of stellar atmospheres. However, such measurements are very challenging with traditional techniques and have usually been restricted to relatively nearby stars or extremely large supergiants. As a result, only a few attempts have been made to measure limb darkening to date. The classical method of tracing the stellar surface brightness profile is the analysis of the light curves of eclipsing binaries \\citep{wi71, tw80}. However, the current practice in eclipsing-binary studies usually takes the opposite approach to limb darkening \\citep{cl98a} -- constructing models of light curves using theoretical predictions of limb darkening. This came to dominate after \\citet{po84} demonstrated that the uncertainty of limb darkening measurements from eclipsing binaries is substantially larger than the theoretical uncertainty. Since the limb-darkening parameter is highly correlated with other parameters of the eclipsing binary, fitting for limb darkening could seriously degrade the measurement of these other parameters. Multi-aperture interferometry and lunar occultation, which began as measurements of the angular sizes of stars, have also been used to resolve the surface structures of stars \\citep{ho98}. In particular, a large wavelength dependence of the interferometric size of a stellar disk has been attributed to limb darkening, and higher order corrections to account for limb darkening have been widely adopted in the interferometric angular size measurement of stars. Several recent investigations using optical interferometry extending beyond the first null of the visibility function have indeed confirmed that the observed patterns of the visibility function contradict a uniform stellar disk model and favor a limb-darkened disk \\citep{qu96, ha98} although these investigations have used a model prediction of limb darkening inferred from the surface temperature rather than deriving the limb darkening from the observations. However, at least in one case, \\citet{bu97} used interferometric imaging to measure the stellar surface brightness profile with coefficients beyond the simple linear model. In addition, developments of high resolution direct imaging in the last decade using space telescopes \\citep{gi96} or speckle imaging \\citep{kl97} have provided a more straightforward way of detecting stellar surface irregularities. However, most studies of this kind are still limited to a few extremely large supergiants, such as $\\alpha$ Ori. Furthermore, they seem to be more sensitive to asymmetric surface structures such as spotting than to limb darkening.  By contrast, microlensing can produce limb-darkening measurements for distant stars with reasonable accuracy. To date, limb darkening (more precisely, a set of coefficients of a parametrized limb-darkened profile) has been measured for source stars in three events, two K giants in the Galactic bulge and an A dwarf in the Small Magellanic Cloud (SMC). MACHO 97-BLG-28 was a cusp-crossing event of a K giant source with extremely good data, permitting \\citet{al99a} to make a two-coefficient (linear and square-root) measurement of limb darkening. \\citet{af00} used data from five microlensing collaborations to measure linear limb darkening coefficients in five filter bandpasses for MACHO 98-SMC-1, a metal-poor A star in the SMC. Although the data for this event were also excellent, the measurement did not yield a two-parameter determination because the caustic crossing was a fold-caustic rather than a cusp, and these are less sensitive to the form of the stellar surface brightness profile. \\citet{al00a} measured a linear limb-darkening coefficient for MACHO 97-BLG-41, a complex rotating-binary event with both a cusp crossing and a fold-caustic crossing. In principle, such an event could give very detailed information about the surface brightness profile. However, neither the cusp nor the fold-caustic crossing was densely sampled, so only a linear parameter could be extracted.  In this paper, we report a new limb-darkening measurement of a star in the Galactic bulge by a fold-caustic crossing event, OGLE-1999-BUL-23, based on the photometric monitoring of PLANET. ",
        "conclusions": ""
    },
    "0004/astro-ph0004133_arXiv.txt": {
        "abstract": "A popular interpretation of recent microlensing studies of the line of sight  towards the Large Magellanic Cloud invokes a population of old white dwarf stars in the Galactic halo. Below I review the basic properties of old white dwarf stars and the ongoing efforts to detect this population directly. ",
        "introduction": "Other authors in this volume cover the microlensing motivations much better than I can, so I shall suffice to remind you that one possible explanation of the microlensing events towards the LMC invokes a population of objects in the range $0.3-0.8$M$_{\\odot}$. Potentially these could be either normal hydrogen-burning stars or white dwarfs, the burnt-out remnants of stellar evolution. To distinguish these populations, we turn to direct searches at optical wavelengths, since the latter population is $10^{-3}-10^{-4}$ as bright as the former. The number counts of faint red stars suggest that hydrogen burning stars cannot account for the microlensing population (Bahcall et al 1994; Graff \\& Freese 1996). The question I wish to address is how well one can constrain the white dwarf hypothesis by similar means. ",
        "conclusions": ""
    },
    "0004/astro-ph0004305_arXiv.txt": {
        "abstract": "The increase in spin frequency as the burning atmospheres of Type~I X-ray bursts cool provides a strong constraint on the radius of the underlying neutron star.  If the change in spin frequency is due to a change in the thickness of the atmosphere, the radius of the star must exceed $3 G M/c^2$ for any equation of state and approximately $3.5 G M/c^2$ for most physically reasonable equations of state.  This constraint arises because the direction of the Coriolis force for radial motion reverses for $R <  G / c^2 ( M +  I / R^2 )$. Furthermore, the marked change in the magnitude of the Coriolis force near compact stars provides a straightforward explanation for why the frequency of the quickly rotating bursters shifts by the same amount as the slow rotators; they are slightly more massive, 1.6~M$_\\odot$ versus 1.4~M$_\\odot$. ",
        "introduction": "If material accretes onto the surface of a neutron star sufficiently slowly, a layer of fuel develops and then suddenly ignites producing a burst of x-rays known as a Type~I X-ray burst.  These thermonuclear flashes each release about $10^{39}$~ergs and repeat on a timescale from hours to days (\\cite{Lewi95,Bild98}).  Although only one source (SAX~J1808.4-3658) exhibits periodic variation in its quiescent emission (\\cite{Wijn98,Chak98}), during the bursts themselves the emission is quasiperiodic (\\eg \\cite{Stro97a}), apparently due to rotational modulation of the inhomogeneities in the thermonuclear burning. \\jcite{1999ApJ...516L..81S} find that the oscillations in the cooling tails of the X-ray bursts from 4U~1702-429 and 4U~1728-34 are nearly coherent with $Q\\sim 4000$ and consistent with an increase in the modulation frequency as the burning layer cools.  \\jcite{Joss78} calculated the first numerical models of thermonuclear burning on the surface of a 1.4~M$_\\odot$ neutron star with a radius of 6.6~km.  He found that the initial surface layer of helium expands from a thickness of about three meters to one of thirty meters during the onset of the burst (also \\cite{Bild95}).  \\jcite{Stro97b} argue that this increase in radius is sufficient to account for the frequency shifts observed during the bursts due to the conservation of angular momentum.  To use the conservation of angular momentum to explain the change in frequency of the emission as the burning region of the atmosphere (or hotspot) expands and then contracts, one assumes that the specific angular momentum of a fluid element is conserved during the burst. The magnetic fields of the neutron stars are likely to be weak (otherwise there would be periodicities in the persistent emission); consequently, this is a viable assumption.  However, the relationship between the specific angular momentum of a fluid element and its position is more complicated in general relativity than in Newtonian theory (\\eg \\cite{AbramowiczMillerStuchlik1993}).  In the case of the Schwarzschild spacetime, the angular momentum of a fluid element is \\be \\d\\! L = \\d m \\Omega r_g^2 = \\d m \\Omega \\frac{r^2}{1-2 M / r} \\label{eq:rgyration} \\ee where $\\Omega$ is the angular velocity of the element as measured by an observer at infinity.  This expression and its generalizations result in the reversal of the centrifugal forces (\\eg \\cite{AbramowiczPrasanna1990,Abramowicz1993}) and gyroscopic precession (\\eg \\cite{1998GReGr..30..593N}) in curved spacetimes.  Rather than using \\eqref{rgyration} to obtain the change in the angular velocity of the hotspot, \\S~\\ref{sec:coriolis} calculates the change the angular velocity in a frame that rotates with the neutron star. The reasons for this treatment are two-fold.  First, the calculation is quite straightforward and illustrative.  Second, the effects of frame dragging are likely to be important for neutron stars, so rather than deriving a new definition for the radius of gyration to include frame dragging, frame dragging may simply be included in the metric.  \\S~\\ref{sec:xrayburst} applies these results to Type~I X-ray bursts, and finally \\S~\\ref{sec:conclus} summarizes the conclusions. ",
        "conclusions": "\\label{sec:conclus}  The Coriolis force in the curved spacetime near a neutron star is substantially smaller than the Newtonian value and reverses its direction as the radius of the star approaches $3 G M/c^2$.  For typical neutron-star parameters of $R=10$~km and $M=1.4$~M$_\\odot$, the fully relativistic Coriolis force is only 30\\% of the Newtonian value.  This contrasts with the assertion of \\jcite{Stro99} that the correction to the Coriolis force is the same order as the gravitational redshift, \\ie 20 -- 30\\% due to time dilation alone. Although time dilation does contribute,  it does not dominate.  Time dilation alone would not predict, the change in the sign of the Coriolis force for $r\\approx 3 M$.  The bulk of the effect arises from the fact that the specific angular momentum of a fluid element has a minimum at the photon circular orbit, which has $r=3 M$ for a non-rotating star.  This also manifests itself through the reversal of the centrifugal force also at $r=3M$ (\\eg\\ \\cite{AbramowiczPrasanna1990,1990MNRAS.245..733A}, \\cite{Abramowicz1993,1996MNRAS.281..659S}) and the fact that compact stars may have moments of inertia which exceed $\\frac{2}{5} M R^2$, the Newtonian value for a sphere with constant density (\\eg\\ \\cite{1967ApJ...150.1005H,1974MNRAS.167...63C}).  Because the Coriolis force sensitively depends on the ratio of the mass to the radius of a relativistic star, Type~I X-ray bursts provide a strong constraint on the equation of state of the underlying neutron star.  Specifically, since the oscillation frequency of the burst increases as the atmosphere cools and presumably shrinks, the radius of the star must exceed $3\\; G M/c^2$.  If frame dragging is included, the constraint is stronger, $R>3.49\\; G M/c^2$. Furthermore, this sensitivity provides a simple explanation for the fact that the quickly rotating bursters, 4U~1636-54, KS~1731-26, Aql X-1 and the galactic center source, spin up by the same amount as the slowly rotating bursters, 4U~1702-43 and 4U~1728-34.  The fast rotators are simply slightly more massive, 1.6~M$_\\odot$ versus 1.4~M$_\\odot$.   Further studies of the thermonuclear burning in the atmospheres of Type~I X-ray bursts will make them a precise probe of the spacetime geometry surrounding rotating neutron stars.  General relativity presents gravity as an inertial force; therefore, it is not surprising that Newtonian notions of inertial forces do not apply in strong gravitational fields.  Gravity strongly affects the Coriolis force which is important in the evolution of Type~I X-Ray bursts; consequently, these bursts provide a strong constraint on the spacetime geometry surrounding accreting neutron stars."
    },
    "0004/cond-mat0004104_arXiv.txt": {
        "abstract": "We consider the Falicov-Kimball model in two dimensions in the neutral case, i.e, the number of mobile electrons is equal to the number of ions. For rational densities between $1/3$ and $2/5$ we prove that the ground state is periodic if the strength of the attraction between the ions and electrons is large enough. The periodic ground state is given by taking the one dimensional periodic ground state found by Lemberger and then extending it into two dimensions in such a way that the configuration is constant along lines at a 45 degree angle to the lattice directions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004149_arXiv.txt": {
        "abstract": "We analyze the three catalogs of nearby loose groups by Garcia (1993). She identified groups in a magnitude--limited redshift galaxy catalog, which covers about $\\sim 2/3$ of sky within $cz =5500$ \\kss, by using two methods, a percolation and a hierarchical method. The free parameters of the group-selection algorithms were tuned to obtain similar catalogs of groups. The author also proposed a third catalog of groups defined as a combination of the two.  Each catalog contains almost 500 groups.  In agreement with previous works on earlier catalogs, we find that groups can be described as collapsing systems. Their sampled size is in general considerably larger than their expected virialized region. We compute the virial masses and correct them by taking into account the young dynamical status of these groups.  We estimate corrected group masses, $M$, for two reference cosmological models, a flat one with a matter density parameter $\\Omega_0=1$ and an open one with $\\Omega_0=0.2$.  For each of the three catalogs we calculate the mass function.  We find that the amplitude of the mass function is not very sensitive to the choice of the group-identification algorithm. The number density of groups with $M> 9 \\times 10^{12}$ \\msunn, which is the adopted limit of sample completeness, ranges in the interval $1.3$--$1.9 \\times 10^{-3} h^{3}Mpc^{-3}$ for $\\Omega_0=1$, and it is about a factor of $15\\%$ lower for $\\Omega_0=0.2$. The mass functions of the hierarchical and combined catalogs have essentially the same shape, while the mass function of the percolation catalog shows a flattening towards large masses.  However, the difference decreases if we do not consider the most massive groups, for which reliable results come from galaxy cluster studies.  After having estimated the mass contained within the central, presumably virialized, regions of groups by adopting a reduction in mass of $\\sim 30$--$40\\%$, we do a comparison with the results coming from the virial analysis of nearby rich clusters (Girardi et al. 1998a).  All three group mass functions turn out to be a smooth extrapolation of the cluster mass function at $M<4 \\times 10^{14}$\\msunn, which is the completeness limit of the cluster sample. The resulting optical virial mass function of galaxy systems, which extends over two orders of magnitude, is fitted to a Schechter expression with a slope of $\\sim -1.5$ and a characteristic mass of $M^*\\sim 3\\times 10^{14}$\\msunn.  We also verify that our group mass function reasonably agrees with the Press--Schechter predictions of models which at large masses describe the virial mass function of clusters.   \\vspace*{6pt}  \\noindent {\\em Subject headings: } galaxies: clusters: general - cosmology: observations - cosmology: theory - large scale structure of universe. ",
        "introduction": "Most galaxies in the local universe belong to loose galaxy groups. Groups seem to be the natural continuation of galaxy clusters at smaller mass scales. Indeed, there is a continuity of properties from rich clusters to poor clusters and to groups (e.g., Ramella, Geller, \\& Huchra 1989; Burns et al. 1996; Mulchaey \\& Zabludoff 1998; Ramella et al. 1999; Girardi, Boschin, \\& da Costa 2000).  Zabludoff \\& Mulchaey (1998) used multi--fiber spectroscopy to obtain velocities for a large number of group members (i.e. 280 galaxies for a total of 12 groups) and Mahdavi et al. (1999) measured several hundreds of redshifts to obtain a sample of 20 groups, each one having , on average, 30 galaxies .  For these well--sampled groups Zabludoff \\& Mulchaey (1998) and Mahdavi et al. (1999) performed refined analyses, i.e. the rejection of interlopers, the study of the internal galaxy distribution and velocity dispersion profiles, and the separation of different galaxy populations (see, e.g.,  Biviano et al. 1997; Carlberg et al. 1996, 1997b; den Hartog \\& Katgert 1996; Dressler et al. 1999; Girardi et al. 1996, 1998b; Mohr et al. 1996; Koranyi \\& Geller 2000; for recent relevant results on rich and poor clusters).  However, all these analyses are so far restricted to a limited number of groups since they require a strong observational effort. Therefore, to analyze group dynamical properties in a statistical sense, one must resort to wide group catalogs where groups are extracted from three--dimensional galaxy catalogs and typically contain $\\lesssim$ five member galaxies (e.g., Huchra \\& Geller 1982; Tully 1987; Ramella et al. 1999).  Here, we focus our attention on the determination of group mass function from wide catalogs of nearby loose groups.  The observational determination of group mass function is plagued by several problems. Some of them concern the estimate of mass and are mainly due to the small number of group members and to uncertainties in the dynamical stage.  In fact, although group cores are virialized or close to virialization (Zabludoff \\& Mulchaey 1998), the size of groups identified in three--dimensional galaxy catalogs, i.e. $\\sim 0.5$--$1$ \\hh, is appreciably greater than their expected virialized region, i.e. $\\sim 0.2$--0.4 \\h for systems with a line--of--sight velocity dispersion of 100--200 \\ks (according to the relations found for galaxy clusters, e.g. Carlberg, Yee, \\& Ellingson 1997; Girardi et al. 1998b).  Indeed, there is a strong indication that these groups are not virialized systems over the whole sampled region, but can be rather described as being in a phase of collapse (e.g., Giuricin et al. 1988; Mamon 1994; Diaferio et al. 1993).  Therefore, usual estimates of velocity dispersion and virial mass are not easily connected to physical quantities such as group potential and mass.  The small number of data and the uncertainties on dynamical status prevent one to use refined methods to reject interlopers in each individual group (e.g., Zabludoff \\& Mulchaey 1998, Mahdavi et al. 1999).  Instead, one must rely on member galaxies as assigned by the group-selection algorithm, while checking a posteriori the presence of spurious groups in a statistical sense (e.g.,  Ramella, Pisani, \\& Geller 1997; Diaferio et al. 1999). Indeed the results could depend on the choice of the group-selection algorithm and its free parameter (e.g., Pisani et al. 1992 -- hereafter P92; Ramella et al. 1997).  For instance, Frederic' s (1995b) analysis of cosmological N--body simulations suggested that the estimated median mass depends on the algorithm and that the resulting bias is sensitive to the depth of the galaxy survey. However, even the analysis of simulated groups is not an easy task and, indeed, the results on mass can depend on the treatment of halos (cf. Frederic 1995b).  A further uncertainty is connected to cosmic variance.  In fact, group catalogs are recovered from local galaxy catalogs which may not be fair samples of the universe.  In view of these difficulties, few statistical distributions of group dynamical properties are available in the literature and they are often discrepant.  The cumulative distributions of internal velocity--dispersion, as computed by Moore, Frenk, \\& White (1993) and by Zabludoff et al.  (1993), are strongly discrepant (the number densities of groups with line--of--sight velocity dispersion larger than 200 \\ks differ by a factor of 100, see Fig.~6 of Fadda et al. 1996 for a comparison). Moreover, analyzing nearby groups ($cz\\le 2000$ \\kss) of three different group catalogs, P92 found a significant dependence of the distribution of mass and other dynamical group parameters on the group-identification algorithm.  The availability of new group catalogs has prompted us to derive a new group mass function, whose connection with the recent determination of the optical virial mass function of nearby rich galaxy clusters (Girardi et al. 1998a, hereafter G98) deserves to be investigated.   The work by Garcia (1993, hereafter G93), who constructed two group catalogs using two different identification algorithms (the percolation and hierarchical ones) and proposed a third catalog which is a combination of the two, represents a good data base for facing the effect of identification algorithms.  So far, G93 catalogs are the largest catalogs of groups presently published. They are largely superior to those analyzed by P92 both for the number of groups (450--500 groups for each of the three catalogs) and the encompassed volumes ($\\sim 2/3$ of sky, $cz\\le 5500$ \\kss).  Moreover, these group catalogs were selected from the same parent galaxy sample, thus allowing us to better investigate on the effects due to differences in the selection algorithm.  Furthermore, the improved statistics in the high-mass range (less than ten groups analyzed by P92 have masses larger than $10^{14}$\\msunn) permits an interesting comparison with cluster mass function and a determination of the virial mass function over an unprecedently large range of masses.  In \\S~2 we briefly describe the data. In \\S~3 we calculate group masses. In \\S~4 and 5 we compute group mass function and verify its stability, respectively. In \\S~6 we compare the results of groups and clusters, recovering the mass function of galaxy systems for a mass range which extends over two orders of magnitude. In \\S~7 we give our discussions.  In \\S~8 we summarize our results and draw our conclusions.  Throughout the paper, errors are given at the $68\\%$ confidence level and the Hubble constant is $H_0=100\\ \\rm{h}\\ $Mpc$^{-1}$ \\kss. ",
        "conclusions": ""
    },
    "0004/astro-ph0004239_arXiv.txt": {
        "abstract": "We report the discovery of a bright (J = 13.83$\\pm$0.03) methane brown dwarf, or T dwarf, by the Two Micron All Sky Survey. This object, 2MASSI J0559191-140448, is the first brown dwarf identified by the newly commissioned CorMASS instrument mounted on the Palomar 60-inch Telescope.  Near-infrared spectra from 0.9 - 2.35 $\\micron$ show characteristic CH$_4$ bands at 1.1, 1.3, 1.6, and 2.2 $\\micron$, which are significantly shallower than those seen in other T dwarfs discovered to date. Coupled with the detection of an FeH band at 0.9896 and two sets of K I doublets at J-band, we propose that 2MASS J0559-14 is a warm T dwarf, close to the transition between L and T spectral classes.  The brightness of this object makes it a good candidate for detailed investigation over a broad wavelength regime and at higher resolution. ",
        "introduction": "T dwarfs are brown dwarfs that exhibit methane absorption bands at 1.6 and 2.2 $\\micron$ \\citep{Ki99a}, and thus have effective temperatures T$_{eff}$ $\\lesssim$ 1200-1300 K \\citep{Fe96,Bu99,Ki00}.  The prototype for this class, Gl 229B \\citep{Na95,Op99}, was identified as a cool companion to the nearby M1V star Gl 229A. Recently, seven field objects \\citep{Ss99,Bg99,Cu99,Ts00} and another companion object \\citep{Bg00a} have also been identified as T dwarfs. The rapid discovery of these cool brown dwarfs has been driven by new sky surveys, such as the Two Micron All Sky Survey \\citep[hereafter 2MASS]{Sk97} and the Sloan Digital Sky Survey \\citep{Gu95}; and deep near-infrared surveys, such as the ESO New Technology Telescope Deep Field \\citep{Ar99}.  The T dwarfs identified to date are remarkably similar to Gl 229B, with colors in the range $-$0.2 $\\lesssim$ J-K$_s$ $\\lesssim$ 0.2. Near-infrared spectra are correspondingly similar \\citep{Ss99,Bg00a}, likely due to the saturation of H$_2$O and CH$_4$ bands that dominate this wavelength regime. Subtle differences in the magnitudes and shapes of H- and K-band flux peaks are observed, due to increased CH$_4$, H$_2$O, and H$_2$ collision-induced absorption (CIA) toward cooler effective temperatures \\citep{Bg99,Ts00}, and variations in the depths of near-infrared H$_2$O and CH$_4$ bands are discerned when compared to Gl 229B \\citep{Na00,Bg00a}. Nonetheless, the similarity of the near-infrared spectra suggests that either the objects thus far identified are very similar in temperature, around 1000 K \\citep{Ma96}, or that near-infrared features are fairly insensitive to temperature, making the definition of a T dwarf spectral sequence in this wavelength regime a difficult proposition, at least at low resolution.  We report the discovery of a T dwarf by 2MASS which is unique among its counterparts, as it shows significant differences in its near-infrared features while retaining defining CH$_4$ bands.  This object, 2MASSI J0559191-140448 (hereafter 2MASS J0559-14), is also 0.4 mag brighter than Gl 229B and more than 1 mag brighter than the field T dwarfs discovered thus far.  It is the first brown dwarf to be identified by the newly commissioned Cornell Massachusetts Slit Spectrograph \\citep[hereafter CorMASS]{Wi00}, mounted on the Palomar 60-inch telescope.  In $\\S$2 we discuss the selection of 2MASS J0559-14 from 2MASS data and its spectral identification by CorMASS.  In $\\S$3 we discuss the observed spectral features and argue that 2MASS J0559-14 is a warm T dwarf, possibly close to the transition temperature between L and T spectral classes. We discuss the brightness of this object and its role in future spectroscopic investigation of the T dwarf class in $\\S$4. ",
        "conclusions": "\\subsection{Spectral Features}  Table 2 summarizes the spectral features detected in 2MASS J0559-14, with identifications from \\citet{Pg63}, \\citet{Da66}, \\citet{Ws66}, and \\citet{Ph87}. Only major absorption bands of H$_2$O and CH$_4$ are tabulated.  The characteristic CH$_4$ bands at 1.6 and 2.2 $\\micron$ are present, as are bands at 1.1 and 1.3 $\\micron$ identified from laboratory data \\citep{Fi79} which are blended with H$_2$O bands at the same wavelengths. An FeH feature is seen at 0.9896 $\\micron$ (0-0 band of A$^4$$\\Delta$-X$^4$$\\Delta$) which has also been identified in SDSS 1624+00 \\citep{Bg00b}. We do not detect the higher order 0-1 FeH band at 1.191 $\\micron$, which is seen to weaken in the latest L dwarfs \\citep{Mc00}. Two sets of K I doublets are noted at 1.1690 \\& 1.1773 $\\micron$ (4p $^2$P$_0$ $-$ 3d $^2$D) and 1.2432 \\& 1.2522 $\\micron$ (4p $^2$P$_0$ $-$ 5s $^2$S), as is Cs I at 0.8943 $\\micron$ (6s $^2$S$_{1/2}$ $-$ 6p $^2$P$_{1/2}$). We do not detect the 2-0 X$^1$$\\Sigma$$^+$-X$^1$$\\Sigma$$^+$ band of CO at 2.3 $\\micron$.  Comparison between 2MASS J0559-14 and SDSS 1624+00 reveals significant differences in spectral morphology. In Figure 2, it is apparent that the slope between 0.9 and 1.05 $\\micron$ is shallower in 2MASS J0559-14, likely due to decreased absorption by the pressure-broadened K I doublet at 0.7665 and 0.7699 $\\micron$ \\citep{Li00}. CH$_4$ and H$_2$O features are generally weaker in 2MASS J0559-14, as noted by the significantly weakened 1.1 - 1.2 $\\micron$ trough between $z$- and J-band peaks. The decreased CH$_4$ opacity noticeably affects the shape of the J-band peak near 1.27 $\\micron$, as the weak CH$_4$ wings at 1.24 and 1.30 $\\micron$ carve out less flux on either side of the peak. In Figure 2, a significant flux offset is readily apparent at the base of the 1.6 $\\micron$ CH$_4$ band in 2MASS J0559-14, and relative enhancement of flux at both H- and K-bands in this object is almost certainly due to decreased H$_2$ CIA (1-0 Quadrupole), H$_2$O, and CH$_4$ opacity.  \\placefigure{fig-2} \\placefigure{fig-3} \\placetable{tbl-2}  \\subsection{2MASS J0559-14 is a Warm T dwarf}  The weak CH$_4$ bands seen in 2MASS J0559-14 are unique among the current sample of T dwarfs, and the reduced opacity can be most readily explained if this object is warmer than other known T dwarfs.  At higher effective temperatures, the dominant carbon-bearing species changes from CO to CH$_4$ at small optical depth, so that the CH$_4$ column density will be less than that of cooler T dwarfs, and observed band strengths correspondingly weaker. Increased thermal flux will also be seen at the base of these bands, particularly at 1.6 $\\micron$, which is unaffected by H$_2$O absorption. The lower CH$_4$ column density directly affects the H$_2$O column density, via the reaction CO + 3 H$_2$ $\\rightarrow$ CH$_4$ + H$_2$O, leading to shallower bands at 1.1 and 1.45 $\\micron$. Water can also be heated and dissociated by dust layers deep in the photosphere \\citep{Le98}. Finally, decreased H$_2$ CIA opacity at K-band, congruous with reduced CH$_4$ and H$_2$O opacity, will result in redder J-K$_s$ colors with warmer T$_{eff}$. These features are observed in 2MASS J0559-14, and its warm temperature is independently supported by the detection of the 0.9896 $\\micron$ FeH band, which is present but weakening in the latest L dwarfs \\citep{Ki99a}.  A similar argument has been made for SDSS 1624+00 by \\citet{Bg00b}, which was one of only three T dwarfs in that paper to show this feature. SDSS 1624+00 has been shown to be a warm object via optical continuum measurements between broadened Na I and K I features \\citep{Li00}, while \\citet{Na00} argue that this object is both warmer and dustier than Gl 229B based on its shallower CH$_4$ and H$_2$O bands. By analogy, 2MASS J0559-14 should be warmer still.  The detection of excited K I lines at J-band, which are seen to weaken in the latest L dwarfs \\citep{Mc00}, is further evidence of the warmth of this object.  The spectral features in 2MASS J0559-14 suggest that it is close to the L/T transition temperature. The lack of CO detection at 2.3 $\\micron$ is not necessarily contradictory to this hypothesis, as overlying CH$_4$ absorption beyond 2.2 $\\micron$ may mask this weaker feature. Naturally, metallicity, dust, and gravity could also play roles in the band strengths seen in 2MASS J0559-14; however, temperature is likely to be the dominant determinant given the concordance of spectral features as discussed above. We can make a conservative T$_{eff}$ constraint for this object based on the temperature of Gl 229B, which is clearly cooler, and a temperature estimate of the L8V companion dwarf Gl 584C \\citep{Ki00}; this translates into a range of 1000 $\\lesssim$ T$_{eff}$ $\\lesssim$ 1300 K. Parallax and bolometric luminosity measurements of this object would allow a direct determination of temperature."
    },
    "0004/astro-ph0004054_arXiv.txt": {
        "abstract": "In very crowded fields, the modulation of the background by the sea of unresolved faint sources induces centroid shifts.  The errors increase with the number of sources per beam.  Even the most optimistic simulations of imaging data show that position errors can become severe (on the order of the beam size) at flux levels at which images contain 1/50 to 1/15 sources per beam, depending on the slope of the number-flux relation $d\\log N/d\\log S$.  These problems are expected to be significant for recent observations of faint submillimeter sources and may be the reason that some sources appear to lack optical counterparts. ",
        "introduction": "Fainter is (usually) better when it comes to star and galaxy counts. However, there are fundamental limits to faint imaging from confusion which cannot be overcome by increasing exposure times alone.  The sea of unresolved sources fainter than the detection limit creates a noise in the sky, which does not improve with more data.  Many of the issues associated with this confusion noise have been discussed before (eg, Scheuer 1957; Condon 1974; Franceschini 1982; Hacking \\& Houck 1987; Barcons 1992), however the large number of present-day observations that are or soon will be pushing this confusion limit suggests a new discussion.  In particular, in recent years there has been a concerted effort to produce very deep, multi-wavelength studies of blank sky in order to identify extragalactic sources as comprehensively as possible.  These studies have been very successful, identifying populations of radio-, submillimeter-, infrared-, visual- and x-ray-bright galaxies and associating them with their counterparts in other bands (Djorgovski et al 1995; Williams et al 1996; Hogg et al 1996; Rowan-Robinson et al 1997; Richards et al 1998; Hughes et al 1998; Barger et al 1998; Eales et al 1999; Aussel et al 1999; Elbaz et al 1999; Gardner et al 2000; Brandt et al in preparation; Dickinson et al in preparation).  Some of the faintest sources in some of the most crowded fields (in the sense of number of sources per resolution element) have not shown clearly distinguished counterparts at other wavelengths (Hughes et al 1998; Smail et al 1998; Barger et al 1999b).  This raises the question ``could confusion be playing a role?''  This paper is a first attempt at characterizing position shifts due to confusion in astronomical images.  Simulated images of crowded fields are presented, made in the most optimistic way: no photon noise, a perfectly understood gaussian point spread function or beam shape, pointlike sources, a power-law number--flux relation of known slope, and no angular clustering.  Even with these optimistic inputs, the resulting images show that it is impossible to accurately measure positions and fluxes of sources that are more than an order of magnitude brighter than the flux level corresponding to one source per beam (a beam being one resolution element in the image).  Recent work making use of the limit ``one source per beam'' (eg, Blain et al 1998) is therefore overly optimistic.  The standard rule-of-thumb is that confusion becomes important at 1/30 of a source per beam.  It is possible to get information from the statistics of the background noise fainter than the level of 1/30 source per beam (eg, Scheuer 1957; Condon 1974), but in terms of identifying and measuring individual sources, 1/30 is regarded as the limit.  This paper tests the rule-of-thumb for confusion-induced astrometry errors, which are particularly important for deep, multi-wavelength studies, in which counterpart identification across multiple data sets is important.  Astrometric shifts due to confusion have been predicted and observed in the context of microlensing data (Goldberg 1998; Goldberg \\& Wo\\'zniak 1998) and are expected to limit future stellar astrometry experiments (Yu et al 1993; Rajagopal \\& Allen 1999).  For the purposes of this paper a ``beam'' is taken to be the solid angle of the $1\\,\\sigma$-radius circle of the gaussian point spread function, or $\\Omega_{\\rm beam}=\\pi\\,\\sigma^2$.  Note that for a gaussian, $\\sigma\\approx \\theta_{\\rm FWHM}/2.35$ where $\\theta_{\\rm FWHM}$ is the angular full width at half-maximum of the point spread function.  The number of sources per beam $s/b$ at a given flux level $S$ is the integrated number of sources $N(>S)$ in an image brighter than flux $S$ divided by the number of beams in the solid angle of the image or $(\\Omega_{\\rm image}/\\Omega_{\\rm beam})$. ",
        "conclusions": "For typical faint imaging in the visual and near-infrared, in which number counts have the form $d\\log N/d\\log S=S^{-\\beta}$ with $\\beta\\approx 0.75$, the confusion limit rule-of-thumb that imaging should not be pursued much fainter than $s/b\\sim 1/30$ sources per beam is essentially correct, both for obtaining good positions and good photometry.  Optimistic simulations show that positions and fluxes of sources more numerous than this condition are likely to have large uncertainties.  When number counts are steep, with Euclidean $\\beta=1.5$ or steeper, the problem is even more severe and a better rule-of-thumb is something like $s/b\\sim 1/50$.  Source identifications in one set of imaging data based on detections in another set will be affected by these confusion-induced astrometry errors.  It is essential that surveys working near the confusion limit perform realistic simulations (which include the sources well faint of any detection limits) in order to draw conservative positional error boxes for source identification."
    },
    "0004/astro-ph0004324_arXiv.txt": {
        "abstract": "The primary difficulty with using transits to discover extrasolar planets is the low probability a planet has of transiting its parent star.  One way of overcoming this difficulty is to search for transits in dense stellar fields, such as the Galactic bulge.  Here I estimate the number of planets that might be detected from a monitoring campaign toward the bulge.  A campaign lasting 10 nights on a 10 meter telescope (assuming 8 hours of observations per night and a 5'x5' field of view) would detect about $100$ planets with radius $\\rp=1.5~\\rjup$, or about $30$ planets with $\\rp=1.0~\\rjup$, if the frequency and distribution of planets in the bulge is similar to that in the solar neighborhood.  Most of these planets will be discovered around stars just below the turn-off, i.e. slightly evolved G-dwarfs.  Campaigns involving 1- or 4-m class telescopes are unlikely to discover any planets, unless there exists a substantial population of companions with $\\rp > 1.5~\\rjup$. ",
        "introduction": "The search for extrasolar planets has garnered enormous attention in recent years, due primarily to the successful implementation of radial velocity searches (\\cite{mandq1995}, \\cite{mandb1996}).  These searches have led to the discovery of a population of massive, close-in planets with orbital separations of $a \\la 0.1~\\au$.  Recently, it was discovered that one such planet, the companion to HD 209458, also transits its parent star (\\cite{charbon2000}; \\cite{henry2000}), yielding a measurement of the mass, radius, and density of the companion.  Clearly, transit observations can be used to extract additional information about known companions.  The {\\it discovery} of an extrasolar planet using transits, however, has remained elusive.  There are two primary difficulties with detecting planets with transits.  First, the photometric requirements are quite stringent: a planet of radius $\\rp\\le\\rjup$ (where $\\rjup$ is the radius of Jupiter) transiting an primary of radius $\\rs=\\rsun$ would produce a fractional deviation of $\\la 1\\%$ during the course of the transit.  Second, the probability that a planet will transit its parent is small: for a planet with separation $\\ge 0.05~\\au$ orbiting a star with $\\rs=\\rsun$, the probability is $\\la 10\\%$.  Several methods for dealing will the small probability have been proposed.  For instance, one can monitor eclipsing binary stars, where the orbital plane is known to be (nearly) perpendicular to the sky (\\cite{deeg}).  Another way of overcoming this small probability is to simply monitor many stars simultaneously.  This can be done by employing a camera with a large field-of-view, or by monitoring very dense stellar fields. Here I focus on the latter possibility.  Specifically I determine the number of planets that might be detected in a campaign monitoring stars toward the Galactic bulge. ",
        "conclusions": "In this Letter, I have estimated the number of planets that may be detected by transits in a monitoring campaign toward the Galactic bulge.  An investment of a relatively modest amount of telescope resources, 10 clear nights of 8 hours per night on a 10m telescope at a site with excellent ($0.75''$) median seeing, would result in the detection of $\\sim 30$ planets of Jupiter size, if the frequency and distribution of planetary companions to stars in the Galactic bulge is similar to those of G-dwarfs in the solar neighborhood.  Most of these planets will be found at orbital separations of $a\\sim 0.02~\\au$ around stars slightly fainter than the turn-off, i.e.\\ evolved G or early K dwarfs.  Modifications to the observing strategy, such as decreasing the number of nights to 5 instead of 10, will not result in substantially fewer detections.  However, if the seeing is substantially worse than $0.75''$, the number of detections will be considerably smaller.  Therefore an excellent site is required. Similar campaigns involving 1m- or 4m-class telescopes are unlikely to result in any detections toward the bulge. Thus, collaborations currently monitoring the Galactic bulge for microlensing events are unlikely to serendipitously detect any transits."
    },
    "0004/astro-ph0004112_arXiv.txt": {
        "abstract": "We present the first unambiguous case of external variability of a radio gravitational lens, CLASS B1600+434. The {\\sl Very Large Array} (VLA) 8.5--GHz difference light curve of the lensed images, taking the proper time-delay into account, shows the presence of external variability with 14.6--$\\sigma$ confidence.  We investigate two plausible causes of this external variability: scattering by the ionized component of the Galactic interstellar medium and microlensing by massive compact objects in the bulge/disk and halo of the lens galaxy. Based on the tight relation between the modulation-index (fractional rms variability) and variability time scale {\\sl and} the quantitative difference between the light curves of both lensed images, we conclude that the observed short-term variability characteristics of the lensed images are incompatible with scintillation in our Galaxy. This conclusion is strongly supported by multi--frequency {\\sl Westerbork Synthesis Radio Telescope} (WSRT) observations at 1.4 and 5\\,GHz, which are in disagreement with predictions based on the scintillation hypothesis.  Several arguments against scintillation might need to be reevaluated if evidence is found for significant scatter-broadening of lensed image B seen through the lens galaxy. However, the frequency-dependence and time scale of variability from image A are not affected by this and remain strong arguments against scintillation.  On the other hand, a single superluminal jet-component in the source, having an apparent velocity 9$\\la$$(v_{\\rm app}/c)$$\\la$26, a radius of 2--5\\,$\\mu$as and containing 5--11\\% of the observed 8.5--GHz source flux density, can reproduce the observed modulation-indices and variability time scale at 8.5\\,GHz, when it is microlensed by compact objects in the lens galaxy. It also reproduces the frequency-dependence of the modulation-indices, determined from the independent WSRT 1.4 and 5--GHz observations. The difference between the modulation-indices of the lensed images (i.e. 2.8\\% and 1.6\\% at 8.5\\,GHz in 1998 for images A and B, respectively), if not affected by scatter-broadening of image B by the ionized ISM of the lens galaxy, can be explained through a different mass-function for the compact objects in the bulge/disk and halo of the lens galaxy. Comparing the observations with microlensing simulations, we place a tentative lower limit of $\\ga$0.5\\,M$_\\odot$ on the average mass of compact objects in the halo line-of-sight. The above-mentioned set of mass-function and source parameters is consistent, although not unique, and should only be regarded as indicative.  The only conclusion fully consistent with the data gathered thus far is that we have indeed detected {\\sl radio microlensing}. The far reaching consequence of this statement is that a significant fraction of the mass in the dark--matter halo at $\\sim$6\\,kpc ($h$=0.65) above the lens--galaxy disk in B1600+434 consists of massive compact objects. ",
        "introduction": "\\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics{fig1_9329.ps}} \\hfill \\parbox[b]{\\hsize}{ \\caption{The normalized light curves of B1600+434 A (upper) and B (lower), corrected for a long-term gradient (Sect.\\,2). The error on each light-curve epoch is 0.7 to 0.8\\% (1--$\\sigma$). Day 0 corresponds to 1998 Febr. 13} \\label{normlc}} \\end{figure*}   Gravitationally lensed compact radio sources have many astrophysical and cosmological applications.  The foremost being the determination of a time-delay between the individual lensed images in order to constrain the Hubble parameter (e.g. Refsdal 1964). Considerable progress has been made during the last few years in measuring time-delays, both through optical and radio observations (e.g Kundi\\'c et al. 1997; Schechter et al. 1997; Lovell et al. 1998; Biggs et al. 1999; Fassnacht et al. 1999; Koopmans et al. 2000). They also allow a detailed study of the mass distribution of the lens galaxy and sometimes the background source, through a large magnification by the lensing potential.  Absorption lines in the spectrum of the background source allow the study of the ISM in the lens galaxy and the HI distribution along the lines of sight to the source.  Temporal changes in the brightnesses or spectra of the lensed images also allow the study of uncorrelated external variability.  The most important sources of external variability are scintillation at radio wavelengths and microlensing in all wavelength bands. Differencing the image light curves, taking the proper time delay into account, removes intrinsic source variability and leaves only uncorrelated external variability. These difference light curves thus provide valuable information on the compact objects in the lens galaxy (e.g. Schmidt \\& Wambsganss 1998) and/or on the intervening ionized medium.  The study of the ionized component of the Galactic interstellar medium (ISM) through scattering of radio emission from pulsars has had a long tradition (e.g. Rickett 1977, 1990). Scattering by the ionized ISM can explain long-term variability at meter wavelengths (e.g. Condon et al. 1979), as well as large-amplitude variability in very compact extra-galactic radio sources (e.g. Rickett, Coles \\& Bourgois 1984). Low-amplitude variability at shorter wavelengths (about 10\\,cm), called `flickering', has been observed by Heeschen (1982, 1984) and is probably associated with refractive interstellar scattering of an extended source (e.g. Rickett et al. 1984).  Strong intra-day variability of very compact radio sources might result from refractive interstellar scattering as well (e.g. Wagner \\& Witzel 1995).  A power-law model of the plasma-density power spectrum (e.g. Rickett 1977, 1990), combined with some distribution of this plasma in our galaxy (e.g. Taylor \\& Cordes 1993 [TC93]) is able to explain most of the observed dispersion measures and variability in pulsars at low frequencies, as well as the variability of extra-galactic radio sources at both low and high frequencies. However, especially for compact flat-spectrum radio sources it remains exceedingly difficult to separate intrinsic variability from scintillation by the Galactic ionized ISM.  Gravitationally lensed (i.e. multiply-imaged) flat-spectrum compact radio sources could offer a solution to this problem. As mentioned previously, these systems provide two or more lines-of-sight through the Galactic ionized ISM. For typical image separations of a few arcseconds, one is looking through parts of the Galactic ionized ISM separated by a few hundred AU. One can expect the scattering of radio waves to be independent for the different lines-of-sight. Differencing the image light curves, after a correction for the appropriate time delay and flux-density ratio, produces a difference light curve that only shows uncorrelated external variability. This difference light curve can be studied to obtain information on the Galactic ionized ISM independent from intrinsic source variability.  However, uncorrelated external variability of the lensed images might also originate from microlensing in the lens galaxy (e.g. Chang \\& Refsdal 1979). This offers the additional opportunity to study the properties of compact objects in the lens galaxy, if microlensing variability dominates or can be separated from scintillation.  Optical microlensing in the lens galaxy of Q2237+0305 has unambiguously been shown (e.g. Irwin et al. 1989; Corrigan et al.  1991; Ostensen et al. 1996; Lewis et al. 1998; Wo\\'zniak et al. 2000). In the radio, several suggestions of microlensing variability have been made (e.g. Stickel et al. 1988; Quirrenbach et al. 1989; Schramm et al. 1993; Romero et al. 1995; Chu et al. 1996; Wagner et al. 1996; Lewis \\& Williams 1997; Takalo et al. 1998; Quirrenbach et al. 1998; Kraus et al. 1999; Watson et al. 1999). In none of these cases, however, has one really been able to convincingly distinguish between intrinsic and external variability. Claims of external variability in singly-imaged radio sources through microlensing should therefore be regarded with some caution.  In this paper, we report the first unambiguous case of external variability of a radio gravitational lens, CLASS B1600+434 (Jackson et al. 1995; Jaunsen \\& Hjorth 1997; Koopmans, de Bruyn \\& Jackson 1998 [KBJ98]; Koopmans et al. 2000 [KBXF00]). The system consists of two compact flat-spectrum radio images, separated by 1.4\\,arcsec.  The background source, at a redshift of $z$=1.59, is lensed by an edge-on disk galaxy at a redshift of $z$=0.41 (Fassnacht \\& Cohen 1998). A time delay of $47^{+12}_{-9}$\\,days (95\\% statistical confidence) was recently found (KBXF00).  What is furthermore of interest is that this system offers two distinct lines-of-sight through the lens galaxy. Image~A only passes mainly through the dark-matter halo around the edge--on lens galaxy, whereas image~B passes predominantly through its disk and bulge (Koopmans et al. 1998; Maller et al. 2000; CASTLE Survey, Munoz et al. 1999).  This makes image~A especially sensitive to microlensing by massive compact objects in the halo and image~B to microlensing by stars in the disk and bulge. This might even offer an opportunity to study compact objects in the dark--matter halo around the lens galaxy of B1600+434.  \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics{fig2_9329.ps}} \\hfill \\parbox[b]{\\hsize}{ \\caption{The normalized difference light curve between the two lensed images, corrected for both the time-delay and flux density ratio (Sect.\\,2.1).  The shaded region indicates the expected 1--$\\sigma$ (1.1\\%) region if all variability were due to measurement errors. The dash--doted lines indicate the observed modulation-index of 2.8\\%. The dotted and dashed curves indicate the normalized difference curves for a time delay of 41 and 52\\,days, respectively. Obviously most variations can not be explained by any reasonable error in the time delay.  There seems to be no evidence of any change in the typical time scale of variability over 6 months. Day 0 corresponds to 1998 Febr. 13.} \\label{normdiff}} \\end{figure*}  The outline of the paper is as follows. In Section 2, we present the VLA 8.5--GHz data from KBXF00 in a different way, unambiguously showing the presence of external variability. We also present additional WSRT 1.4 and 5--GHz monitoring data of B1600+434. In Section 3, we investigate whether Galactic scintillation can explain the fractional rms variabilities (modulation-indices) and time scales of the short-term variability seen in the VLA 8.5--GHz light curves. Similarly, in Sections 4 and 5 the possibility of microlensing by compact objects in the lens galaxy is studied. In Section 6, we present microlensing simulations of a more complex jet structure and compare the results to B1600+434. In Section 7, we discuss a critical test (i.e. the frequency-dependence of the modulation-index) to discriminate between scintillation and microlensing {\\sl and} compare predictions from the VLA 8.5--GHz light curves with the independent multi-frequency WSRT data. In Section 8 our results and conclusions are summarized. ",
        "conclusions": "We have shown {\\sl unambiguous} evidence of external variability in the CLASS gravitational lens B1600+434. The difference between the 8.5--GHz VLA light curves of the two lensed images shows external variability at the 14.6--$\\sigma$ confidence level. The modulation indices of the short-term variability are 2.8\\% for image~A and 1.6\\% for image~B. The difference light curve has an rms scatter of 2.8\\%, indicating that the short-term variability in both light curves is mostly of external origin (Sect.\\,2).  We have investigated two plausible sources of this external variability: (i) scattering by the ionized component of the Galactic interstellar medium (ISM) and (ii) microlensing by massive compact objects in the bulge/disk and halo of the lens galaxy.  Based on the `standard' theory of scintillation (e.g. Narayan 1992; Rickett et al. 1995) there should be a considerable increase in the modulation-index with wavelength (Sections 3 and 7). From simultaneous WSRT 1.4 and 5--GHz observations we find, however, that $m_{21}$=1.2\\% and $m_{6}$=3.7\\% (Table~1), which is a considerable decrease. Scintillation theory predicts $m_{21}$=9.0\\% for $m_{6}$=3.7\\% (Sect.\\,7). If the 1.4 and 5--GHz short-term variability is intrinsic, it is hard to reconcile with the fact that in 1998 the VLA 8.5--GHz light curves were dominated by external variability during the full eight months of monitoring (Sect.\\,2), although it can not be fully excluded yet. Moreover, from microlensing simulations, we expect that $m_{21}$=1.2--2.4\\% if $m_{6}=3.7\\%$ (Fig.\\,\\ref{predict}), based on constraints on the source structure and mass function of compact objects found from the VLA 8.5--GHz light curves (Sections 4, 5 and 7).  This range agrees remarkably well with the observed modulation index $m_{\\rm part}$=1.2\\% at 21 cm.  Supplementary to this argument, the difference in modulation-index between the lensed images would, in the case of scintillation, argue for either a very different Galactic ionized ISM (SM$_{\\rm A}/$SM$_{\\rm B}\\approx$3.1; Section 3.1--2) towards the lensed images or a different image size ($\\Delta\\theta_{\\rm B}/\\Delta\\theta_{\\rm A}$$\\approx$1.75; Sect.\\,3.2), although the latter might result from scatter-broadening. Furthermore, the longer variability time scales at 8.5\\,GHz ($\\gg$1 day; Figs \\ref{normlc}--\\ref{normdiff}) are also difficult to explain in terms of scintillation, as well as the absence of variability with short time scales in several 12\\,h WSRT observations at 5\\,GHz (Koopmans et al. in prep.).  However, the strongest argument against scintillation remains the dominant presence of short-term external variability at 8.5\\,GHz in 1998, combined with the fact that in 1999 significant short-term variability is seen at 5\\,GHz, but almost none at 1.4\\,GHz.  Under the microlensing hypothesis, we find a consistent, although not unique set of jet-component parameters. A core plus a single-jet-component with a size of 2--5 $\\mu$as, containing 5--11\\% of the flux density and moving superluminally with 9$\\la$$\\beta_{\\rm app}$$\\la$26, can explain the modulation-index and variability time scale in both lensed images (Sections 4--5). For image~A we find a significantly higher average mass of compact objects ($\\ga$0.5 M$_\\odot$), compared with those near image~B.  A much lower mass of compact object would result in a finer magnification pattern and thus in less variability. If image~B is scatter-broadened, its microlensing modulation-index is reduced, which might change the lower-limit on the compact object mass.  If one, based on the evidence gathered thus far, accepts that the 1.4, 5 and 8.5--GHz short-term variability in B1600+434--A and B is dominated by microlensing, the profound consequence is that the dark-matter halo at $\\sim$6 kpc above the plane of the disk-galaxy lens in B1600+434 is partly filled with massive compact objects. New WSRT, VLA and VLBI multi-frequency data is being obtained at the moment, which combined with a more comprehensive statistical analysis should provide us with refined constraints on the mass function of compact objects and the source structure (Koopmans et al. in prep.)."
    },
    "0004/astro-ph0004397_arXiv.txt": {
        "abstract": "We investigate the structure of dark matter halos by means of the kinematics of a very large sample of spiral galaxies of all  luminosities.  The observed rotation curves show  an  universal profile which is the sum of an exponential thin disk term and a spherical halo term with a flat  density core. We find that the Burkert  profile proposed to  describe  the dark matter halo density distribution of dwarf galaxies also provides an excellent mass model for the dark halos around  disk systems up to 100 times more massive. Moreover, we find that spiral  dark matter core densities $\\rho_0$ and core radii $r_0$ lie in  the same scaling relation $\\rho_0 = 4.5 \\times 10^{-2} (r _0/kpc)^{-2/3} M_{\\odot} pc^{-3} $ of dwarf galaxies with core radii up to  10 times smaller.  At  the highest masses $\\rho_0$   decreases with $r_0$ faster than the   $-{2\\over {3}}$ power law implying     a  lack of objects with disk masses  $> 10^{11}M_\\odot$ and  central densities $> 1.5 \\times 10^{-2}(r_0/kpc)^{-3}~M_{\\odot} pc^{-3}$  that can be explained by the existence of  a {\\it maximum} mass  of $\\approx 2 \\times 10^{12} M_{\\odot}$ for an halo   hosting a spiral galaxy. ",
        "introduction": "It is now well established that spiral galaxies have universal rotation curves (URC) that can be characterized by one single free parameter, the luminosity (e.g. Rubin et al. 1980; Persic  \\& Salucci, 1991, Persic, Salucci \\& Stel 1996 (PSS)). For instance,  low-luminosity spirals show  ever-rising rotation curves  (RC) out to the optical radius \\footnote{ $R_{opt} = 3.2 R_d$, with $R_d$ the exponential disk length-scale}, while,  in the same region, the RC of   high-luminosity spirals are flat or even decreasing.  It has been demonstrated  by Persic \\& Salucci (1988, 1990) and Broeils (1992) that, as the  galaxy luminosity decreases,  the light is progressively  unable to trace the radial distribution of the gravitating  matter (see also Salucci, 1997). This  discrepancy  is, in general, interpreted as the  signature of an  invisible mass component (Rubin et.al. 1980; Bosma 1981). As pointed out  by PSS the universality of the rotation curves, in combination with the invariant  distribution of the luminous matter, implies  an universal dark matter distribution with luminosity-dependent scaling properties.  On  the theoretical side,  recent high-resolution cosmological N-body simulations have shown that cold dark matter halos achieve a specific equilibrium density profile (Navarro, Frenk \\& White 1996, NFW; Cole \\& Lacey 1997; Fukushige \\& Makino 1997; Moore et al. 1998; Kravtsov et al. 1998). This can be characterized by one free parameter, e.g.  $M_{200}$, the dark mass enclosed within    the radius  inside which the average over-density is 200 times the critical density of the Universe. In the innermost regions  the dark matter profiles show some scatter around an average  profile which  is characterized by a power-law cusp $\\rho \\sim r^{-\\gamma} $,  with $\\gamma =1-1.5$ (NFW, Moore et al. 1998, Bullock et al., 1999).  Until recently,   due  to both  the limited number of suitable rotation curves and  a poor knowledge   of the exact  amount of luminous matter present in the innermost  regions of spirals, it has been  difficult to investigate    the  internal structure of dark matter halos. The situation is  more favorable  for (low surface brightness) dwarf galaxies which are   strongly dark matter dominated even at small radii.  The kinematics  of these systems  shows  an  universality of the dark halo density profiles, but, it results in disagreement with that predicted by CDM, in particular because of the existence of  dark halo density  cores. (Moore 1994; Burkert 1995). The origin of these features  is not yet understood (see e.g. Navarro, Eke \\& Frenk 1996; Burkert \\& Silk 1997, Gelato \\& Sommer-Larsen 1999),  but it is likely  that it  involves more physics  than  a simple  hierarchical assembly of cold structures. To cope with this observational  evidence, Burkert (1995) proposed an empirical profile that successfully  fitted  the halo  rotation curves of four dark matter dominated dwarf galaxies \\begin{equation} \\rho_b(r) = \\frac{\\rho_0 r^3_0}{(r+r_0)(r^2+r_0^2)} \\end{equation} \\noindent where $\\rho_0$ and $r_0$ are free parameters which represent the central dark matter density and the scale radius. This sample has been extended,  more recently,  to 17 dwarf  irregular and low surface brightness galaxies (Kravtsov et al. 1998, see however van den Bosch et al. 1999) which all  are found to  confirm  equation (1). Adopting spherical symmetry, the mass distribution of the Burkert halos is given by \\begin{equation} M_b(r) = 4~M_0 \\Big\\{ ln \\Big( 1 + \\frac{r}{r_0} \\Big) -tg^{-1} \\Big(   \\frac{r}{r_0} \\Big) +{1\\over {2}} ln \\Big[ 1 +\\Big(\\frac{r}{r_0} \\Big)^2 \\Big] \\Big\\} \\end{equation} \\noindent with $M_0$,  the dark mass within the core given by: \\begin{equation} M_0 = 1.6 \\rho_0 r_0^3 \\end{equation} \\noindent The halo contribution to the circular  velocity is then: \\begin{equation} V^2_b(r)=GM_b(r)/r. \\end{equation} Although the dark matter core parameters $r_0$, $\\rho_0$ and $M_0$ are in principle independent,   the observations reveal a clear  correlation (Burkert 1995): \\begin{equation} M_0  =  4.3\\times 10^7 \\left( \\frac{r_0}{kpc} \\right)^{7/3} M_{\\odot} \\end{equation} \\noindent which indicates that dark  halos represent a 1-parameter family which is completely specified,  e. g. by the core mass.  \\vbox{ \\vskip 1.truecm \\centerline{\\epsfxsize=4.8truecm \\epsfbox{deco.ps}} \\vskip 0.00truecm \\figcaption[]{ Synthetic rotation curves (filled circles with error bars) and  URC   (solid line) with  its  separate dark/luminous contributions (dotted line: disk;  dashed line: halo.) See PSS for details.  } \\vskip 0.7truecm }  The analysis of a   recently published  large sample of  RC's (Persic \\& Salucci, 1995) has  provided a suitable framework    to investigate the dark halo density  distribution in spirals. The starting  points of this study  are: {\\it a)} the mass in spirals is distributed according to   the  Inner Baryon Dominance (IBD) regime: there is a characteristic transition  radius $R_{IBD} \\simeq 2 R_d  \\Big({V_{opt}\\over{220 km/s}}\\Big)^{1.2}$ for which,  at  $r\\leq R_{IBD}$,    the luminous matter  totally accounts for the mass distribution,   while, for   $r> R_{IBD}$, the DM {\\it rapidly} becomes the  dominant dynamical component (Salucci and Persic, 1999a,b;  Salucci et al, 2000;  Ratnam and Salucci, 2000;  Borriello and Salucci, 2000). Then,  although the  dark halo might  extend  down to the galaxy center, it is only  for   $r>R_{IBD}$ that it  gives  a non-negligible contribution to the circular velocity. {\\it b)}  The dark matter  is distributed in a   different way with respect to any of  the various    baryonic components  (PSS, Corbelli and Salucci, 2000), and {\\it c)}  The HI  contribution to the circular velocity, for $r< R_{opt}$, is negligible (e.g.  Rhee, 1996 ; Verheijen, 1997).    The main aim of this  letter is to expand   the above  results   to  derive the luminosity-averaged density profiles  of the  dark halos   and   to relate   them  with the   Burkert profiles. Section 2 presents the  analysis of  a homogeneous sample of about 1100 rotation curves in which the dark halo contribution to the circular velocity is first derived and then  matched to the Burkert halo  mass models. In section 3 we discuss  the results. We take $H_0=75 km/s/Mpc$ and $\\Omega_0=0.3$, however no result depends on these choices. ",
        "conclusions": "Out to two optical radii,  the  Burkert density profile  reproduces, for the whole  spiral  luminosity sequence, the DM halos mass distribution. This density   profile, though  at very large radii  coincides with the NFW profile, approaches a constant, finite density value at the center,  in a way consistent with an isothermal distribution. This is in contradiction to cosmological models (e.g. Fukushige and Makino 1997) which predict that the velocity dispersion $\\sigma$ of the dark matter particles decreases towards the center to reach $\\sigma \\rightarrow 0$ for $r \\rightarrow 0$. After the result of this study, the dark halo inner regions, therefore, cannot be considered as kinematically cold structures but rather as \"warm\"  regions with size $ r_0 \\propto  \\rho_0^{-1.5}$. The halo core sizes  are  very large:  $r_0 \\sim 4-7 R_d$. Then, the boundary of the core region is  well beyond  the region where the stars are located and, as in  Corbelli and Salucci (2000), even at  the  outermost observed  radius there is not  the  slightest  evidence that dark  halos  converge to a    $\\rho \\sim r^{-2}$ (or a steeper) regime.  \\vbox{ \\vskip 0.5truecm \\centerline{\\epsfxsize=7.9truecm \\epsfbox{bur3.ps}} \\vskip 0.00truecm \\figcaption[]{   {\\it top} Disk mass (in solar units) {\\it vs}   central halo density  $\\rho_0$ (in $g/cm^3$) for   normal spirals ({\\it filled circles}). The straight line is the extrapolation  to high luminosities  of the   relation of  dwarfs. {\\it bottom }  Central density    {\\it vs}  core radii  (in kpc)  for  normal spirals ({\\it filled circles}),   compared with  the extrapolation of the relationship of  dwarfs ({\\it dotted line}, the point with errorbar represents a typical  object of Burkert,  1995). The solid line is the   eye-ball fit; $\\rho_0=5\\times 10^{-24}r_0^{-2/3}e^{-(r_0/27)^2} g/cm^3$. The effect of a limiting halo mass is also  shown ({\\it dashed line}).} \\vskip 0.5truecm } We find  that the dark  halos  around  spirals  are essentially an one-parameter family. It is relevant that the order parameter  (the central density or the core radius) correlates   with the luminous  mass  (see Fig 4). We do however not know how it is related to the global structural properties of the dark halo, like the virial radius or the virial mass. In fact, the RC out to $6R_D$ is completely determined by the core parameters, i.e. the central core density and the core radius,  both of which are not defined  in the CDM scenario.  The location of spiral galaxies in the parameter space of virial mass, halo  central density and baryonic mass is determined by different processes that occur on different scales and at different red-shifts. Yet, this  3D space degenerates into  a single  curve (see Figure 4,remind that: $\\rho_0 = {\\pi\\over {24}} ~ M_{200}/r_0^{3}$ and remind that: $M_d=G^{-1}\\beta V^2_{opt}R_{opt})$ which describes the  dark-luminous coupling.    Let us  discuss the limitations   of the present results. First,  here we have considered the luminosity dependence of the dark halo structure.  Although  this   is  probably the most relevant one, other dependences (Hubble type and surface brightness) should also be investigated.  Moreover, the existence of a (weak) cosmic variance   in   the halo structural properties cannot be excluded  until we analyze individual objects (Salucci, 2000, Borriello and Salucci, 2000). Secondly, we have derived the profile of DM halos  out to  about six disk-scale lengths, i.e. out to a distance much smaller than the virial radius.  To assess the  {\\it global}  validity of the  proposed mass model data at larger radii are obviously required."
    },
    "0004/astro-ph0004168_arXiv.txt": {
        "abstract": "We calculate the temperature profiles of (thin) accretion disks around rapidly rotating neutron stars (with low surface magnetic fields), taking into account the full effects of general relativity.  We then consider a model for the spectrum of the X--ray emission from the disk, parameterized by the mass accretion rate, the color temperature and the rotation rate of the neutron star.  We derive constraints on these parameters for the X--ray source Cygnus X--2 using the estimates of the maximum temperature in the disk along with the disk and boundary layer luminosities, using the spectrum inferred from the EXOSAT data.  Our calculations suggest that the neutron star in Cygnus X--2 rotates close to the centrifugal mass--shed limit. Possible constraints on the neutron star equation of state are also discussed. ",
        "introduction": "\\label{sec: I}  The soft X--ray spectra of luminous low--mass X--ray binaries (LMXBs) are believed to originate in geometrically thin accretion disks around neutron stars with weak surface magnetic fields (see for e.g. White 1995). An important parameter in modeling these spectra is the maximum value of the effective temperature in the accretion disk. The effective temperature profile in the disk can be estimated (assuming the disk to radiate from its surface like a blackbody) if one knows the accretion energy released in the disk. In a Newtonian treatment, the innermost region of an accretion disk surrounding a neutron star with weak magnetic field will extend rather close to the neutron star surface. The amount of energy released in the disk will be one--half of the total accretion energy, the other half being released in the thin boundary layer between the disk's inner edge and the neutron star's surface. This then gives the disk effective temperature $(T_{\\rm eff})$ varying with the radial distance $(r)$ as $T_{\\rm eff} \\propto r^{-3/4}$ and the maximum effective temperature $(T^{\\rm max}_{\\rm eff})$ will depend on the (nonrotating) neutron star mass $(M)$ and radius $(R)$ as $T_{\\rm eff}^{\\rm max} \\propto (M \\dot{M}/R^3)^{1/4}$, where $\\dot{M}$ is the steady state mass accretion rate. The value of $(T^{\\rm max}_{\\rm eff})$ in the disk, in this approach, occurs at a radial distance $1.36~R$.  Mitsuda et al. (1984) parameterized the disk spectrum by the maximum temperature of the disk, using the above formalism and assuming the mass of the neutron star is equal to $1.4$~\\msun.  These authors assumed that the inner parts of the disk do not contribute to the X--ray spectrum, and suggested a multi--color spectrum for the X--ray emission from the disk. It was shown by these authors, that the observed spectra of Sco X--1, 1608--52, GX 349+2 and GX 5--1, obtained with the {\\it Tenma} satellite, can be well fitted with the sum of a multi--color spectrum and a single blackbody spectrum (presumably coming from the boundary layer). White, Stella \\& Parmar (1988) (WSP) suggested that the simple blackbody accretion disk model should be modified to take into account the effects of electron scattering. Using {\\it EXOSAT} observations, these authors compared the spectral properties of the persistent emission from a number of X--ray burst sources with various X--ray emission models. This work suggests that either the neutron star (in each system considered) rotates close to equilibrium with the Keplerian disk, or that most of the boundary layer emission is not represented by a blackbody spectrum.  For accretion disks around compact objects, one possibility is that of the accretion disk not being Keplerian in nature.  For e.g. Titarchuk, Lapidus \\& Muslimov (1998) have formulated a boundary problem in which the Keplerian accretion flow in the inner disk is smoothly adjusted to the neutron star rotation rate.  The generality of such a formulation permits application even to black holes, but only for certain assumed inner boundary conditions.  These authors demonstrate that there exists a transition layer (having an extent of the order of the neutron star radius) in which the accretion flow is sub-Keplerian.  An attractive feature of this formalism is that it allows super-Keplerian motion at the outer boundary of the transition layer, permitting the formation of a hot blob that ultimately bounces out to the magnetosphere. This formalism (Titarchuk \\& Osherovich 1999; Osherovich \\& Titarchuk 1999a; Osherovich \\& Titarchuk 1999b; Titarchuk, Osherovich \\& Kuznetsov 1999) therefore provides a mechanism for the production of high frequency quasi--periodic oscillations (QPOs) observed in the X--ray flux from several LMXBs.  Such effects, when taken into account, can modify the Newtonian disk temperature profile (Chakrabarti \\& Titarchuk 1995).  There are several other effects which will modify the Newtonian disk temperature profile, such as the effects of general relativity and of irradiation of the disk by the central neutron star. The wind mass loss from the disk and the residual magnetic field near the disk's inner edge may also play a part in modifying the effective temperature (Knigge 1999). Czerny, Czerny \\& Grindlay (1986) calculated LMXB disk spectra assuming that a disk radiates locally as a blackbody with the energy flux detemined by viscous forces, as well as irradiation by the boundary layer, and took into account relativistic effects, some of them in an approximate way. The possible effects of general relativity were also discussed by Hanawa (1989), using the Schwarzschild (nonrotating) metric, assuming that the neutron star radius is less than the radius of the innermost stable circular orbit ($r_{\\rm in} = 6 G M/c^2$), which they identified as the disk inner boundary. The color temperature was assumed to be higher than the effective temperature by a factor of 1.5.  It was found by Hanawa (1989) that the observations are consistent with a geometrically thin, optically thick accretion disk, whose inner edge is at $r=r_{\\rm in}$, $r$ being the Schwarzschild radial coordinate.  An important dynamical aspect of disk accretion on to a weakly magnetized neutron star is that the neutron star will get spun up to its equilibrium period, which is of the order of milliseconds (see Bhattacharya \\& van den Heuvel 1991, and refereces therein). The effect of rotation is to increase the equatorial radius of the neutron star, and also to relocate the innermost stable circular orbit (for a corotating disk) closer to the stellar surface (as compared to the Schwarzschild case). These effects will be substantial for rapid rotation rates in a fully general relativistic treatment that includes rotation. Therefore, for accreting neutron stars with low magnetic fields,  the stellar radius can be greater or less than the radius of the innermost stable orbit, depending on the neutron star equation of state and the spacetime geometry.  The effect of magnetic field will be to constrain the location of the inner--edge of the accretion disk to the magnetospheric (Alf\\'{v}en) radius. In such a case, $r_{\\rm in}$ would lose the astrophysical relevance as discussed here.  However, this will be so only if the magnetic field strength ($B$) is large.  The problem addressed in this paper refer to LMXBs which contain old neutron stars which are believed to have undergone sufficient magnetic field decay (Bhattacharya \\& Datta 1996).  Clearly, for low magnetic field case, a number of different disk geometries will be possible if general relativistic effects of rotation are taken into account. These structural differences influence the effective temperature profile and the conclusions derived by Czerny, Czerny \\& Grindlay (1986) and Hanawa (1989) are likely to be modified.  In this paper, we attempt to highlight the effects brought in due to general relativity and rotation of the neutron star on the accretion disk temperature profile and then apply this to the particular case of the X--ray source Cygnus X--2.  For simplicity (unlike Titarchuk, Lapidus \\& Muslimov 1998), we assume the accretion disk to be fully Keplerian, geometrically thin and optically thick. We first give a theoretical estimate of the modifications in $T_{\\rm col}^{max}$ that would result if inclusion is made of the rotational effects of general relativity, and illustrate these by taking representative neutron star equations of state. We then consider a model for the spectrum parameterized by the mass accretion rate, the color factor, and the rotation rate of the accreting neutron star (assumed to be weakly magnetized).  We derive constraints on these parameters for the X--ray source Cygnus X--2, for which we take the estimates of $T_{\\rm eff}^{\\rm max}$, the disk luminosity and boundary layer luminosity from the analysis of WSP.  A conclusion of our work is that the neutron star in Cygnus X--2 has a rapid spin rate close to the centrifugal mass--shed limit.  The format of this paper is as follows. In Section~\\ref{sec: II}, we discuss the rotational general relativistic effects on the disk temperature, using a formalism given by Page \\& Thorne (1974) and the disk irradiation by the neutron star.  The theoretical predictions  for the temperature profiles with these effects taken into account are presented in Section~\\ref{sec: III}. Section~\\ref{sec: IV} deals with comparison with observations, and its implications for parameters of our model for Cygnus X--2. A summary and discussions are presented in Section~\\ref{sec: V}. ",
        "conclusions": "\\label{sec: V}   In this paper, we have calculated the temperature profiles of accretion disks around rapidly rotating and non--magnetized neutron stars, using a fully general relativistic formalism.  The maximum temperature and its location in the disk are found to differ substantially from their values corresponding to the Schwarzschild space--time, depending on the rotation rate of the accreting neutron star.  We have considered a model for the spectrum of the X--ray emission from the accretion disk, parameterized by the mass accretion rate, the color temperature, and the rotation rate of the star.  We have compared the maximum effective temperature in the disk and the accretion luminosities (corresponding to the disk and the boundary layer) with the results of spectral fitting for the X--ray source Cygnus X--2 (WSP), and derived constraints on these parameters for the neutron star in this X--ray binary.  The main conclusion of our analysis is that the neutron star in Cygnus X--2 has a rapid spin rate (close to the centrifugal mass shed value), and that the system has a fairly large accretion rate (several times $10^{18}$~g~s$^{-1}$).  The low luminosity of the boundary layer compared to that of the disk for Cygnus X--2 is consistent with the above conclusion that the neutron star in this system has a rapid rotation rate. The low value of the ratio $L_{\\rm BL}/L_{\\rm D}$ justifies our assumption that the radiation pressure is negligible in the disk, so that the geometrically thin approximation for the disk is reasonable. According to Shimura \\& Takahara (1988), the spectrum from the disk can be represented as a multi-color blackbody only if $\\dot{M} >0.1 \\dot{M}_{\\rm edd}$. Our results for Cygnus X--2 are in accord with this. Interestingly, if we take the lower value $1.7$ for the color index $f$ (Shimura \\& Takahara 1988), we obtain a consistent set of results, except for the stiffest EOS model (D).  This suggests that the comparatively lower values of $f$ would disfavor stiff EOS for neutron star matter. However, if we take the value of $f=2.6$, as reported by Borozdin et al. (1999), one would require an EOS model that is stiffer than the stiffest used here, or a mass greater than $M=1.78$~\\msun (if one uses the narrower limits on the luminosity and color temperature).  On the other hand, if one were to use the broader limits, the hardening factor $f=2.6$ is disallowed only by the softest EOS model.  We have assumed here that the magnetic field of the neutron star is weak, which implies that the radius of the last orbit of the accretion disk should be much greater than the Alfv\\'{e}n radius ($r_{\\rm a}$) (e.g., Shapiro \\& Teukolsky 1983), \\begin{equation} R >> r_{\\rm a} = 2.9 \\times 10^7  {({\\dot {M}\\over \\dot {M}_{\\rm edd}})}^{-2/7} \\mu_{30}^{4/7} ({M\\over M_\\odot})^{-3/7} \\end{equation} \\noindent where $M$ is the mass of the neutron star, $\\mu_{30}$ is the magnetic moment in units of $10^{30}$ G cm$^3$ and $r_{\\rm a}$ is in cm. The above condition implies that for $R \\approx 15$ km., $\\dot{M}/\\dot {M}_{\\rm edd} \\approx 20$ and $M = 1.4 M_\\odot$, the magnetic moment  $\\mu_{30} << 3.4 \\times 10^{-2} $ or the magnetic field in the surface should be less than $10^{10}$ G. So the conclusions presented by us will be valid for the neutron star magnetic field upto a few times $10^{9}$~G.  In our analysis, we have assumed that the the boundary layer between the disk and the neutron star surface does not affect the inner regions of the disk. This will be a valid approximation when the boundary layer luminosity is smaller than the disk luminosity, and the boundary layer extent is small compared to the radius of the star. The flux received at earth  from this region is  \\begin{equation} F_{\\rm BL} = ({2 \\pi R \\frac{\\Delta R}{D^2}}) \\cos{\\theta} ({\\sigma T_{\\rm BL}^4 \\over \\pi}) \\end{equation}  \\noindent where $\\Delta R$ is the width of the boundary layer, $D = 8$ kpc is the distance to the source, $\\theta = 60^o$ is the inclination angle and $T_{\\rm BL}$ is the effective temperature. Spectral fitting gives a best fit value for $F_{\\rm BL} \\approx 4 \\times 10^{-9}$ ergs sec$^{-1}$ cm$^{-2}$ and $T_{\\rm BL} = T_{\\rm col(BL)}/f_{\\rm BL}$ = 2.88/$f_{\\rm BL}$~keV, where $f_{\\rm BL}$ is the color factor for the boundary layer and $T_{\\rm col(BL)}$ is the color temperature of the boundary layer. Using these values, $\\Delta R \\approx $ 0.2 $f_{\\rm BL}^4$ km, which is indeed smaller than $R$ provided the boundary layer color factor $f_{\\rm BL}$ is close to unity. This is supported by the work of London, Taam \\& Howard (1986) and Ebisuzaki (1987), who obtain $f_{\\rm BL}\\approx 1.5$.  A few comments regarding the validity of the Page \\& Thorne (1974) formalism for accreting neutron star binaries are in order here. Unlike for the case of black holes, neutron stars possess hard surface that could be located outside the marginally stable orbit. For neutron star binaries, this gives rise to a possiblity of the disk inner edge coinciding with the neutron star surface.  We have assumed that the torque (and hence the flux of energy) vanishes at the disk inner edge even in cases where the latter touches the neutron star surface.  In the case of rapid spin of the neutron star (as we infer for Cygnus X--2), the angular velocity of a particle in Keplerian orbit at disk inner edge will be close to the rotation rate of the neutron star. Therefore, the torque between the neutron star surface and the inner edge of the disk is expected to be negligible.  Independently of whether or not the neutron star spin is large, Page \\& Thorne (1974) argued that the error in the calculation of $T_{\\rm eff}$ will not be substantial outside a radial distance $r_{\\rm o}$, where $r_{\\rm o}$ is given by $r_{\\rm o} - r_{\\rm in} = 0.1 r_{\\rm in}$. In our calculation, we find that $r_{\\rm eff}^{\\rm max}$ (which is the most important region for the generation of X--rays) is greater than $r_{\\rm o}$ by several kilometers for all the cases considered.  The shortest time-scale of the system is given by the frequency in the innermost stable circular orbit ($\\nu_{\\rm in}$, Table 2: column 5). A periodic oscillation in the system should be at a frequency lower than $\\nu_{\\rm in}$ (unless the model invoked to explain the temporal behavior predicts substantial power in the second harmonic, i.e.,  $\\nu_{\\rm QPO} \\approx 2 \\nu_{\\rm in}$). The maximum frequency of the kHz quasi-periodic oscillation (QPO) observed for Cygnus X-2 is 1005 Hz (Wijnands et al. 1998). The stiffest EOS, model (D), will then be disfavored since $\\nu_{\\rm in} < 1$~kHz for this model. Further, the neutron star mass estimate in Cygnus X-2 ($ \\approx  1.78 M_\\odot$, Orosz \\& Kuulkers 1998) is not consistent with the soft EOS model (A). Our analysis, therefore, favors neutron star EOS model which are intermediate in the stiffness parameters.  We have not attempted to model the observed temporal behavior of the source, and in particular, the QPO observations. Beat frequency models identify the peak separation of the two kHz QPO observed with the neutron star spin rate. For Cygnus X-2 the observed peak separation is $\\Delta \\nu = 346 \\pm 29$ Hz (Wijnands et al. 1998) which is smaller than the typical rotation frequencies calculated here. However, a pure beat-frequency model has been called into question due to several observations. For instance, $\\Delta \\nu$ has been observed to vary by about 40\\% for Sco X-1 (van der Klis et al. 1997) and the kHz QPO frequencies have been found to be correlated with the  break frequency ($\\approx$ 20 Hz) of the power spectrum density.  An alternate model, where the QPOs are suggested to originate due to non--Keplerian motion of matter in the disk (Titarchuk \\& Osherovich 1999; Osherovich \\& Titarchuk 1999a; Osherovich \\& Titarchuk 1999b; Titarchuk, Osherovich \\& Kuznetsov 1999) have been proposed.  These authors have also demonstrated the model by applying it to particular sources.  Inclusion of this Newtonian model into the framework of the calculations mentioned in this paper require a parallel formulation within the space--time geometry chosen herein.  X-ray binaries like Cygnus X-2 are believed to be the progenitors of the millisecond pulsars. Therefore, the discovery of a pulsar with a period $\\approx 1 $ ms will strengthen the model presented in this paper, in terms of a rapidly rotating accreting neutron star. X-ray spectral analysis of Cygnus X-2 and similar sources using data from recent satellites (e.g. BeppoSAX, ASCA, Chandra) are required to provide further support to the model presented in this paper."
    },
    "0004/astro-ph0004383_arXiv.txt": {
        "abstract": "We present imaging and spectro- polarimetric observations of the ultraluminous infrared galaxy IRAS P09104+4109 using the Keck 10-m Telescope. We detect the clear presence of broad \\hb, \\hg, and Mg~II~\\wave 2800 emission lines in the polarized flux spectra of the nucleus and of an extranuclear emission region $\\sim$ 4\\arcsec~away, confirming the presence of a hidden central quasar. The polarization of the broad Mg II emission line is high ($\\sim$ 29\\%), consistent with the remarkably high polarization ($\\sim$ 30\\%--40\\%) observed in the extended continuum emission. This indicates that the off-nuclear continuum is dominated by light scattered from the hidden quasar, most probably by dust mixed with the line emitting gas. The high polarizations, combined with the ``foreshortened'' morphology of the polarized brightness distribution allow us to constrain the scattering biconical structure to be at inclination $i~\\approx~50$\\arcdeg~with a half-opening cone angle $\\theta_c~\\approx$~40\\arcdeg.  The narrow emission lines are polarized in a stratified fashion, with the high ionization lines (\\oiii, \\nev, \\fevii) being polarized 0.7\\%--1.7\\% and \\oii~essentially unpolarized. The line polarizations are positively correlated with critical density, ionization potential, and velocity width of the emission lines. This indicates that, as is the case with the narrow-line radio galaxies, which also often contain powerful quasars, the narrow-emission line region may be partially shadowed by the putative torus, with the higher ionization lines originating closer to the nucleus.  One notable characteristic of the extranuclear knot is that all species of Fe are markedly absent in its spectrum, while they appear prominently in the nucleus. In addition, narrow Mg II is observed to be much weaker than predicted by ionization models. Our favored interpretation is that there is a large amount of dust in the extranuclear regions, allowing gaseous refractory metals to deposit. Near the nucleus, dust is destroyed in the strong radiation field of the quasar, inhibiting metal depletion onto grains. The extended emission regions are most likely material shredded from nearby cluster members and not gas condensed from the cooling flow or expelled from the obscured quasar.  The higher temperature inferred from [O III] lines compared to that from [N II] and the general better agreement with models of line ratios, especially \\oiii~\\wave 5007/\\wave 4363 and He~II/\\hb, provide strong evidence for matter-bounded clouds in addition to ionization-bounded clouds in the NLR. Ionization by pure velocity shocks can be ruled out. Shocks with photoionizing precursors may be present, but are probably not a dominant contributor to the energy input. ",
        "introduction": "Discovered by the $IRAS$ all-sky survey, ultraluminous infrared galaxies (ULIRGs, see review by \\cite{sm96}) emit most of their energy in the infrared ($L_{IR} > 10^{12}$ \\lsun), and could harbor infant quasars enshrouded in a large amount of dust (\\cite{san88}). On the other hand, they may also represent energetic, compact starbursts (\\cite{con91}; \\cite{gen98}). Much recent research effort has been devoted to understanding the dominant energy source in these ULIRGs -- whether it is obscured quasars or intense bursts of star formation.  IRAS P09104+4109 ($z$ = 0.44; \\cite{k88}) is one of the few exceptionally luminous ULIRGs sometimes referred to as the ``hyperluminous'' ($L_{IR}~\\simgt~10^{13}$~\\lsun) infrared galaxies. The others include F15307+3252 ($z$ = 0.926; \\cite{c94}) and F10214+4724 ($z$ = 2.286; \\cite{rr91}). All of these galaxies have been shown to harbor a quasar nucleus obscured from direct view (P09104+4109, \\cite{hin99}; F15307+3252, \\cite{hin95}; F10214+4724, \\cite{goo96}). This has led to the suggestion that perhaps all warm\\footnote{``Warm'' ULIRGs are those having $f_{25}/f_{60}~>~0.2$, \\cite{low88}; \\cite{san88}); $f_{25}$ and $f_{60}$ are the $IRAS$ flux densities in units of Jy at 25 $\\mu$m and 60 $\\mu$m, respectively.} ULIRGs contain buried QSOs and that they may be the misdirected type 2 QSOs (AGN with Seyfert 2-like emission-line characteristics but QSO-like luminosities). A diagnostic diagram (\\cite{t99}) involving \\oiii~emission line and infrared color $f_{25}/f_{60}$ does indeed show that essentially all warm ULIRGs with sufficiently high ionization harbor energetic quasars in their centers.  IRAS P09104+4109 has been identified with a central cD galaxy in a rich cluster (\\cite{k88}; Hall \\etal~1997). It has a high-ionization spectrum characteristic of a Seyfert 2 galaxy. Ground-based and $Hubble~Space~Telescope$ ($HST$) images show that \\objn~has an off-nuclear northern extension or ``plume'' of ionized gas (\\cite{k88}; \\cite{hn88}; \\cite{armus99}) that may serve as a scattering mirror of the light originating from the nucleus. We wish to probe the origin of this extension and constrain the scattering geometry for \\objn~by obtaining polarimetric observations of this knot.  In this paper, we confirm the high polarization and the broad emission lines of \\hb, \\hg~and Mg II in polarized flux of the nucleus. In addition to the $B$-band imaging polarimetry, we also present new spectropolarimetric data for the nucleus and off-nuclear emission regions of \\objn. We study in detail the emission-line spectra of the nuclear and extra-nuclear regions in order to derive an understanding of the ionization mechanism of the narrow-line region (NLR) of this galaxy. Throughout this paper, we assume $H_o$ = 75 \\hubu, $q_o$ = 0 and $\\Lambda$ = 0. At the redshift $z=0.44$ of P09104+4109, 1\\arcsec~corresponds to a projected size of 5 kpc. ",
        "conclusions": "Our high-quality Keck spectropolarimetric and imaging polarimetric data confirm the high polarization and broad Balmer emission lines in the polarized flux spectra of \\objn. In addition, we also detect broad Mg II in both total and polarized flux spectra of the nuclear and extension regions, indicating the presence of a hidden quasar visible only in scattered light. The narrow-line polarizations exhibit a strong positive correlation with line width, critical density and ionization potential of the transition, indicating that the line emission arise from a radially stratified gas. The lack of a clear fan-like morphology in the polarization image suggests that our viewing angle ($i~\\approx$~50\\arcdeg) is not far outside the ionization cone ($\\theta_c~\\approx$~40\\arcdeg) of \\objn.  The emission-line spectra of \\objn~are consistent with AGN photoionization with contribution from both ionization-bounded and matter-bounded clouds. Under the assumption that photoionization by the central hidden quasar is the main operating mechanism, \\fevii~should be detected in the extended gas and Mg II should be much stronger than observed. The absence of all Fe lines and weakness of Mg II in the extension can be readily explained if dust is mixed with the ionized gas of the NLR. Furthermore, the fact that both the nuclear and the extended continua are highly polarized with relatively blue polarized flux spectra compared to that of the average quasar also supports the existence of dust, which could efficiently scatter and bluen the light. The other possibility is that the extended gas simply is highly deficient in iron, or is metal poor. This does not seem very likely, since the strengths of the other metal lines appear to be in good agreement with solar abundance. Since grains remain intact, this in turns implies that shocks are probably not important in the extra-nuclear regions. Although the diagnostic diagrams offer somewhat inconclusive assessment of the shocks+precursor models for the extension region, there is no strong evidence for shocks in \\objn~from the line ratios. Ionization by pure shocks can be ruled out by the data. The existence of dust and lack of strong shocks in the extended gas also suggest that it is probably dismembered remnants of a cluster neighbor being disrupted or in the process of merging with the central cD galaxy.  Two predictions follow directly from our interpretation of the current data: The near-IR line [Fe II] 1.257 $\\mu$m has been observed in the nucleus of \\objn~by Soifer \\etal~(1996), but should be absent from the extension, confirming the absence of gaseous Fe, and hence the lack of shocks and presence of dust there. Likewise, [Ca II] \\waves 7291, 7324 should be absent in EXT spectrum, since Ca is more than 10 times more sensitive to grain depletion than Fe or Mg (\\cite{fer97b})."
    },
    "0004/astro-ph0004330_arXiv.txt": {
        "abstract": "We report the discovery of two young isolated radio pulsars with very high inferred magnetic fields.  PSR~J1119$-$6127 has period $P=0.407$\\,s, and the largest period derivative known among radio pulsars, $\\dot P=4.0\\times10^{-12}$.  Under standard assumptions these parameters imply a characteristic spin-down age of only $\\tau_{\\rm c} = 1.6$\\,kyr and a surface dipole magnetic field strength of $B=4.1\\times10^{13}$\\,G.  We have measured a stationary period-second-derivative for this pulsar, resulting in a braking index of $n=2.91\\pm0.05$.  We have also observed a glitch in the rotation of the pulsar, with fractional period change $\\Delta P/P = - 4.4\\times10^{-9}$.  Archival radio imaging data suggest the presence of a previously uncataloged supernova remnant centered on the pulsar.  The second pulsar, PSR~J1814$-$1744, has $P=3.975$\\,s and $\\dot P=7.4\\times10^{-13}$.  These parameters imply $\\tau_{\\rm c} = 85$\\,kyr, and $B=5.5\\times10^{13}$\\,G, the largest of any known radio pulsar.  Both PSR~J1119$-$6127 and PSR~J1814$-$1744 show apparently normal radio emission in a regime of magnetic field strength where some models predict that no emission should occur.  Also, PSR~J1814$-$1744 has spin parameters similar to the anomalous X-ray pulsar (AXP) 1E~2259+586, but shows no discernible X-ray emission.  If AXPs are isolated, high magnetic field neutron stars (``magnetars''), these results suggest that their unusual attributes are unlikely to be merely a consequence of their very high inferred magnetic fields. ",
        "introduction": "\\label{sec:intro}  The pulsar in the Crab nebula (PSR~B0531+21), with period $P = 33$\\,ms, was born in a type~II supernova observed in 1054~AD, supporting the view that at least some core-collapse supernovae (SNe) form pulsars. Based largely on studies of the Crab and a few other young objects, a picture has emerged where pulsars are born spinning rapidly (with initial period $P_0 \\approx 20$\\,ms in the case of the Crab), and spin down due to their large magnetic moments according to $\\dot{\\nu} \\propto - \\nu^n$.  In this spin-down law $\\nu = 1/P$ is the pulsar rotation frequency, $\\dot \\nu$ is its derivative, and $n = \\nu \\ddot \\nu/(\\dot \\nu)^2$ is the ``braking index.''  Integration of the spin-down law with constant magnetic moment gives the age of the pulsar, \\begin{equation} \\label{eq:tau} \\tau = \\frac{P}{(n-1) \\dot{P}} \\left[ 1 - \\left(\\frac{P_0}{P} \\right)^{n-1} \\right]. \\end{equation} Braking indices have been measured for only four pulsars, namely PSRs~B0531+21, B0540$-$69, B0833$-$45, and B1509$-$58, with values for $n$ of $2.51\\pm0.01$, $2.2\\pm0.1$, $1.4\\pm0.2$, and $2.837\\pm0.001$ respectively (\\cite{lps93}; \\cite{dnb99}; \\cite{lpgc96}; \\cite{kms+94}).  In other cases, an oblique rotating vacuum dipole model is typically assumed, for which $n=3$ (\\cite{mt77}), and if $P_0 \\ll P$, equation~(\\ref{eq:tau}) reduces to $\\tau = P/(2 \\dot P) \\equiv \\tau_{\\rm c}$, the characteristic age of a pulsar.  With a neutron star radius of $10^6$\\,cm and moment of inertia of $10^{45}$\\,g\\,cm$^2$, the surface magnetic field strength is \\begin{equation} \\label{eq:B} B = 3.2 \\times 10^{19} \\sqrt{P \\dot{P}} \\;\\; {\\rm G}. \\end{equation} The luminosity generated in the braking of the pulsar rotation, $\\dot E = 4\\pi^2 I \\nu \\dot \\nu$, is emitted in the form of magnetic dipole radiation and a relativistic particle wind.  The vast majority of this luminosity may be deposited in the ambient environment, powering a plerionic supernova remnant (SNR) such as the Crab synchrotron nebula, while a very small portion may be observed as pulsed electromagnetic radiation.  Despite the above, many questions remain regarding the outcome of type~II SNe and the manifestation of young neutron stars.  Although Galactic SNe and pulsar formation rates are both notoriously difficult to estimate (see, e.g., \\cite{vt91}; \\cite{tls94}; \\cite{wol98}, and \\cite{no90}; \\cite{lbdh93}; \\cite{lml+98}), it is quite plausible that type~II SNe occur significantly more often than radio pulsars of the kind already known are born (see \\cite{vt91}; \\cite{wol98}).  If this is the case, perhaps some young neutron stars are being ``missed.''  A possible example is SNR 3C58, the likely outcome of a type~II SN observed about 820\\,yr ago, with no detectable pulsar.  Studying the energetics and morphology of the remnant, Helfand, Becker, \\& White (1995)\\nocite{hbw95} make a compelling case for the presence of an unseen pulsar with higher magnetic field than any previously known. Having a short period like the Crab at birth, such a pulsar would have spun down rapidly to a present long period.  Maybe yet other pulsars are born spinning slowly and never generate the large $\\dot E$ required to power an easily detectable nebula.  In addition, some neutron stars may never manifest themselves as radio pulsars at all.  It has been suggested that there exists a class of isolated rotating neutron stars with ultra-strong magnetic fields, the so-called ``magnetars'' (\\cite{dt92a}).  The observational properties of radio pulsars and magnetar candidates are very different.  Radio pulsars rarely exhibit X-ray pulsations, and when they do, their X-ray power is small compared to their $\\dot E$.  By contrast, magnetars emit pulsed X-rays with luminosities far in excess of their spin-down power (\\cite{vg97}; \\cite{ksh+99}) but remain undetected at radio wavelengths.  The dichotomy is thought to result from the much larger magnetic fields in magnetars (\\cite{td93a}; \\cite{hh97}).  In this paper we report the discovery of two isolated radio pulsars with some properties that are unusual and interesting in the context of the above questions. ",
        "conclusions": "\\label{sec:disc}  \\subsection{PSR~J1119$-$6127}\\label{sec:p1119}  PSR~J1119$-$6127 has the largest period derivative known among radio pulsars.  Partly for this reason it was relatively straightforward to measure a stationary $\\ddot \\nu$ with a phase-connected timing solution (i.e., through absolute pulse numbering), only the third pulsar for which this has been possible.  The resulting value of braking index is $n=2.91\\pm0.05$, including possible contamination by timing noise (see \\S~\\ref{sec:obs}), and is in good agreement with that predicted by a model treating the pulsar as an oblique rotator with a current-starved outer magnetosphere (\\cite{mel97}).  For the four other pulsars for which it has been measured, $n$ ranges between 1.4 and 2.8 (see \\S~\\ref{sec:intro}). That observed braking indices are smaller than 3 can be explained in a variety of ways (see \\cite{mel97} for a review), including a kinetic energy-dominated flow at the light cylinder, or an increase in the magnetic moment of the star over time (\\cite{br88}).  None of these scenarios are consistent with all the observations (\\cite{aro92}). Measurement of $\\dddotnu$ would constrain these possibilities further. At present the upper limit in Table~\\ref{tab:parms} is 30 times the value expected from a simple spin-down law (\\cite{br88}).  Whether this can be measured, and how much the measurement of $n$ can be improved with further observations, will depend on the level of timing noise and glitch activity displayed by the pulsar.  Assuming that $P_0 \\ll P$, but using the measured values of $P$, $\\dot P$, and $n$ (Table~\\ref{tab:parms}) in equation~(\\ref{eq:tau}), the age of PSR~J1119$-$6127 is $1.7\\pm0.1$\\,kyr, including possible biases due to timing noise and glitches (see \\S~\\ref{sec:obs}). Of course, if the pulsar were born spinning slower, it would be younger.  For $P_0 = 0.2$\\,s, half the present period, the age is 1.2\\,kyr.  In any case it is clear that PSR~J1119$-$6127 is among the very youngest neutron stars known.  Three other pulsars with characteristic ages under 2\\,kyr are known: the Crab pulsar ($\\tau_{\\rm c} = 1.3$\\,kyr), PSR~B1509$-$58 in G320.4$-$1.2 ($\\tau_{\\rm c}=1.6$\\,kyr), and PSR~B0540$-$69 in the Large Magellanic Cloud ($\\tau_{\\rm c}=1.7$\\,kyr).  All three are associated with SNRs.  We have searched for evidence of an SNR near PSR~J1119$-$6127.  Although none is cataloged (\\cite{gre96}), data from the Molonglo Observatory Synthesis Telescope obtained at a radio frequency of 843\\,MHz (\\cite{gcly99}) reveal a faint ring of radius 7$'$ centered on the pulsar.  This could be the expanding blast wave of the SNR.  Its size would imply an expansion velocity of $\\sim 10^4$\\,km\\,s$^{-1}$, for a distance of 8\\,kpc and age of 1.6\\,kyr, reasonable for the blast-wave interpretation if the surrounding medium is of low density and relatively uniform.  Additional ATCA data show that the shell has a non-thermal radio spectrum (Crawford et al., in preparation).  This possible SNR is also X-ray-bright, with its spectrum described by either a power-law or thermal model, and additional observations are required to further constrain its properties (\\cite{pkc00b}).  Although the supernova that gave birth to this pulsar occurred in an era in which celestial events were recorded by some civilizations, this explosion may have been too far south and/or too distant or too obscured to have been detected by these observers.  The glitch observed in PSR~J1119$-$6127 is small compared to most glitches in most pulsars, with $\\Delta \\nu/\\nu = (4.4\\pm0.4)\\times10^{-9}$ (Table~\\ref{tab:parms}), but it is of similar fractional size as three of the five glitches observed in the Crab pulsar over 23 years (Lyne et al.~2000b\\nocite{lsg00}).  It remains to be seen whether at least some of the change measured in $\\dot \\nu$ is permanent, as seen in the Crab glitches.  Unless we were unreasonably lucky, PSR~J1119$-$6127 glitches more often than the Crab pulsar, but it is curious that its glitches share some characteristics with those of the Crab: while its period and magnetic field are approximately 10 times larger than the Crab's, its age, and perhaps therefore its internal temperature, are similar.  Finally, we compare the PSR~J1119$-$6127 system with some young pulsar/SNR systems.  For ages $\\la 2000$\\,yr, the radio luminosity $L_{\\rm R}$ of a synchrotron nebula (``plerion'') with a central pulsar is a measure of the energy output of the pulsar over its lifetime, due to the relatively long lifetime of the radiating electrons.  The plerion X-ray luminosity $L_{\\rm X}$, on the other hand, reflects the current $\\dot E$ of the pulsar.  For the Crab and PSR~B0540$-$69, $L_{\\rm X} \\sim 0.05 \\dot E$, while for PSR~B1509$-$58, $L_{\\rm X} \\sim 0.01 \\dot E$ (see Helfand et al.~1995, and references therein\\nocite{hbw95}).  For SNR 3C58, Helfand et al.\\nocite{hbw95} find that all available data can be reconciled with a (candidate) pulsar having $P \\sim 0.2$\\,s, $\\dot P \\sim 4\\times 10^{-12}$ (parameters similar to those of PSR~J1119$-$6127 --- see Table~\\ref{tab:parms}), and with $L_{\\rm X} \\sim 5 \\times 10^{-4} \\dot E$.  For PSR~J1119$-$6127, with a current $\\dot E$ 200 times smaller than the Crab's, the limit on plerionic X-ray emission is $L_{\\rm X} \\la 10^{-3} \\dot E$ (Pivovaroff et al.~2000a\\nocite{pkc00b}).  If PSR~J1119$-$6127 were born with a small period, it would have had a much larger $\\dot E$ within the past $\\sim 1700$\\,yr, possibly larger than the Crab's initially.  That energetic past might be reflected in plerionic radio emission near the pulsar, depending on the local environment.  A measurement of $L_{\\rm X}$ and $L_{\\rm R}$ may in principle provide information about whether PSR~J1119$-$6127 was born with a rapid spin rate, as commonly assumed for most pulsars, or whether it was born a slow rotator.   \\subsection{PSR~J1814$-$1744}\\label{sec:p1814}  Figure~\\ref{fig:ppdot} is a plot of $\\dot P$ versus $P$ for the radio pulsar population.  PSRs~J1119$-$6127 and J1814$-$1744 are indicated, and we infer $B = 4.1 \\times 10^{13}$ and $5.5 \\times 10^{13}$\\,G respectively, using equation~(\\ref{eq:B}).  These are the highest magnetic field strengths yet observed among radio pulsars.  The pulsars with the next largest values of $B$ are PSRs~J1726$-$3530 $(P=1.1\\,\\mbox{s}; B = 3.7 \\times 10^{13}\\,\\mbox{G})$ and J1632$-$4818 $(P=0.8\\,\\mbox{s}; B = 2.3 \\times 10^{13}\\,\\mbox{G})$, also discovered in the multibeam survey\\footnote{See http://www.atnf.csiro.au/$\\sim$pulsar/psr/pmsurv/pmwww/pmpsrs.db.}. Prior to this survey the largest value was $B = 2.1 \\times 10^{13}$\\,G for the 2.4\\,s PSR~B0154+61 (\\cite{antt94}).  Also shown in Figure~\\ref{fig:ppdot} are the sources usually identified as magnetars, namely the five anomalous X-ray pulsars (AXPs) and two soft gamma repeaters (SGRs) for which $P$ and $\\dot{P}$ have been measured.  AXPs are characterized by X-ray periods in the range 5--12\\,s and extremely rapid spin down (\\cite{gv98}), while the SGRs exhibit occasional enormous bursts of $\\gamma$-radiation and AXP-like X-ray pulsations during quiescence.  Most models of the radio emission physics (\\cite{mt77}) depend on pair-production cascades above the magnetic poles and hence on the magnitude of the magnetic field.  However, at field strengths near or above the so-called ``quantum critical field,'' \\begin{equation} \\label{eq:Bc} B_{\\rm c} \\equiv \\frac{m_e^2 c^3}{e \\hbar} = 4.4 \\times 10^{13} \\;\\; {\\rm G}, \\end{equation} the field at which the cyclotron energy is equal to the electron rest-mass energy, processes such as photon splitting may inhibit pair-producing cascades.  It has therefore been argued (\\cite{bh98b}) that a radio-loud/radio-quiet boundary can be drawn on the $P$--$\\dot P$ diagram, with radio pulsars on one side, and AXPs and SGRs on the other (see dotted line in Fig.~\\ref{fig:ppdot}).  \\medskip \\epsfxsize=8truecm \\epsfbox{ppdot.eps} \\figcaption[ppdot.eps]{\\label{fig:ppdot} Plot of $\\dot P$ versus $P$ for radio pulsars (dots), anomalous X-ray pulsars (AXPs), and soft gamma-ray repeaters (SGRs).  PSRs~J1119$-$6127 and J1814$-$1744 are identified by large filled circles, and sources plausibly associated with supernova remnants (SNRs) are noted.  Lines of constant characteristic age and surface magnetic field strength are drawn.  The dotted line shown between the lines for $B=10^{13}$ and $10^{14}$\\,G indicates a hypothesized approximate theoretical boundary (\\cite{bh98b}) separating radio-loud and radio-quiet neutron stars due to effects relating to magnetic fields close to the critical field $B_{\\rm c}$ (see discussion following equation~[\\ref{eq:Bc}]). } \\bigskip  The existence of PSRs~J1119$-$6127, J1726$-$3530, and J1814$-$1744 demonstrates that radio emission can be produced in neutron stars with $B \\ga B_{\\rm c}$.  The radio luminosities of these objects (Table~1) are typical for observed radio pulsars.  Thus, photon splitting does not appear to inhibit radio emission at these magnetic fields, in agreement with Usov \\& Melrose (1995)\\nocite{um95} who argue that this process is inhibited by polarization selection rules.  Also, there are both astrophysical and instrumental selection effects which bias searches against the detection of long-period ($P \\ga 5$\\,s) radio pulsars such as J1814$-$1744:  evidence suggests their beams are narrower (e.g., \\cite{ymj99}), so the chances of one intersecting our line-of-sight are smaller, and instrumental high-pass filtering intended to remove baseline variations reduces the sensitivity of searches for such pulsars.  Pulsars such as J1814$-$1744 could therefore be more prevalent than present numbers suggest.  Especially noteworthy is the proximity of PSR~J1814$-$1744 to the cluster of AXPs and SGRs at the upper right corner of Figure~\\ref{fig:ppdot}.  In particular, this pulsar has a very similar $\\dot P$ to that of the well-known AXP 1E~2259+586 (\\cite{fg81}; \\cite{bsss98}), which has a period of 7\\,s.  The disparity in their emission properties is therefore surprising.  The absence of X-ray emission from the direction of PSR~J1814$-$1744, inferred from archival ASCA and ROSAT observations, implies that it must be significantly less luminous than 1E~2259+586 (\\cite{pkc00}).  The radio emission upper limit for 1E~2259+586 (\\cite{cjl94}; \\cite{llc98}) implies an upper limit on the radio luminosity at 1400\\,MHz of 0.8\\,mJy\\,kpc$^2$, $10^{-2}$ that of PSR~J1814$-$1744, assuming a distance of 4\\,kpc (\\cite{rp97}).  This limit is comparable to the lowest values observed for the radio pulsar population (\\cite{tnj+94}).  That the radio pulse may be unobservable because of beaming cannot of course be ruled out.  The radio-loud/radio-quiet boundary line displayed in Figure~\\ref{fig:ppdot} is more illustrative than quantitative (\\cite{bh98b}).  However, the apparently normal radio emission from PSRs~J1119$-$6127, J1726$-$3530, and J1814$-$1744, and the absence of radio emission from 1E~2259+586, located very close to PSR~J1814$-$1744 on a $P$--$\\dot P$ diagram (Fig.~\\ref{fig:ppdot}), suggests that it may be difficult to delineate any such boundary.  The two sources are also similar in their levels of rotational stability, at least on time scales of $\\sim 2$\\,yr: PSR~J1814$-$1744 displays timing noise in the amount expected for a radio pulsar with its $\\dot P$ (see \\S~\\ref{sec:obs}), as is the upper limit on timing noise for 1E~2259+586 over a 2--3\\,yr span (\\cite{kcs99}).  However, longer term incoherent timing of 1E~2259+586 has revealed significant deviations from a simple spin-down model.  These have been interpreted as being evidence for radiative precession of the neutron star, due to its physical distortion by the strong magnetic field (\\cite{mel99}). Alternatively, Heyl \\& Hernquist~(1999)\\nocite{hh99} suggest the deviations are due to extremely large glitches.  In either model, similar behavior might be expected of PSR~J1814$-$1744; continued radio timing will be sensitive to it.  The similar spin parameters for these two stars and, in turn, many common features between 1E~2259+586 and some other AXPs and SGRs, suggest that very high inferred magnetic field strengths cannot be the sole factor governing whether or not an isolated neutron star is a magnetar or a radio pulsar.  Other possible factors include heavy-element atmospheric composition and youth (\\cite{td93a}; \\cite{hh97}; see also Pivovaroff et al.~2000b\\nocite{pkc00}).  The age of PSR~J1814$-$1744, if $P_0 \\ll P$ and $n=3$, is 85\\,kyr.  It is unlikely that any associated supernova remnant would still be observable and indeed there is none known in the vicinity (\\cite{gre96}).  We also note that the recently proposed accretion model for AXPs (\\cite{chn00}), in which they are accreting from a fall-back disk formed from material remaining after the supernova explosion, is challenged by PSR~J1814$-$1744.  In this model, the neutron star should not be a radio pulsar, but rather an AXP progenitor in a ``dim propeller phase,'' its rotational frequency being still too high for the accreting material to overcome the centrifugal barrier.  Of course, it is always possible that in this one case no fall-back disk formed.  Proof that AXPs or SGRs are isolated high-magnetic-field neutron stars would come from either the discovery of magnetar-like emission from a radio pulsar, or radio pulsations from a putative magnetar.  While such radio emission was not expected due to theoretical considerations, because of the high inferred magnetic fields, the discovery of PSR~J1814$-$1744 shows that this emission does occur at magnetic field values characteristic of at least some magnetars, opening the possibility that magnetars also emit observable radio waves."
    },
    "0004/astro-ph0004089_arXiv.txt": {
        "abstract": "We propose a direct method to detect close-in giant planets orbiting stars in the Galactic bulge.  This method uses caustic-crossing binary microlensing events discovered by survey teams monitoring the bulge to measure light from a planet orbiting the source star.  When the planet crosses the caustic, it is more magnified than the source star; its light is magnified by two orders of magnitude for Jupiter size planets.  If the planet is a giant close to the star, it may be bright enough to make a significant deviation in the light curve of the star.  Detection of this deviation requires intensive monitoring of the microlensing light curve using a 10-meter class telescope for a few hours after the caustic.  This is the only method yet proposed to directly detect close-in planets around stars outside the solar neighborhood. ",
        "introduction": "One of the scientific goals of microlensing searches towards the Galactic bulge is the detection of planets orbiting the primary lenses.  These searches are conducted in the following manner: One of the microlensing searches, EROS (\\cite{erostrigger}) or OGLE (\\cite{udalski1994}) or the now terminated MACHO program (\\cite{alcock1996}) launches an electronic alert of an ongoing microlensing event.  These events are then monitored by follow-up groups such as the PLANET (\\cite{albrow1998}), MPS (\\cite{mps}), or GMAN (\\cite{alcock1997}) collaborations.  While the searching teams typically monitor stars $\\sim$ once per day, the follow-up campaigns, with a network of telescopes around the globe, sample much more frequently.  A planet orbiting the primary lens with semi-major axis $a$ in the ``lensing zone'', $0.6-1.5 R_E$ (\\cite{mandp1991}; \\cite{gandl1992}), where the Einstein radius, $R_E$ is \\begin{equation} R_E=\\left ( \\frac{4 G M}{c^2}\\frac{D_{L}D_{LS}}{D_S} \\right )^{1/2} \\, , \\end{equation} may cause detectable deviations from the standard microlensing light curve (see \\cite{sackett1999} for a review).  Here, $D_L$, $D_S$, and $D_{LS}$ are respectively the distances to the lens, the source star, and between the lens and source stars. For a bulge source lensed by a 0.3 $\\msun$ lens, the lensing zone, where one is most sensitive to planets, is in the range $1.2 - 3.2$ AU.  In this paper, we discuss another means by which a microlensing follow-up experiment may detect a planet, in this case a giant planet close ($a\\la0.1~{\\rm AU}$) to the {\\it source} star.  Confounding prior expectations, such close-in planets are found to be relatively common.  They have been detected by several collaborations (\\cite{mandq1995}, \\cite{cochranetal1997}, \\cite{noyesetal1997}; see \\cite{mandb1998} for a review) and can be found around $\\sim 1\\%$ of all stars (\\cite{mandb2000}).  Since direct detection is most sensitive to planets close to the source ($a\\la0.1~{\\rm AU}$), it is complementary to the traditional microlensing light curve deviation method which can only detect planets in the lensing zone ($a\\simeq 1-3~{\\rm AU}$).  Unlike other methods of planet detection, such as radial velocity measurements or astrometric shifts\\footnote{It may also be possible to indirectly detect planets orbiting stars in the Galactic bulge using occultation.}, this method can be used to directly detect planets in the bulge of our Galaxy, thus allowing us to compare planet formation under conditions different from those in the solar neighborhood.  \\subsection{Caustics}  All strong gravitational lenses produce ``caustics,'' regions in which a point source is (formally) infinitely magnified.  In reality, no source is truly point-like, and the passage of the source through the caustic allows us to spatially resolve the source.  Already, limb-darkening profiles have been measured in sources as far away as the Small Magellanic Cloud (\\cite{afonso2000}), and it has been proposed that star-spots could be imaged when a source passes through a caustic (\\cite{hands2000}; \\cite{hkcp2000}; \\cite{bryce2000}).  In this paper, we examine the possibilities of directly detecting light from a planet as it passes through a caustic.  There are several different types of caustics depending on the lens configuration.  A single point lens has a point caustic corresponding to perfect alignment between the source and the lens, when the image of the source is the Einstein ring.  These events are rare since they require such perfect alignment.  In the case of binary lenses, the caustics form a network of ``folds'' and ``cusps.''  The caustic structure of binary lenses is cataloged in Schneider \\& Weiss (1986) and Erdl \\& Schneider (1993).  Because caustics form closed curves, a caustic light curve generically has pairs of crossings.  The time of the second caustic crossing can often be predicted days in advance, allowing for scheduling of detailed monitoring of the caustic crossing.  The best studied caustic crossing event, which allows us to illustrate the technology, is MACHO-98-SMC-1 (\\cite{alcock1999}), an event which took place in the Small Magellanic Cloud.  The MACHO group first sent an alert indicating that a microlensing event was taking place.  After the first caustic crossing, the MACHO group launched a second alert with a rough prediction of when the second caustic crossing would occur.  As the second caustic crossing approached, predictions of its time became more accurate so that when it occurred, several groups were able to devote all their telescope time for one night to the observation of this one event.  Notably, the passage of the source through the caustic was imaged every $\\sim 5$ minutes by the PLANET collaboration (\\cite{albrow1999}) from South Africa, followed immediately, and with similar sampling frequency by the EROS collaboration (\\cite{afonso1998}) from Chile.  A combined analysis of this event involving data from the MACHO/GMAN, EROS, PLANET, OGLE and MPS collaborations was published in (\\cite{afonso2000}).  Several other caustic crossings have been studied with similar sampling frequency, by the MACHO/GMAN and PLANET collaborations (\\cite{alcock2000}; \\cite{albrow2000}).  Roughly 7\\% of all microlensing events are caustic crossing events (\\cite{udalski2000}).  Near a fold caustic, the magnification of a single point is (\\cite{bible}) \\begin{equation} A=A_0 + \\Theta(-u_\\perp) (u_\\perp/u_r)^{1/2} \\end{equation} where $u_\\perp$ is the distance of the source to the caustic normal to the caustic in units of the projected Einstein radius, $u_r$ is the length scale of the caustic in units of the projected Einstein radius, $A_0$ is the magnification not associated with the caustic, and $\\Theta$ is a step function.  For typical binary lenses, $u_r$ and $A_0$ are ${\\cal O}(1)$.  The magnification of a uniform disk by a fold caustic is discussed in (\\cite{sandw1987}).  The maximum magnification is $1.4 (u_r/\\rho_*)^{-1/2}$, where $\\rho_*$ is the radius of the stellar disk in units of the projected Einstein radius, while the average magnification during the crossing is $(u_0/\\rho_*)^{1/2}$.  Note that since the planet is smaller than the source star, it is more highly magnified. ",
        "conclusions": "We have shown that the presence of a giant close-in planet around the source star of a caustic crossing microlensing event could generate an order 1\\% deviation in the light curve.  This deviation would be detectable with a 10 meter class telescope making measurements every few minutes over the course of a night.  In some cases, the planet will transit the caustic more than once, allowing for complete solution of the orbit and a determination of the radius and albedo of the planet. Follow-up with 100 meter class telescopes may allow for resolution of limb darkening and spots or bands on the surface of these planets.  \\smallskip This paper profitted from useful discussions with Pierre Vermaak, Andrew Gould and Penny Sackett.  This work was supported in part by grant AST 97-27520 from the NSF. B.S.G. acknowledges the support of a Presidential Fellowship from the Ohio State University."
    },
    "0004/astro-ph0004276_arXiv.txt": {
        "abstract": "Which sample of objects can give strong constraints on single-star evolution theory ? Whilst star cluster members share the same age and the same metallicity, many questions (e.g. field stars contamination, stellar rotation, presence of unresolved binaries) are difficult to clarify properly when using their colour-magnitude diagram to compare with theoretical isochrones. Alternatively, binary stars can be used to put constraints on theoretical predictions. However, while the stellar mass is accurately known for some sample of well-detached binary stars, their metallicities are often poorly known. Then it appears that a better test could be obtained by combining both advantages (well-detached binaries members of a star cluster). This idea is applied in this work to binaries in the Hyades and one binary in the Cepheus OB3 association to test the validity of three independent sets of theoretical tracks. A detailed comparison of theoretical vs. observational masses (and radii when possible) are presented. ",
        "introduction": "Binary systems are the main source of fundamental data on stellar masses and radii. These data give stringent constraints that should be fitted by any set of theoretical stellar models. This work is exclusively focused on the mass estimate (and radius when available) of well-detached binary systems which can be presumed typical of single stars properties. Moreover three of them should share same age, chemical composition and distance due to their membership to the same open cluster. The 6 selected Hyades stars have individual stellar masses known with an accuracy of about 10\\% for 51 Tauri and $\\theta^2$ Tauri and better than 1\\% for V818 Tauri. The mass accuracy is better than 3\\% for the CW Cephei eclipsing binary (Table~\\ref{tab:par}). Moreover, these stars cover a wide mass range which is useful to obtain some interesting tests between $\\sim$0.77 and 13.5$M_{\\odot}$.  \\begin{table}[htb] \\caption[]{Cross identification and masses for the 4 selected binary systems.} \\label{tab:par} \\smallskip \\begin{tabular}{lrcl} \\hline \\hline Stars            &  HIP   & Mass            & Ref. \\\\ &        & ($M_{\\odot}$)   &      \\\\ \\tableline 51 Tau A         &  20087 &  1.80 $\\pm$ 0.13     & [TSL97a] \\\\ 51 Tau B         &  ...   &  1.46 $\\pm$ 0.18     & [TSL97a] \\\\ V818 Tau A       &  20019 &  1.072 $\\pm$ 0.010   & [PS88] \\\\ V818 Tau B       &  ...   &  0.769 $\\pm$ 0.005   & [PS88] \\\\ $\\theta^2$ Tau A &  20894 &  2.42 $\\pm$ 0.30     & [TSL97b]$^{(1)}$ \\\\ $\\theta^2$ Tau B &  ...   &  2.11 $\\pm$ 0.17     & [TSL97b]$^{(2)}$ \\\\ CW Cep A         & 113907 &  13.52 $\\pm$ 0.39    & [A91] \\\\ CW Cep B         &  ...   &  12.08 $\\pm$ 0.29    & [A91] \\\\ \\hline \\hline \\end{tabular}  $^{(1)}$ The value originally quoted by [TSL97a], 2.10$\\pm$0.60 $M_{\\odot}$, is the determination of Tomkin et al. (1995) adjusting the error upward by a factor of two. \\\\ $^{(2)}$ The value quoted by [TSL97a], 1.60$\\pm$0.40 $M_{\\odot}$, is also from Tomkin et al. (1995) adjusting the error upward by a factor of two.\\\\ \\end{table}  Comparisons between measured masses and predictions of three widely used models from the Geneva group (Mowlavi et al. 1998 and references therein), the Padova group (Fagotto et al. 1994 and references therein) and from Claret \\& Gim\\'enez (1992) (CG92 thereafter) are presented in the next sections. The theoretical quantities are derived from the isochrone technique in the colour magnitude diagram as described in Lastennet et al. (1999). ",
        "conclusions": "The main conclusions of this work are the following: \\\\ 1/ The masses predicted by the three sets of theoretical tracks, widely used in the literature, are in good agreement with the measured individual masses of each system in HIP 20087, HIP 20894 and HIP 113907. \\\\ 2/ Padova isochrones can not fit simultaneously the two stellar components of V818 Tau in the mass-radius diagram, whatever the age and metallicity assumed. The radius of the more massive component is overestimated by $\\sim$4$\\sigma$, and the radius of the less massive component is underestimated by more than 7$\\sigma$. \\\\ 2a/ No conclusion on this point for the two other models because the low mass star ($\\sim$ 0.77 $M_{\\odot}$) of V818 Tau does not allow us to test either the CG92 or the Geneva models whose lower mass limit is 0.8 $M_{\\odot}$. \\\\ 3/ The result 2/ points out that the Padova models have to be used cautiously for accurate studies in the mass range $\\sim$0.7-1.1 $M_{\\odot}$.  \\\\ 4/ It is of interest to keep in mind that mass and radius can be compared with theoretical predictions only for V818 Tau and CW Cep, the only double-lined eclipsing binaries of the sample studied in this work. If Padova tracks predict the true stellar masses of V818 Tau (but only at a 5$\\sigma$ level), they are not able to reproduce simultaneously its true masses and radii with a great accuracy. Unfortunately, this result cannot be checked with the two other Hyades binaries, hence the importance of double-lined eclipsing binaries to fully constrain stellar tracks. \\\\ In this context, accurate data of well-detached double lined eclipsing binaries (e.g. Kurpinska-Winiarska et al. 2000) are highly needed to perform further detailed comparisons. Since 1996, Oblak and collaborators (see Oblak et al. 1999) started an observational campaign of radial velocity curves for a sample of new eclipsing binaries discovered by Hipparcos, and these new data will be of great interest for a better understanding of stellar evolution."
    },
    "0004/astro-ph0004199_arXiv.txt": {
        "abstract": "We present red spectra in the region $\\sim\\lambda$7000--8300\\AA\\ of the eclipsing dwarf nova IP Peg, with simultaneous narrow-band photometry centered at 7322\\AA. We show that by placing a second star on the slit we can correct for the telluric absorption bands which have hitherto made the TiO features from the secondary star unusable. We use these TiO features to carry out a radial velocity study of the secondary star, and find this gives an improvement in signal-to-noise of a factor two compared with using the Na{I} doublet. In contrast with previous results, we find no apparent ellipticity in the radial velocity curve. As a result we revise the semi-amplitude to $K_{2}=331.3\\pm 5.8$ km s$^{-1}$, and thus the primary and secondary star masses to $1.05^{+0.14}_{-0.07}$M$_{\\odot}$ and $0.33^{+0.14}_{-0.05}$M$_{\\odot}$ respectively. Although this is the lowest mass yet derived for the secondary star, it is still over-massive for its observed spectral type. However, the revised mass and radius bring IP Peg into line with other CVs in the mass-radius-period relationships.  By fitting the phase resolved spectra around the TiO bands to a mean spectrum, we attempt to isolate the lightcurve of the secondary star. The resulting lightcurve has marked deviations from the expected ellipsoidal shape. The largest difference is at phase 0.5, and can be explained as an eclipse of the secondary star by the disc, indicating that the disc is optically thick when viewed at high inclination angles. ",
        "introduction": "IP Peg is a dwarf nova which brightens by $\\sim$2 mag every 100 days or so, with each outburst lasting $\\sim$2 weeks, and is important because it is the brightest known eclipsing dwarf nova above the period gap. The eclipses give tight constraints on the geometry of the system, and in combination with other data, such as radial velocity studies, allow the masses of the two stars to be determined. We should thus be able to compare the mass, radius and density of the secondary star with those of main sequence stars, and learn something of its structure and evolution. Unfortunately the radial velocity curve of the secondary star in IP Peg is problematical, being apparently elliptical (Martin et al. 1989, henceforth M89). The star is believed to be in a circular orbit, and the apparent ellipticity is thought to be due to irradiation of the secondary star (although we shall present an alternative explanation in Section \\ref{ecc_dis}), and an uncertain correction must be applied before the binary parameters can be determined. The result is acutely embarrassing. Smith \\& Dhillon (1998) collect together all the available masses, radii and spectral types for CVs. IP Peg is one of the few objects that does not rely on the dubious method of using the accretion disc lines to determine the white dwarf radial velocity, and so should give reliable parameters. However, in all the relationships IP Peg is a persistent offender, lying well clear of the supposedly less reliable points in both the mass vs. orbital period and mass vs. spectral type diagrams.  A further complication is presented by the lightcurve of the system. High-speed photometric observations show a light curve dominated by a bright spot and a deep eclipse. The eclipse is that of the bright spot and white dwarf by the secondary star, superposed on a more gradual disc eclipse. Unfortunately the  bright spot ingress overlaps with that of the white dwarf, making the ephemeris determination more problematical. Wood \\& Crawford (1986) were able to use their high-speed photometric observations to derive an ephemeris for IP Peg based on white dwarf egress timings. Later observations by Wood et al. (1989) and Wolf et al. (1993) show that the white dwarf egress varies markedly in strength and duration (10s to 300s).  In this paper we present red spectroscopy and photometry of IP Peg. The observations and their reduction are explained in Section \\ref{obs}. After that the paper divides into two main threads. The first is the radial velocity study and its results. The extraction of the velocities is described in Section \\ref{rad_vel} after which we discuss the summed spectra and new ephemeris (Sections \\ref{disentangle} and \\ref{ephemeris}). In Section \\ref{circularity} we discuss the absence of apparent eccentricity in our data, and conclude the system has changed in some way. We use our radial velocity semi-amplitude to derive new system parameters in Section \\ref{masses}. The second thread, isolating the lightcurve of the secondary star, is explored in Section \\ref{flux_deficits}. ",
        "conclusions": "The main conclusions from this work are as follows. \\hfil\\break (1) There is some form of variability which means that radial velocity studies of the infrared Na{I} doublet in IP Peg can sometimes return an apparently elliptical orbit. We suspect this is contamination from disc emission. \\hfil\\break (2) We strongly recommend the use of TiO rather than Na{I} for future radial velocity studies. Not only does it avoid the contamination problem described above, but it gives a factor two gain in signal-to-noise. \\hfil\\break (3) Our TiO study suggests the radial velocity semi-amplitude for the secondary star should be revised to $K_{2}=331.3\\pm 5.8$km s$^{-1}$, and as a result the primary and secondary star masses become $M_{1}=1.05^{+0.14}_{-0.07}$M$_{\\odot}$ and $M_{2}=0.33^{+0.14}_{-0.05}$M$_{\\odot}$ respectively. \\hfil\\break (4) Despite this downwards revision of the secondary star mass, it is still over-massive for its observed spectral type, but now agrees with the current semi-empirical and theoretical mass-radius-period relationships. \\hfil\\break (5) The accretion disc eclipses the secondary star in quiescence, implying that, when viewed at high inclination, it is optically thick."
    },
    "0004/astro-ph0004150_arXiv.txt": {
        "abstract": "We study the orbits in the MOND theory within a dwarf galaxy of mass $M_d\\sim 10^8M_\\odot$ at a distance of $\\sim 100$kpc from a neighboring galaxy of mass $M_g=5\\times 10^{11} M_\\odot$, such as ours. It is assumed that a second mass $m<<M_d$ is gravitationally bound to $M_d$ by a previously calculated potential for the MOND theory. This potential is obtained for a free falling mass $M_d$ in a constant external gravitational acceleration field $\\nabla\\phi_g$. The numerical technique of surfaces of section is used to study the stability of the phase-space orbits in the dwarf galaxy. Equatorial orbits with sufficiently small eccentricities $e<0.65$ are found to be stable with respect to small changes in the initial conditions. (The equatorial plane is perpendicular to the direction of $\\nabla\\phi_g$, which is along the line joining $M_d$ and $M_g$.) For decreasing values of the conserved component of the angular momentum, in the direction of $\\nabla\\phi_g$, equatorial stability is lost. ",
        "introduction": "The inconsistency between luminous and dynamical mass measurements is well known. An alternative way to explain this inconsistency without dark matter, is by Modified Nonrelativistic Dynamics (MOND) for galaxy and galactic systems \\cite{M1} \\cite{M2}, \\cite{M3}. This theory was later developed into an alternative theory of gravitation by \\cite{BM} (1984), which we hereafter  refer to as the Bekenstein-Milgrom (BM) theory. MOND agrees with the \\cite{TF} and \\cite{FJ} laws, and explains the dynamics in elliptical and spiral galaxies, without the need for dark matter. It satisfies conservation of energy, momentum, angular momentum, as well as the weak equivalence principle (\\cite{BM} 1984).  Recent studies of rotation disks of Low Surface Brightness (LSB) Galaxies by \\cite{SG}, \\cite{Block} and \\cite{McG} are in agreement with MOND.  Departure of MOND from Newtonian theory is not connected with a distance scale. In MOND, deviations from Newtonian theory become significant for very small accelerations. The main argument presented is that Newtonian gravitation has not been tested for very weak fields and it could be that the theory is not valid in this regime. Other theories modifying Newtonian theory were reviewed by \\cite{B1} (1987).  Stability of galactic disks in BM theory was first studied in a WKB approximation by \\cite{MM}. The WKB scheme deals with perturbation wavelengths $\\lambda=2\\pi/k$, much smaller than the involved distances $\\rho$ on the disk ($|k\\rho|>>1$). In Newtonian gravitation, the presence of a dark halo is important in stabilizing the disk against violent bar formation \\cite{OP}. Recent N-body simulations for the BM theory, indicate that disks are more stable in MOND than in Newtonian dynamics with dark halos (\\cite{BrdM} 1999).  We obtain the conservative Hamiltonian that describes the motion of a particle in a previous calculated potential for the BM theory in Section \\ref{2}. This Section also includes a summary of the stability theory for phase-space orbits. Orbits in a dwarf galaxy are discussed in Section \\ref{3}. Our conclusions are presented in Section \\ref{4}. ",
        "conclusions": "}  There exists a previously calculated potential $\\varphi$ in the MOND theory, given by equation (\\ref{pot}), for a free falling sphere of mass $M_d$ in a constant external gravitational acceleration $\\nabla\\phi_g$. $\\varphi$ is valid when $|\\nabla\\phi_g|>>|\\nabla\\varphi|$. We assume that the existence of a second mass $m<<M_d$ does not modify the potential.  The system $M_d-m$, bounded by the potential $\\varphi$, is described by the Hamiltonian given in equation (\\ref{H}). Since the potential is axially symmetric, only the component of the angular momentum in the direction of $\\nabla\\phi_g$, $l_z$, is conserved. As a consequence, when $z=p_z=0$ and $l_z$ is the total angular momentum, motion is periodic and occurs in a plane perpendicular to $\\nabla\\phi_g$, which we call the equatorial plane  (Fig. \\ref{esq}). Kepler's theory is then recovered with an effective gravitational constant (eq. (\\ref{lbd})), and $l_z$ is related to the eccentricity by equation (\\ref{exc}).  We analyze the phase space of $H$, using the numerical technique of surfaces of section. It is found that for eccentricities $\\epsilon<0.65$, the equatorial orbits are stable with respect to small variations in the initial conditions (Fig. \\ref{f4}). For smaller values of $l_z$, the stability of the equatorial orbit is lost and there is an increase in the size of the chaotic regions (Fig. \\ref{fg2}). There is a disk perpendicular to $\\nabla\\phi_g$, as well as some regions not in the disk, which are stable against variations of initial conditions."
    },
    "0004/astro-ph0004366_arXiv.txt": {
        "abstract": "An intrinsically narrow line emitted by an accretion disk around a black hole appears broadened and skewed as a result of the Doppler effect and gravitational redshift. The fluorescent iron line in the X-ray band at 6.4 -- 6.9~keV is the strongest such line and is seen in the X-ray spectrum of many active galactic nuclei and, in particular, Seyfert galaxies. It is an important diagnostic with which to study the geometry and other properties of the accretion flow very close to the central black hole. The broad iron line indicates the presence of a standard thin accretion disk in those objects, often seen at low inclination. The broad iron line has opened up strong gravitational effects around black holes to observational study with wide-reaching consequences for both astrophysics and physics. ",
        "introduction": "In recent years, observations with both ground-based and space-based instruments have led to realization that most, if not all, nucleated galaxies harbor a massive black hole at their center (Kormendy \\& Richstone 1995; Magorrian et al. 1998).  While many of these black holes appear to be relatively isolated, some fraction accrete significant amounts of material from the surrounding galaxy.  The angular momentum of the incoming material leads to the formation of a flattened rotating disk --- the accretion disk.  The gravitational potential energy of material flowing through the accretion disk is converted into radiative (i.e. electromagnetic) and kinetic energy. These powerful and compact energy sources, observed in approximately 1--10\\% of galaxies are termed {\\it active galactic nuclei} (AGN). AGN are also observed to be copious X-ray emitters.  These X-rays are thought to originate from the innermost regions of an accretion disk around a central supermassive black hole.  Since the accretion disk itself is expected to be an optical/UV emitter, the most likely mechanism producing the X-rays is inverse Compton scattering of these soft photons in a hot and tenuous corona that sandwiches the accretion disk.  Thus, in principle, the study of these X-rays should allow the immediate environment of the accreting black hole as well as the exotic physics, including strong-field general relativity, that operates in this environment to be probed.  This review discusses how, in the past decade, X-ray astronomy has begun to fulfill that promise.  Guided by observations with the {\\it Ginga}, ASCA, RXTE and BeppoSAX satellites, there is a broad concensus that X-ray irradiation of the surface layers of the accretion disk in a class of AGN known as Seyfert 1 galaxies gives rise to fluorescent K$\\alpha$ emission line of cold iron via the process of ``X-ray reflection''. Since this line is intrinsically narrow in frequency, the observed energy profile of the line is shaped by both special relativistic (i.e. Doppler shifting) and general relativistic (i.e.  gravitational redshifting and light bending) effects into a characteristic skewed profile predicted over a decade ago (Fabian et al 1989) and first clearly seen in ASCA data (Tanaka et al 1995).  Since these lines are typically broadened to a full-width half maximum of $5\\times 10^4\\kmps$ or more, they are often referred to as ``broad iron lines''. After discussing the physical processes responsible for the production of these spectral signatures, we will summarize the current observational status of broad iron line studies.  We will show how current observations are already addressing the nature of the accretion disk within a few gravitational radii of the black hole.  Observations of the broad iron line also provide valuable insights into the physical differences behind AGN of differing luminosities and type.  Finally, we discuss and attempt to predict the results that will emerge from high throughput X-ray spectroscopy with {\\it XMM--Newton}, {\\it Constellation-X} and {\\it XEUS}. We argue that these future data will provide unprecedented constraints on the spacetime geometry near the black hole (thereby measuring the spin of the black hole), well as the physical nature of the accretion disk. ",
        "conclusions": "In this section we summarize some of the key points of this review.  \\begin{list}{$\\bullet$}{}  \\item \\emph{X-ray continuum.} The hard X-ray continuum in AGN and GBHCs is thought to be produced in active or flaring regions in a corona above the accretion disk. Thermal electrons multiply inverse Compton scatter optical and UV photons from the disk to X-ray energies. The hard X-ray power law that results irradiates the accretion disk and produces a ``reflection'' component in the spectrum.  \\item \\emph{Reflection component.} The reflection component causes the observed spectrum to flatten above 10\\,keV as Compton recoil reduces the backscattered flux and also results in a strong iron fluorescence line at approximately 6.4\\,keV. The precise energy and strength of the line depend on a number of factors such as the iron abundance, the inclination of the disk and its ionization state. An ionized disk may also produce a strong iron absorption edge.  \\item \\emph{Iron line profile.} The line profile is determined by Doppler shifts and relativistic boosting due to the motion of the disk and the gravitational redshifting of the black hole. This produces a broad, skewed line profile. Since the line originates from the innermost regions of the accretion disk, these effects are very pronounced. From observations of the line profile the black hole spin and the inclination angle of the accretion disk may be determined. In most Seyfert~1 galaxies (i.e. AGN in which we can view the accreting black hole directly) the accretion disk is inclined at about 30$^\\circ$ to the observer. This is consistent with the standard model in which the Seyfert nucleus is surrounded by an optically thick torus with an opening angle of 30--40$^\\circ$.  \\item \\emph{Observations.} The iron line was first clearly detected by \\emph{Ginga} and a line profile subsequently resolved by \\emph{ASCA} confirming the broad and skewed shape expected from an accretion disk around a Schwarzschild black hole. \\emph{RXTE} has been able to study the line and continuum variability on much shorter timescales, although with reduced energy resolution.  \\item \\emph{Black hole spin.} The radius of the smallest stable circular orbit around the black hole decreases with the spin of the black hole. Since the line profile is sensitive to the innermost radius of fluorescent emission this may be used (with some assumptions about the astrophysics of this region) to estimate the spin of the black hole. With present time-averaged observations, however, such measurements may be ambiguous as alternative models with very different values for the black hole spin may produce almost identical line profiles.  \\item \\emph{Alternative models for the production of the broad iron line.} Models for the broad iron line that do not require a black hole accretion disk appear to fail. In particular the line width cannot be entirely due to Comptonization.  Hybrid models in which both Comptonization and Doppler/gravitational effects produce the line profile are heavily constrained.  \\item \\emph{Variability.} Rapid X-ray continuum variability is observed in most AGN and the iron line is expected to vary in response to this with a short time lag. Whilst these timescales are too short to be probed with present instruments, significant and complex iron line variability has been observed. Curiously, the line flux is seen to remain constant whilst the continuum changes, and there appears to be an anti-correlation between the reflected fraction and the equivalent width of the line.  In another study the reflected fraction and the photon index of the power law are correlated, both for an individual object, and between different objects (including both AGN and GBHC). Such observations need to be explained, especially since they appear contrary to our simple model of reflection. Flux-correlated changes in the ionization state of the disk may explain some of these facts.  \\item \\emph{Reverberation mapping.} The rapid X-ray variability is associated with the activation of new flares in the corona above the accretion disk. X-ray reverberation mapping is the technique of using observations of the iron line response, or ``echo'', to sudden changes in the continuum to study the accretion disk and black hole. In principle this may be used to determine the geometry of the X-ray emission and the black hole spin and mass. Such observations will be within the capabilities of the next generation of X-ray observatories.  \\item \\emph{Other classes of object.} Iron lines are observed in other classes of object in addition to Seyfert galaxies. In quasars the strength of the iron line decreases with increasing luminosity. This may be because the more luminous sources accreting closer to the Eddington limit and more highly ionized. The observation of iron lines in LLAGN may determine whether the accretion with low rates is an ADAF or a thin disk.  Weak iron lines have also been seen in LLAGN suggesting their low accretion rate flows are thin disks as opposed to geometrically thick ADAFs.  In radio loud AGN, broad iron lines and reflection humps are weak or absent, perhaps because the reflection signature is swamped by a beamed continuum.  All of these require further detailed observations. In GBHC the accretion disk is ionized and the reflection spectra show smeared absorption and emission feature, and there is debate as to the precise nature of the accretion flow within a few tens of Schwarzschild radii of the black hole.  \\end{list}  Over the past decade observations of the broad iron line have provided an unprecedented probe of the region within a few tens of Schwarzschild radii of the black hole event horizon. The next generation of X-ray observatories, beginning with \\emph{XMM--Newton}, will address many of the puzzling questions we have, and significantly enhance our understanding of these enigmatic objects."
    },
    "0004/astro-ph0004016_arXiv.txt": {
        "abstract": " ",
        "introduction": "The origin of the highest energy cosmic rays is a problem that has persisted for 4 decades since the pioneering measurements at Volcano Ranch \\cite{L-S,Linsley}.  There is some consensus that, above the spectrum's ankle at about $10^{18.5}$ eV, they originate outside the disk of the Galaxy.  For particles of such high magnetic rigidity, sources in the Galaxy's disk would presumably cause an obvious anisotropy in arrival directions that is not observed.  Evidence for a composition changing to lighter particles at the ankle \\cite{FE} strengthens this argument (particles having even greater rigidity because of lesser charge) and supports the view that cosmic rays with energies above the ankle are of extragalactic origin.  The sources of those particles remain to be identified.  The observations of cosmic rays with energies above the expected GZK cutoff \\cite{Greisen,Z-K} should be a powerful clue to the nature of the sources.  The Fly's Eye \\cite{Big1} and AGASA \\cite{Big2} measured air showers with energies well above the GZK threshold.  Recent reports \\cite{AGASA,Jui,Watson} suggest that the spectrum might continue without a strong GZK effect.  These super-GZK results have posed several related, but distinguishable, puzzles:  \\begin{enumerate} \\item{How are particles produced with such prodigious energy?} \\item{Why do the arrival directions of those particles not point back to recognizable sources in our local part of the universe?} \\item{Why is the intensity of particles above $6\\times 10^{19}$ eV not more strongly suppressed?} \\end{enumerate}  New attempts to improve the observational data include the recently commissioned High Resolution Fly's Eye (HiRes) \\cite{HiRes}, the Pierre Auger Observatory \\cite{Auger}, the proposed Telescope Array Project \\cite{TA}, and Airwatch/OWL \\cite{OWL} plans for future space-based detectors of atmospheric air showers.  The Auger Observatory and the potential space-based detectors will have exposure to the entire sky, which will open new possibilities for anisotropy analysis.  These methods will be explored here in the context of the better controlled exposure of the Auger surface arrays.  Hillas \\cite{Hillas1} pointed out years ago that there are few astrophysical sites that can produce a large enough ``electrical potential'' $\\beta B L$ to accelerate even highly charged nuclei to $10^{20}$ eV.  (Here $\\beta=v/c$ is the relative velocity of moving media with magnetic field strength $B$ and size $L$.  In the case of statistical acceleration, including Fermi shock acceleration, the same product $\\beta B L$ governs the maximum particle rigidity even though particles do not pass through a monotonic change in electrical potential.)  Puzzle \\#1 is exacerbated in many contexts by synchrotron radiation and/or pion photoproduction.  Puzzle \\#2 can be resolved by invoking stronger-than-expected extragalactic magnetic fields, but that does not readily simplify puzzle \\#3.  The expected suppression of particles above the GZK cutoff is based on travel time (or total distance traveled) rather than the straight-line distance to the sources.  Particles below the GZK threshold have been accumulating over billions of years, whereas the mean age of particles well above the threshold cannot be greater than tens of millions of years \\cite{Elbert}.  One possible inference from the lack of an observable GZK spectral break is that sub-GZK particles may not be much older than 30 million years either, in which case the GZK effect would not significantly suppress the cosmic rays above the threshold relative to those below.  It is not reasonable to suppose that the sources of high energy cosmic rays turned on so recently, so the young age of sub-GZK cosmic rays would require an intergalactic mechanism for dissipating their energy.  Known mechanisms (e.g. nuclear interactions, synchrotron radiation, $e^{\\pm}$ production, pion photoproduction via infrared and visible background photons, etc.) do not rob energy from sub-GZK particles rapidly enough.  Perhaps high energy cosmic rays are attenuated through interactions with the unknown dark matter of the universe \\cite{Sommers}.  More conventional approaches to puzzle \\#3 are to defeat the GZK cutoff with a very hard extragalactic spectrum (e.g. from topological defect annihilation \\cite{Schramm}) or to evade it by invoking neutrinos \\cite{Yoshida,Weiler} or non-standard particles \\cite{Farrar1,Farrar2} that are immune to the microwave background radiation.  Others \\cite{Berezinsky,Hillas2} conclude that the sources must be localized to the Galaxy, but distributed in a halo large enough that galactic anisotropy has not become obvious.  The cumulative cosmic ray observations at this time are not sufficient to sort out the possibilities.  AGASA and HiRes are currently building up the world's total exposure at the highest energies.  With better statistics and better measurements, the observations could soon lead to a breakthrough that identifies the sources of the highest energy cosmic rays.  This same hope has been expressed for decades through the course of numerous experiments, however, and the puzzles have only become deeper mysteries.  The answers may not come easily, and we should prepare the best possible analyses of the energy spectrum, particle mass distribution, and arrival directions.  Careful determinations of the energy spectrum and mass composition can be used to weed out classes of theories, but these tools are not likely to yield a clear signature for picking out a unique theory.  A positive identification of the cosmic ray sources requires seeing their fingerprint in the sky.  This may come in the form of arrival direction clusters \\cite{Clusters,Chi} that identify discrete sources, or it may come as a large-scale celestial pattern that characterizes a particular class of potential sources.  In the worst case, we might discover that the arrival directions are isotropic and the sources still elude positive identification.  In that case, observers must strive for the best possible upper limit on anisotropy.  The Auger Project's surface arrays will provide the best search for anisotropy fingerprints.  Their combined exposure function on the celestial sphere will be unambiguous because they operate continuously and are not sensitive to atmospheric variability. Continuous operation means the celestial exposure function is uniform in right ascension.  By having observatories appropriately located in both the southern and northern hemispheres, the exposure does not vary strongly with declination either.  The methods described in this paper are applicable to any observatory with full-sky coverage.  They are not limited to the Auger Project, although the specific Auger site locations are used in the example simulations that are reported here.  Coverage of the full sky could be achieved piecemeal by combining results from different experiments.  There is serious risk of spurious results from such meta-analyses, however, unless the exposures, energy resolutions, and detector systematics are perfectly understood and correctly incorporated in the analysis. The reliable approach is to use identical detectors in both hemispheres or the same (orbiting) detector for both hemispheres.  This paper seeks to evaluate the sensitivity of a full-sky observatory to large-scale anisotropy patterns and how that sensitivity depends on the number of arrival directions in a data set.  Large-angle fingerprinting will be needed if there are many contributing sources or if the flux from each single source is diffused over a large solid angle due to magnetic deflection of the charged cosmic rays.  If, instead, there are point sources to be detected, then the advantage of a full-sky observatory is in mapping the entire celestial sphere with comparable sensitivity in all regions.  Full-sky coverage is crucial for large-scale anisotropy analysis. It makes it possible to do integrals over the sky, so the powerful tools of multipole moments and angular power spectra are available.  With full sky coverage, cosmic ray anisotropy analysis will be similar to gamma ray burst anisotropy analysis. The numbers of events will be comparable, the direction error boxes will be comparable, the exposure non-uniformities will be comparable, and in both cases events come from all parts of the sky.  All of the techniques that were employed to search for anisotropy in the BATSE data \\cite{Batse1,Batse2} can be applied to a full-sky cosmic ray data set.  With a cosmic ray detector in only one hemisphere, there is a solid angle hole in the sky where the detector has zero exposure despite the Earth's rotation.  A zero-exposure hole makes it impossible to do integrals over the whole celestial sphere.  No matter how many events the detector collects overall, it will never determine any multipole moment.  A single-hemisphere detector can test hypotheses like, ``Does the observed distribution match better what would be accepted from the clustering of radio galaxies toward the supergalactic plane or what would be accepted from an isotropic distribution?''  It can also make qualified measurements like, ``Assuming the anisotropy is a perfect quadrupole with axial symmetry, fit for the axis orientation that best explains the observed celestial distribution.''  The role of an observatory, however, should be to map the sky and make results available in a form which is readily usable without knowledge of the detector properties and which is independent of any theoretical hypothesis.  Low-order multipole tensors (or spherical harmonic coefficients) can summarize the large-scale information.  The angular power spectrum reveals if there is clumpiness on smaller scales.  These results can be tabulated so that theorists can test arbitrary models quantitatively without privileged access to the data.  With approximately uniform exposure, even eyeball inspection of arrival direction scatter plots can show large-scale patterns that are hidden when steep exposure gradients dominate the scatter plots.  While the role of an observatory should be to map the sky and determine the patterns without preconceived expectations, it is nevertheless worthwhile to consider what might be learned by measuring the low order multipoles or the angular power spectrum.  \\vspace{.3cm} {\\it Monopole.}  There is no information about anisotropy patterns in the monopole scalar by itself.  It is simply the sky integral of the cosmic ray intensity.  That is information already present in the energy spectrum.  A pure monopole intensity distribution is equivalent to isotropy.  The strength of other multipoles relative to the monopole is a measure of anisotropy.  \\vspace{.3cm} {\\it Dipole.}  A pure dipole distribution is not possible because the cosmic ray intensity cannot be negative in half of the sky.  A ``pure dipole deviation from isotropy'' means a superposition of monopole and dipole, with the intensity everywhere $\\geq 0$.  A predominantly dipole deviation from isotropy might be expected if the sources are distributed in a halo around our Galaxy, as has been suggested \\cite{Berezinsky,Hillas2}.  In this case, there is a definite prediction that the dipole vector should point toward the galactic center.  An approximate dipole deviation from isotropy could be caused by a single strong source if magnetic diffusion or dispersion distributes those arrival directions over much of the sky.  In general, a single source would produce higher-order moments as well.  A dipole moment is measurable in the microwave radiation due to Earth's motion relative to the universal rest frame \\cite{COBE}.  If we are moving relative to the cosmic ray rest frame, a dipole moment should exist also in the cosmic ray intensity (the Compton-Getting effect).  At lower energies, this may occur if the sun and Earth are moving relative to the galactic magnetic field or if the cosmic rays are not at rest with respect to the galactic field.  For extragalactic cosmic rays, a Compton-Getting dipole is expected if the Galaxy is moving relative to the intergalactic field or if the cosmic rays themselves are streaming in intergalactic space.  In any case, the expected velocities would be small ($v/c \\simlt 10^{-3}$), and the Compton Getting anisotropy \\cite{Longair} ($\\frac{I_{max}-I_{min}}{I_{max}+I_{min}}=(\\gamma+2)(v/c)$) should be $\\simlt 0.005$.  (Here $\\gamma$ is the differential spectral index, which is roughly 3.)  An anisotropy of one-half percent would require high statistics for detection (see section 3).  A larger dipole anisotropy might be produced by a cosmic ray density gradient.  If the magnetic field is disorganized, the gradient produces streaming by diffusion and the Compton-Getting dipole vector is parallel to the density gradient.  If there is a regular magnetic field, however, the expected dipole vector $\\vec{D}$ can be perpendicular to both the gradient and the field direction, $\\vec{D} \\propto \\vec{\\nabla} \\rho \\times \\vec{B}$. The direction of strongest intensity corresponds to the arrival direction of particles whose orbit centers are located in the direction of increased density.  \\vspace{.3cm} {\\it Quadrupole.}  An equatorial excess in galactic coordinates or supergalactic coordinates would show up as a prominent quadrupole moment.  A measurable quadrupole is expected in many scenarios of cosmic ray origins, and is perhaps to be regarded as the most likely result of a sensitive anisotropy search.  In general, a quadrupole tensor is characterized by 3 relative eigenvalues with associated orthogonal eigenvectors.  In the case of axial symmetry, there is a single non-degenerate eigenvector that gives the symmetry axis.  An axisymmetric ``prolate'' distribution would be hot spots at antipodal points of the sky, whereas an ``oblate'' distribution has the excess concentrated toward the equator that is perpendicular to the symmetry axis.  The axis of an oblate quadrupole distribution might differ from the galactic axis or the supergalactic axis if we are embedded in a magnetic field that systematically rotates the arrival directions.  \\vspace{.3cm} {\\it The angular power spectrum.}  Spherical harmonic coefficients for a function on a sphere are the analogue of Fourier coefficients for a function on a plane.  Variations on an angular scale of $\\theta$ radians contribute amplitude in the $\\ell=1/\\theta$ modes just as variations of a plane function on a distance scale of $\\lambda$ contribute amplitude to the Fourier coefficients with $k=2\\pi/\\lambda$.  For cosmic ray anisotropy, we might look for power in modes from $\\ell = 1$ (dipole) out to $\\ell \\sim 60$, higher order modes being irrelevant because the detector will smear out any true variations on scales that are smaller than its angular resolution.  For charged cosmic rays, magnetic dispersion will presumably smear out any point source more than the detector's resolution function.  Even at the highest observed particle energies, there is unlikely to be any structure in the pattern of arrival directions over angles smaller than $3^{\\circ}$.  The interesting angular power spectrum is therefore probably limited to $\\ell \\simlt 20$. ",
        "conclusions": "There is great advantage in a cosmic ray observatory having exposure to the entire celestial sphere, especially if the relative exposure is nearly uniform.  In that case, scatter plots of arrival directions are immediately interpretable, and eyeball evaluations can readily identify discrete sources or large-scale patterns.  Discrete sources will be identified with equal sensitivity anywhere in the sky.  If no such sources are found, the flux upper limits will be uniform over the sky.  At the highest energies there is no proven anisotropy.  Unlike the COBE anisotropy analysis, it is not necessary to subtract a large known dipole pattern and a myriad of uninteresting foreground sources.  Any cosmic ray deviations from isotropy will be of immediate interest.  The search for cosmic ray anisotropy is more similar to the case of gamma-ray bursts, where expectations and early claims of anisotropy were not supported by additional data.  The role of an observatory is to map the sky and make the results available to the scientific community.  This is highly challenging for an observatory without full-sky coverage. Measurements in that case are made with different sensitivity in different parts of the sky, and nothing at all can be said about a large hole where the exposure is zero.  Certainly it is not possible to perform the full-sky integrations that are required to measure the multipoles of the celestial cosmic ray intensity.  In this paper, frequent use is made of the inverse of the relative exposure, $1/\\omega$.  Such methods obviously fail if the relative exposure anywhere becomes infinitesimal or zero.  To underscore the difficulty of anisotropy analysis without full-sky coverage, one can cite the work by Wdowczyk and Wolfendale \\cite{Wolfendale}.  In that paper, the authors argue that the cosmic ray intensity measurements support a model of excess arrivals from equatorial galactic latitudes.  The argument is based on the same data that two experimental groups had previously used in support of a gradient in galactic latitude that suggested an excess from southern latitudes relative to northern latitudes.  In effect, because those northern detectors had poor exposure for southern galactic latitudes, Wdowczyk and Wolfendale were able to argue that the data supported a quadrupole distribution rather than a dipole distribution. Neither a dipole moment nor a quadrupole moment can be measured without full-sky coverage.  The techniques outlined in this paper pertain to any full-sky detector.  Non-uniformity in celestial exposure is not hard to handle, provided it is well determined and there is adequate exposure to all parts of the sky.  The true cosmic ray intensity is mapped with a sensitivity that depends primarily on the total number of detected arrival directions.  This number $N$ is related to particle energy and observing time for any detector of known acceptance.  This relationship for the Auger Observatory was given in the Introduction.  While complete information about anisotropy is encoded in the $a_{\\ell m}$ coefficients (tied to some specified coordinate system), important gross properties of the anisotropy are characterized by the (coordinate independent) angular power spectrum $C(l)$.  One can tell, for example, if there is a strong dipole or quadrupole moment.  Such large-scale patterns are expected in many theories.  It should be noted, however, that $C(1)$ gives the dipole moment but not its direction.  It is obviously important whether the dipole points toward the galactic center, toward Virgo, toward Cen A, or in some unexpected direction. Similarly, all components of the quadrupole tensor are of interest, not just the average of their squares, $C(2)$. Nevertheless, the angular power spectrum provides a powerful tool for discriminating between viable and non-viable theories without detailed investigation.  Also, the higher-order moments of the angular power spectrum can quantitatively characterize whatever clumpiness may exist in a map of arrival directions.  The techniques and examples mentioned in this paper are only representative of the powerful analysis methods that become possible with full-sky observatories.  Data sets from such observatories will open a rich field of anisotropy study.  The primary goal seen from the present time is the discovery of the highest energy cosmic ray origins.  If that objective is accomplished (perhaps even before full-sky data sets are available), then the observed patterns from a known source distribution will be analyzed to infer properties of magnetic fields in the galactic halo and in intergalactic space.  A full-sky observatory has the ability to summarize completely the anisotropy information with the use of the spherical harmonic expansion coefficients $a_{\\ell m}$.  A table of coefficients (perhaps with multiple columns for different energy cuts) will provide the whole story.  As has been done in this paper, those coefficients will be reliably corrected for the observatory's non-uniform exposure.  Detailed anisotropy analysis will no longer require privileged access to detector data.  The published anisotropy fingerprint encoded in the $a_{\\ell m}$ spherical harmonic coefficients can be matched against any theoretical suspect by any interested investigator.  \\vspace{1.5cm} \\noindent {\\it Acknowledgement.}  It is a pleasure to thank Brian Fick for many helpful discussions about these matters.  \\vspace{1.5cm}"
    },
    "0004/astro-ph0004220_arXiv.txt": {
        "abstract": " ",
        "introduction": "VLBI Exploration of Radio Astrometry (VERA), approved to start its construction in 2000, is a new Japanese VLBI array of 20m antennas dedicated to astrometry (Honma et al. 2000; Kawaguchi et al. 2000). VERA will undertake astrometry of Galactic maser sources based on the differential VLBI technique, in which two adjacent sources are observed simultaneously to remove the atmospheric fluctuation. A position accuracy of $\\sim 100$ $\\mu$arcsec has been already achieved by switching differential VLBI in which an object and a reference are observed by turns (e.g., Lestrade et al. 1999), but the atmospheric fluctuation is not completely suppressed because of the time lag between observations of an object and a reference. In order to obtain a higher position accuracy, one needs to observe an object and a reference source at the same time. VERA, with its dual beam antenna system, will enable us to observe a Galactic maser and an extra-galactic reference source simultaneously. With such a system, positions of maser sources relative to reference sources can be measured with 10 $\\mu$arcsec level accuracy, allowing us to obtain parallaxes of maser sources in the whole Galaxy.  To remove the atmospheric fluctuation effectively, a reference source must be sufficiently close to the maser source for which one intends to measure the parallax and proper motion. In order for VERA to achieve a 10 $\\mu$arcsec level accuracy, it is required that a maser and a reference source must be located within 2$^\\circ$ from each other. This requires a large number of reference sources spreading over the whole sky area. Currently, there exist several catalogs that contain hundreds of continuum sources detected with VLBI (e.g., Preston et al. 1985; Morabito et al. 1986; Johnston et al. 1995; Ma et al. 1998; Peck \\& Beasley 1998). In total, more than 2000 VLBI sources are known. However, the number of VLBI sources is considerably smaller in the Galactic plane than that in the off-plane region. This is because several surveys of VLBI sources were conducted in the off-plane regions: for example, the VLBI survey by Preston et al.(1985) are for $|b|\\ge 10^\\circ$, and the VLBA calibrator survey by Peck \\& Beasley (1998) are mainly based on the Jodrell Bank-VLA Astrometric Survey (Patnaik et al. 1992; Browne et al. 1998; Wilkinson et al. 1998), which are conducted at $|b|\\ge 2.5^\\circ$. On the other hand, H$_2$O and SiO masers, which are the main targets of VERA, are mostly concentrated in the Galactic plane because they are emitted from Galactic star forming regions and Mira-type variable stars: for example, about 40\\% of H$_2$O maser sources listed in Palagi et al.(1993) are located within $b=\\pm 5^\\circ$. Therefore, a large number of VLBI sources are necessary as phase and position reference for VERA, and hence the search for VLBI sources in the Galactic plane is indispensable to the success of VERA.  A search for VLBI sources in the Galactic plane has its own scientific merits even apart from the necessity for VERA. First, many of the continuum radio sources are likely to be radio galaxies and QSOs, and so a search for radio sources in the Galactic plane will provide a basis for studying the distribution of galaxies and large scale structures behind the Milky Way. Second, a search for VLBI sources in the Galactic plane may lead us to discoveries of new active Galactic sources like X-ray binaries or flare stars which are of significant interest in stellar physics and high energy astrophysics. Third, the radio sources can be used to study the cold interstellar matter in the Galaxy by observing the absorption line systems toward them, and also can be used to investigate the ionized interstellar matter by evaluating the interstellar scintillation effect on these sources. Hence, a search for radio sources in the Galactic plane is useful not only for VERA, but also for a number of interesting studies in astronomy and astrophysics. For these reasons, we have performed a survey of continuum radio sources in the Galactic plane using the Japanese VLBI Network (hereafter J-Net; Kameya et al. 1993; Shibata et al. 1994; Omodaka et al. 1994) at 22GHz, and here we report on the results. ",
        "conclusions": "We have searched for new VLBI sources in the Galactic plane using J-Net. We have observed 267 sources that are selected from single dish and low resolution interferometer observations, and detected 93 sources including 51 sources that are detected with VLBI for the first time. We have presented the correlated flux for Nobeyama-Kashima baseline for all sources we detected, and also positions for 21 sources that are detected with 3 or more baselines. These VLBI sources are located within $|b|\\le 5^\\circ$ of the Galactic plane, and there are a large number of maser sources within 2$^\\circ$ of them. Thus, the sources detected in the present study are potential candidates for the reference sources of phase-referencing VLBI astrometry with VERA."
    },
    "0004/astro-ph0004293_arXiv.txt": {
        "abstract": "We present observations of radio recombination lines (RRL) from the starburst galaxy Arp 220 at 8.1 GHz (H92\\a) and 1.4 GHz (H167\\a\\- and H165\\a) and at 84 GHz (H42\\a), 96 GHz (H40\\a) and 207 GHz (H31\\a) using the Very Large Array and the IRAM 30 m telescope, respectively. RRLs were detected at all the frequencies except at 1.4 GHz where a sensitive upper limit was obtained. We also present continuum flux measurements at these frequencies as well as at 327 MHz made with the VLA. The continuum spectrum, which has a spectral index $\\alpha \\sim -0.6$ ($S_\\nu \\propto \\nu^{\\alpha}$) between 5 and 10 GHz, shows a break near 1.5 GHz, a prominent turnover below 500 MHz and a flatter spectral index above 50 GHz.  We show that a model with three components of ionized gas with different densities and area covering factors can consistently explain both RRL and continuum data.  The total mass of ionized gas in the three components is $3.2\\times 10^7$ \\msun\\- requiring $3\\times 10^5$ O5 stars with a total Lyman continuum (Lyc) production rate \\nlyc $\\sim 1.3\\times 10^{55}$ photons s$^{-1}$.  The ratios of the expected to observed Br\\a\\- and Br$\\gamma$ fluxes implies a dust extinction $A_V \\sim 45$ magnitudes. The derived Lyc photon production rate implies a continuous star formation rate (SFR) averaged over the life time of OB stars of $\\sim$ 240 \\msun\\- $yr^{-1}$.  The Lyc photon production rate of $\\sim$ 3\\% associated with the high density HII regions implies similar SFR at recent epochs ($t<10^5$ yrs).  An alternative model of high density gas, which cannot be excluded on the basis of the available data, predicts ten times higher SFR at recent epochs. If confirmed, this model implies that star formation in Arp 220 consists of multiple starbursts of very high SFR (few $\\times 10^3$ \\msun\\- yr$^{-1}$) and short durations ($\\sim 10^5$ yrs).  The similarity of IR-excess, $L_{IR}/L_{Ly\\alpha}\\sim 24$, in Arp 220 to values observed in starburst galaxies shows that most of the high luminosity of Arp 220 is due to the on-going starburst, rather than produced by a hidden AGN. A comparison of the IR-excesses in Arp 220, the Galaxy and M33 indicates that the starburst in Arp 220 has an IMF which is similar to that in normal galaxies and has a duration longer than $10^7$ yrs. If there was no infall of gas during this period, then the star formation efficiency (SFE) in Arp 220 is $\\sim$50\\%. The high SFR and SFE in Arp 220 is consistent with their known dependences on mass and density of gas in star forming regions of normal galaxies. ",
        "introduction": "Arp 220 is the nearest ($d$ = 73 Mpc for $H_0$ = 75 \\kms Mpc$^{-1}$) example of an ultra-luminous infrared galaxy (ULIRG). ULIRGs are a class of galaxies with enormous infrared luminosities of $L_{IR} \\ge 10^{12}L_\\odot$, believed to have formed through mergers of two gas-rich galaxies (Sanders et al 1988).  What powers the high luminosity of ULIRGs is not well understood, although there is considerable evidence that intense starburst which may be triggered by the merger process may play a dominant role (Sanders and Mirabel 1996, G\\\"{e}nzel et al 1998). There is also evidence in many ULIRGs for a dust-enshrouded AGN which could contribute significantly to or in some cases dominate the observed high luminosity (G\\\"{e}nzel et al 1998). In the case of Arp 220, there is clear evidence for a merger in the form of extended tidal tails observed in the optical band (Arp 1966) and the presence of a double nucleus in the radio and NIR bands with a separation of $\\sim 1''$ (370 pc)(Norris 1985, Graham et al 1990). Prodigious amounts of molecular gas ($M_{H_2} \\sim 10^{10} M_\\odot$; Scoville, Yun \\& Bryant 1997, Downes \\& Solomon 1998) are present in the central few hundred parsecs.  Detection of CS emission (Solomon, Radford \\& Downes 1990) indicates that the molecular densities may be high ($n_{H_2}\\sim 10^5$ \\cmthree).  Recent high resolution observations reveal that the two nuclei in Arp 220 consists of a number of compact radio sources - possibly luminous radio supernovae (Smith et al 1998) - and a number of luminous star clusters in the near IR (Scoville et al 1998). High resolution observations of OH megamasers (Lonsdale et al 1998) reveal multiple maser spots with complex spatial structure. These observations have led to the suggestion that the nucleus of Arp 220 is dominated by an intense, compact starburst phenomena rather than the presence of a hidden AGN (Smith et al 1998, Downes \\& Solomon 1998). The presence of AGNs in the two nuclei is, however, not ruled out since large velocity gradients, $\\sim 500$ \\kms\\- over $\\sim 1''$ at different position angles across each nuclei (Sakamoto et al 1999) have been observed, indicating high mass concentrations at the location of both nuclei. Obscuration due to dust in the nuclear region of Arp 220 is extremely high. Ratios of fine structure lines observed with the ISO satellite (Lutz et al 1996, G\\\"{e}nzel et al 1998) indicate a lower limit of $A_V \\sim$ 45 mag. This high extinction seriously hampers measurements even at IR wavelengths (see also Scoville et al 1998).   Radio recombination lines (RRLs) do not suffer from dust obscuration and thus may potentially provide a powerful method of studying the kinematics, spatial structure and physical properties of ionized gas in the nuclear region of galaxies.  RRLs have been detected, to date, in about 15 Galaxies (Shaver, Churchwell and Rots 1977, Seaquist and Bell 1977, Puxley et al 1991, Anantharamaiah et al 1993, Zhao et al 1996, Phookun, Anantharamaiah and Goss, 1998).  Arp 220 is the most distant among the galaxies that are detected in recombination lines (Zhao et al 1996). The detection of an RRL in Arp 220 (Zhao et al 1996) opens up the possibility of studying the ionized gas in the nuclear region.  Zhao et al (1996) presented VLA observations of the H92\\a\\- line ($\\nu_{rest}$ =8.309 GHz) with an angular resolution of $\\sim 4''$ and a velocity coverage of $\\sim 900$ \\kms.  This paper presents new H92\\a\\- VLA observations with an angular resolution of $\\sim 1''$ and a velocity coverage of $\\sim 1700$ \\kms.  We also present observations of millimeter wavelength RRLs H42\\a\\- ($\\nu_{rest}$ =85.688 GHz), H40\\a ($\\nu_{rest}$ = 99.022 GHz) and H31\\a\\- ($\\nu_{rest}$ = 210.501 GHz) made using the IRAM 30 m radio telescope in Pico Veleta, Spain. Additionally, we report a sensitive upper limit to lower frequency RRLs H167\\a\\- ($\\nu_{rest}$ = 1.40 GHz) and H165\\a ($\\nu_{rest}$ = 1.45 GHz ) obtained using the VLA. These observations together with the observed continuum spectrum are used to constrain the physical properties and the amount of ionized gas in the nuclear region of Arp~220. Since the Lyman continuum photons which ionize the gas are generated by young O and B stars with a finite life time, the derived properties of the ionized gas are  used to constrain the properties of the starburst in Arp 220.   The larger velocity coverage used for the present H92\\a\\- observations showed that the previous observation with half the bandwidth (Zhao et al 1996) had underestimated both the width and the strength of the H92\\a\\- line. The integrated line intensity is found to be more than a factor of two larger than reported earlier and thus the deduced number of Lyman continuum photons that are required to maintain the ionization is also larger.  The millimeter wavelength lines are found to be much stronger than  expected based on models for the H92\\a\\- line. The observed strength of the millimeter wavelength RRLs indicate the presence of a higher density ($n_e > 10^5$ \\cmthree) ionized gas component and leads to information about star formation rates at recent epochs ($t < 10^5$ yrs).  The upper limits to the lower frequency lines provide significant constraints on the amount of low density ($n_e < 10^3$ \\cmthree) ionized gas in the nuclear region of Arp 220. A comparison of the predictions (based on RRLs) of the strengths of NIR recombination lines (e.g. Br\\a\\- and Br\\g lines) with observations (e.g. Goldadder et al 1995, Sturm et al 1996) leads to a revised estimate of the dust extinction.  This paper is organized as follows. In Section 2, observations with the the VLA and the IRAM 30 m telescope are described and the results are presented. Section 3 summarizes the available radio continuum measurements of Arp 220 in the frequency range 150 MHz to 230 GHz including new measurements at 327 MHz, 1.4 GHz, 8.3 GHz, 99 GHz and 206 GHz. In Section 3, we discuss the various possible models to explain the observed RRLs and continuum in Arp 220 and arrive at a three-component model, consistent with the observations.  Section 4  discusses the implications of the three-component model to several aspects of the starburst in Arp 220.  We attempt to answer the question of whether the deduced star formation rates in Arp 220 can power the observed high luminosity  without invoking the presence of an obscured active nucleus. The paper is summarized in Section 5. ",
        "conclusions": "\\subsection{Density}  As discussed in the previous section, only models with densities $> \\sim 10^3$ \\cmthree\\- can fit the recombination line data. Lower density models are inconsistent with the upper limit to the recombination line strength near 20 cm. While densities in the range $ 10^3 - 2.5\\times 10^4$ \\cmthree\\- fit the H92\\a\\- line, even higher densities (\\ne$\\sim 1-5\\times 10^5$ \\cmthree) are required to account for the millimeter wavelength recombination lines. These results confirm the point made by Zhao et al (1996) that recombination lines at different frequencies act as ``density filters''. Thus multi-frequency RRL observations provides an excellent method of determining the various density components in starburst regions.  That a substantial  fraction of the gas in the nuclear regions of Arp 220 is at densities $>10^4$ \\cmthree\\- is evident from the high ${\\rm L_{HCN}/L_{CO}}$ ratio observed by Solomon, Downes and Radford (1992). The total mass of H$_2$ gas at densities higher than $10^4$ \\cmthree\\- is $>10^9$ \\msun\\- or more than 25\\% of the dynamical mass of the two nuclear components (Downes \\& Solomon 1998). Less than $10^{-4}$ of this gas needs to be in ionized form to account for the high density ionized gas deduced from millimeter wavelength recombination lines. The bulk of the ionized gas ($\\sim 3\\times 10^7$ \\msun) which is in component A1 (Table 9) is at a density $\\sim 10^3$ \\cmthree\\- or higher.  Electron densities inferred from [S {\\small III}] and [Ne {\\small III}] IR line ratios observed with the ISO satellite typically yield lower values ($n_e \\sim 100 - 300$ \\cmthree) (Lutz et al 1996).  We suggest that this is a selection effect caused by the insensitivity of the [S {\\small III}] line ratios to densities outside the range $10^2 - 10^{3.5}$ \\cmthree\\- (Houck et al 1984). The upper limit to the H167\\a\\- line near 1.4 GHz presented in Table 2, implies severe constraints on the amount of low-density (i.e. \\ne $<$ 500 \\cmthree) ionized gas present in Arp 220. Fig 11 shows the expected strength of radio recombination lines as a function of frequency from ionized gas with $n_e = 500$ \\cmthree\\- which is ionized by a Lyman continuum luminosity equal to that deduced from Br\\a\\- and Br\\g\\- recombination lines observed with ISO (i.e. \\nlyc = $1.3 \\times 10^{55}$ s$^{-1}$, G\\\"{e}nzel et al 1998). The predicted strength of the H167\\a\\- recombination line is five times the 3$\\sigma$ upper limit in Table 2. Thus the only way to consistently explain the ISO and radio results is if the IR recombination lines (Br\\a\\- and Br\\g) also arise in higher density gas (i.e. $n_e > 10^3$ \\cmthree).  The three component model in Table 9 shows that the total number of Lyman continuum photons, \\nlyc, is $\\sim 1.3\\times 10^{55}$ photons s$^{-1}$.  This value of \\nlyc\\- is equal to the intrinsic Lyman continuum luminosity inferred from IR recombination lines observed with the ISO satellite (Sturm et al 1996, Lutz et al 1996, G\\\"{e}nzel et al 1998). The simplest reason for this correspondence is that almost all of the IR recombination lines also arise in the same gas with density \\ne $ > 10^3$ \\cmthree.  That still leaves the question why [S III] line ratios observed with ISO indicates a density of only a few hundred \\cmthree, although they are sensitive to up to ten times higher densities. It is possible that if dust extinction is very high (see next section), [SIII] lines are not reliable indicators of density.   \\subsection{Extinction}  The predicted flux of the Br\\a\\- line from the three component model in Table 9 is $1.0\\times 10^{-15}$ W~m$^{-2}$. Sturm et al (1996) report a measured flux of $2.1 \\times 10^{-16}$ W~m$^{-2}$ from ISO data. Thus extinction of Br\\a\\- is $A_{Br\\alpha} \\sim 1.7$ which correspond  to a $V$ band extinction $A_V \\sim 45$. Similar calculation based on the predicted Br\\g flux of $3.2 \\times 10^{-16}$ W~m$^{-2}$ and a measured flux of $5.9\\times 10^{-18}$ W~m$^{-2}$ (Goldader et al 1995) yield $A_{Br\\gamma} \\sim 4.3$ and $A_V \\sim 41$. These derived values of $A_V$ are comparable to the values obtained using various IR line ratios by Sturm et al (1996). This similarity in the derived $A_V$ values provides further evidence that the IR recombination lines of hydrogen observed by ISO arise in high density gas. The radio recombination line data confirms the high extinction values found from ISO spectroscopy (Sturm et al 1996, G\\\"{e}nzel et al 1998) and provides support to the idea that Arp 220 is powered by a massive starburst (see below).  \\subsection{\\nlyc and Star Formation Rate (SFR)}  One of the derived quantities of the three component model in Table 9 is the production rate of Lyman continuum photons \\nlyc. Lyman continuum (Lyc) photons are produced by massive O and B stars during their main sequence phase.  Since these stars have a relatively short main sequence life time of $< 10^7$ years, the derived value of \\nlyc can be related to the formation rate of O and B stars if direct absorption of Lyc photons by dust is not significant. If an initial mass function (IMF) is adopted, then the formation rate of OB stars can be used to derive the total star formation rate (SFR) (e.g. Mezger 1985). The derived SFR is both a function of the adopted IMF and assumed upper and lower  mass limits of stars that are formed ($m_u$ and $m_l$). It is thought that in the case of induced star formation, triggered for example by an external event such as a galaxy-galaxy interaction or a merger, the lower mass cut off in the IMF may be a few solar masses (Mezger 1987). In spontaneous star formation in a quiescent molecular cloud, the lower mass cut off may be $<0.1$ \\msun\\- set by theoretical considerations (Silk 1978). In the following discussion we use $m_l$ = 1 \\msun\\- and $m_u$ = 100 \\msun, the latter corresponding to an O3.5 star. Using these mass limits and assuming a constant rate of star formation (over the life time of O stars), we get from the formulae given by Mezger (1985) \\begin{equation} N_{Lyc} = 5.4 \\times 10^{52} \\times {\\bf \\Psi}_{OB}~~ {\\rm s^{-1}}, \\end{equation} where ${\\bf \\Psi}_{OB}$ (\\msun\\- yr$^{-1}$) is the SFR averaged over the lifetime of the OB stars. In eqn 1, the IMF proposed by Miller and Scalo (1978) has been used. Using a single power-law IMF of Salpeter (1955) results in a factor of $\\sim 3$ increase in \\nlyc. A reduction in the lower mass limit $m_l$ to 0.1 \\msun\\- decreases \\nlyc by $\\sim 2$. The value of \\nlyc is less sensitive to the upper mass limit in the range 50-100 \\msun.  The sum of \\nlyc for the three components in Table 9 is $1.3\\times 10^{55}$ s$^{-1}$. An examination of the various models discussed in the previous section indicate that the uncertainty in this value is likely to be less than a factor of two. This value of \\nlyc is consistent with the results from ISO based on NIR recombination lines (Sturm et al 1996 and G\\\"{e}nzel et al 1998) and also with the lower limit of $1 \\times 10^{55}$ s$^{-1}$ deduced by Downes and Solomon (1998) based on their estimated thermal continuum at 113 GHz.  The derived SFR based on equation (1) and the total \\nlyc in Table 9 is $\\sim$ 240 \\msun\\- yr$^{-1}$. The SFR in Arp 220 is thus about two orders of magnitude higher than in the Galaxy and may be the highest SFR derived in any normal, starburst or ultra-luminous galaxy (e.g. Kennicutt et al 1998). If the Salpeter-IMF is used and the upper mass limit is reduced to 60 \\msun, then the SFR is reduced to 90 \\msun\\- yr$^{-1}$.  The high SFR in Arp 220 is likely a consequence of the large gas content in the nuclear region with high volume and surface densities. Downes and Solomon (1998) have estimated that within a radius of about 500 pc, which includes the east and west nuclei as well as a portion of a gas disk that surrounds them, the total mass of H$_2$ gas is $\\sim 4 \\times 10^9$ \\msun. For the same region, Scoville et al (1997) obtain a higher mass of $\\sim 8 \\times 10^9$ \\msun . For the discussion here, we take a mean of the two, $6\\times 10^9$ \\msun. The average volume density of the H$_2$ is $\\sim 250$ \\cmthree\\- and the average gas surface density $\\mu_g \\sim$ 6000 \\msun\\- pc$^{-2}$. A detailed study of star formation rates in normal and ultra-luminous galaxies by Kennicutt (1998) has shown that SFR per unit area can be fitted to a Schmidt law of the form $\\Sigma_{SFR} = 2.5\\times 10^{-4} \\mu_g^{1.4}$ \\msun\\- yr$^{-1}$ kpc$^{-2}$. This relation yields an SFR of 38 \\msun\\- yr$^{-1}$ for $\\mu_g$ = 6000 \\msun\\- pc$^{-2}$, which is a factor of six lower than deduced above. In his compilation of starburst properties of luminous galaxies, Kennicutt (1998) uses $\\mu_g$ = 5.8$\\times 10^4$ \\msun\\- pc$^{-2}$ for Arp 220 and an SFR of 955 \\msun\\- yr$^{-1}$ kpc$^{-2}$. This high value of $\\mu_g$ is the peak surface density in the inner most region obtained by Scoville et al (1997) and it is about a factor of 10 higher than the average surface density over a $\\sim$kpc$^2$ area.  Thus, when averaged over a kpc$^2$ region, the SFR predicted by Schmidt law with an exponent of 1.4 falls short of the value deduced above (240 \\msun\\- yr$^{-1}$) by about a factor of six.  Although other parameters such as the upper and lower mass limits and the shape of the IMF could be adjusted to lower the deduced SFR, it is unlikely to reduce it to the value predicted by the empirical Schmidt law obtained by Kennicutt (1998).   In a study of star forming regions in normal galaxies, Viallefond et al (1982) found a relation between $M_{gas}$, \\nlyc and density $n$ of the form \\nlyc = $7.9\\times 10^{44}~M_{gas}~n^{0.36}$. Using $M_{gas} = 6\\times 10^9$ \\msun\\- and $ n= 250$ \\cmthree\\-  (the average molecular density), this relation predicts \\nlyc = $3.3\\times 10^{55}$ s$^{-1}$. Although this value is about a factor of 2.5 higher than the derived total \\nlyc in Table 9, these values are comparable given the uncertainties. This result shows that the starburst in Arp 220 behaves  as expected from scaling the known properties of star forming regions in normal galaxies.       \\subsection{Star Formation at Recent Epochs}   As shown in Section 4, the observed intensities of the millimeter wavelength RRLs H31\\a, H40\\a\\- and H42\\a\\- in Arp 220 can only be explained by ionized gas at a high density of $\\sim 2.5\\times 10^5$ \\cmthree. The presence of these high density HII regions which account for about 3\\% of the total \\nlyc of $1.3\\times 10^{55}$ s$^{-1}$, can be used to derive the  star formation rates at recent epochs. Since the high density compact HII regions are relatively short-lived ($\\tau_{HII} \\le 10^5$ yrs), the \\nlyc value corresponding to these regions will indicate the SFR averaged over the life time of  the HII regions rather than the main-sequence life time of O, B stars which ionize them. Based on equation (1), and approximating the Lyc photon production rate of OB stars to be constant during their life time, we can write \\begin{equation} {\\bf \\Psi}_{HII} = {\\bf \\Psi}_{OB} \\times \\left[\\frac{\\tau_{OB}}{\\tau_{HII}}\\right] \\times \\left[\\frac{N_{Lyc,HII}}{N_{Lyc,OB}}\\right]. \\end{equation} The compact, high density phase of an HII region is short lived because the HII region expands as it is over-pressured with respect to the surroundings. The expansion proceeds at the sound speed $c_i$ ($\\sim 10$ \\kms) and approximately follows the relation \\begin{equation} r(t) = r_i\\left[1+\\frac{7c_it}{4r_i}\\right]^{4/7} \\end{equation} Spitzer (1968), where $r_i$ is the initial size of the HII regions and $r(t)$ is its size after a time $t$. The HII regions of component A2 in Table 9 have a density of $2.5\\times 10^5$ \\cmthree. The density of these HII regions will drop to about 1000 \\cmthree\\- and become indistinguishable from component A1, if they expand to $\\sim 6$ times their initial size. Taking the initial size to be 0.1 pc, as given in Table 9, equation (3) gives $t =10^5$ yrs. If the initial size is smaller, then the time scale is also shorter. (For a single O5 star, the initial size of the Stromgren sphere is $\\sim$0.05 pc if the density is  $2.5\\times 10^5$ \\cmthree.) Using $\\tau_{HII} = 10^5$ yrs, $\\tau_{OB} = 3\\times 10^6$ yrs, ${\\bf\\Psi}_{OB} $ = 240 \\msun\\- yr$^{-1}$ and using \\nlyc values from Table 9, equation 2 gives the SFR averaged over the life time of HII regions  ${\\bf \\Psi}_{HII} =$ 194 \\msun\\- yr$^{-1}$. This rate is similar to the SFR averaged over $\\tau_{OB}$. This result indicates that the starburst in Arp 220 is an on-going process with a minimum age exceeding $\\tau_{OB}$. Other evidence (see below) indicates that the age of the starburst is probably much longer.  The above calculation demonstrates that a reliable measurement of the number of Lyman continuum photons in high density HII regions can lead to a determination of the ``instantaneous'' SFR. Such a measurement seems possible through RRLs and continuum at millimeter wavelengths which are sensitive to the dense component.  The actual value of the instantaneous SFR determined above (194 \\msun\\- yr$^{-1}$) can be almost a factor of ten  higher or about a factor of two lower  since the value of ${N_{Lyc,HII}}$ in equation 2 is uncertain by those factors (see Table 8). If the ten times higher value of $N_{Lyc,HII} \\sim 4.6\\times 10^{54}$ s$^{-1}$ corresponding to model B2 in Table 8 is established through further measurements, then the implied SFR at recent epochs will be ten times the average value. Such a result will fundamentally alter our picture of the starburst process in Arp 220. Instead of being bursts of constant but high SFR (approximately 10-100 times the Galactic rate) over $10^6 - 10^7$ yrs, starburst in Arp 220 will be multiple events of short duration ($10^5 - 10^6$ yrs) but at extremely high SFR (100-1000 times the Galactic rate). Further continuum and RRL measurements at millimeter wavelengths are required to clarify this aspect.  \\subsection{\\nlyc and IR-Excess}  In dense galactic HII regions, the ratio $\\frac{L_{IR}}{L_\\alpha}$ is in the range 5-20 which is known as the IR-excess (IRE) (Mezger et al 1974, Panagia 1978, Mathis 1986). $L_{IR}$ is the observed infrared luminosity. $L_\\alpha$ is the luminosity of Lyman-$\\alpha$ photons deduced from observations of the ionized gas using recombination lines or radio continuum and assuming that all the Lyman continuum photons which ionize the gas are eventually converted to L\\a\\- photons (i.e. $N_{L_\\alpha}$=\\nlyc).  These L\\a\\- photons are then absorbed by dust. If this is the only source of heating the dust in an ionization bounded HII region, and since dust re-radiates most of the absorbed energy in the IR, then the expected IRE is $\\sim$1. If IRE $>1$, then the dominant mechanism for heating the dust in HII regions contains a major contribution other than  $L\\alpha$ heating. Two possible sources of heating are: (1) direct absorption of Lyc photons by dust and (2) absorption of photons long ward of the Lyman limit. Lyc photons which are directly absorbed by dust are not counted in the \\nlyc derived from RRL or continuum observation of the ionized gas.  Therefore, if direct absorption of Lyc photons is the main contributor to the IRE, then the SFR derived using eqn 1 will be an underestimate by a factor $\\sim$IRE. On the other hand, heating of dust by photons long ward of the Lyman limit (which are produced mainly by lower mass, non-ionizing stars) can account for the IRE without a corresponding increase in the SFR. Viallefond  (1987) has shown that a population of non-ionizing stars formed continuously or in a burst with a Miller-Scalo IMF can indeed explain an IRE of $\\geq$10. In Arp 220, $L_{IR} = 1.3 \\times 10^{12} L_\\odot$ and $L_{L_\\alpha} = 5.5\\times 10^{10}L_\\odot$ using $N_{L_\\alpha}$ = \\nlyc = $1.3\\times 10^{55}$ s$^{-1}$. Thus for Arp 220, IRE $\\sim$ 24. Based on this value of IRE we make a few deductions about the star formation history in Arp 220.  \\subsubsection{Starburst or AGN?}   In giant and super-giant HII regions of the normal galaxy M33, the value of IRE is in the range 3 to 8 and the global value of IRE is $\\sim 14$ (Rice et al 1990). In the Galaxy, IRE is typically $\\sim$ 7 in compact HII regions and $\\sim 5$ in extended HII regions (Mezger 1987). G\\\"{e}nzel et al (1998) have computed a modified form of IRE (they use $L_{Bol}/L_{Lyc}$) in starburst galaxies observed by ISO. Converting their ratios to the above definition of IRE (i.e.$\\frac{L_{IR}}{L_\\alpha}$), the mean value of IRE in starburst galaxies is $\\sim 25$ with values ranging from 12 to 45.  Similar computation made for AGN dominated sources show that IRE$ \\sim 45-65$ (G\\\"{e}nzel et al 1998). Thus IRE in Arp 220 ($\\sim 24$) is very similar to the values observed in starburst galaxies and slightly higher than in star forming regions in the Galaxy and in M33. IRE in Arp~220 is significantly lower than in AGN dominated sources.  The high infrared luminosity of Arp 220 can therefore be entirely accounted for by processes that operate in starburst regions. We therefore conclude that Arp 220 is powered entirely by a starburst. There is no significant contribution from an ``active'' nucleus to the observed high IR luminosity in Arp 220.  \\subsubsection{IMF, Age of Starburst and Star Formation Efficiency in Arp 220}  The fact that the IRE for Arp 220 is not significantly different from the IRE in star forming regions in the Galaxy and in M33 also implies that the starburst in Arp 220 has an IMF which is not unusual. If, for example, IRE$\\sim$1, then there is no contribution to heating of dust by non-ionizing photons which would imply that the burst is mainly producing massive stars and thus $m_l > \\sim$ 5 \\msun. This is not the case in Arp 220. On the other hand if IRE$>>$10, then it would imply one of the following: (1) the IMF is truncated at the upper end ($m_u < 10-20$ \\msun\\- or so) which makes the heating of dust by non-ionizing photons even more pronounced, leading to a much larger IRE or (2) the starburst ended about $5\\times 10^6$ years ago which would reduce \\nlyc in relation to non-ionizing photons, again resulting in a larger IRE or (3) an AGN is contributing predominantly to the heating and IR emission. These three possibilities are also ruled out in Arp 220.  Another implication of IRE $\\sim$ 24 in Arp 220, which is significantly higher than in young star forming regions, is that the starburst in Arp 220 must be much longer in duration than the lifetime of OB stars (i.e. $t_{SB} > 10^7$ yrs). As mentioned above, in younger star forming regions such as the super-giant HII region NGC 604 in M33 and in compact HII regions in the Galaxy, the IRE is significantly smaller ($\\sim 5-7$, Mezger 1985, Rice et al 1990).  To illustrate that a relatively high value of IRE implies a longer duration of starburst, we consider a simple model of continuous star formation discussed in Viallefond and Thuan (1983) and the tabulation of the derived parameters in Viallefond (1987). In this model, for an assumed IMF with lower ($m_l$) and upper ($m_u$) stellar mass limits, \\nlyc and \\lbol\\- are computed as a function of duration of star formation ($\\tau_{csf}$) for a SFR of 1 \\msun\\- yr$^{-1}$.  These values, taken from Viallefond (1987), are tabulated in Table 10 for a Miller-Scalo IMF with $m_u = 100$ \\msun\\- and $m_l = 1$ \\msun. The results would be similar for $m_l = 0.1$  \\msun. In this model, a steady state production rate of Lyc photons is reached after $\\sim 5\\times 10^6$ yrs. The steady state value of \\nlyc is $3.7 \\times 10^{52}$ ph s$^{-1}$/(\\msun\\- yr$^{-1}$) and implies a SFR of $\\sim$ 350 \\msun\\- yr$^{-1}$ since the derived value of \\nlyc = $1.3\\times 10^{55}$ ph s$^{-1}$. This SFR is consistent with the value derived in Section 5.3. Table 10 also shows that in this model, the observed IRE$\\sim$ 24 in Arp 220 is reached at $\\tau_{csf} \\sim 1.5\\times 10^7$ yrs.  Finally, a continuous SFR of 350 \\msun\\- yr$^{-1}$ for a duration of $1.5\\times 10^7$ yrs converts $\\sim 5\\times 10^9$ \\msun\\- of gas into stars. Thus the mass of stars formed is approximately equal to the present mass of gas ($\\sim 6\\times 10^9$ \\msun\\-).  In a closed-box model (i.e. with no fresh gas added to the system from outside), the star formation efficiency (SFE) in Arp 220 is 50\\%, much higher than the average SFE of a few \\% in Galactic star forming regions (Elmegreen 1983, Myers et al 1986). The high SFE ($\\sim 50$\\%) and the high SFR ($\\sim$300\\msun\\- yr$^{-1}$) in Arp 220 and its corresponding high luminosity ($L_{IR} = 1.3 \\times 10^{12}$ \\lsun) must be a result of confinement of a large quantity of gas (M$_{gas} \\sim 6\\times 10^9$ \\msun) in a relatively small volume ($\\sim$kpc$^3$) with a high average density of 250 \\cmthree. At the East and West peaks of Arp 220, about $10^9$ \\msun\\- of gas (i.e. about the mass of gas in the Galaxy) is confined to a region $\\sim$100 pc in size leading to a much higher {\\em average} density of  $\\sim 10^4$ \\cmthree. The high concentrations of large quantities of gas in Arp 220 have given rise to a starburst with an efficiency and SFR that is consistent with known properties of star forming regions in normal and starburst galaxies."
    },
    "0004/astro-ph0004308_arXiv.txt": {
        "abstract": "The discrimination capabilities of a 70 g heat and ionization Ge bolometer are studied. This first prototype has been used by the EDELWEISS Dark Matter experiment, installed in the Laboratoire Souterrain de Modane, for direct detection of WIMPs. Gamma and neutron calibrations demonstrate that this type of detector is able to reject more than 99.6\\% of the background while retaining 95\\% of the signal, provided that the background events distribution is not biased towards the surface of the Ge crystal. However, the 1.17 kg.day of data taken in a relatively important radioactive environment show an extra population slightly overlapping the signal. This background is likely due to interactions of low energy photons or electrons near the surface of the crystal, and is somewhat reduced  by applying a higher  charge-collecting inverse bias voltage ( -6 V instead of -2 V) to the Ge diode. Despite this contamination, more than 98\\% of the background can be rejected while retaining 50\\% of the signal. This yields a conservative upper limit of 0.7 ${\\rm event.day^{-1}.kg^{-1}.keV_{recoil}^{-1}}$ at 90\\% confidence level in the 15-45 keV recoil energy interval; the present sensitivity appears to be limited by the fast ambient neutrons. Upgrades in progress on the installation are summarized.\\\\ \\vskip0.2truecm\\noindent PACS: 95.35.+d\\par\\noindent {\\it Keywords:} Dark Matter, WIMP, Cryogenic Detector ",
        "introduction": "The search for dark matter is one of the major challenges of contemporary physics. Despite extensive scans for MAssive Compact Halo Objects (MACHOs) initiated by several groups \\cite{alc,aub,pac}, it appears at present that the local density of dark matter can hardly be composed of baryonic matter essentially. Indeed, recent results from these experiments indicate that a significant fraction of the MACHOs observed in the direction of the Small Magellanic Cloud (SMC) are probably due to deflectors internal to the SMC, further reducing the possible amount of dark matter in the form of MACHOs in the halo itself \\cite{alc2}.  The coincidence of the electroweak interaction scale (SUSY theory) with that required for Weakly Interacting Massive Particles (WIMPs) to contribute significantly to the solution of the dark matter problem is a further motivation for attacking the WIMPs direct detection challenge. In addition these hypothetical particles represent an attractive solution to the problem of galaxy formation. Experimental efforts are underway to detect these particles, either indirectly, by searching for products of their annihilation in the core of the Sun or of the Earth, or directly, by detecting the interactions of the WIMPs themselves in ordinary matter, as first suggested in the mid 1980s \\cite{stodol,good} .  The sensitivity of present experiments appears to be limited by the radioactive background rate of the detectors \\cite{baud1}  and by the systematics of the rejection scheme using statistical identification methods based on pulse shape discrimination techniques (PSD) in NaI crystals \\cite{dama,ukdmc,gerb} . The motivation for the development of cryogenic detectors is the possibility, when the heat measurement is coupled with a measurement of ionization \\cite{shutt,schne,drain,berge} (or of scintillation \\cite{meunier}), to discriminate much more reliably the main source of radioactive background, producing electron recoils, from the nuclear recoils expected from WIMP interactions. Indeed, WIMPs scatter off nuclei, which ionize (or produce scintillation) less efficiently than electron recoils for a given energy deposit in the absorber.  In the following, we present the discrimination capabilities of this new generation of detectors in the EDELWEISS (Exp\\'erience pour DEtecter Les Wimps En SIte Souterrain)  underground experiment and we show that a 70 g high-purity Ge crystal can already reach sensitivity levels typical of the best existing experiments despite its small mass and a relatively high radioactivity level in the detector environment. The present limitations of these detectors are then explored and the first attempts to counteract them are discussed. ",
        "conclusions": "Our prototype heat and ionization bolometer provides an excellent separation of the calibrated background and signal down to an energy of 4 keV e.e.. However data taken during background runs contain an important population of off-axis events (electron surface events) essentially not seen during the calibrations. Despite this, the rejection capabilities of the detector remain satisfactory, still reducing the number of electron recoil events by nearly two orders of magnitude. Sensitivity limits on WIMPs \\cite{drain,berge} obtained with this prototype are only a factor of 10 higher than the best existing limits obtained with longer-established techniques \\cite{baud1,dama} despite the small mass (70 g) of our detector and the relatively high radioactivity level in the detector environment. Moreover the present sensitivity of the experiment appears to be limited, principally for the runs with a -6V bias voltage, not by the gamma background or by parasitic electron surface events, but by the ambient neutrons  in the underground laboratory. Given the systematic uncertainties in the Monte Carlo neutron simulations  a background subtraction is not attempted but, clearly, shielding against neutrons is required.  A new generation of detectors is being tested in the LSM, with  new implantation  schemes for the electrodes and a rigorous selection procedure of the crystal holder material \\cite{XFN3}. Several improvements have been undertaken in the detector environment: nitrogen flushing for radon removal, improved passive protection against gamma radioactivity (including  close Roman lead shielding) and 30 cm removable paraffin shielding against neutrons. In an initial test run, two 70 g prototypes show background levels before rejection of $\\approx 2\\  {\\rm event.keV^{-1}.kg^{-1}.day^{-1}}$  below 50 keV, a factor of 10 improvement over previous data, with a comparable progression for electron surface events. The number of events remaining in the neutron zone is now extremely limited and could be due to residual neutron interactions inside the present shielding. Assuming a good performance of the neutron shielding and that the energy resolution and background rejection are the same as those reported here, a few months of data taking with these detectors should allow a test of the entire parameter region claimed by the DAMA experiment \\cite{dama} and the entrance into the relevant supersymmetry parameter space."
    },
    "0004/astro-ph0004414_arXiv.txt": {
        "abstract": "The {\\it Extreme Ultraviolet Explorer\\/} ({\\it EUVE\\/}) satellite was used for three days beginning on 1997 Mar 26.96 UT to obtain photometric and spectroscopic observations of the eclipsing SU~UMa-type dwarf nova OY~Carinae in superoutburst. Because of the longer time on source (143 ks), the larger number of eclipses observed (17), and the higher count rate in the detector (0.5--2.2 $\\rm counts~s^{-1}$), we are able to significantly strengthen previous reports that there is little or no eclipse by the secondary of the EUV emission region of OY~Car in superoutburst. The mean {\\it EUVE\\/} spectrum extends from 70 to 190~\\AA \\ and contains broad ($\\rm FWHM\\approx 1$~\\AA ) emission lines of \\ion{N}{5}, \\ion{O}{5}--\\ion{O}{6}, \\ion{Ne}{5}--\\ion{Ne}{7}, \\ion{Mg}{4}--\\ion{Mg}{6}, \\ion{Fe}{6}--\\ion{Fe}{8}, and possibly \\ion{Fe}{23}. Good fits of the observed spectrum are obtained with a model (similar to that of Seyfert 2 galaxies) wherein radiation from the boundary layer and accretion disk is scattered into the line of sight by the system's photoionized accretion disk wind. It is possible to trade off continuum luminosity for wind optical depth, but reasonable models have a boundary layer temperature $T_{\\rm bl}\\approx 90$--130~kK and a boundary layer and accretion disk luminosity $L_{\\rm bl}= L_{\\rm disk}\\lax 4\\times 10^{34}~\\rm erg~s^{-1}\\approx 10\\> \\Lsun $, corresponding to a mass-accretion rate $\\Mdota\\lax 10^{-8}~\\rm \\Msun~yr^{-1}$; an absorbing column density $N_{\\rm H}\\approx 1.6$--$3.5\\times 10^{19}~\\rm cm^{-2}$; and a wind mass-loss rate $\\Mdotw\\lax 10^{-10}~{\\rm \\Msun~yr^{-1}} \\approx 0.01\\, \\Mdota $. Because radiation pressure alone falls an order of magnitude short of driving such a wind, magnetic forces must also play a role in driving the wind of OY~Car in superoutburst. ",
        "introduction": "Cataclysmic variables (CVs) are a diverse class of semidetached binaries composed typically of a low-mass main-sequence secondary, an accreting white dwarf, and a luminous accretion disk. There are three generalized classes of CVs: novae, novalike variables, and dwarf novae. In dwarf novae, a mass-transfer instability causes the disk to brighten by 3--5 magnitudes in the optical at quasi-periodic intervals of tens to hundreds of days. Based on their light curve morphologies, most dwarf novae are classified as either U~Gem systems with ``normal'' outbursts only, Z~Cam systems with ``normal'' outbursts and prolonged states of intermediate brightness (``standstills''), or SU~UMa systems with ``normal'' and ``super'' (brighter, longer-duration) outbursts. Nonmagnetic novalike variables behave like dwarf novae ``stuck'' in outburst. \\citet{war95} provides an excellent background on all of the various CV types and subtypes, while \\citet{can93} and \\citet{osa96} provide recent reviews of the theory of dwarf nova outbursts.  While the white dwarf and accretion disk are the dominant sources of optical through FUV light in dwarf novae, the boundary layer between the disk and the surface of the white dwarf is the dominant source of higher-energy emission. Modulo a correction factor for white dwarf rotation, the boundary layer luminosity $L_{\\rm bl}= G\\Mwd\\Mdota/ 2\\Rwd\\sim 4\\times 10^{34}\\, (\\Mdota/10^{-8}~\\rm \\Msun~yr^{-1})~erg~s^{-1}$, where $\\Mdota $ is the accretion rate and $\\Mwd $ and $\\Rwd $ are respectively the mass and radius of the white dwarf. When $\\Mdota $ is low (e.g., dwarf novae in quiescence), the boundary layer is optically thin and quite hot (of order the virial temperature $T_{\\rm vir} = G\\Mwd m_{\\rm H}/3k\\Rwd\\sim {\\rm tens}$ of keV; \\citealt{pri79, tyl81, nar93}); when $\\Mdota $ is high (e.g., novalike variables and dwarf novae in outburst), the boundary layer is optically thick and quite cool (of order the blackbody temperature $T_{\\rm bb} =[G\\Mwd\\Mdota/8\\pi\\sigma\\Rwd ^3]^{1/4} \\sim 100$ kK; \\citealt{pri77, tyl77, pop95}).  Observationally, the location and extent of the source of the X-rays in nonmagnetic CVs is constrained by (1) the X-ray light curves of eclipsing systems, (2) the variation of apparent emission measure with binary inclination \\citep{tes96}, and (3) the relative strength of the line and continuum spectrum reflected off the white dwarf and accretion disk \\citep[e.g.,][]{don97}. All these constraints imply that all or most of the hard X-rays in low-$\\Mdot $ CVs are emitted very close to the white dwarf, but the eclipse observations provide the most direct and unambiguous evidence. In HT~Cas ($i\\approx 81^\\circ $; \\citealt{woo95a, muk97}), Z~Cha ($i\\approx 82^\\circ $; \\citealt{tes97}), and OY~Car ($i\\approx 83^\\circ $; \\citealt{pra99a}) in quiescence, the X-rays are fully eclipsed for an orbital phase interval comparable to that of the optical eclipse of the white dwarf.  In contrast to the compact source of hard X-rays in low-$\\Mdot $ CVs, there is strong evidence for an extended source of soft X-rays in high-$\\Mdot $ CVs. {\\it EXOSAT\\/} and {\\it ROSAT\\/} light curves of OY~Car in superoutburst \\citep{nay88, pra99b} and {\\it ROSAT\\/} light curves of the novalike variable UX~UMa ($i\\approx 71^\\circ $; \\citealt{woo95b}) rule out an eclipse by the secondary of a compact source of X-rays centered on the white dwarf. In both cases, it is argued that the boundary layer is obscured from view by the accretion disk at all orbital phases, and that the observed X-rays come from a corona or wind above the disk.  The situation is far more complex in U~Gem ($i\\approx 70^\\circ $). In quiescence, the hard X-rays are partially eclipsed at orbital phase $\\phi \\sim 0.7$ \\citep{szk96}, and during outburst the EUV and soft X-rays are partially eclipsed at orbital phases $\\phi\\sim 0.7$ and 0.1 \\citep{mas88, lon96}. These dips in the X-ray and EUV light curves are interpreted as partial eclipses of the boundary layer by vertical structure at the edge of the accretion disk ($r\\approx 4\\times 10^{10}$~cm) or at the circularization radius of the accretion stream ($r\\approx 5\\times 10^9$~cm); in either case, the heights above the orbital plane are distressingly large: $h\\approx 0.5\\, r$, whereas the disk thickness $H\\approx c_{\\rm s}/\\Omega_{\\rm K}\\ll r$. The {\\it EUVE\\/} spectra of U~Gem in outburst supply additional diagnostic information. The phase-averaged spectrum contains emission lines of Ne VI--VIII, Mg VI--VII, and Fe VII--X superposed on a $T= 110$--140~kK blackbody continuum \\citep{lon96}. The phase-resolved spectra demonstrate that the eclipses affect the continuum more strongly than the lines (most strikingly, the Ne~VIII $\\lambda 88.1$ line is nearly unaffected by the eclipses; see \\citealt{mau97}), implying that the lines are produced in a region of larger extent than that of the continuum, which presumably is formed in the boundary layer.  To gain a better understanding of the extended source of soft X-rays in high-$\\Mdot $ CVs, we obtained {\\it EUVE\\/} observations of the eclipsing SU~UMa-type dwarf nova OY~Car in superoutburst. The resulting data are superior to those from the previous {\\it EXOSAT\\/} and {\\it ROSAT\\/} observations because of the longer time on source, the larger number of eclipses observed, the higher count rate in the detector, and the ability to obtain dispersed EUV spectra. We discuss the {\\it EUVE\\/} observations in \\S 2, the DS light curve in \\S 3, and the SW spectrum in \\S 4. We find that there is little or no eclipse of the EUV emission region in OY~Car in superoutburst (\\S 3), and that the EUV spectrum contains emission lines of intermediate ionization stages of N, O, Ne, Mg, and Fe (\\S 4). The spectrum is not fit well in detail by a model of a collisionally ionized plasma (\\S 5), and although it is fit well by a model wherein radiation from the boundary layer and accretion disk is scattered into the line of sight by OY~Car's accretion disk wind (\\S 6), radiation pressure alone falls an order of magnitude short of driving the wind (\\S 7). We close in \\S 8 with a discussion and summary of our results. ",
        "conclusions": "Because of the longer time on source, the larger number of eclipses observed, and the higher count rate in the detector, the {\\it EUVE\\/} observations discussed here significantly strengthen previous reports that there is little or no eclipse by the secondary of the EUV emission region of OY~Car in superoutburst. The mean {\\it EUVE\\/} spectrum extends from 70 to 190~\\AA \\ and contains broad ($\\rm FWHM\\approx 1$~\\AA ) emission lines of \\ion{N}{5}, \\ion{O}{5}--\\ion{O}{6}, \\ion{Ne}{5}--\\ion{Ne}{7}, \\ion{Mg}{4}--\\ion{Mg}{6}, \\ion{Fe}{6}--\\ion{Fe}{8}, and possibly \\ion{Fe}{23}. The gross details of the spectrum are reproduced by the spectrum of a collisionally ionized plasma with a temperature $T\\approx 320$~kK, an absorbing column density $N_{\\rm H}\\approx 3.4\\times 10^{19}~\\rm cm^{-2}$, and an emission measure $EM\\approx 3.2\\times 10^{55}~\\rm cm^{-3}$, but such a model fails to reproduce the observed spectrum in detail. Better fits are obtained with a model (similar to that of Seyfert 2 galaxies) wherein radiation from the boundary layer and accretion disk is scattered into the line of sight by the system's cool ($T\\sim 30$ kK) photoionized accretion disk wind. The likely range of acceptable parameters for the boundary layer are $T_{\\rm bl}\\approx 90$--130 kK and $N_{\\rm H}\\approx 1.6$--$3.5\\times 10^{19}~\\rm cm^{-2}$. This range for the absorbing column density is typical of the interstellar column densities of nearby CVs beyond the Local Bubble \\citep{mau88, pol90, lon96}, while the range of boundary layer temperatures is significantly less than that of SS~Cyg \\citep[$T_{\\rm bl} \\approx 230$--350~kK;][]{mau95} but comparable to that of U~Gem \\citep[$T_{\\rm bl}= 110$--140~kK;][]{lon96} and VW~Hyi \\citep[$T_{\\rm bl}\\approx 120$~kK;][]{mau96}. The corresponding range of the boundary layer luminosity is less clear because it is possible to trade off the intensity of the EUV continuum for the wind optical depth (i.e., the wind mass-loss rate). A lower limit to the fractional emitting area of the boundary layer is $f\\approx 0.1$, and a lower limit to the boundary layer luminosity is $L_{\\rm bl}\\approx 1\\times 10^{34}~\\rm erg~s^{-1} \\approx 3\\> \\Lsun $. Adding a ``minimal'' Shakura-Sunyaev accretion disk\\footnote{The observed values of $L_{\\rm disk}/L_{\\rm bl}$ range from $\\approx 2$ for U~Gem \\citep{lon96}, $\\gax 10$ for SS~Cyg \\citep{mau95}, and $\\sim 20$ for VW~Hyi \\citep{mau96}.} with $L_{\\rm disk}=G\\Mwd \\Mdota/2\\Rwd =L_{\\rm bl}=4\\pi\\Rwd^2f\\sigma T_{\\rm bl}^4$ provides further constrains. Requiring that the maximum disk temperature be less than or equal to the boundary layer temperature requires that $f\\le 5.8$ and $L_{\\rm bl}\\le 6.5\\times 10^{35}~\\rm erg~s^{-1}\\approx 170\\> \\Lsun $. Adding reasonable brightness constraints in the UV [$f_{1500}\\lax 1\\times 10^{-10}\\, (d/85~{\\rm pc})^{-2}~\\rm erg~cm^{-2}~s^{-1}~\\AA^{-1}$] and optical ($M_V\\gax 5$) is equivalent to the requirement that the mass-accretion rate $\\Mdota\\lax 10^{-8}~\\rm \\Msun~yr^{-1}$, consistent with observations of other high-$\\Mdot $ CVs, hence $L_{\\rm bl}\\lax 4\\times 10^{34}~\\rm erg~s^{-1}\\approx 10\\> \\Lsun $. Such models have wind mass-loss rates of $\\Mdotw\\lax 10^{-10}~{\\rm \\Msun~yr^{-1}}\\approx 0.01\\, \\Mdota $ and fall an order of magnitude short of driving the wind by radiation pressure alone. An alternative is that the wind of OY~Car is driven by a combination of radiation pressure and magnetic forces (a magnetocentrifugal wind). Such winds have been investigated recently by \\citet{pro00}."
    },
    "0004/astro-ph0004287_arXiv.txt": {
        "abstract": "In 1998 The Optical Gravitational Lensing Experiment (OGLE) successfully implemented automated data reductions for QSO 2237+0305. Using a new image subtraction method we achieved a differential photometry scatter of 1--5\\% for images A--D respectively. Combined with a time sampling of 1--2 times a week this is sufficient for early detection of caustic crossings. Nearly real time photometry of QSO 2237+0305 is available from the OGLE web site http://www.astro.princeton.edu/$\\sim$ogle/ogle2/huchra.html. During the 1999 observing season, the apparent $V$ magnitude of the A, B, C and D images changed by 0.50, 0.15, 0.65 and 0.35 mag, respectively. Most likely however, none of the microlensing events involved a caustic crossing. The most rapid variation was 0.25 mag in 30 days, observed for image C. The alert system will continue to be active in the next observing season from late April until September 2000, when OGLE suspends operation for an upgrade. Observations will resume for season of 2001. ",
        "introduction": "The Optical Gravitational Lensing Experiment is a long term project focusing on detection and monitoring of microlensing events in crowded fields. During the second phase of the experiment, data are collected with the dedicated 1.3 m Warsaw Telescope at the Las Campanas Observatory (LCO), Chile (Udalski, Kubiak and Szyma\\'nski 1997). In 1997, OGLE started a monitoring program for two multiply imaged lensed quasars: QSO 2237+0305 and HE 1104-1805. In QSO 2237+0305 (Huchra's lens) one expects individual quasar images to vary due to microlensing in the bulge of the macrolensing galaxy (Webster et al. 1991, Schneider et al. 1988). Theoretical models predict high magnification events at caustic crossings. The sharpest features in the light curve provide in principle a size measurement of the source, in this case a quasar (e.g. Wambsganss et al. 1990). This letter presents observations made in 1999, previously unpublished except for a web site release. Also we would like to call the attention of astronomical community to the fact that good photometry of QSO 2237+0305 is available practically in real time, and encourage follow up programs. ",
        "conclusions": ""
    },
    "0004/astro-ph0004078_arXiv.txt": {
        "abstract": "We examine the origin of clustercentric gradients in the star formation rates and colors of rich cluster galaxies within the context of a simple model where clusters are built through the ongoing accretion of field galaxies. The model assumes that after galaxies enter the cluster their star formation rates decline on a timescale of a few Gyrs, the typical gas consumption timescale of disk galaxies in the field. Such behaviour might be expected if tides and ram pressure strip off the gaseous envelopes that normally fuel star formation in spirals over a Hubble time. Combining these timescales with mass accretion histories derived from N-body simulations of cluster formation in a $\\Lambda$CDM universe, we reproduce the systematic differences observed in the color distribution of cluster and field galaxies, as well as the strong suppression of star formation in cluster galaxies and its dependence on clustercentric radius.  The simulations also indicate that a significant fraction of galaxies beyond the virial radius of the cluster may have been within the main body of the cluster in the past, a result that explains naturally why star formation in the outskirts of clusters (and as far out as two virial radii) is systematically suppressed relative to the field.  The agreement with the data beyond the cluster virial radius is also improved if we assume that stripping happens within lower mass systems, before the galaxy is accreted into the main body of the cluster.  We conclude that the star formation rates of cluster galaxies depend primarily on the time elapsed since their accretion onto massive virialized systems, and that the cessation of star formation may have taken place gradually over a few Gyrs. ",
        "introduction": "Extensive observational work has established that the star formation properties of galaxies in rich clusters differ significantly from those of field galaxies (e.g., \\cite{O60,DTS,B+97,KK}). In the field, galaxies form stars at rates several times higher than systems of similar luminosity in the cores of clusters. This is partly a result of the well-known morphology-density relation, since ellipticals and S0 galaxies are more abundant in clusters (\\cite{Dressler,WGJ}), but there is evidence that even later type galaxies in clusters form stars at lower rates then in the field.  For example, Balogh et al (\\cite*{PSG}) report that field galaxies of given bulge-to-disk ratio and luminosity have, on average, much larger [OII] equivalent widths than their counterparts in rich clusters, suggesting that the cluster environment somehow curbs the star formation rates of all galaxies, regardless of morphology.  Various physical mechanisms have been proposed to explain this and other systematic differences between the field and cluster galaxy populations. Ram pressure stripping by the intracluster medium, for example, has been suggested as a means of removing the gaseous component of disk galaxies and of dramatically altering their morphology and subsequent star forming history (e.g., \\cite{GG,FN99,A+99}). Tides, either by the main cluster potential (e.g. \\cite{BV,F98}) or by fly-by encounters with other cluster galaxies (``galaxy harassment'', \\cite{Moore+96}), have also been proposed to explain the observations.  Although the above processes are all plausible, the actual changes in star formation rate induced by them remain a matter of debate. For example, some authors have suggested that the cluster environment triggers intense bursts of star formation that rapidly consume the gas of an infalling cluster galaxy (\\cite{DG83,Barger,P+99}), while others have favored a scenario where star formation is truncated in a galaxy almost immediately after it is accreted into the cluster (\\cite{A2390,NBK,1621,JSC}).  The virtues of these scenarios can in principle be tested observationally, since they are expected to produce a population of galaxies with unusually strong Balmer absorption lines in which star formation has been recently terminated (e.g., \\cite{DG83,CS87}). Recently, Balogh et al. (\\cite*{PSG}) have applied this idea to galaxies in the CNOC1 survey of X-ray luminous clusters (\\cite{YEC}). From the paucity of galaxies with strong H$\\delta$ absorption in their spectra they conclude that the decline of star formation in these clusters may actually be a fairly gradual process. This is in contradiction to other authors' conclusions, which are derived from datasets with relatively high numbers of H$\\delta$-strong objects (e.g. \\cite{Barger,P+99}); whether this discrepancy is due to selection effects or to real differences in the clusters studied has not yet been fully elucidated.  A natural timescale for a gentler reduction in star formation rate may be gleaned from the long--known observation that, given their present disk gas content, normal field spirals currently form stars at rates that cannot be sustained over a Hubble time. At face value, it appears that most spirals would use up their disk gas supply in a few Gyrs (e.g., \\cite{G+89}), though this timescale may be considerably extended for some choices of the IMF and if the effects of non-instantaneous mass recycling is taken into account (\\cite{K+94}). Star formation lifetimes can be significantly longer if galaxies continuously accrete fresh star formation fuel from their surroundings (e.g., \\cite{HVC}). If, as proposed by Larson, Tinsley \\& Caldwell (\\cite*{LTC}), this extended reservoir is stripped off a galaxy when it first enters the cluster, its star formation rate may decay significantly within a few Gyr, leading to large differences in the cluster and field populations. In clusters that are built hierarchically (as in the ``bottom-up'' scenario favored by cold dark matter cosmogonies) this mechanism would also establish a radial gradient in the star formation properties of cluster galaxies, reflecting the relation between the clustercentric radius of a galaxy and the time of its accretion into the cluster. This is an important ingredient in the success of semianalytic models, which have been shown to match global cluster properties such as the morphological composition and the blue galaxy fraction (\\cite{semianal,KC}).  We explore here a simple model based on this interpretation where the mass accretion history of a cluster obtained from numerical simulations of a universe dominated by cold dark matter is coupled with a simple star formation prescription for galaxies following accretion. Once calibrated to reproduce the properties of local field galaxies, the model has no free parameters and its results can be compared directly with observations. We shall focus our analysis on a quantitative discussion of the clustercentric radial gradients in star formation properties and galaxy colors expected in this model.  In \\S\\ref{sec-sample} we describe the observational dataset we use, selected from the CNOC1 survey. The numerical simulations used to derive mass accretion histories of different clusters are described in \\S\\ref{sec-sims}.  Our model prescriptions are presented in \\S\\ref{sec-model} and our results in \\S\\ref{sec-results}.  We discuss the implications of our results in \\S\\ref{sec-discuss}, and list our conclusions in \\S\\ref{sec-conc}. ",
        "conclusions": "\\label{sec-conc}  We present a simple model to account for the systematic differences in the star formation properties of galaxies in CNOC1 clusters and the field. The model assumes that the cluster is built through the accretion of field galaxies whose star formation rates gradually decline after entering the cluster as a result of the removal, through tides or ram pressure, of the gaseous envelopes needed to supply normal spirals with the fuel needed to sustain their present star formation rates over a Hubble time. Once calibrated to reproduce observations of nearby spirals the model has no free parameters.  Using cluster mass accretion rates determined from N--body simulations of cluster formation in a $\\Lambda$CDM universe, our model is able to reproduce qualitative and quantitative differences in the mean star formation rates and colors between clusters and the field. The model demonstrates that the origin of radial gradients in these properties is the natural consequence of the strong correlation between radius and accretion times which results from the hierarchical assembly of the cluster.  Interestingly, the model also explains why star formation in the outskirts of clusters is found to be almost a factor of two below the field average as far out as twice the virial radius of the cluster. This is a result of cluster members being pushed onto highly eccentric, loosely bound orbits during major merger events. Our results are robust to evolution in the star formation properties of field galaxies, but somewhat more sensitive to the manner in which fading of cluster galaxies is modeled.  Realistic modeling of the field population star formation histories is required to fully understand the consequences of this fading.  We conclude that the stripping of extended gaseous reservoirs by the cluster environment and the gradual decline that follows gas consumption is likely to be the main mechanism that differentiates the star formation properties of cluster and field galaxies."
    },
    "0004/astro-ph0004234_arXiv.txt": {
        "abstract": "If a black hole formed in a core-collapse supernova is accreting material from the base of the envelope, the accretion luminosity could be observable in the supernova light curve. Here we continue the study of matter fall back onto a black hole in the wake of a supernova and examine realistic supernovae models which allow for an early emergence of the accretion luminosity. Such cases may provide a direct {\\it observational} identification of the black hole formed in the aftermath of the explosion. Our approach combines analytic estimates and fully relativistic, radiation-hydrodynamic numerical computations. We employ a numerical hydrodynamical scaling technique to accommodate the diverse range of dynamical time scales in a single simulation. We find that while in typical Type II supernovae heating by radioactive decays dominates the late-time light curve, low-energy explosions of more massive stars should provide an important exception where the accretion luminosity will emerge while it is still relatively large. Our main focus is on the only current candidate for such an observation, the very unusual SN1997D. Due to the low energy of the explosion and the very small ($2\\!\\times\\!10^{-3}\\;M_\\sun$) inferred mass of \\Coa in the ejected envelope, we find that accretion should become the dominant source of its luminosity during the year 2000. The total luminosity at emergence is expected to lie in the range $0.5-3\\!\\times\\!10^{36}\\;$\\ergss, potentially detectable with HST. We also discuss the more favorable case of explosions which eject negligible amounts of radioactive isotopes and find that the black hole is likely to emerge a few tens of days after the explosion, with a luminosity of $\\sim 10^{37}\\;$\\ergss. ",
        "introduction": "\\label{Sect:Intro}  Core-collapse supernovae mark the death of a massive star and the birth of a compact remnant. Theory suggests that the remnant can be either a neutron star or a black hole, depending on the character of the progenitor and the details of the explosion. Radio pulsar emission has allowed to compile a variety of observational evidence associating neutron stars with sites of known supernovae, but similar evidence for a black hole - supernova connection is still mostly unavailable. Interestingly, indirect evidence that the black hole candidate in the X-ray binary system GRO J1655-40 was formed in a supernova explosion has recently been reported by \\citet{Israelianal99}. The inference is based on detection of high abundances of nitrogen and oxygen on the surface of the companion, which has too low a mass to have produced them by thermonuclear burning, hence indicating that they were deposited there by the supernova that created the black hole.  The formation of a compact object in a supernova can be inferred {\\it directly} if the object causes an observable effect on the total luminosity which follows the explosion -- i.e., the {\\it light curve}. In particular, some material from the base of the expanding envelope may remain bound to the compact object and continuously fall back onto it, thus generating an accretion luminosity. \\citet{PaperI} have recently performed a self-consistent investigation of the accretion luminosity generated by spherically symmetric ``fallback'' of matter onto a black hole in the wake of a supernova. With the aid of a general-relativistic, radiation hydrodynamic Lagrangian code they showed that, at late times, the fallback evolves quasi-stationarily with the accretion luminosity obeying the analytic expression found by \\citet{Blondin86} for steady state, spherical hypercritical accretion\\footnote{We use the term ``hypercritical'' to describe an accretion rate $\\dot{M}$ that largely exceeds the Eddington limit for the accreting object, i.e., $\\dot{M}\\gg \\dot{M}_{Edd}\\equiv L_{Edd}/c^2$. The actual luminosity that arises would still be sub-Eddington if the efficiency of converting energy into radiative luminosity is low enough.}. The accretion luminosity declines secularly with a time dependence of $t^{-25/18}$. This decay is driven by the decrease in the accretion rate due to the continuous expansion of the supernova envelope (Colpi, Shapiro \\& Wasserman 1996),  Such an accretion luminosity will produce a distinct signature on the total light curve if and when it becomes comparable to the output of other power sources, namely the initial internal energy of the envelope and decays of radioactive isotopes synthesized in the explosion. It is well established that the efficiency of converting spherical accretion onto a black hole into radiative luminosity may be quite low (Shapiro 1973; Blondin 1986; Wandel, Yahil \\& Milgrom 1986; Park 1990; Nobili, Turolla \\& Zampieri 1991; Zampieri, Miller \\& Turolla 1996), and, in general, the accretion luminosity will be undetectable in comparison with the luminosity generated by radioactive decays. This is particularly important when compared to the case of accretion onto a neutron star, which is expected to be far more efficient in generating an accretion luminosity. For the last decade this distinction has been specifically applied to the case of SN1987A: \\citet{Chevalier89} and \\citet{HouckChev91} found that accretion onto a newly formed neutron star in SN1987A would generate an accretion luminosity at the Eddington-rate ($\\ga 10^{38}\\;$\\ergss) within a few months after the explosion, and the luminosity would persist at this rate for several years\\footnote{It is noteworthy that magnetic dipole emission of a Crab-like pulsar would deposit energy in the envelope at a similar rate \\citep{Woosleyal89}.}. The proximity of SN1987A has allowed observations to follow the light curve for over 3000 days (to a present luminosity of $\\sim10^{36}\\;$\\ergss), and the light curve has been completely consistent with heating by decays of \\Coa, \\Cob and \\Ti. The absence of any evidence for accretion luminosity in the case of SN1987A has been used to argue that the compact remnant in SN1987A is a low mass black hole, implying significant consequences regarding the nuclear equation of state \\citep{BetheBrown95}. It has also been argued, however, that this absence could still be consistent with a weakly magnetized neutron star, if the accretion flow has become dynamically unstable (Fryer, Colgate \\& Pinto 1999).  A positive indication of the presence of a black hole would be an actual detection of the accretion luminosity. The power-law time dependence would make the accretion luminosity easily discernible from luminosity generated by radioactive heating, which decreases exponentially with time. Furthermore, the difference in the time dependence implies that the accretion luminosity {\\it must} eventually become dominant over radioactive decays, so in principle, the black hole always ``emerges'' in supernovae light curves. However, the potential for observation is usually quite slim: for example, \\citet{PaperI} estimated that if a low mass black hole was formed in SN1987A, the accretion luminosity will not become dominant over that of heating by \\Ti decays until $\\sim900$ years after the explosion, when the luminosity will have dropped to a mere $10^{32}\\;$\\ergss. The abundance of radioactive elements and the amount of fallback in SN1987A are rather typical of core collapse supernovae of low to intermediate mass progenitors (12--25 $M_\\sun$) so that, if a black hole is formed, the ``emergence'' of the accretion luminosity occurs only at such late times and low luminosities that it is of no observational consequence.  An important exception may exist, however, in the explosions of more massive stars. In their survey of supernova explosions and nucleosynthesis, \\citet{WoosWeav95} find that more massive stars tend to produce more massive compact remnants: the original iron core is larger, leading to a more massive proto-neutron star and to larger amounts of fallback (see also Fryer 1999). Correspondingly, larger mass progenitors are more likely to produce black holes (as the remnant mass would be larger than the maximum mass of stable neutron stars), and to have larger amounts of material that remain bound to the black hole as a reservoir for late time fallback. Just as important, the final abundance of the key radioactive isotopes (\\Ni, $^{57}$Ni and \\Ti) in the expanding envelope is quite sensitive to the location of the cutoff that defines the material which settles on the black hole during the early stages (while the explosion is still proceeding). For the more massive progenitors the cutoff is further out, and since the radioactive elements are synthesized in the innermost layers above the proto-neutron star, the total mass of these elements in the ejected envelopes of higher mass stars can be significantly smaller than in ``standard'' supernovae. In fact, progenitor of masses in the range $30-40\\;M_\\sun$, are expected to be practically free of radioactive elements that can power the bolometric light curve \\citep{WoosWeav95}. For such stars, black hole accretion luminosity should emerge immediately after the decay of the recombination peak (tens of days after the explosions), while it is still relatively powerful, and hence potentially observable.  Naturally, the more massive stars are a small minority in the stellar population, as are their explosions among core collapse supernovae. However, one recent supernova, SN1997D in NGC 1536 \\citep{deMello97}, may present a border-line candidate for a direct detection of black hole emergence. Its observed light curve indicates that that the total mass of \\Coa (the daughter of \\Ni) in the envelope is only $\\sim2\\!\\times\\!10^{-3}\\;M_\\sun$, a factor of $\\sim35$ lower than SN1987A. By analyzing the light curve and spectra, \\citet{Turatto98} suggested that SN1997D was a low energy ($E_{tot}\\sim 4\\!\\times\\!10^{50}\\;$ergs) explosion of a $26\\;M_\\sun$ star. The expected mass of the remnant is about $3\\;M_\\sun$ (formed by a $1.2\\;M_\\sun$ early fallback on a $1.8\\;M_\\sun$ collapsed core), hence most likely a black hole. Based on the inferred average properties of the material available for late time fallback and on the late time behavior of the accretion luminosity found by \\cite{PaperI}, Zampieri, Shapiro \\& Colpi (1998b) estimated that accretion luminosity will emerge in SN1997D in the optical band within {\\it only a few years} after the explosion.  In this work we revisit the issue of black hole emergence in supernova due to spherically-symmetric late-time accretion. In particular, we investigate a more robust estimate for our ``proto-type'' case, SN1997D. We furnish revised analytic estimates and a detailed numerical investigation of the radiation-hydrodynamic evolution of the ejecta, incorporating the effects of a realistic chemical composition, opacities and radioactive heating on the accretion history and luminosity. The numerical study is based on an improved version of the radiation-hydrodynamic code used by \\citet{PaperI}, where the main modifications are the inclusion of realistic envelope composition and opacities (rather than pure hydrogen) and of heating by radioactive decays. We find that a $3\\;M_\\sun$ black hole should emerge in the SN1997D light curve at about 1000-1500 days after the explosion -- i.e., {\\it during late 1999 and 2000} -- when the total bolometric luminosity is $\\sim 0.5-3\\!\\times\\!10^{36}\\;$\\ergss, and is marginally detectable with HST.  In the course of our computational study we were compelled to introduce a rescaling scheme in order to numerically follow the accretion history and light curve for several years. The huge dynamical range of pertinent time scales renders a full scale time-integration numerically impractical. Our solution has been to study a scaled model which evolves more quickly and allows a simple relation between the properties and history of the rescaled model and the true one. We note that this scheme, which is presented in detail in the appendix, is potentially suitable for other problems which involve time-dependent radiation hydrodynamics spanning a large dynamic range.  We begin our discussion in \\S~\\ref{Sect:basics} with a summary of the characteristic features of supernova light curve evolution when accretion onto a central black hole is occurring. Quantitative relations between the light curve magnitude and the relevant time scales are also reviewed. The principal features of the numerical code of \\citet{PaperI} and the modifications incorporated for this work are described in \\S~\\ref{Sect:code}. The light curve of SN1997D is examined in detail in \\S~\\ref{Sect:SN1997D}, where we present analytic and numerical results, including our estimates regarding black hole emergence. Prospects of observing black hole emergence in other supernovae are discussed in \\S~\\ref{Sect:otherSN}, and further discussion and conclusions are offered in \\S~\\ref{Sect:conc}. ",
        "conclusions": "\\label{Sect:conc}  In this work we have examined the emergence of luminosity due to spherical accretion onto black holes formed in realistic supernovae. The study is based on a combination of analytic estimates and numerical simulations. Our main focus has been on SN1997D, where the low explosion energy and very low inferred abundance of \\Coa in its ejecta provide a possible (and currently, the only) candidate where the emergence of a black hole in a supernova may actually be observable.  We confirmed the main features of the accretion history and luminosity as derived by \\citet{PaperI}. At late times, the accretion flow settles on a dust-like motion, and the accretion rate falls off as a $t^{-5/3}$ in time \\citep{CSW96}. Accretion proceeds in a sequence of quasistationary states, so the accretion luminosity achieves the expected magnitude of a hypercritical, spherical accretion flow \\citep{Blondin86}, accounting for the secular decay in time of the accretion rate. The resulting time dependence of the late-time accretion luminosity is $L(t)\\propto t^{-25/18}$. This relation marks a fundamental feature regarding emergence: since the main competitor to accretion in powering the late time light curve is radioactive heating which decays exponentially with time, emergence of the accretion luminosity is inevitable if spherical accretion does indeed persist. While in typical Type II supernovae the abundance of radioactive isotopes is expected to be large enough to prevent a practical observation of the emergence, more rare explosions where this abundance is significantly reduced and fallback is enhanced provide important exceptions.  The case of SN1997D offers a unique opportunity to examine black hole emergence \\citep{SN97Dlet}, in view of its inferred explosion energy of $\\sim 4\\!\\times\\! 10^{50}\\;$\\ergss~ and very low abundance of \\Coa, i.e., only $2\\!\\times\\!10^{-3}\\;M_\\sun$ \\citep{Turatto98, Benetti00}. We have have examined the expected emergence of a black hole in SN1997D analytically and numerically. In the analytic study we noted that the early-time accretion is almost certain to generate an Eddington-rate luminosity, which will modify the accretion history and decrease the late-time accretion rate. In the numerical investigation we have included the effects of a variable envelope composition. In particular, we incorporated heating due to radioactive decays, including a finite (time-dependent) $\\gamma-$ray optical depth. The wide diversity of time scales involved compelled us to use a rescaling scheme in the numerical simulations (see in the appendix).  Our various analyses offer a consistent assessment of the emergence of the black hole in SN1997D. We confirm the preliminary estimate of \\citet{SN97Dlet} that a $3\\;M_\\sun$ should emerge in SN1997D about $1000\\;$days after the explosion, in contrast to hundreds of years in the case of ``standard'' Type II supernovae (like SN1987A) where the abundances of radioactive elements are significantly higher. We show that the dominant parameter in determining the luminosity at emergence is the amount of kinetic energy carried in the post-explosion helium-rich layer (which is the source of late-time accretion). For realistic values of this energy, we estimate that the total luminosity at emergence will lie in the range $5\\!\\times\\! 10^{35}\\la L_{tot} \\la 3\\!\\times\\! 10^{36}\\;$\\ergss, emerging at $1000-1500\\;$days after the explosion. In our numerical study we have examined the best fit model of \\citet{Turatto98} for SN1997D. The simulations suggest that, for this particular model, the black hole emerges at about $1050\\;$days after the explosion, with $L_{tot}\\approx 6.5\\!\\times\\!10^{35}\\;$\\ergss. The results of the simulation, which is based on a consistent, fully-relativistic, radiation-hydrodynamics code, provides significant support to the analytic approach. Most notably, it demonstrates that, after passing through an early radiation-pressure-limited stage which reduces the accretion rate, the accretion flow does settle on the dust-like solution and that the influence of the radioactive decays on the hydrodynamic evolution is negligible for the low abundances inferred in SN1997D.  The time of emergence and the corresponding luminosity depend on the abundances of \\Cob and \\Ti in the envelope, which are unknown a-priori. However, as long as the accretion luminosity is significantly larger than the threshold set by positron emission in \\Ti decays, both emergence time and luminosity will not vary substantially. At the time of emergence $\\gamma-$ray escape from the envelope is non-negligible and increases rather rapidly with time, as the envelope's optical depth drops. Correspondingly, the bolometric luminosity due to radioactive heating decreases even faster than exponentially with time.  Although our results appear to be quite robust within the context of the present study, some uncertainties arise due to effects not considered here. Most critical is the assumption of perfectly spherical accretion: if the bound material can maintain sufficient angular momentum, a centrifugal barrier would drive the flow into a disk structure \\citep{Chevalier96}. The radiative efficiency of hypercritical disk accretion is uncertain, complicating any a-priori estimate regarding the character of the late-time accretion luminosity. The residual angular momentum is likely to be dependent on the details of the progenitor and the explosion (we note that for SN1987A the observed neutrino burst seemed consistent with nonrotating models \\citep{Burrows88}, indicating that angular momentum of matter close to the collapsed core did not play a significant role). In principle, late time accretion occurs many viscous time scales after the explosion, and much of the initial angular momentum could be lost, but the minimum specific angular momentum required to force disk accretion is quite small (about $(GM_{BH} R_{ISCO}c)^{1/2}$, where $R_{ISCO}=6GM/c^2$ is the innermost-stable circular orbit near a black hole). Furthermore, even if at late times angular momentum is negligible, some residual effect may arise if the character of the early-time accretion has been sufficiently altered (Mineshige et al. 1997 found through simulations that sufficient angular momentum strongly affects the magnitude and time dependence of the accretion during the first {\\it hours} after collapse). We defer further investigation of fallback including angular momentum to future work.  Additional effects could arise due to convection, mixing and clumping of envelope material (and radioactive isotopes in particular), and dust extinction, although we expect that at a time of $\\gtrsim3\\;$years after the explosion, any dust formed will have become transparent to the underlying luminosity. This may not be true if photoionization by the central accreting black hole (not included in our calculation) is important.  Our results illuminate the uniqueness of SN1997D. It is the low abundance of radioactive elements, combined with the low explosion energy in general, that allows the black hole to emerge only a few years after the explosion, at which time the absolute magnitude of the accretion luminosity is still quite large (recall that for ``standard'' explosions we expect emergence luminosities of $10^{32}-10^{33}\\;$\\ergss). Since gas and radiation are in LTE in the inner accreting region and the photosphere is well localized in radius, the emerging spectrum is expected to be very similar to a blackbody at the photospheric temperature ($\\sim 7500$ K) with the superposition of heavy line elements. We therefore expect that emission will be mainly in the optical band. At a distance of $\\sim 14\\;$Mpc, the corresponding apparent magnitude for the range of emergence luminosities predicted for SN1997D is $m_V\\approx 28-30$ and $m_R\\sim m_I \\approx 27-29$. The higher luminosity (lower magnitude) end of this range coincides with the detection threshold of the HST STIS camera, indeed making emergence in SN1997D potentially {\\it observable}. We note that the capabilities of HST are needed both because of the faintness of the source at this time, and the resolution, required to distinguish the supernova over the background of the parent galaxy NGC 1536. A recent estimate (M.~Turatto 1999, private communication) suggests that the required exposure time for the HST STIS to resolve the apparent magnitude expected at emergence in SN1997D at a signal to noise (S/N) ratio of 10 would be $\\sim 25000\\;$seconds. The accretion luminosity should be distinguishable from radioactive heating (or from circumstellar interaction, which has not been observed in SN1997D so far, Benetti et al.~2000) due to its unique power law dependence. This would be especially true if the abundances of \\Cob and \\Ti are negligible, since after emergence the luminosity would be due only to accretion. But, even if there is a finite contribution from these two isotopes to the light curve, the presence of a power-law source may still be identifiable if several measurements with reasonable accuracy could be made. The novel consequences of a positive observation are self-evident - the first {\\it direct} observational evidence that a black hole can be formed by an otherwise ``successful'' supernova. Although our main focus has been SN1997D, other theoretical Type II supernova models offer more favorable conditions for an actual observation of black hole emergence. The survey of \\citet{WoosWeav95} suggests that progenitors with masses in range $30-40\\;M_\\sun$ are quite likely to leave behind a remnant of mass $3-10\\;M_\\sun$ and a practically radioactive-free envelope, as all the iron group isotopes are engulfed by the early fallback. Our calculations in \\S~\\ref{Sect:otherSN} show that accretion luminosity will emerge after recombination depletes the initial internal energy of the envelope, hence after {\\it a few tens of days}. The corresponding luminosity at emergence will be $\\sim 10^{37}\\;$\\ergss, which with the present capabilities of HST would allow for detection at emergence perhaps out to $20-25\\;$Mpc. The luminosity at emergence is higher than for SN1997D so that the galactic background would be less important, and ground based instruments might be relevant as well.  How often can we expect an opportunity to observe a black hole emerging in a supernova? An upper limit is, of course, the rate at which black hole forming supernovae occur. Adopting the recent results of \\citet{FryKal99} that about $10\\%$ of core collapse supernovae make black holes, we expect a rate of about one event per thousand years per galaxy. There are somewhat less than one thousand galaxies at a distance $\\la 20\\;$Mpc from the Milky way, and hence there may be approximately one observable event (equal or better than SN1997D) per several years. This estimate is an upper limit, since a significant fraction of high mass progenitors will end their lives as Type Ib/Ic supernovae after losing their hydrogen envelope, and such supernovae are unfavorable for a practical observation at emergence. Furthermore, even ``perfect'' candidates may not allow for an actual observation, as seems to have been the case in SN1994W. Following its recombination peak, the tail of the SN1994W light curve failed to settle on an exponential decay that would arise due to \\Coa decay \\citep{Sollerman94W}. Interestingly, the last data point led to a limit of $\\le2\\!\\times\\!10^{-3}$ on the mass of \\Coa in the ejected envelope. However, the light curve showed a very steep power-law in time until it dropped below the detection threshold some 220 days after the explosion. The luminosity during this phase was much too high to be due to accretion ($10^{40}-10^{42}\\;$\\ergss) and \\citet{Sollerman94W} found that it was most likely caused by interaction with circumstellar material, with dust formation also playing an important role. We can only speculate that in the absence of circumstellar interaction and dust extinction, the light curve would have declined exponentially after the recombination peak and the accretion luminosity may have emerged while still marginally detectable. The case of SN1994W implies that the rate of potential observable black hole emergence may be significantly smaller than that of black hole forming Type II supernovae.  Nonetheless, we urge that any Type II supernovae which shows a diminished \\Coa abundance should be monitored, and its light curve be tracked for several months after the explosion. The prospects of observing black hole emergence may improve significantly with potential future instruments, such as NGST or a dedicated faint-supernovae project."
    },
    "0004/astro-ph0004002_arXiv.txt": {
        "abstract": "Current submillimetre surveys are hindered in their ability to reveal detailed information on the epoch of galaxy formation and the evolutionary history of a high-redshift starburst galaxy population. The difficulties are due to the small primary apertures ($\\rm{D < 15}$-m) of existing submillimetre telescopes and the limited sensitivities of their first generation of bolometer cameras.  This situation is changing rapidly due to a variety of powerful new ground-based, airborne and satellite FIR to millimetre wavelength facilities. Improving our understanding of the luminosity and clustering evolution provides the motivation for conducting cosmological submillimetre and millimetre surveys.  It is therefore important that we quantify the limitations of the future surveys and the significance of the results that can be drawn from them.  In this paper we present simulated surveys which are made as realistic as possible in order to address some key issues confronting existing and forthcoming surveys.  We discuss the results from simulations with a range of wavelengths (200$\\mu$m -- 1.1\\,mm), spatial resolutions (6\\arcsec -- 27\\arcsec) and flux densities (0.01--310 mJy).  We address how the measured source-counts could be affected by resolution and confusion, by the survey sensitivity and noise, and by the sampling variance due to clustering and shot-noise. ",
        "introduction": "By early 2001 the initial extensive programme of extragalactic SCUBA (850$\\mu$m) surveys conducted on the 15-m JCMT will be completed.  The targets of these sub-mm surveys include lensing clusters (Smail \\etal\\ 1997), the Hubble Deep Field (Hughes \\etal\\ 1998), the Hawaii Deep Survey fields (Barger \\etal\\ 1998, 1999), the Canada-France Redshift Survey fields (Eales \\etal\\ 1999, Lilly \\etal\\ 1999) and the UK 8\\,mJy SCUBA surveys of the ISOPHOT ELAIS and Lockman Hole field (in prep.)  These blank-field sub-mm surveys, which cover areas of $\\sim 0.002 - 0.2$\\,deg$^{2}$, allow a measure of the $850\\mu$m source-counts between flux densities of 1--15mJy.  The relatively small survey areas are due to the low mapping speed (a combination of the field-of-view and sensitivity of the bolometer array, and the necessary overheads to fully sample the focal-plane with $\\rm >0.5F\\lambda$ feedhorns, {\\it e.g.} 16 separate secondary-mirror positions are required to fully sample the SCUBA array). The restricted 850$\\mu$m flux range in the measured source-counts is the result of a source-confusion limit of $\\sim 2$\\,mJy (due to a small primary aperture or map resolution) at faint flux densities and the small survey area at the bright end, given the intrinsic low source-density of bright sub-mm sources ({\\it e.g.}  $N[S_{850_{\\mu m}} > 20\\rm \\,mJy] < 70~~ sources/deg^{2}$).     The limitations on the accuracy of our understanding of high-z galaxy evolution from sub-mm surveys has been described elsewhere (Hughes 2000). To improve the constraints on the competing evolutionary models, future sub-mm and mm surveys must extend their wavelength coverage, increase their survey area and sensitivity. To quantify the advantages of these forthcoming surveys we have simulated the extra-galactic sky at various sub-mm and mm wavelengths (200-3000$\\mu$m), spatial resolutions ($\\theta_{FWHM} \\sim 1-120$\\arcsec) and sensitivities ($S >0.01$ mJy). We include in the simulations contributions from CMB primary fluctuations, S-Z clusters, an evolving extra-galactic population of star-burst galaxies and high-latitude galactic-plane cirrus emission.  These simulated surveys allow us to compare the measured source-counts for a wide variety of current and future FIR and sub-mm/mm telescopes ({\\it e.g.} SIRTF, FIRST, MAP, PLANCK, BLAST, SOFIA, CSO, JCMT, LMT, GBT, ALMA).  We discuss these multi-wavelength simulations in more detail elsewhere (Hughes \\& Gazta\\~{n}aga 2000). In this paper we concentrate on illustrating the effects of galaxy clustering and low spatial resolution on the measured source-counts at 1100$\\mu$m, 850$\\mu$m, 350$\\mu$m and 200$\\mu$m. ",
        "conclusions": "We have presented simulated surveys which are made as realistic as possible in order to address some key issues confronting existing and forthcoming millimetre and submillimetre surveys. These simulations assume a model for the luminosity and clustering evolution which is poorly constrained by current data. Improving our understanding of this evolution provides the motivation for conducting submillimetre and millimetre surveys. It is therefore important that we quantify the limitations of these surveys and the significance of the results drawn from them.  In this paper we discuss the results from survey simulations with a range of wavelengths (200$\\mu$m -- 1.1\\,mm), spatial resolutions (6\\arcsec -- 27\\arcsec) and flux densities (0.01--310 mJy). Several issues have been addressed related to the estimation of the measured source-counts in future surveys: resolution and confusion; survey sensitivity and noise; and sampling variance due to clustering and shot-noise.  The redshift information, determined from the millimetre and submillimetre colours of individual sources (Hughes 2000), is an essential ingredient to discriminate between the possible evolutionary models that describe the star formation history of galaxies.  Clustering in future wide-area submillimetre surveys might also help to understand some of these issues. Both the number counts and the two-point clustering pattern depend strongly on the cosmological model, but its combination can break some degeneracies. For example if galaxy formation occurs in the rare peaks of the underlying matter field then the resulting  galaxy clustering should be stronger than if galaxy formation occurred in random places. The measure of the coherence of the clustering pattern, as described by the higher-order statistics (which is apparent in our simulated maps), can also be used to test how star forming galaxies trace the underlying mass ({\\it e.g.} Gazta\\~naga 2000, and references therein). This issues will be address in more detail elsewhere."
    },
    "0004/astro-ph0004372_arXiv.txt": {
        "abstract": "We describe and evaluate a new instrument design called a Tunable Echelle Imager (TEI). In this instrument, the output from an imaging Fabry-Perot interferometer is cross-dispersed by a grism in one direction and dispersed by an echelle grating in the perpendicular direction. This forms a mosaic of different narrow-band images of the same field on a detector. It offers a distinct wavelength multiplex advantage over a traditional imaging Fabry-Perot device.  Potential applications of the TEI include spectrophotometric imaging and OH-suppressed imaging by rejection. ",
        "introduction": "\\label{sec:intro}  One of the most versatile instruments for astronomical narrow-band imaging is the gap-scanning Fabry-Perot (FP) etalon. Light passing through a FP produces a series of wavelength passbands called orders (examples of the transmission are shown in Figure~\\ref{fig:gap-sp}). \\begin{figure} \\plotone{fg1.ps} \\caption{Transmission profiles for Fabry-Perot filters with two different gap spacings.  The upper plot shows orders from 14 on the right to 23 on the left; the lower plot shows orders 24 to~40.} \\label{fig:gap-sp} \\end{figure} The positions of the passbands can be moved by varying the separation of the cavity surfaces \\citep[e.g.,][]{ARRH81,Ath95,BJ98pasa}. However, in order to image a field with a FP, broad- or intermediate-band filters are usually used to block out the light from all the passbands except the one of interest.  An alternative method is to use a dispersive element to separate the images formed by different FP orders. Recently this principle has been used in the PYTHEAS instrument \\citep{LBC95} and the GraF instrument (\\citeauthor{CL95} \\citeyear{CL95}; \\citeauthor{CLL99} \\citeyear{CLL99}, \\citeyear{CLR99}). In this paper, we consider the use of a cross-dispersed echelle grating\\footnote{We use the term `echelle' for any multi-order diffraction grating used above about the fifth order, rather than distinguishing between an `echellette' and an `echelle' grating.} crossed with a gap-scanning FP called a Tunable Echelle Imager (TEI). This forms a mosaic of narrow-band images with different wavelengths on a detector, and by scanning the plate separation of the FP, a complete data cube can be obtained (2D image plus 1D spectral).  \\subsection{The basic physics}  The wavelength positions of the FP orders are described by the equation for constructive interference:\\footnote{Note that Equation~\\ref{eqn:orders-fp} does not include the effect of phase changes on reflection \\citep{ARRH81,JB98}.  Such phase changes only have a significant impact for low-order etalons which are not considered here.} \\begin{equation} m_{\\rm fp} \\lambda = 2 \\mu d \\cos \\theta \\: , \\label{eqn:orders-fp} \\end{equation} where $m_{\\rm fp}$ is the FP order, $\\mu$ is the refractive index of the medium in the cavity gap, $d$ is the plate separation (or gap spacing) and $\\theta$ is the angle of incidence.  The bandwidth of each FP order is dependent on the reflectivity and defects of the coated surfaces (non-parallelism). The ratio between inter-order wavelength spacing ($\\Delta \\lambda$) and the FWHM bandwidth ($\\delta \\lambda$) is called the effective finesse: \\begin{equation} N = \\frac{\\Delta \\lambda}{\\delta \\lambda} \\: . \\label{eqn:finesse} \\end{equation} The finesse is also related to the contrast between the peak transmission of an order and the minimum transmission between orders. For contrast values greater than 100, 250, 500 or 1000, the finesse needs to be greater than 16, 25, 35 or 50, respectively. At high finesse, the efficiency can drop significantly due to defects \\citep{ARRH81}. A finesse between 20 and 40 is optimal for a high-throughput imaging FP \\citep{Bla95}.  The echelle grating (EG) disperses the FP orders using multi-order diffraction. The diffracted angle ($\\beta$) for gratings that are not immersed between prisms is determined from the grating equation: \\begin{equation} m_{\\rm eg} \\lambda = \\mu \\Lambda (\\sin \\alpha + \\sin \\beta) \\: , \\label{eqn:orders-eg} \\end{equation} where $m_{\\rm eg}$ is the EG order, $\\mu$ is the refractive index of air, $\\Lambda$ is the grating period and $\\alpha$ is the incidence angle in air.  With an echelle grating, several wavelengths corresponding to different EG orders are diffracted at or near the same angle. For this reason, a low-dispersion element is needed to separate the EG orders. This could be another grating, one or more prisms, or a grism. ",
        "conclusions": "\\subsection{Observing and data reduction}  To fully sample the spectral dimension at the resolving power of the TEI setup, it will be necessary to take a number of exposures (at different FP gap spacings) equal to or greater than the finesse of the FP.  The reason for a greater number of exposures is because the requirements for sampling at the shortest wavelength may be different from the longest wavelength.  In many cases, there will be double sampling of some wavelengths near the shortest wavelength, these could be used as a self-calibration check in reconstructing the spectral dimension.\\footnote{\\cite{BT89} noted that it was easier to reconstruct the $\\lambda$ dimension than an image dimension (e.g., as with a scanning slit or IFU).} Spectrophotometric measurements could be obtained in photometric conditions by calibrating with a standard star.  The TEI can mimic a true 3D spectrograph by scanning, as described above. Some applications may not need scanning or only partial scanning (a few exposures) if the sampling of the spectral dimension is adequate. For example, multiple narrow-band imaging could be used to selectively combine images in order to suppress sky-emission lines, thus mimicking an OH-suppression instrument.  Wavelength calibration of a TEI would be made using arc lamps to define the wavelength for each order and gap spacing for the etalon. The calibration is then uniquely defined for a given gap spacing and should be stable during an exposure using capacitance micrometry \\citep{HRA84}. Advantageously, the spectral stability of the TEI is not subject to instrument flexure as with a normal Cassegrain spectrograph.  In some of the example image formats, it appears that optimal use is not made of the detector in the sense that there are repeated orders (e.g., Figure~\\ref{fig:im5}). However, if every order is repeated twice on the detector, it may increase the efficiency and allow the possibility of removing image defects (e.g., hot pixels, cosmic rays, bad columns) by comparing the same narrow-band image.  \\subsection{Aperture width and the PSF}  The dispersion of the gratings spreads out the light from a single FP order, contributing to the point spread function (PSF) of the image. The FWHM of this one-dimensional contribution is equal to the separation (on the detector) between adjacent FP orders divided by the finesse.  The lines in the bottom-left corner of each image in the Figures~\\ref{fig:im1}--\\ref{fig:im6} represent this PSF.  If the aperture size is maximized for a particular echelle grating and FP setup then the FWHM of the dispersive PSF is equal to the width of the aperture divided by the finesse, for the lowest wavelength. Therefore, if a certain image resolution is required in both spatial dimensions (the TEI only degrades resolution in one direction), then there is a limit on the setups to those with certain aperture widths. For example, in the case of adaptive optics a resolution of 0.1\\,arcsec may be required limiting the aperture width to 3\\,arcsec (with a finesse of 30). Alternatively, if 0.5\\,arcsec is satisfactory then an aperture width of 15\\,arcsec would be suitable. If good image resolution is only required in one direction then the aperture width can be larger still.  We note that the elongated PSF is a common occurrence in radio interferometry where many software tools exist to treat such data \\citep{PSB94}. The TEI situation is somewhat better because an instrument rotation allows us to sample the field with orthogonal PSF orientations.  \\subsection{Anamorphic factor}  The ratio of the beam sizes after and before diffraction by a grating is called the anamorphic factor. For a reflection grating or a transmission grating without prisms, the anamorphic factor \\begin{equation} A = \\frac{\\cos \\beta}{\\cos \\alpha} \\: , \\label{eqn:anamorphic} \\end{equation} where $\\alpha$ and $\\beta$ are the incident and diffracted angles in air.  The change in beam size is inverse to the change in angular size (the angle between light rays in the collimated beam). A grating setup which has an anamorphic factor above unity can increase the resolution because the angular size of the aperture (in the dispersion direction) is reduced. However, the increase in resolution is less than the anamorphic factor because the dispersion is also reduced (though not by as much).  Figures~\\ref{fig:im4} to~\\ref{fig:im5} show the effect of varying the anamorphic factor with the same resolution and EG orders (see Table~\\ref{tab:setup-para}). Larger apertures can be used with the blaze-to-collimator setup (increasing beam size) even though the images appear smaller on the detector (Fig.~\\ref{fig:im5}). Using the example of a 2K\\,x\\,2K CCD and a 6\\,x\\,6\\,arcminute field, the pixel sampling on the detector would change from 0.18 to 0.28\\,arcsec/pixel across the width of the aperture (14\\,arcsec). The change in sampling is not critical in this case because the dispersion PSF is about 0.47\\,arcsec (FWHM). In quadrature with 0.5\\,arcsec seeing, we find that the sampling is still adequate. In other cases a large anamorphic factor could cause undersampling of the PSF (instrumental and seeing) on the detector. Ideally, the change in sampling caused by an increase in beam size could be matched to the decrease in image resolution. A decrease in beam size cannot cause undersampling on the detector.  With a limited number of gratings, it would be useful to be able to use either a transmission or a reflection grating setup, with an increase or decrease in beam size. In this way, the aperture width could be matched more easily to the desired spectral resolving power.  A decrease or no change in beam size after dispersion has the advantage of instrumental compactness, in that, the camera optics could be smaller than if an increase in beam size is allowed.  \\subsection{Finesse}  Choosing the finesse is tradeoff between a number of factors. A lower finesse has the advantage of needing less scans to complete the data cube. In effect, the total throughput of the instrument is approximately proportional to $1/N$. A higher finesse has the advantage of higher spectral and image resolution, and better suppression of the light between the narrow bands.  \\subsection{Comparison with a mono-order FP}  The advantages of the TEI over a FP, where a single order has been selected with a blocking filter, include: imaging in many narrow bands simultaneously, and; imaging at any wavelength and spectral resolution without the need for many order-sorter filters. The disadvantages include: a smaller field of view; compromised image resolution (in one direction) because of the PSF caused by the dispersion, and; efficiency loss for a single order due to the extra dispersive elements.  A significant aspect of the TEI is that the field, although smaller, will be monochromatic or nearly monochromatic in many cases. In fact, the size of the field decreases as the spectral resolving power ($R$) is increased which is in the same sense as the \\citeauthor{Jac54} (\\citeyear{Jac54}, \\citeyear{Jac60}) criterion for a monochromatic field: \\begin{equation} R \\, \\Omega = 2 \\pi \\: , \\end{equation} where $\\Omega$ is the solid angle at the FP. The TEI parameters can be chosen to give maximum spectral coverage for the monochromatic field.  \\subsection{Comparison with the GraF concept}  The GraF concept uses a single dispersive element to separate the FP orders \\citep{CL95}. Such a concept has been used in the GRAF/ADONIS instrument on the ESO 3.6-m telescope (\\citeauthor{CLL99} \\citeyear{CLL99}, \\citeyear{CLR99}). This instrument operates in the near-IR with adaptive optics used to obtain 0.1\\,arcsec image resolution. As an example, it can provide 10 images of 9'' by 0.9'' in the $H$ band at a spectral resolution of 10000 \\citep{CLR99}. The field is thus small and narrow for high image resolution.  Here, we compare the TEI with a more general GraF instrument. A lower image resolution GraF instrument could, for example, have a field of 6\\,arcmin by 20\\,arcsec.  The advantage of the TEI is the ability to image in more FP orders simultaneously, i.e., providing a larger wavelength coverage except for the lowest spectral resolutions. The disadvantages include: efficiency losses due to the extra dispersive element, and; a shorter field of view.  In general, the shorter field of view is not a significant disadvantage because (i) the field will not be monochromatic over a longer field of view and (ii) astrophysical objects of interest may not be extended in one direction only, over the length of field possible with a GraF instrument (arcminutes).  The TEI could be converted to a GraF instrument by replacing the echelle grating with a mirror, or positioning the camera for direct imaging.  \\subsection{Comparison with an echelle grating}  The advantages of the TEI over a standard echelle grating include: imaging, and; better spectrophotometry because a wider aperture can be used without compromising wavelength calibration or resolution. A disadvantage is that it requires scanning (tens of exposures) to fully sample a wavelength region.  The TEI can reach the same resolution as a standard echelle with a wider aperture and a smaller beam. The first factor allows better sampling of extended sources and the second allows the TEI to be a Cassegrain instrument with a smaller camera. These two factors partially offset the disadvantage of the scanning requirement in terms of efficiency. In certain cases, the TEI could be comparable to a slit echelle spectrograph for high-resolution spectroscopy of extended sources. For spectrophotometry and imaging, the TEI is significantly better.  \\subsection{Comparison with an IFS}  In recent years, there has been great interest in the use of integral-field spectrographs (IFS).  These include image slicers \\citep[e.g.,][]{Con00}, integral-field units and the TIGER approach \\citep{Cour82,CGM88,BAB95}.  Both TEI and IFS instruments are true 3D spectrographs. However, they differ in one key respect. For the IFS, the incident wavefront is divided and dispersed, in contrast to the TEI where the wavefront is only dispersed.  The advantages of the TEI compared to an IFS include: better imaging and spectrophotometry because images do not need to be reconstructed; a larger field of view for a given detector size, and; potentially higher spectral resolution. The disadvantages include: the scanning requirement to fully sample a wavelength region, and; compromised image resolution (in one direction) because of the PSF.  Integral-field spectrographs divide the focal plane and therefore image integrity can be compromised (i.e., they are not spectrophotometric). So while an IFU is more efficient than a TEI device, for the same image size and spectral resolution requirements, there are certain cases where the TEI's undivided approach to 3D spectro-imaging could be more beneficial.  \\subsection{Some applications}  The potential applications include: spectro-imaging and/or high-resolution spectroscopy of extended sources, through production of a full data cube; OH suppression in the 600--1600\\,nm wavelength range ($RIzJH$), and; targeted line imaging.  The OH-suppression method is to selectively combine the images produced by a TEI, rejecting those images that contain OH-line emission (or applying a lower weight). This will considerably reduce the sky background especially in the $J$ and $H$ bands. Figure~\\ref{fig:oh-suppress} shows an example of a FP transmission profile plotted over the $J$-band sky emission spectrum.  \\begin{figure} \\plotone{fg12.ps} \\caption{Plots of the transmission profile of a FP (dotted line) with a gap spacing of 220\\micr\\ and a finesse of 30, and an OH sky emission spectrum (solid line), in the $J$ band. The TEI images produced by the different FP orders can be combined selectively to exclude most of the OH emission.  The sky emission spectrum is the same as that used by \\cite{OB98}.} \\label{fig:oh-suppress} \\end{figure}  Evaluation of OH-suppression methods has been done by \\cite{Her94,JBB96,OB98}.  Equation~5 of \\citeauthor{JBB96} gives the gain ($G$) in performance of an OH-suppression filter for the broad-band S/N.\\footnote{The analysis of \\cite{JBB96} uses Poisson-noise statistics. They point out that even $G \\sim 1$ could provide a significant S/N gain, if most of the OH emission is removed from the background, because the continuum background is more stable and therefore systematic effects could be reduced.} For background limited conditions, this equation reduces to \\begin{equation} G = \\frac{\\tau}{\\sqrt{\\beta}} \\: , \\end{equation} where $\\tau$ is the fraction of object flux remaining, and $\\beta$ is the fraction of background remaining. Simple numerical calculations, for the broad-band S/N in the $J$ band, show that $G$ can be larger or about unity, when selectively combining TEI images at resolving powers of above 10000 using a finesse of 30. Note that this broad-band gain only applies to dark time because the TEI is imaging between the OH lines where the background flux is moon-phase dependent. Additionally, a suitable TEI instrument could image in the $J$ and $H$ bands simultaneously and provide a spectrum (low-resolution if the FP is not scanned).  As with a mono-order FP, the TEI could be used for targeted line imaging but, in addition, the TEI can simultaneously provide an OH-suppressed broad-band image using the non-targeted FP orders. The reduction of field of view in comparison with a standard FP mode is not always significant because, if the FP is not scanned, it is only the monochromatic field of view that images the targeted emission line."
    },
    "0004/astro-ph0004144_arXiv.txt": {
        "abstract": "Stars appearing in CCD images obtained over 224 nights during the course of 69 observing runs have been calibrated to the Johnson/Kron-Cousins {\\it BVRI\\/} photometric system defined by the equatorial standards of Landolt (1992, AJ, 104, 340).  More than 15,000 stars suitable for use as photometric standards have been identified, where ``suitable'' means that the star has been observed five or more times during photometric conditions and has a standard error of the mean magnitude less than 0.02\\mag\\ in at least two of the four bandpasses, and shows no significant evidence of intrinsic variability.  Many of these stars are in the same fields as Landolt's equatorial standards or Graham's (1982, \\pasp, 94, 244) southern E-region standards, but are considerably fainter.  This enhances the value of those fields for the calibration of photometry obtained with large telescopes.  Other standards have been defined in fields containing popular objects of astrophysical interest, such as star clusters and famous galaxies, extending Landolt-system calibrators to declinations far from the equator and to stars of sub-Solar chemical abundances.  I intend to continue to improve and enlarge this set of photometric standard stars as more observing runs are reduced.  The full current database of photometric indices is being made freely available via a site on the World-Wide Web, or by direct request to the author. Although the contents of the database will evolve in detail, at any given time it should represent the largest sample of precise {\\it BVRI\\/} broad-band photometric standards available anywhere. ",
        "introduction": "Accurate photometry with modern detectors on large telescopes is hampered by the scarcity of suitable photometric standard stars.  At present, the largest and most definitive collection of fundamental standard stars in the Johnson {\\it UBV\\/} Kron-Cousins {\\it RI\\/} broad-band photometric system is that of Landolt (1992), which consists of 526 stars that are mostly quite close to the celestial equator.  However, if one restricts oneself only to those stars that were observed a minimum of five times each (for instance), with standard errors of less than 0.02\\mag\\ in both $V$ and \\bmv\\ (say), then the total number of Landolt's ``good\" standards is reduced to 318.  Of these, perhaps of order 200 are appropriate for use with a 2.5-m telescope ($V \\gtsim 12$); maybe $\\sim\\,$130 can be used with a 4-m telescope ($V \\gtsim 13)$; and $\\ltsim\\,$40 are suitable for use with an 8-m telescope or a 2.4-m telescope in space ($V \\gtsim 14.5$). Graham (1982) has published a list of some 103 stars with {\\it UBVRI\\/} photometry in nine fields at declination --45\\deg; if one again considers only those stars having at least five observations and standard errors in $V$ and \\bmv\\ less than 0.02\\mag, the number of ``good'' Graham standards is reduced to some 61, of which only 11 are fainter than $V = 12$.  There are places on the sky where several standards can be imaged onto a CCD at the same time, but many of the Landolt and Graham stars are comparatively isolated, so that trying to observe a diverse sample of standards over a range of colors and airmasses with a CCD can be quite inefficient.   Furthermore, since these standards are primarily equatorial or far south, they never reach the zenith at many good terrestrial observing sites, and cannot cover the same range of azimuth as many scientifically interesting targets.   Observers trying to make the most of their large-telescope time are often reluctant to undertake large slews from the science target to one or more standard fields more than a few times per night.  Another drawback of Landolt's and Graham's standards is that few or no Population~II stars are included.  However, when one does observe fundamental standards like Landolt's or Graham's with a CCD, one usually gets for free the images of nearby stars, most of which are fainter than the official standards.  I expect that most CCD photometrists have toyed with the notion of combining these serendipitous observations of neighbor stars for the purpose of defining new, fainter standards, and this is what I have begun to do.  As of this date (Spring 2000) I have combined photometric data from a total of 69 observing runs consisting of 224 individual nights, of which 135 nights were completely clear, while on the remaining 89 nights observations were obtained through thin cloud during at least part of the night.  (CCD observations made through cloud can contribute to the precision of photometric indices provided that each image contains either fundamental standards or secondary standards that have also been observed under photometric conditions on numerous occasions.  Differential photometry relative to the brighter, well-established stars reduces the random errors of the mean magnitudes estimated for the fainter stars in the same field.)  These observations have been made by many different observers using ten telescopes at five sites (Kitt Peak National Observatory:  4-m, 2.1-m, 0.9-m; Cerro Tololo Inter-American Observatory: 4-m, 1.5-m, 0.9-m; La Palma:  Isaac Newton Telescope, Jacobus Kapteyn Telescope; Canada-France-Hawaii Telescope; Wyoming Infrared Observatory) over the period 1983--1999.  Many of these observations were made by me or my collaborators, but I have also obtained data for many of these observing runs through the excellent services of the Isaac Newton Group Archive and the Canadian Astronomy Data Centre.  In addition to photometric measurements for faint neighbors of Landolt and Graham standards, I have defined new standard sequences on the same photometric system in fields where the presence of an astrophysically interesting object (\\eg, a star cluster or a nearby galaxy) has led to the field's being observed several times during the observing runs at my disposal.  In the case of star clusters or dwarf galaxies very near the Milky Way, many of new standards will actually be members of the science target.  In the case of more distant galaxies and other types of extragalactic object, the new standard stars obviously belong to the Galactic foreground.  At the present moment, the available data permit the definition of more than 15,000 primary and secondary standards in 198 fields, where the following criteria are satisfied:  at least five independent observations under photometric conditions {\\it and\\/} standard errors of the mean magnitude smaller than 0.02\\mag\\ in {\\it at least two\\/} of {\\it BVRI\\/}, and no evidence for intrinsic variation in excess of 0.05\\mag, root-mean-square, based upon consideration of all available bandpasses.  The Johnson $U$ bandpass is not much observed with CCDs due to a variety of inconveniences, such as the low and highly wavelength-dependent relative quantum efficiency of many CCDs at these short wavelengths.  Although I do have and have tabulated some $U$-band observations for a number of these stars, I have not considered the availability of $U$ data to be relevant in making the decision whether a given star warrants being considered a photometric standard for my present purposes.  Lists of these standards are available to interested photometrists via the World-Wide Web or by direct communication with me.  The available data are:  digital finding charts (FITS-format images) on a common half-arcsecond-per-pixel scale, with $x$ increasing east and $y$ increasing north; ASCII files with astrometric positions---both absolute right ascensions and declinations, and relative (x,y) positions in the finding charts; and lists of photometry, consisting of mean apparent magnitudes in {\\it UBVRI\\/}, the standard errors of those quantities, the number of independent observations in each filter (the number of observations made on photometric occasions and the total number of observations, including those made through thin cloud, are both tabulated), and a measure of the intrinsic root-mean-square photometric variation.  All of these observations have been placed on the system of Landolt (1992) with an accuracy of order 0.001\\mag\\ in the mean.  It is my intention to keep the database up to date as additional observing runs become reduced, so the random photometric errors should go down and the number of individual standards and independent fields may be expected to grow with time.  However, at any given moment the instantaneous state of the database should represent the largest and most precise sample of {\\it BVRI\\/} broad-band photometric standards available anywhere. ",
        "conclusions": "At the moment, the total set of CCD observations considered here consists of some 1,092,401 individual magnitude measurements for 28,552 stars.  The instrumental magnitudes are based entirely on synthetic aperture photometry (bright, isolated stars) or profile-fitting photometry with aperture growth-curve corrections (fainter stars, or those with neighbors less than a few arcseconds away) obtained with CCDs and extracted by means of software written by me (Stetson 1987, 1990, 1993, 1994).  The instrumental magnitudes are transformed to the standard system using nightly equations that generally include linear and quadratic color terms as well as linear extinction terms.  Whenever practical, mean color coefficients are determined for all the nights of a given observing run with a particular instrumental setup.  However, extinction coefficients and photometric zero points are determined on a night-by-night basis, except for a few cases where the range of airmass spanned by the observations is too small for a meaningful extinction measurement; in such cases mean extinction coefficients for the site are imposed.  The equations for non-photometric nights do not model the effects of extinction.  Instead a separate photometric zero point for each frame is determined from measurements of at least two standard stars included within that frame; color terms determined from photometric nights during the same run and/or from individual frames containing standards that span a broad range of color are employed just as for the photometric nights. After the transformation equations for all nights have been determined, all the observations for each star are collected and transformed to the best possible magnitudes in $U$, $B$, $V$, $R$, and $I$ based on a simultaneous least-squares optimization involving all available data for the star.  The whole process is iterated.  Initially transformation equations are determined only from observations of the fundamental standard stars.  Then standard-system magnitudes can be derived for other stars contained within the same fields as the fundamental standards, and for stars in other program fields that were observed on photometric nights.  The subset of these stars that meet the criteria mentioned above, {\\it viz.\\/} at least five observations made under photometric conditions {\\it and\\/} standard errors smaller than 0.02\\mag\\ in at least two of the four {\\it BVRI\\/} filters, and no significant evidence of intrinsic variability, may now be considered to be additional standards.  Improved transformations are then re-determined using this enlarged set of standard stars.  Starting with this second iteration, the newly defined ``standards'' allow the inclusion of non-photometric observations for the former program fields, increasing the precision (but not the accuracy) of their derived photometric magnitudes.  Another iteration of this process is undertaken every time a new observing run is added to the database, resulting in some new standard stars, more precise mean magnitudes for the previously existing standard stars, and occasionally the loss of a putative standard if the new observations suggest intrinsic variability.  The fundamental basis for the photometric system employed here is that of Landolt (1992) consisting of (mainly) equatorial standards observed in {\\it UBVRI\\/} with photomultipliers.  I have augmented this primary set of reference stars with the data in Landolt (1973; photomultiplier-based {\\it UBV\\/} observations that are apparently independent of those of Landolt 1992, unlike the observations in Landolt 1983, which appear to be a subset of those included in the 1992 catalog); Landolt (1983; a very few stars that were not republished in the 1992 paper); Graham (1982; photomultiplier {\\it UBVRI\\/} photometry of stars in the E regions at declination --45\\deg) and Graham (1981; photomultiplier {\\it UBVRI\\/} photometry of a standard sequence near the spiral galaxy \\ngc{300}); W.~E.~Harris (unpublished; photomultiplier {\\it UBV\\/} photometry of stars in the equatorial open cluster M11 = \\ngc{6705}); and L.~Davis (unpublished; CCD {\\it UBVRI\\/} photometry of stars in the Kitt Peak consortium fields in the star clusters \\ngc{4147}, 2419, 6341 = M92, 7006, and 7790; Christian \\etal\\ 1985).  All of these data must be assumed {\\it a priori\\/} to be on effectively the same photometric system as Landolt (1992) --- within the errors --- with two exceptions.  (1)~There are enough stars in common between Landolt (1973) and Landolt (1992) that a direct comparison of the two systems can be undertaken in $U$, $B$, and $V$.  In fact, I base this comparison on {\\it only\\/} those stars that are common among Landolt (1973), Landolt (1992), {\\it and\\/} the set of Landolt stars included among my observations.  This restriction is made just in case any difference between Landolt (1973) and Landolt (1992) might depend in some systematic way on the stars' magnitudes, colors, right ascensions, or other properties; if such should be the case, obviously we want to know the value for any (1992) {\\it minus\\/} (1973) difference that would be appropriate specifically for the sorts of stars considered here.  When the comparison is made, I find that the Landolt (1992) magnitudes differ from those of Landolt (1973) by --0.0034$\\pm$0.0011\\mag\\ (standard error of the mean difference) in $V$, --0.0026$\\pm$0.0013\\mag\\ in $B$, and +0.0022$\\pm$0.0023\\mag\\ in $U$, based upon 81 stars common to all three data sets.  Landolt's 1973 {\\it UBV\\/} magnitudes have been adjusted by these offsets and combined with his 1992 data.  (2) According to Davis (private communication) her data for \\ngc{7790} were taken under dubious photometric conditions. The way in which these data are included will be described below.  The assumption that the remaining Graham, Harris, and Davis data are on essentially the same system as Landolt (1992), at least within the standard errors of the available data sets, can be tested {\\it a posteriori}, as I will now describe.  Specifically, after each iteration I compare my photometry with Landolt's for those stars where (a)~Landolt has at least four observations and a standard error of the mean magnitude less than 0.03\\mag\\ in a given filter, {\\it and\\/} (b)~I have at least four observations and a standard error of the mean magnitude less than 0.03\\mag\\ in the same filter, {\\it and\\/} (c)~the star shows no evidence in my data for intrinsic variability greater than 0.05\\mag, root-mean-square, in all filters considered together.  (Selection criteria more restrictive than these resulted in a sample size too small to be very meaningful.) Any net difference remaining between my weighted average results and the combined results of Landolt (1992) and (1973) for stars meeting these criteria is evaluated and added to all my magnitudes, forcing my photometric system to be identically equal, in the mean, to that of Landolt with a high level of accuracy.  After the most recent iteration these corrections were all less than 0.0005\\mag\\ in $B$, $V$, $R$, and $I$, with standard errors of the correction better than 0.0013\\mag\\ in each case, based on 144 stars in $B$, 144 stars in $V$, 30 stars in $R$, and 79 stars in $I$; in $U$ the correction  was $0.0009\\,\\pm\\,0.0084$\\mag\\ based on only 3 stars.  Figures 1--4 show the differences between my photometry and Landolt's for these stars versus magnitude and color.  The observed root-mean-square magnitude residuals between Landolt's results and mine exceed the quadrature sum of both our estimated standard errors by less than 10\\%.  This leaves very little room for systematic errors due to neglected high-order transformation terms occasioned by, for instance, filter-bandpass mismatch.  In fact, to the naked eye, some seemingly systematic differences between my photometry and Landolt's may be seen in Figs.~1--4.  For instance, in Fig.~1 it seems that for $10\\,\\ltsim\\,B\\,\\ltsim\\,11.5$ my $B$-band magnitudes are fainter than Landolt's, while for $11.5\\,\\ltsim\\,B\\,\\ltsim\\,12.0$ my $B$ magnitudes are brighter.  Similarly, my $B$ magnitudes for the bluest stars ($\\bmv < 0.00$) may be slightly fainter, on average, than Landolt's.  If such behavior is real, it would imply a subtle systematic nonlinearity either Landolt's photometry or mine.  In either case, the nonlinearity would have to be a {\\it collective\\/} property of many devices, since Landolt used a number of different photomultipliers and cold boxes, while my results have certainly been based on a large number of different CCD and filter combinations.  In each case, all data for each detector were placed on a common photometric system using appropriate transformation models.  It is noteworthy that the apparently systematic differences in the $B$-band photometry are {\\it not\\/} duplicated in either the $V$ band or the $I$ band, while the plots for each of these other bandpasses have idiosyncracies likewise not reflected in the other filters.  I cannot come up with a plausible physical mechanism that would produce this variety of effects systematically across an ensemble of detectors of either technology.  In the absence of more definitive data, it seems most likely that these seeming deviations are the result of small-number statistics and the propensity of the human eye for finding patterns even in random data.  After I have thus forced my mean results onto Landolt's system, the net differences between Davis's unpublished magnitudes for \\ngc{7790} stars and my results for the same stars are determined and applied to her mean magnitudes; these corrections, which are of order a few hundredths of a magnitude, place Davis's \\ngc{7790} data on the same system as mine, in the mean, which is---via the previous step---the same as Landolt's. Finally, a weighted average of my derived photometric magnitudes and the previous ones is determined for all stars in common.  To the extent that the Graham results, the Harris results, and the rest of Davis's results may {\\it not\\/} be inherently on the Landolt system, these weighted mean magnitudes will not be on exactly the Landolt system either.  However, they will be much {\\it closer\\/} to the Landolt system inasmuch as my observations generally greatly outnumber the previous ones.  In fact, these other data sets turn out to be fairly close to the Landolt system in comparison to their standard errors, as Table~1 shows.  Here I have tabulated the robust mean magnitude differences and standard errors of the mean differences  for {\\it all\\/} stars common to my and the previous data sets, without regard to the number of observations, the standard errors, or any evidence of variablility.  Only two elements of the table reveal systematic differences as large as 0.01\\mag, and it is to be expected that the overall set of my observations combined with the previous ones will differ from the mean Landolt system by amounts much less than this.  The line of mean differences for the comparison to Landolt's photometry exemplifies the ultimate uncertainty of placing my photometry on his system:  it represents the distinction between (a)~comparing only those stars that he and I both measured ``well,'' in some sense, which has zero net difference after the procedure of the previous paragraph, and (b)~comparing {\\it all\\/} stars common to the two data sets, which yields the differences in Table~1."
    },
    "0004/astro-ph0004329_arXiv.txt": {
        "abstract": "We present the ultraviolet spectrum of the SW~Sex star and nova-like variable DW~UMa in an optical low state, as observed with the {\\em Space Telescope Imaging Spectrograph} on board the {\\em Hubble Space Telescope} (HST). The data are well described by a synthetic white dwarf (WD) spectrum with $T_{eff} = 46,000 \\pm 1000$~K, $\\log{g} =  7.60 \\pm 0.15$, $v\\sin{i} = 370 \\pm 100$~km~s$^{-1}$ and $Z/Z_{\\odot} = 0.47 \\pm 0.15$. For this combination of $T_{eff}$ and $\\log{g}$, WD models predict $M_{WD} = 0.48 \\pm 0.06$~M$_{\\odot}$ and $R_{WD} = (1.27 \\pm 0.18) \\times 10^9$~cm. Combining the radius estimate with the normalization of the spectral fit, we obtain a distance estimate of $d = 830 \\pm 150$~pc.  During our observations, DW~UMa was approximately 3 magnitudes fainter in V than in the high state. A comparison of our low-state HST spectrum to a high-state spectrum obtained with the {\\em International Ultraviolet Explorer} shows that the former is much bluer and has a higher continuum level shortward of 1450~\\AA. Since DW~UMa is an eclipsing system, this suggests that an optically thick accretion disk rim blocks our view of the WD primary in the high state. If self-occulting accretion disks are common among the SW~Sex stars, we can account for (i) the preference for high-inclination systems within the class and (ii) their V-shaped continuum eclipses. Moreover, even though the emission lines produced by a self-obscured disk are generally still double-peaked, they are weaker and narrower than those produced by an unobscured disk. This may allow a secondary line emission mechanism to dominate and produce the single-peaked, optical lines that are a distinguishing characteristic of the SW~Sex~stars. ",
        "introduction": "\\label{introduction}  The SW~Sex stars are a sub-class of nova-like cataclysmic variables (NLCVs). Like all NLCVs, they are semi-detached, close binary systems, in which a late-type main sequence secondary star loses mass to a WD primary via Roche-lobe overflow onto an optically thick accretion disk. However, they also exhibit a number of peculiar characteristics that are not shared by ``ordinary'' NLCVs (e.g. Thorstensen et al. 1991; Dhillon, Jones \\& Marsh 1994; Rutten, van Paradijs \\& Tinbergen 1992): (i) 7 of the 10 current members of the class are eclipsing systems (for an up-to-date list with references, see Hoard, Thorstensen \\& Szkody 2000); (ii) their continuum eclipses appear to be more V-shaped (as opposed to U-shaped) than those of other NLCVs; (iii) their optical emission lines are single-peaked (instead of double-peaked, as expected for lines formed in a rotating accretion disk); (iv) their low-excitation optical emission lines are only weakly eclipsed at primary minimum, but appear to undergo absorption events at the opposite orbital phase; (v) the radial velocity curves derived from their optical emission lines exhibit significant phase shifts relative to photometric minimum. The V-shaped eclipses noted in (ii) are often interpreted as evidence for flatter-than-expected radial temperature distributions and non-stationary accretion in SW~Sex star disks (e.g. Rutten et al. 1992; Horne 1999).  At least four different models have been proposed to explain the SW~Sex stars' ``mysterious, but consistent behavior'' (Thorstensen et al. 1991). The physical mechanisms invoked in these models are: (i) a WD-anchored magnetic field that dominates the geometry of the accretion flow (Williams 1989; Dhillon, Marsh \\& Jones 1991); (ii) an accretion disk wind that alters the radiative transfer of the optical emission lines and the energy budget of the accretion disk (Honeycutt, Schlegel \\& Kaitchuck 1986; Murray \\& Chiang 1996; Dhillon \\& Rutten 1995); (iii) an accretion stream from the secondary star that overflows the outer disk edge (Hellier \\& Robinson 1994), possibly in combination with either non-axisymmetric, vertical disk structure (Hoard 1998; Hoard et al. 2000) or a flared, but not self-occulting disk (Hellier 1998); (iv) a disk-anchored magnetic field that acts as a propeller and ejects a significant amount of material out of the system (Horne 1999).  No real consensus has emerged to date regarding which, if any, of these models is correct. In our view, this is partly because the models have been tested almost exclusively against the known, optical characteristics of the SW~Sex stars -- which they were designed to reproduce in the first place. A promising, complementary approach is to compare the models to observation in a qualitatively different observational regime. With this idea in mind, we have obtained high time- and medium spectral resolution ultraviolet (UV) spectroscopy of the SW~Sex star DW~UMa with the {\\em Space Telescope Imaging Spectrograph} (STIS) on board the HST. Much to our initial disappointment, our observations found DW~UMa in a deep (V $\\simeq$ 17.6) low state (in the high state, V $\\simeq$ 14.5). However, a comparison of our HST data to {\\em International Ultraviolet Explorer} (IUE) observations of DW~UMa in its ``normal'', high state reveals a surprising result: DW~UMa's far-UV continuum (shortward of 1450~\\AA) is {\\em higher} in the low state than in the high state. As we will discuss, the simplest interpretation of this result is that the WD in DW~UMa is geometrically occulted by the rim of the accretion disk in the high state. ",
        "conclusions": "For comparison with our low state HST observations, we also show in Figure~1 a high state, out-of-eclipse spectrum of DW~UMa obtained approximately 13 years earlier with IUE (SWP27097; Szkody 1987). The IUE spectrum is much redder than the HST one, and, unlike the latter, contains strong emission lines. Most importantly, however, the UV continuum shortward of 1450~\\AA~is actually higher in the low state HST spectrum than in the high state IUE one. Thus if we are correct in attributing the low-state UV light to the WD primary, our view of this system component must be blocked in the high state. \\footnote{The only alternative we can think of is that the WD is much cooler in the high state than in the low state. We discard this as physically implausible, since the temperature of the WD is expected to increase with increasing accretion rate (Sion 1999).} We have also inspected the five other (2 SWP and 3 LWP) UV spectra of DW UMa that were taken by IUE. All of them share the flat UV continuum slope of the high-state IUE spectrum shown in Figure~1, and none exhibits a significantly higher UV flux level. Since the combined IUE observations cover the entire orbital cycle (Szkody~1987), the WD in DW~UMa appears to be obscured at most or all binary phases during the high state.  The simplest interpretation of these findings is that, in the high state, our view of the WD is blocked by the rim of a flaring, optically thick accretion disk. \\footnote{A flat, geometrically thin, optically thick disk can occult no more than half of the projected area of the WD. The presence of such a disk can therefore reduce the high state WD flux by at most a factor of two. The maximum discrepancy between low-state and high-state continua is larger than this. Moreover, the hot, inner regions of a geometrically thin disk would themselves contribute substantially to the far-UV flux, offsetting the effect of the WD occultation.} For a disk with semi-opening angle $\\alpha$ (measured at its outer edge), this requires that $i+\\alpha \\geq 90^o$. In a deeply eclipsing system like DW~UMa, this condition may well be met (see below). Our results can then be explained as follows. In the high state, the inner disk, the WD and the front of the disk are all hidden from view by the disk's front rim. The only elements that remain visible are the cool, outer regions on the back of the disk, as well as the occulting disk rim itself. The observed UV flux in the high state is therefore much lower (and the UV spectrum much redder) than if the hot, inner disk, and the WD at its center, were visible. In the low state, the mass transfer rate from the secondary star is severely reduced or completely shut off. The accretion disk then becomes optically thin (or disappears completely), and the WD emerges. In order to be consistent with the observations, this picture requires that the hot, compact WD should match or exceed the far-UV -- but not the optical -- output of the much cooler, visible regions of the high-state accretion disk. This is plausible, given the differences in effective temperature and projected area associated with these system components. For example, for two blackbodies at temperatures $T_1=45,000$~K and $T_2=15,000$~K, the ratio $B_\\lambda(T_1)/B_\\lambda(T_2)$ is approximately 92 at $\\lambda=1450$~\\AA~but only about 6 at $\\lambda=5550$~\\AA. So if the visible portions of the high state disk have a characteristic temperature near $T_2$ and present a projected area around 100 times that of the WD, the UV flux in the low state will be similar to that in the high state, whereas the optical flux will be 3 magnitudes (15$\\times$) lower, as is observed.  Thus the simple hypothesis that DW~UMa in the high state harbors a self-occulting accretion disk can account for: (i) the red UV colors in the high state; (ii) the blue UV colors in the low state; (iii) the approximate equality of the high- and low-state UV flux levels; (iv) the lack of such an equality in the optical waveband. Given this success, it is interesting to explore the impact of  the self-occultation scenario more generally, i.e. beyond its specific application to our HST observations of DW~UMa. First, it is obvious that the preference for eclipsing systems in the SW~Sex class is explained quite naturally if at least some aspects of the SW~Sex syndrome are related to the presence of a self-occulting accretion disk. Second, Smak (1994) has pointed out that a flared accretion disk will produce increasingly V-shaped continuum eclipse light curves as it is viewed increasingly edge-on. This effect is due to the increasing relative contribution of the facing disk rim to the total light. But an additional effect is at work in a {\\em self-occulting} disk: since the hot disk center is permanently hidden from view, even the top of the disk produces a more V-shaped eclipse in this case (Rutten 1998). Third, the line profiles produced by a self-obscured, flared disk will be significantly different from those produced by a fully visible disk. This may have a bearing on the peculiar optical emission line behaviour characteristics of the SW~Sex stars.  To illustrate and explore the last two points, we show in Figure~2 some synthetic eclipse light curves and line profiles. In calculating these models, the following parameters were adopted: (i) primary mass: $M_{WD} = 0.5M_{\\odot}$; (ii) mass ratio and inclination: $q=M_{2}/M_{WD} = 0.4$ and $i=82^o$; (iii) secondary radius: $R_2=0.27 R_{\\odot}$; (iv) accretion disk radius: $R_{disk} = 0.8 L_1$; (v) radial surface brightness distribution: $S \\propto R^{-3/2}$. The combination of $q$ and $i$ in (ii) gives an eclipse width roughly consistent with that of DW~UMa (Dhillon et al. 1994). $R_2$ is based on Smith \\& Dhillon's (1998) mass-radius relation for CV secondaries, taking $M_2 = q M_{WD}$. $L_1$ is the distance from the WD to the inner Lagrangian point, and $R_{disk}$ is based on DW~UMa's total eclipse width in the high state (Dhillon et al. 1994). The radial surface brightness distribution in (v) is applied to both lines and continuum. To explore the effect of the shape of the disk on line profiles and continuum light curves, we have calculated two types of models. In one, the disk is described as flat, but with a raised rim at its outer edge; in the other, the disk is taken to have a constant aspect ratio H/R. These two shapes are the limiting cases of concave disks. In both cases, we have set $\\alpha=\\arctan{[H(R_{disk})/R_{disk}]} = 10^o$. We have also considered two limiting cases for emission from the facing rim: for each disk shape, we present models with and without rim emission. In the former, the surface brightness of the rim is the same as at the top edge of the disk. All model line profiles were calculated in both optically thin and thick limits. For the latter, we generalized slightly the treatment given in Horne \\& Marsh 1986 to account for the flare of the disk.  The third row of panels in Figure~2 shows that self-obscured disks do indeed produce much more V-shaped continuum eclipses than a flat disk, regardless of the precise disk shape. As expected, the effect is particularly pronounced when the rim contributes significantly to the emission. Thus self-occultation can easily account for the V-shaped continuum eclipses of the SW~Sex stars and for the correspondingly flat T(R) distributions obtained in eclipse mapping studies of SW~Sex stars (Rutten et al. 1992). The bottom row of panels in Figure~2 shows that the line profiles produced by self-obscured disks are weaker and narrower than those produced by a flat disk. This is because the bright, quickly rotating inner disk regions are no longer visible in a self-occulted disk. However, the line profiles are generally still double peaked. Thus self-occultation alone can probably not account for the single-peaked profiles of the optical lines in the SW~Sex stars. Self-occultation may nevertheless be key to the production of single-peaked line profiles in some of these systems. In particular, the weakness of lines produced by a self-obscured disk may allow a secondary, normally negligible, line emission mechanism to dominate. This mechanism could actually be at work in all NLCVs, but may only be observable in the SW~Sex stars.  \\vspace*{0.1cm}  We are grateful to Joe Patterson, Jonathan Kemp, John Thorstensen and Cindy Taylor for their contributions to this project and to Leandro Althaus for providing his WD mass-radius relations in electronic form. Support for this work was provided by NASA through grant number GO-7362  from the Space Telescope Science Institute (STScI), which is operated by AURA, Inc.,under NASA contract NAS5-26555. Additional support for CK was provided by NASA through Hubble Fellowship grant HF01109.0198A, also awarded by STScI."
    },
    "0004/astro-ph0004023_arXiv.txt": {
        "abstract": "I compare various semi-analytic models for the bias of dark matter halos with halo clustering properties observed in recent numerical simulations. The best fitting model is one based on the collapse of ellipsoidal perturbations proposed by Sheth, Mo \\& Tormen (1999), which fits the halo correlation length to better than 8 per cent accuracy. Using this model, I confirm that the correlation length of clusters of a given separation depends primarily on the shape and amplitude of mass fluctuations in the universe, and is almost independent of other cosmological parameters. Current observational uncertainties make it difficult to draw robust conclusions, but for illustrative purposes I discuss the constraints on the mass power spectrum which are implied by recent analyses of the APM cluster sample. I also discuss the prospects for improving these constraints using future surveys such as the Sloan Digital Sky Survey. Finally, I show how these constraints can be combined with observations of the cluster number abundance to place strong limits on the matter density of the universe. ",
        "introduction": "Recent years have seen dramatic increases in the size of the largest cosmological N-body simulations, with giga-particle runs now able to track the non-linear gravitational evolution of a significant fraction of the observable universe (Colberg et al.\\markcite{C98} 1998). Correspondingly, it is becoming possible to make more and more detailed predictions for the non-linear properties of dark matter in any cosmological model, with convergent results being obtained for such features as the profiles and correlation functions of dark matter halos (Moore et al.\\markcite{M99} 1999). In spite of these advances, the one great challenge remaining for the simulators is to realistically track the evolution of cosmic gas through its complex cycles of star-formation and feedback.  Given the current state of the art for cosmological simulations, rich clusters of galaxies provide one of the most powerful tools with which to make detailed tests of cosmology. Firstly, cluster potential wells are so deep that their formation and evolution is largely unaffected by the uncertain history of the cosmic gas, making dissipationless N-body simulations an ideal place in which to study their properties. Second, clusters correspond to rare peaks in the primordial density field, so their statistics are extremely sensitive to detailed features of the cosmology.  In this paper, I will show that a simple model proposed by Sheth, Mo \\& Tormen\\markcite{SMT} (1999) can give an extremely good fit to the correlation function of massive halos observed in recent state of the art numerical simulations. Using this model I will compare the predictions of a range of cosmological models with observations of the cluster correlation length. In particular, I will show that correlation length observations place strong constraints on the power spectrum of mass fluctuations in the universe, constraints that are almost independent of other cosmological parameters. In section~\\ref{sec-model} I describe the Sheth, Mo \\& Tormen model for the cluster correlation length, together with a similar semi-analytic model due to Mo \\& White (1996), and other bias models discussed in the literature. In section~\\ref{sec-simulations} I compare the predictions of these models with recent numerical simulations, and in section~\\ref{sec-data} I show the constraints which can be derived when the models are compared with observations. In section~\\ref{sec-conclusions} I discuss the relationship of this work to previous studies, and draw my conclusions. ",
        "conclusions": "\\label{sec-conclusions} I have shown that the halo correlation length measured in recent numerical simulations can be well fit (to better than 8 per cent accuracy) by a semi-analytic model based on the collapse of ellipsoidal perturbations due to Sheth, Mo \\& Tormen\\markcite{SMT} (1999). A similar analytic model due to Mo \\& White\\markcite{MW} (1996) on the other hand gives a poorer fit to the simulations (roughly 25 per cent accuracy). Applying the SMT result to present data, I have shown that correlation length observations place strong, almost cosmology independent constraints on the shape and amplitude of the matter power spectrum in the universe. By far the greatest source of uncertainty is systematic discrepancies between current datasets, but if the cluster correlation length is at least as high as is implied by the APM survey, then an interesting region of power spectrum parameter space (everything to the right of the ``Croft'' region in Fig.~\\ref{fig-square}) can be excluded. Future surveys, including the Sloan Digital Sky Survey, should greatly reduce the observational uncertainties, with 2$\\sigma$ confidence limits from SDSS being as tight as the 1$\\sigma$ limits shown for the C97 data in Fig.~\\ref{fig-square}. One particularly interesting result is that the analysis is almost unaffected when a significant scatter (up to 50 per cent) is introduced between true cluster mass and the richness property by which clusters are ranked. Such a scatter is inevitable in any cluster survey, so it is extremely useful to know that even an effect this large does not influence the results. Finally, I have shown how the correlation length can be combined with other cluster observations to place limits on the matter density of the universe.  A number of previous studies have discussed constraints from the cluster correlation length. Bahcall \\& Cen\\markcite{BC} (1992) and Croft \\& Efstathiou\\markcite{CE} (1994) carried out numerical simulations showing that cluster correlation length was a strong discriminator amoung models, and obtained results which are consistent with those discussed in this paper. Mo, Jing \\& White\\markcite{MJW} (1996) and later Robinson, Gawiser \\& Silk (1998) made use of the MW formalism to compare models and observations for a wider range of parameters, and owing to the inaccuracy of the MW model, their results are not consistent with those found here (the inconsistency of the MW model is actually larger than is apparent in MJW due to an error in Fig.~8 of that paper). The work described here improves on previous studies by utilizing a formula which accurately fits the most recent numerical simulations, and which can be used to compute predictions for a large range of model parameters very quickly.  Lastly, it should be noted that the analysis discussed here assumes that the primordial fluctuations in the universe are Gaussian. Non-gaussianity also influences the value of the cluster correlation length, and a number of authors, including Robinson, Gawiser \\& Silk (1998) and Koyama, Soda \\& Taruya \\markcite{KST} (1999) have attempted to exploit this fact to use cluster observations to place constraints on primordial non-gaussianity. Both these works have modeled the correlation function for non-gaussian models using an extension of the MW formalism, and therefore their conclusions should be modified in the light the results discussed here, which show that MW does not accurately predict the halo correlation function even for Gaussian models.  To summarize, cluster correlation length observations place strong cosmology independent constraints on the matter power spectrum in the universe, constraints which future surveys such as SDSS will allow us to fully exploit."
    },
    "0004/astro-ph0004215_arXiv.txt": {
        "abstract": "For all 209 RXTE observations of the 7-Crab X-ray nova XTE~J1550--564, we analyzed the X-ray power spectra, phase lags, and coherence functions.  These observations cover the entire 250 day outburst cycle of the source and are the most complete and arguably the richest data set for any black hole X-ray nova.  We find that there are three fundamental types of QPO behavior--one more than reported by Wijnands, Homan, \\& van der Klis (1999, \\apj, 526, 33).  The new type occurred during the first half of the outburst.  These three types of QPO behavior can be grouped according to the relative contributions of the disk and power-law components to the total flux. ",
        "introduction": "XTE~J1550--564 is an X-ray nova and black hole candidate discovered with the All Sky Monitor (ASM; Levine et al.~1996) onboard the {\\it Rossi X-ray Timing Explorer} (RXTE) in 1998 September (Smith et al.~1998).  The outburst lasted until 1999 May; its duration was approximately 250 days, during which time RXTE performed an almost daily series of pointed observations.  A summary of these observations can be found in Sobczak et al.~(2000b).  Extensive spectral and timing studies of this source have been performed using these data.  The source has been observed in the very high, intermediate, and high/soft outburst states of black hole X-ray novae, with an X-ray spectrum that can be well modeled by multicolor blackbody disk and power-law components (Sobczak et al.~1999,2000b; Homan et al.~2000).  The power spectra of XTE~J1550--564 exhibit quasiperiodic X-ray oscillations (QPOs) at both high frequencies ($\\sim100$--285~Hz) and low frequencies (0.08--18~Hz) (Remillard et al.~1999; Sobczak et al.~1999; Cui et al.~1999; Homan et al.~2000). Herein we are interested only in the latter `low-frequency' QPOs.  These QPOs can have large amplitudes, with peak to trough ratios as high as 1.5 (Remillard et al.~1999), which indicates that they are a fundamental aspect of the accretion process.  A sample lightcurve of XTE~J1550--564 can be seen in Figure~1 of Sobczak et al.~(2000a), which explored the correlations between the QPO frequency and amplitude and the X-ray spectral parameters.  They reported that both the power-law and the disk components are linked to the QPO phenomenon: QPOs are observed only when the power-law component contributes more that 20\\% of the 2--20~keV flux, and the QPO frequency generally increases as the disk flux increases.  Our focus in this paper is on the properties of the phase lags associated with the low-frequency QPOs.  The Fourier time delay, or lag, can be computed to measure the phase lags between soft and hard X-ray energy bands.  Recent studies have shown that the phase lags associated with low-frequency QPOs and their harmonics are quite complex.  In one observation of GRS~1915$+$105, Cui (1999) found that the lower energy (soft) photons lagged behind the higher energy (hard) photons (``soft lag'') for the fundamental QPO and it's 2nd harmonic; however, the higher energy photons lagged behind the lower energy photons (``hard lag'') for the 1st and 3rd harmonics. Wijnands, Homan, \\& van der Klis (1999) reported even more complicated behavior from 14 observations of XTE~J1550--564 near the end of the outburst.  However, Wijnands et al.~were able to show that the complicated phase lag behavior observed in these 14 observations actually consists of variations on only two fundamental types.  Phase lags for 13 observations during the initial rise of XTE~J1550--564 are discussed by Cui, Zhang, \\& Chen~(2000).  In this paper we analyze the phase lags for all 209 RXTE observations of XTE~J1550--564.  These observations cover the entire 250~day outburst cycle of the source and are the most complete and arguably the richest data set for any black hole X-ray nova.  We find that there are three fundamental types of phase lag behavior -- one more than reported by Wijnands et al.~(1999).  The new type occurred during the first half of the outburst.  We discuss correlations between the phase lags and the X-ray spectral parameters, QPO frequencies, and amplitudes. ",
        "conclusions": "In conclusion, we summarize the observed QPO properties for XTE~J1550--564.  First, whenever QPOs are observed the power-law component contributes more than 20\\% of the 2--20~keV flux (Sobczak et al.~2000a).  During the very high state of the outburst, when the QPOs have the largest amplitude, the disk component contributes as little as 3\\% of 2--20~keV flux (Sobczak et al.~1999,2000b).  However, the QPO frequency generally increases as the disk flux increases (Fig.~7a).  These results demonstrate that both the disk and the power-law components are linked to the QPO phenomenon. Secondly, there are three fundamental types of QPO behavior which can be grouped according to the relative contributions of the disk and power-law components to the total flux (Fig.~6).  The QPOs display both hard and soft lags, as well as near unity coherence in two of the QPO types."
    },
    "0004/astro-ph0004165_arXiv.txt": {
        "abstract": "We examine   a new paradigm to account for the exponential decay seen in the light curves of some of the bright X-ray novae. These systems show an exponential decay in soft X-rays with an $e-$folding time constant of $\\sim30$ d. We investigate a scenario in which evaporation of  matter into a corona is the dominant mass removal mechanism from the accretion disk. We utilize the  thermal evaporative instability discovered by Shaviv and Wehrse. First we  parametrize  local mass loss rates from the disk (fitted to vertical structure computations of the optically thin structure using the photoionization code CLOUDY), and then we utilize the scalings  in our numerical time dependent model for the decay. Both the $\\sim30$ d $e-$folding time scale for the decay and the secondary maximum with its rapid rise time $\\sim1-3$ d which is seen in the X-ray nova  light curves can be produced by adjusting the strength of the evaporation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004353_arXiv.txt": {
        "abstract": "The contribution of the Source Cosmic Rays (SCRs), confined in Supernova Remnants, to the diffuse high energy \\gr emission above 1 GeV from the Galactic disk is studied. \\grs produced by the SCRs have a much harder spectrum compared with those generated by the Galactic Cosmic Rays which occupy a much larger residence volume uniformly. SCRs contribute less than 10\\% at GeV energies and become dominant at \\gr energies above $100$~GeV. The contributions from $\\pi^0$-decay and Inverse Compton \\grs have comparable magnitude and spectral shape, whereas the Bremsstrahlung component is negligible. At TeV energies the contribution from SCRs increases the expected diffuse \\gr flux almost by an order of magnitude. It is shown that for the inner Galaxy the discrepancy between the observed diffuse intensity and previous model predictions at energies above a few GeV can be attributed to the SCR contribution. ",
        "introduction": "Observations of the diffuse Galactic \\gr emission give information about the Galactic Cosmic Rays (GCRs), the interstellar gas and diffuse photon fields, and about the interactions between them. The observational results obtained with the Energetic Gamma Ray Experiment Telescope (EGRET) on the Compton Gamma Ray Observatory can be described fairly well by a suitable model for the diffuse interstellar gas, Cosmic Ray (CR) and photon distributions (e.g. Hunter et al. 1997a; Hunter et al. 1997b).  However, above 1~GeV the observed average diffuse \\gr intensity in the inner Galaxy, $300^{\\circ}<l<60^{\\circ}$, $|b|\\leq 10^{\\circ}$, exceeds the model prediction significantly. There are at least two possible explanations for this discrepancy (e.g. Hunter et al. 1997b; Weekes et al. 1997). The high-energy $\\gamma$-ray excess may indicate that the GCR spectrum observed in the local neighborhood is not representative of the diffuse CR population in the Galactic disk; a harder average diffuse proton spectrum is required to explain the $\\gamma$-ray excess if it is due to $\\pi^0$-decay. An unresolved distribution of CR sources is the other possibility.  The physical picture which we consider in this paper corresponds to the second possibility. The idea that CRs, after leaving their sources, could in principle produce \\grs in ambient dense clouds with a harder spectrum than those produced by the average GCRs, was proposed by Aharonian and Atoyan (1996). Another proposed possibility, also invoking transport effects, is the local hardening of the CR energy spectrum in the direction perpendicular to the Galactic disk above strong CR sources, especially above the inner region of the Galaxy, due to faster CR convection in a faster Galactic Wind (V\\\"olk 1999). This contributes a principally observable harder than average \\gr component in such regions. In contrast, we consider here the accelerating particles inside their sources, where they are much more strongly scattered than in the ISM, neglecting the contributions invoked by Aharonian and Atoyan, and by V\\\"olk (see also below).  We assume that SNRs are the dominant sources of the Galactic CRs. On this premise we find that CRs, accelerated and confined in SNRs, give an important contribution to the high-energy $\\gamma$-ray emission from the Galactic disk. Since the CR energy spectrum inside SNRs is much harder than on average in the Galaxy --- the average spectrum being softened by rigidity-dependent escape from the Galaxy in the diffusion region above the disk --- the relative SNR contribution increases with energy and becomes in fact dominant at \\gr energies $\\epsilon_{\\gamma} \\gsim 100$~GeV. It may substantially increase the diffuse TeV $\\gamma$-ray emission from the Galactic disk so as to constitute a significant and spatially variable observational background which must also be taken into account in the search for spatially extended Galactic CR sources in this energy region. A physically analogous problem is the contribution of CR electrons in SNRs to the {\\it radio synchrotron spectrum} of normal galaxies, without an Active Galactic Nucleus. It has been recently discussed by Lisenfeld \\& V\\\"olk (1999).  In this paper we shall investigate pion-decay $\\gamma$-ray emission from CR nuclei as well as Inverse Compton (IC) radiation and Bremsstrahlung due to CR electrons. It will be shown that the average IC $\\gamma$-ray background from SNRs is comparable in magnitude and spectral form to the pion-decay background at high energies, whereas the corresponding Bremsstrahlung component is negligible. ",
        "conclusions": "We note that of order ten SNRs of age younger than $10^5$~yr can on average lie within a 1 degree field of view of a detector directed towards the Galactic Center. Therefore a moderate fluctuation of the measured \\gr intensity is expected due to variations of the actual number of SNRs within such a detector's field of view. At the same time, for directions perpendicular to the Galactic plane, the chance to observe the contribution of SCR \\gr emission is quite negligible. A question is then how one might best study the nonuniformities of this background experimentally. Clearly this is an investigation of its own. Therefore we would like to restrict ourselves to a few comments here. Due to the spectral form of the background its graininess is most pronounced at high \\gr energies. This is even more true due to the fact that for individual sources with a very low magnetic field the IC emission could be much stronger than the \\gr emission due to hadronic interactions; in addition the IC emission has a harder spectrum. This suggests the use of imaging atmospheric Cherenkov telescopes. Their resolution in angle and energy is as good or better than that of other ground-based detectors. However, the study of extended sources is not an easy task with imaging telescopes which have a very limited field of view, even employing the stereoscopic method. For low brightness extended sources a satellite instrument like the future GLAST detector is well suited since it does not have to deal with the charged CR background due to its use of an anticoincidence shield. On the other hand, the statistics achievable with a small area space detector gets very low above a few tens of GeV. Thus one should probably attempt such a study with both types of instruments due to their complementary properties.  A different question concerns the limitations of our approach due to its concentration on SNRs as the sources of the GCRs. In fact the considerations in this paper can be applied to any alternative class of dominant GCR sources. The most important aspect, which leads to the dominance of SCRs in high-energy \\gr production, is that the GCR sources should generate SCRs with a spectrum that is significantly harder than the GCR spectrum. Eq. (15) is valid for an arbitrary class of CR sources if we substitute some other value of the SCR confinement time $T_S$ instead of $T_{SN}$, since the grammage $x$ is an experimentally fixed quantity. Let us then assume that some class of compact CR sources produces an energy $E_C$ in the form of CRs with spectrum $N_{SCR}\\propto \\epsilon^{-\\gamma_{SCR}}$ with average frequency $\\nu_S$, and let us further assume that this spectrum remains unchanged inside the source regions for some period of time $T_S$ after which these CRs are released into the ISM as the GCRs. It is obvious that due to the general energy requirement the production rate $\\nu_S E_C$ should be about the same as $\\nu_{SN}\\delta E_{SN}$. The SCR energy $E_C$, deposited in some initial volume $V$, will produce a dynamically significant disturbance in the background ISM if we assume that the initial SCR energy density $E_C/V$ is much greater than the thermal ISM energy density which in turn is of order $e_{GCR}$. This will inevitably lead to the confinement of these SCRs inside an expanding, disturbed volume for some period of time $T_S$ before the SCRs will be released. It is difficult to give a general relation between $E_C$ and $T_S$ and there may exist only lower bounds on $T_S$, given $E_C$. Therefore, we cannot exclude speculative source classes with many weak but long-lived sources. The opposite case of many weak and short-lived sources is excluded to the extent that the present explanation of the hard \\gr spectrum by the contribution of the SCRs is unique."
    },
    "0004/astro-ph0004347_arXiv.txt": {
        "abstract": "The rotation rates of six weakly-magnetic neutron stars accreting in low-mass X-ray binaries have most likely been measured by Type I X-ray burst observations with the {\\it Rossi X-Ray Timing Explorer} Proportional Counter Array.  The phenomenology of the nearly coherent oscillations detected during the few seconds of thermonuclear burning is most simply understood as rotational modulation of brightness asymmetries on the neutron star surface. We show that, as suggested by Strohmayer and colleagues, the frequency changes of 1--2 Hz observed during bursts are consistent with angular momentum conservation as the burning shell hydrostatically expands and contracts during the burst. We calculate how vertical heat propagation through the radiative outer layers of the atmosphere and convection affect the coherence of the oscillation.  We show that the evolution and coherence of the rotational profile depends strongly on whether the burning layers are composed of pure helium or mixed hydrogen/helium.  Our results help explain the absence (presence) of oscillations from hydrogen-burning (helium-rich) bursts that was found by Muno and collaborators.  We also investigate angular momentum transport within the burning layers and address the recoupling of the burning layers with the star. We show that the Kelvin-Helmholtz instability is quenched by the strong stratification, and that mixing between the burning fuel and underlying ashes by the baroclinic instability does not occur. However, the baroclinic instability may have time to operate within the differentially rotating burning layer, potentially bringing it into rigid rotation. ",
        "introduction": "\\label{sec:Intro}  Type I X-ray bursts have been long understood as thermonuclear flashes on the surfaces of neutron stars accreting at rates  of $10^{-11} M_\\odot \\ {\\rm yr^{-1}} < \\dot M < 10^{-8} M_\\odot \\ {\\rm yr^{-1}}$ in low mass X-ray binaries (LMXBs). The accreted hydrogen and helium accumulates on the surface of the neutron star and periodically ignites and burns. Thermonuclear flash models successfully explain the burst recurrence times (hours to days), energetics ($\\sim 10^{39}\\,{\\rm ergs}$), and durations ($\\sim 10$--$100\\,{\\rm s}$) (Lewin, van Paradijs, \\& Taam 1993; Bildsten 1998), though many quantitative comparisons to observations are less successful (for example, see Fujimoto et al. 1987; Bildsten 2000).   \\begin{deluxetable}{llccccccl} \\tiny \\footnotesize \\tablecaption{Properties of Observed Type I Burst Oscillations\\label{tab:obs}} \\tablewidth{0pt} \\tablehead{ \\colhead{Object} & \\colhead{Time} & \\colhead{Radius} & \\colhead{$\\nu_0$} & \\colhead{$\\Delta\\nu$} & \\colhead{$\\tau$} & \\colhead{$\\Delta\\nu/\\nu_0$} & \\colhead{Oscillations} & \\colhead{References} \\\\ \\colhead{} & \\colhead{(UT)} & \\colhead{Expansion?} & \\colhead{(Hz)} & \\colhead{(Hz)} & \\colhead{(s)} & \\colhead{($10^{-3}$)} & \\colhead{During Rise?} & \\colhead{} } \\startdata 4U 1636-54  & 1996 Dec 28 (22:39:22) & Y & 580.5 & $\\approx 1$--$2.5$ & \\nodata & $\\approx 2$--$4$ &Y &1,2\\tablenotemark{b}\\\\ & 1996 Dec 29 (23:26:46) & Y& 581.5 & $\\approx 2$ & \\nodata & $\\approx 3$ &Y &2,3\\tablenotemark{b}\\\\ & 1996 Dec 31 (17:36:52)\\tablenotemark{c} & Y& 581 & $\\approx 2$ & \\nodata & $\\approx 3$ &Y &2,3\\tablenotemark{b}\\\\ 4U 1702-43  & 1997 Jul 26 (14:04:19) & ? & $329.85\\pm 0.1$ & 2.5 & $1.88\\pm 0.25$ & $7.7\\pm 0.3$ & Y & 4\\tablenotemark{a}\\\\ & 1997 Jul 30 (12:11:58) & ? & $330.55\\pm 0.02$ & 1.6 & $4.02\\pm 0.07$ & $4.8\\pm 0.3$ & N & 4\\tablenotemark{a}\\\\ 4U 1728-34  & 1996 Feb 16 (10:00:45) & N & $364.23\\pm 0.05$ & 2.4 & $3.52\\pm 0.28$ & $6.6\\pm 0.1$ & Y &4,5\\tablenotemark{a}\\\\ & 1997 Sep 9 (06:42:56) & ? & $364.10\\pm 0.05$ & 2.1 & $1.84\\pm 0.15$ & $5.9\\pm 0.2$ & ? &4\\tablenotemark{a}\\\\ KS 1731-26 & 1996 Jul 14 (04:23:42) & Y & $524.61^{+0.13}_{-0.07}$ & $0.96\\pm 0.04$ & $4.1^{+1.3}_{-0.9}$ & 1.8 & N &6,7\\tablenotemark{a}\\\\ & 1999 Feb 27 (17:23:01) & Y & $524.48^{+0.05}_{-0.03}$ & $1.06\\pm 0.05$ & $2.6^{+0.4}_{-0.3}$ & 2.0 &N & 7\\tablenotemark{a}\\\\ Galactic Center & 1996 May 15 (19:32:10) &Y & 589 & $\\approx 1$ & \\nodata & $\\approx 2$& N &8\\tablenotemark{b} \\\\ (MXB 1743-29?) &&&&&&&&\\\\ Aql X-1 & 1997 Mar 1 (23:27:39) & N & $549.76\\pm 0.04$ & $2.38^{+0.03}_{-0.07}$ & $2.69\\pm 0.15$ & 4.3 & N & 9,10\\tablenotemark{a}\\\\ \\enddata \\tablenotetext{a}{In these cases, the frequency evolution was fitted by an exponential model $\\nu(t)=\\nu_0-\\Delta\\nu e^{-t/\\tau}$.} \\tablenotetext{b}{For these bursts, $\\nu_0$ is the frequency seen in the tail, and $\\Delta\\nu$ how much the frequency changes during the burst.} \\tablenotetext{c}{This burst observed an episode of spin down in the tail (Strohmayer 1999b).} \\tablenotetext{}{References.---(1) Strohmayer et al. 1998c; (2) Miller 2000; (3) Strohmayer 1999a; (4) Strohmayer \\& Markwardt 1999; (5) Strohmayer et al. 1996; (6) Smith, Morgan \\& Bradt 1997; (7) Muno et al.  2000; (8) Strohmayer et al. 1997a; (9) Zhang et al. 1998; (10) Fox et al. 2000.} \\end{deluxetable}   Evolutionary scenarios connecting the neutron stars in LMXBs to the millisecond radio pulsars (see Bhattacharya 1995 for a review) predict that neutron stars in LMXBs should be spinning rapidly. This has been confirmed for one system, the $\\nu_s=401$ Hz accreting pulsar SAX J1808.4-3658 (Wijnands \\& van der Klis 1998; Chakrabarty \\& Morgan 1998), in which the neutron star magnetic field ($B\\sim 10^8$--$10^9\\ {\\rm G}$; Psaltis \\& Chakrabarty 1999) channels the accretion flow onto the magnetic polar caps, creating an asymmetry which is modulated by rotation. However, most neutron stars in LMXB's show no evidence for coherent periodicity in the persistent emission, implying that the neutron stars do not possess magnetic fields strong enough to make a permanent asymmetry ($B\\ll 10^{10}\\,{\\rm G}$).  Type I X-ray bursts have provided a new way to determine the spin of these neutron stars. Observations with the Proportional Counter Array (PCA) on the {\\it Rossi X-Ray Timing Explorer (RXTE)} of neutron stars in six LMXBs have shown coherent oscillations during Type I X-ray bursts, with frequencies $\\nu_0$ that range from 300 to 600 Hz (see Table \\ref{tab:obs}). The simplest interpretation is that the burning is not spherically symmetric, providing a temporary asymmetry on the neutron star that allows for a direct measurement of rotation. The coherent nature of the periodicities ($Q\\gtrsim 300$), large modulation amplitudes and stability of the frequency over at least a year support its interpretation as the neutron star spin (Strohmayer 1999b, and references therein).  It is expected theoretically that the burning should not be spherically symmetric. Joss (1978) and Shara (1982) suggested that ignition of a burst occurs at a local spot on the star, and not simultaneously over the whole surface. This is because it takes hours to days to accumulate the fuel, but only a few seconds for the thermal instability to grow. Simultaneous ignition thus requires synchronization of the thermal state of the accreted envelope to one part in $10^3$--$10^4$ over the surface. More likely is that ignition is local and a burning front then spreads laterally (Fryxell \\& Woosley 1982; Nozakura, Ikeuchi, \\& Fujimoto 1984; Bildsten 1995), burning the rest of the accreted fuel and creating a temporary brightness asymmetry on the neutron star surface (Schoelkopf \\& Kelley 1991; Bildsten 1995). Consistent with the picture of a spreading burning front, the pulsation amplitude is observed to decrease during the burst rise while the emitting area increases, reaching a constant value during the decay (Strohmayer, Zhang \\& Swank 1997b; Strohmayer et al. 1998c).  The oscillations during the burst rise are well-explained by the picture of a spreading hotspot. However, several mysteries remain. First, oscillations are often seen in the burst tail, when the whole surface of the star has presumably ignited. The cause of the azimuthal asymmetry at late times is not understood. Secondly, for those objects with $\\nu_0\\approx 550\\ {\\rm Hz}$ in Table \\ref{tab:obs}, there is evidence that the burst oscillation frequency is twice the spin frequency of the star, including observations of a $\\approx 290\\ {\\rm Hz}$ subharmonic (Miller 1999) and the fact that the kHz QPO separation is approximately half the burst oscillation frequency in these sources (see van der Klis 2000 for a review). What might cause an $m=2$ azimuthal asymmetry is unknown. Third, oscillations are not seen in all sources, or all Type I bursts from the same source.  An initially puzzling feature of the observations was that the oscillation frequency often changes during the burst, increasing by $\\Delta\\nu\\approx 1$--$2$ Hz in the burst tail. Strohmayer et al. (1997a) proposed a simple explanation --- that this frequency shift results from angular momentum conservation. The slight hydrostatic radial expansion (contraction) of the burning layers as the temperature increases (decreases) results in spin down (spin up) if angular momentum is conserved. The time for radial heat transport from the burning layers to the photosphere is about one second, which means that the layers are puffed up and spun down by the time the observer sees the burst. As the burning layers cool during the tail of the burst, they contract and spin up. Strohmayer \\& Markwardt (1999) and Miller (1999, 2000) have modelled the observed frequency evolution. They find that, in the evolving frame of the burning shell, the oscillations are coherent, as expected if they are due to rotation. Calculations show that the change in thickness of the burning layers during a burst is $\\Delta z\\approx 20\\,{\\rm m}$ (Hanawa \\& Fujimoto 1984; Ayasli \\& Joss 1982; Bildsten 1998). A simple estimate of the spin down of the burning shell due to this change in thickness is then $\\Delta\\nu\\approx\\nu_s (2\\Delta z/R)\\approx 1\\,{\\rm Hz}\\,(\\nu_s/300\\,{\\rm Hz})(\\Delta z/20\\,{\\rm m})(10\\,{\\rm km}/R)$ where $R$ is the neutron star radius and $\\nu_s$ the spin frequency. This roughly agrees with the observed values (Table \\ref{tab:obs}).  In this paper, we make a first attempt at understanding the evolution of the neutron star atmosphere during a Type I X-ray burst on a rotating neutron star. We use hydrostatic models of the neutron star atmosphere to calculate the expansion and resulting spin evolution during a burst. Our aim is to investigate whether the observed frequency changes during bursts are consistent with spin down of the burning shell, thus lending support to the interpretation that the burst oscillation frequency is intimately related to the neutron star spin. This simple picture demands that the hot burning material decouples from the cooler underlying ashes and conserves its angular momentum, completing a few phase wraps with the underlying star during the burst. We thus examine mechanisms that might couple the burning layers to the neutron star, and ask whether it is plausible that the burning layers can remain decoupled for the $\\approx 10\\,{\\rm s}$ duration of the burst.  We stress that we consider only one-dimensional models for these hydrostatic and coupling calculations. We do not consider the complex question of how the burning front spreads around the star during the burst rise, what determines the number of hotspots on the surface, or what causes the asymmetry at late times during the cooling tail of the burst. We leave these questions for future investigations.  We start in \\S 2 by describing the hydrostatic structure of the atmosphere during fuel accumulation and the X-ray burst. In \\S 3, we calculate the expansion of the atmosphere and show that the expected spin down of a decoupled layer is consistent with observed values. We consider heat transport through the atmosphere, and ask how a single coherent frequency is transmitted to the observer. In \\S 4, we discuss hydrodynamic mechanisms that could transport angular momentum within the burning layers or couple the hot burning material to the underlying colder and denser ashes. In \\S 5, we summarize our results and discuss the many remaining puzzles. Finally, we present our conclusions in \\S 6. ",
        "conclusions": "We have shown that the hydrostatic expansion, spin down, and later spin up during contraction of the neutron star atmosphere naturally explains the magnitude and sign of the observed changes in the nearly coherent oscillations observed in Type I X-ray bursts. Our results support the simple picture proposed by Strohmayer et al. (1997a) to explain the frequency evolution seen in Type I bursts, and thus the identification of the burst oscillation frequency with the neutron star spin.  The amplitude of the oscillation is set not only by the lateral extent of the asymmetry on the neutron star surface, but also by the vertical propagation of heat through the shearing atmosphere. In addition, we find that the spin behavior differs depending on whether the burst results from pure He ignition or He ignition in a H-rich environment. The spin down is smaller by roughly a factor of two for pure He ignition as opposed to mixed H/He ignition, because the pure He models expand outwards less due to the larger mean molecular weight. In addition, the time for heat propagation through the atmosphere is smaller for pure He models.  This might explain why oscillations have not been detected during bursts of long duration ($\\gtrsim 30\\ {\\rm s}$) typical of H burning (for example, the bursts from GS 1826-24 shown in Figures 2 and 5 of Kong et al. 2000). In particular, the recent study of Muno et al. (2000) of bursts from KS 1731-26 found oscillations only during those bursts with short durations $\\approx 10\\ {\\rm s}$, typical of pure He ignition. However, the largest values of $\\Delta\\Omega/\\Omega$ that are observed (for 4U~1702-43 and 4U~1728-34) require base temperatures very close to the limiting value from radiation pressure for pure He bursts, and this may prove to be a problem. More observational studies of the phenomenology of burst oscillations and their relation to burst properties and source properties, in particular accretion rate, are needed. With further theoretical work, these studies could give us important clues as to how the nature of nuclear burning during bursts depends on accretion rate, and help to resolve the disagreement with one-dimensional models first found with EXOSAT (Bildsten 2000 and references therein).  We have not found a robust hydrodynamical mechanism that recouples the burning layers to the underlying star in the $\\approx 10\\ {\\rm s}$ burst duration.  However, the coherence of the observed burst oscillations suggests that little differential rotation occurs within the burning layer itself. As we note in \\S \\ref{sec:coupling}, this might be accomplished through short vertical wavelength baroclinic instabilities, though more theoretical work is needed.  We have not discussed the complex question of the lateral spreading of the burning front during the burst rise, or the cause of asymmetry in the cooling tail of bursts. In the simplest picture of a spreading burning front (Strohmayer et al. 1997b; Strohmayer et al. 1998c) it is a mystery why oscillations are seen in the burst tail, when the whole surface has ignited. One possible explanation is that an asymmetry arises because the atmosphere takes a finite time to cool once ignited, and a finite time to spread around the star. Thus the matter on the opposite side of the star from the ignition point will be a little hotter since it ignited later. If the flux at each point is $\\propto e^{-t/\\tau}$ once ignited, and the burning front propagates around the whole star (starting at the equator) in a time $t_{\\rm burn}$, it is straightforward to show that the peak-to-peak luminosity amplitude in the tail for an observer looking down on the ignition point is $\\Delta L/L=2(1-\\exp(-t_{\\rm burn}/\\tau))/(1+\\exp(-t_{\\rm burn}/2\\tau))^2$, or, for $t_{\\rm burn}\\ll\\tau$, $\\Delta L/L\\approx t_{\\rm burn}/2\\tau$. Thus for a spreading time $t_{\\rm burn}\\approx 0.5\\ {\\rm s}$ and $\\tau\\approx 10\\ {\\rm s}$, the peak-to-peak amplitude in the tail is a few percent, a little less than the 5--20\\% amplitudes observed (Smith, Morgan \\& Bradt 1997; Zhang et al. 1998; Strohmayer et al. 1998a). Another possibility is that the oscillations come and go because of changes in the vertical transport of heat rather than changes in the size of the asymmetry on the stellar surface. For example, Figure \\ref{fig:therm} shows that as the atmosphere heats up during the burst rise, the ratio of thermal time to wrap-around time $t_{\\rm therm}\\Delta\\nu$ increases. As this ratio increases, it becomes more likely that the oscillation will be washed out. It may be that at the burst peak, the oscillations disappear because of this effect.  We have not addressed what mechanism might cause frequency doubling during bursts, as seems to be the case for at least one source (4U 1636-56; Miller 1999). There are two processes that we have discussed in this paper that may play a role. The first is the recoupling mechanism. If the burning layers recouple to the neutron star during the burst, they may do so in an $m=2$ pattern, leading to frequency doubling. Our linear instability analysis of the two-layer beta-plane model (\\S \\ref{sec:baroclinic}) suggests that the baroclinic instability would have $m\\approx 20$ when it goes unstable. However, we have not calculated the correct spherical eigenfunctions. The second possible process is the azimuthal wrapping of the the magnetic field. If an initially dipolar field is wound up, an $m=2$ pattern might emerge in the flow.  A test for the rotational modulation hypothesis might be the distribution of observed $\\Delta\\Omega/\\Omega$ values. For sources undergoing X-ray bursts of the same type (for example, similar accretion rates and same type of ignition) on similar mass neutron stars, we would expect the distribution of $\\Delta\\Omega/\\Omega$ values to be the same for each source, and independent of the spin frequency or number of hotspots on the surface. In fact, the distribution of $\\Delta\\Omega/\\Omega$ so far (Table 1) looks bimodal, though we should be wary of small number statistics. Two sources, 4U 1728 and 4U 1702, which have $\\nu_0\\approx 350\\ {\\rm Hz}$ show frequency changes $\\Delta\\nu/\\nu\\approx 5$--$8\\times 10^{-3}$, whereas the other four objects which have $\\nu_0\\approx 550\\ {\\rm Hz}$ show frequency changes about half as large, $\\Delta\\nu/\\nu\\approx 1$--$4\\times 10^{-3}$. If further observations show this bimodal distribution to be true, this might indicate that the expansion of the atmosphere is less for the $\\nu_0\\approx 550\\ {\\rm Hz}$ objects, for example, because of pure He rather then mixed H/He ignition. Alternatively, we note that the $\\Delta\\nu$ values are very similar for all six objects, and this might point to another explanation for the observed frequency evolution. Care must be taken that the observed $\\Delta\\Omega/\\Omega$ represents the change in burst oscillation frequency from the burst rise to the burst tail. For example, if the oscillation is seen only during the burst tail, the observed $\\Delta\\Omega/\\Omega$ does not represent the full spin down of the atmosphere, as in this case some cooling will have occurred before the oscillation is seen.  We have not discussed other explanations for the frequency drifts. One possibility is that the burst oscillation is a non-radial oscillation (NRO) on the neutron star surface (Livio \\& Bath 1982; McDermott \\& Taam 1987; Bildsten \\& Cutler 1995; Strohmayer \\& Lee 1996). Non-radial oscillations of the frequency observed would be plausible on a rapidly rotating star, as the Coriolis force takes the typically much lower-frequency modes and makes them mugh higher (Bildsten, Ushomirsky and Cutler 1996). The observed frequency would be close to the spin frequency of the star or some multiple $m\\Omega$, and might change by a $1$--$2\\ {\\rm Hz}$ during the burst as the atmosphere cooled. However, it is not clear why only a single mode would be selected, and why that frequency would be stable on timescales of a year or more, for which conditions in the surface layers of the star would most likely change significantly.  The physics of expansion-induced shear during shell flashes may prove important in other contexts. For example, mixing due to shear instabilities during a classical nova explosion could provide the enrichment of CNO elements necessary to achieve a fast nova (see Truran 1982 for a review of the enrichment problem). Previous authors have investigated shear mixing during accumulation of the fuel (see Livio \\& Truran 1990 for a review), but not during the burst itself when shear is induced by expansion of the atmosphere as in a Type I X-ray burst. A similar situation arises in AGB stars during He shell flashes, where mixing of carbon into the proton rich envelope is necessary to allow production of s-process elements (Iben 1991; Sackmann \\& Boothroyd 1991; Langer et al. 1999), and perhaps could be achieved by shear instabilities during shell flashes."
    },
    "0004/astro-ph0004171_arXiv.txt": {
        "abstract": "We examine the supernova remnant (SNR)  0540$-$697 in the Large Magellanic Cloud (LMC) using data from the \\chandra ACIS.  The X-ray emission from this SNR had previously been hidden in the bright emission of nearby X-ray binary LMC X-1;  however, new observations with \\chandra can finally reveal the SNR's structure and spectrum. We find the SNR to be a thick-shelled structure about 19 pc in diameter, with a brightened northeast region.  Spectral results suggest a temperature of 0.31 keV and an X-ray luminosity (0.3-3.0 keV) of 8.4 $\\times 10^{33}$ erg s$^{-1}$. We estimate an age of 12,000-20,000 yr for this SNR, but note that this estimate does not take into account the possibility of cavity expansion or other environmental effects. ",
        "introduction": "Recent observations of the Large Magellanic Cloud (LMC) have opened up a new era in near-extragalactic studies of supernova remnants (SNRs). Thanks to advances in instrumentation over the last decade, we have been able to observe this large and varied sample of remnants without the difficulties in distance determination and obscuration for Galactic remnants, or those in resolution and sensitivity for ones in more distant galaxies. This has allowed us both to examine individual remnants in detail (\\eg\\ Williams \\et\\ 1997, 1999a; Chu \\et\\ 1997, 2000) and as a group (\\eg\\ Williams \\et\\ 1999b; Williams 1999) to explore their interactions with the surrounding interstellar medium (ISM).  The completeness of the LMC sample, however, is a vexing question. Many SNRs lack one or more of the traditional SNR signatures (high \\sii/\\ha\\ ratios, nonthermal radio emission, and X-ray emission) and must be identified through other means (\\eg\\ Chu 1997). New LMC SNRs continue to be found (\\eg\\ Smith \\et\\ 1994; Chu \\et\\ 1993, 1995, 1999).  Each  discovery not only adds to the list of known SNRs, but provides additional insight into how to detect other such SNRs. This has profound implications for both the completeness of the LMC sample and the assemblage of samples from other galaxies, and thus for our understanding of the contribution from SNRs to the energetics and dynamics of the ISM.  One of these discoveries, SNR 0540$-$6944, may prove particularly informative in this regard. The SNR, discovered by Chu \\et\\ (1997), is difficult to observe in optical and radio due to emission from the surrounding \\hii\\ region N159; the X-ray emission was similarly obscured by that from the nearby X-ray binary LMC X-1.  The SNR's expansion was serendipitously uncovered in optical echelle observations, and its SNR nature confirmed by using soft-band (0.1-1.0 keV) \\rosat PSPC images to separate the SNR's thermal emission from the much harder emission of LMC X-1. However, the SNR remained spatially and spectrally confused with LMC X-1 at the resolution of the \\rosat\\ observations.  The unprecedented combination of spatial and spectral resolution and sensitivity from the Chandra (formerly AXAF) X-ray Observatory and its instruments allows us, for the first time, to get a picture of SNR 0540$-$6944 distinct from that of LMC X-1.  Using these data, we are able to present images and spectra from the object itself.  The results confirm the SNR identification of Chu \\et\\ 1997; provide spatial and spectral data on the SNR itself; and illustrate the information that may be gleaned by separating objects whose close proximity confuses their emission. It provides a striking demonstration of the many new areas of investigation made possible by \\chandra's power, even in a brief observation. ",
        "conclusions": "Morphologically, the SNR is without a sharply defined shell in X-rays. The large and distributed X-ray structure, interior to much of the optical emission, implies an older, Sedov-stage SNR.  This picture is strengthened by the relatively low temperature and low luminosity of the X-ray emission.  The velocity of $\\sim$180 \\kms\\ found by Chu et al. (1997) indicates a slow expansion consistent with this picture.  If this velocity is representative of the actual expansion velocity of the remnant, and the SNR is undergoing Sedov-like expansion, we would expect the shock velocity to be $v_{shock} = 4/3 v_{exp} = 240$ km s$^{-1}$.  This shock velocity, in turn, would imply a temperature of $kT = 3/16 \\mu v_{shock}^2$, where $\\mu$ is the mean molecular weight (assumed 1.1$m_H$). For the shock velocity given above, this gives approximately 0.12 keV. The temperature derived from X-ray spectral fits is somewhat higher; the X-ray temperature would imply $v_{exp} = 280$ (220-340) km s$^{-1}$. Similar discrepancies have been noted for other LMC remnants (Williams 1999). One possible explanation for this discrepancy is that the highest-velocity material may be too faint to show up clearly in the echelle spectrum. Another possibility is that the expansion velocity is indeed accurately reflected by the echelle spectroscopy and that this slow shock speed is insufficient to produce X-ray emission at the shock front. In this latter case, the remnant may have entered the shell-forming phase, as evidenced by the pronounced H$\\alpha$ shell.  The observed X-rays in such a case are likely to be ``fossil\"  radiation produced by the cooling of gas shocked to high temperatures earlier in the remnant's expansion.  Using the Sedov equation and the velocity from Chu et al. (1997), and assuming constant external density, we find an age of $\\sim$20,000 yr for this SNR.  The velocity derived from the X-ray temperature would give an age of $\\sim$12,000 yr.  These must be regarded as only approximate figures.  For instance, if the SNR is expanding in the wind-blown bubble formed by its progenitor - a quite plausible scenario, as the remnant is within an \\hii\\ region - it may very well be considerably younger than these estimates. Given this estimate of age, we do not expect a causal connection between SNR 0540$-$6944 and LMC X-1.  The projected distance between the center of the SNR and LMC X-1 is about 1\\ampt75, or a minimum of 26 pc separation.  To reach this distance within the estimated age of the SNR, the compact object would have had to travel at a constant speed of over 1200 km s$^{-1}$.  Our fits to the X-ray spectrum use a normalization constant directly related to the emissivity, $K = 10^{-14} \\int n_e n_H dV / (4 \\pi D^2)$. Here D is the distance to the remnant, $n_e$ and $n_H$ the electron and particle densities (we assume $n_e = 1.1 n_H$), and $V$ the remnant volume.  This allows us to make estimates of the density, energy and pressure within the X-ray emitting gas of this remnant.  Using the fitted value of $K = 5.66 \\times 10^{-4}$ and assuming a volume filling factor for the gas of 10\\% (corresponding to a shell thickness of about 0.33 pc, 3\\% of the radius, a reasonable value for a middle-aged remnant), we calculated the gas density at about $n$=1.2 cm$^{-3}$.  Using the formula $E_{th} = (3/2) N kT$, where N is the total particle number in the hot cavity, we find a thermal energy of about $10^{49}$ erg, suitable for a remnant in which much of the remaining energy is tied up in the kinetic energy of expansion.  The thermal pressure of the hot gas, according to $P=nkT$, is about 6 $\\times 10^{-10}$ dyne cm$^{-2}$.  These figures should, of course, only be regarded as rough estimates, as there is perhaps an order of magnitude uncertainty in the actual volume occupied by the hot gas, as well as additional uncertainties in the fitted temperature and emissivity.  The reason for the brightening of X-rays in the northeast section of the remnant is unclear.  A search for timed emission revealed no significant peaks in the power spectrum; however, given the short exposure and high background for this observation, further investigation is indicated.  The X-ray brightening occurs near an optical ``knot\" of bright emission, and may indicate a region where the SNR shock is encountering denser material.  In summary, we find that Chandra data are sufficient to distinguish SNR 0540$-$6944 from the emission of the nearby X-ray binary LMC X-1.  While these findings are preliminary, as there remain uncertainties in the spatial and spectral responses, they are still informative.  Given that this SNR has a lower luminosity than most LMC SNRs (eg Williams 1999), the results point out the capacity of Chandra to distinguish faint objects in confused regions.  As observations continue, we may uncover an entire population of low-luminosity, older SNRs, in the LMC and even in more distant galaxies.  This has the potential to substantially increase our estimates of SNR rates, which to date have been largely based on the fraction of SNRs that are more readily detected.  SNR 0540$-$6944 also suggests the possibility of finding previously undetected SNRs within the crowded environs of \\hii\\ regions, where we would indeed expect a high population of Type II SNRs."
    },
    "0004/astro-ph0004201_arXiv.txt": {
        "abstract": "We discuss the optical parameter space (luminosity and surface brightness) over which galaxies are known to exist and report on the results from three recent datasets: the Two-degree Field Galaxy Redshift Survey, the Hubble Deep Field north, and the Local Group. These data are combined to provide a unique and comprehensive insight into the luminosity distribution of galaxies over the range $-24 < M_{B} < -8$ mags,  and the surface brightness distribution over the range $18 < \\mu_{e} < 28$ mags$/\\Box''$. The main conclusions are: (1) luminous low surface brightness galaxies are rare; (2) there exists a universal luminosity-surface brightness relation; (3) the most numerous galaxies are low luminosity low surface brightness dwarfs. ",
        "introduction": "The study of galaxy formation and evolution cannot commence in earnest until we have a full and complete description of the local galaxy population - our benchmark if you will. Until we've fully defined the local ``zoo of galaxies'' and their relative abundances we can never be sure that any trend in redshift cannot be interpreted as ``selection bias'' rather than ``evolution''. So how close are we getting to a complete description ? Typically the galaxy population is quantified via the following observable parameters: luminosity, surface brightness, morphology and colour. These are in turn loosely associated with the fundamental properties of: mass, angular momentum, evolutionary history and age, respectively. In this article we review our current insight into the parameter space occupied by galaxies in terms of their luminosity and surface brightness distributions. We argue that in order to unambiguously remove the spectre of selection bias, the luminosity and surface brightness distributions must be considered simultaneously. This is borne out by three recent datasets: the Two-degree Field Galaxy Redshift Survey (2dFGRS), the Hubble Deep Field (HDF), and the latest Local Group (LG) census. Combining these datasets provides the current state-of-the art insight into the low redshift galaxy population. ",
        "conclusions": "Here we have summarised the current understanding of the field galaxy luminosity function and the surface brightness distribution. Because of the clear luminosity-surface brightness relation we conclude that treating these key parameters as separable can lead to confusion and error. We advocate a bivariate brightness distribution which allows us to construct a window into the galaxy population. Using three distinct independent datasets to explore this region we map out the locations in luminosity and surface brightness over which galaxies are known to exist. The three datasets considered here are the Two-degree Field Galaxy Redshift Survey (spanning the high surface brightness luminous regime), the Hubble Deep Field (spanning the medium surface brightness medium luminosity regime), and the Local Group (spanning the low luminosity regime). Combining these three surveys reveals a consistent picture in which {\\it BOTH} the luminosity and surface brightness distribution are flat to the detection boundaries. Across the full optical window over which galaxies are known to exist we have only censured a quarter-panel of this parameter space. Much work remains to be done before we can confidently state that we have a comprehensive and complete understanding of the low redshift universe and all that lies within."
    },
    "0004/astro-ph0004037_arXiv.txt": {
        "abstract": "Recently it has been suggested that the majority of dark matter in the universe resides in  the form of Jupiter mass  black holes distributed cosmologically.  This population makes itself apparent by microlensing high  redshift  quasars and  introducing  pronounced variability  into their observed light curves.   While several arguments dismissing this hypothesis have  been presented,  a conclusive observational  test is, alas, sadly  lacking.  In  this paper we  investigate the effect  of a cosmologically  distributed  population   of  microlensing  masses  on galaxies at low to intermediate redshift.  The magnification of bright stars in  these galaxies leads to small,  but observable, fluctuations in  their   surface  brightness.   The  variability   time  scale  for Jupiter-mass lensing objects is of the  order of a few months and this population  may be  detected through  a future  space-based monitoring campaign of a field containing  $z \\sim 0.5$ galaxies.  The monitoring of galactic surface  brightness will provide an effective  test of the nature of dark matter on cosmological scales. ",
        "introduction": "The nature of dark matter is  one of the major outstanding problems in current  astrophysics, with  the community  being divided  between two broad  camps: those  who propose  particle dark  matter and  those who support  massive  compact objects.   Both  these  ideas have  received recent impetus due to  the measurement of a neutrino mass~\\citep{fu99} and the  detection of MACHOs  in the halo of  our Galaxy~\\citep{al93}, although the  small mass  of the neutrino  and the uncertainty  in the location of the  MACHO lenses suggests that the  nature of majority of the dark matter still remains hidden from us.  One  recent suggestion  proposes  that the  variability  seen in  high redshift quasars, rather than being intrinsic to the source, is due to the  microlensing  action  of  Jupiter mass  black  holes  distributed cosmologically    along    the    line-of-sight    to    the    quasar \\citep{ha93,ha96}.   To   account  for  the   characteristics  of  the variability, a substantial population of black holes is required, with a density  high enough  to account for  a significant fraction  of the total mass  in the universe \\citep{sc93}.  While  there is theoretical support for such  a model, with models of  the quark-hadron transition producing  black  holes  with   a  mass  of  $\\sim10^{-3}{\\rm  \\msun}$ [e.g.~\\citet{ha97}], observations have failed to reveal their presence in     studies    of    the     equivalent    widths     of    distant quasars~\\citep{ca82,da91}, although this conclusion has been contested \\citep{ha96}.  Other  objections to this  microlensing hypothesis have been   raised    \\citep{ba95,al95},   with   corresponding   rebuttals \\citep{ha97}.   More recently, arguments  have become  quite vitriolic with the suggestion that this idea of cosmologically distributed black holes is being  ignored by the astronomical community  even though the apparent evidence is overwhelming~\\citep{ha99}.  The microlensing of cosmological supernovae provides a potential probe of  any   population  of  black   holes  \\citep{sc87,ra91,ra91a,me99}, although their rarity suggests that a long observational period may be required until  an unambiguous  detection is made.   In this  paper we provide a  complementary probe  of a cosmologically  distributed black hole population;  instead of supernovae  as the source  population, we focus here  on microlensing  on the much  more numerous  population of normal   stars   in   galaxies    out   to   ${\\rm   z\\sim0.5}$.    In Section~\\ref{method} we outline the analytic and numerical approach to this problem,  while in  Section~\\ref{distributions} the results  of a suite  of  simulations  are  presented.   Section~\\ref{timescales}  is concerned with the time scale of the microlensing variability, and the observational   feasibility   of    this   study   is   discussed   in Section~\\ref{observational}.   Section~\\ref{conclusions}  presents the overall conclusions of this study. ",
        "conclusions": "Recent  studies have  claimed that  the fluctuations  observed  in the light  curves of  high redshift  quasars are  due to  the microlensing effect  of a  population  of cosmologically  distributed Jupiter  mass black holes. While arguments have been presented to refute this claim, no observational  test of this hypothesis has  been readily available. In this paper we have investigated the effect this putative population would have on  the surface brightness distribution of  galaxies out to intermediate  redshift.   The   results  of  this  study  conclusively demonstrate  that  if   a  significant  population  of  cosmologically distributed compact  objects are present, with enough  mass to account for  the  universal  dark  matter  budget, then  they  will  induce  a `twinkling'  into  the observed  surface  brightness distributions  of galaxies  at low  to  intermediate redshifts.   The  magnitude of  the observed fluctuations  are a  function of the  redshift of  the source galaxy,  ranging from  negligible  values in  the  local Universe,  to $\\sim1$\\% at $z=0.5$, with a time-scale of variability on the order of weeks to months.  While the  optical depth to  microlensing at such low  to intermediate redshifts is still  quite small and the induced  fluctuations are at a relatively  low  level, the  sky  density  of  galaxies out  to  these redshifts vastly  exceeds that of  supernova and quasars,  the current focus  for  the  search  for cosmologically  distributed  microlensing masses.  With sufficiently deep imaging with source pixel luminosities of greater than $10^5L_\\odot$, microlensing induced surface brightness variability can  be detected over galaxies  at intermediate redshifts, and the ubiquitous  nature of this variability means  it will dominate over any contaminating effects, such as supernovae, variable stars and self-lensing.  The scale of  the induced variability make observations conducive over a  single observing season, and with  the advent of the Next  Generation  Space  Telescope,  such observations  will  soon  be possible.  Hence,  a monitoring program to search  for fluctuations of the  surface brightness of  the plethora  of galaxies  at $z\\simlt0.5$ offers a effective test of the existence of cosmologically distributed compact matter."
    },
    "0004/astro-ph0004084_arXiv.txt": {
        "abstract": "The possibility of the magnetic flux expulsion from the Galaxy in the superbubble (SB) explosions, important for the $\\alpha{\\mbox -}\\Omega$ dynamo, is considered. Special emphasis is put on the investigation of the  downsliding of the matter from the top of the shell formed by the SB explosion which is able to influence the kinematics of the shell. It is shown that either Galactic gravity or the development of the Rayleigh-Taylor instabilities in the shell, limit the SB expansion, thus, making impossible magnetic flux expulsion. The effect of the cosmic rays in the shell on the sliding is considered and it is shown that it is negligible compared to Galactic gravity. Thus, the question of possible mechanism of flux expulsion in the $\\alpha{\\mbox -}\\Omega$ dynamo remains open. ",
        "introduction": "The magnetic field in our Galaxy and in other spiral galaxies is usually believed to have been amplified from a weak seed field by a hydromagnetic dynamo, which exists due to the presence  of the large-scale differential rotation and small-scale cyclonic turbulence in the Galaxy (Parker 1970, 1971; Vainshtein $\\&$ Ruzmaikin 1971, 1972; Moffat 1978). It has been  suggested that any primordial magnetic field could be expelled from the Galaxy by the dynamic motions in less than a billion years (Parker 1971) so it would seem that some flux amplification is necessary to explain the Galactic field. The theory of such a dynamo has been formulated in a precise way through the mean field equations and the solutions to these equations indicate that the field would be amplified. The resulting field appears to correspond to the magnetic field patterns in our galaxy and others.  On the other hand, a number of criticisms of this theory have emerged. One of them concerns the intense development of the small scale fields, which could  damp the turbulence and stop the dynamo action unless they saturate at levels which do not interfere with mean field dynamo (Kulsrud $\\&$ Anderson 1992; Parker 1992; Vainshtein $\\&$ Cattaneo 1992).  Another unresolved problem is the expulsion of flux from the galactic disc. A very important point is that the theory of the $\\alpha{\\mbox -}\\Omega$ dynamo predicts the amplification of some small preexisting magnetic field only if some magnetic diffusion is present in the Galaxy. However, too large a diffusion is destructive for the dynamo, because of its dissipative role in the process of generation and it seems attractive to suppose that a dynamo without diffusion at all will be the most effective. But Ruzmaikin, Shukurov, $\\&$ Sokoloff (1988) showed that this is not possible. The physical reason for this is rooted in the very strong flux freezing of the galactic plasma, because magnetic lines cannot break and the number of field lines in the disc can be increased only by toroidal stretching, which is accomplished by the dynamo action. Any stretching creates field of both signs, to conserve the total flux, and, for net amplification to occur, those portions of field lines which are of the wrong sign must be expelled from the Galaxy. In the standard $\\alpha{\\mbox -}\\Omega$ dynamo theory this is done by magnetic diffusion.  Ruzmaikin, Shukurov, $\\&$ Sokoloff (1988) considered  $\\alpha{\\mbox -}\\Omega$-dynamo in the case of the thin disc and demonstrated that temporal evolution of the magnetic flux is governed by the following set of equations: \\begin{equation} \\frac{\\partial}{\\partial t}\\int\\limits^1_0 B_r(t,z)dz=\\beta \\left.\\frac{\\partial B_r}{\\partial z}\\right|^1_0, \\label{flux1} \\end{equation} \\begin{equation} \\frac{\\partial}{\\partial t}\\int\\limits^1_0 B_{\\varphi}(t,z)dz=\\beta \\left.\\frac{\\partial B_{\\varphi}}{\\partial z}\\right|_0^1+G\\int\\limits^1_0 B_r dz, \\label{flux2} \\end{equation} where $G=r\\partial \\Omega/\\partial r$-measure of the differential rotation, $\\beta$-magnetic diffusion, and $0$ and $1$ correspond to the center plane and boundary of the Galactic disc.  The $\\beta$ terms represent the expulsion of flux. It is clear, that setting $\\beta=0$ in (\\ref{flux1}) and (\\ref{flux2}) we immediately get that \\begin{equation} \\int\\limits^1_0 B_r dz=const, \\label{flux3} \\end{equation} and that $\\int\\limits^1_0 B_{\\phi}dz$ may grow only linearly due to the stretching of lines by the galactic differential rotation. Thus, there is no exponential field growth, essential to the dynamo. One can easily see that for the dynamo to operate, there must be some  nonzero   flux escape through the upper boundary of the disc, that is $\\beta\\neq0$.  This is consistent with the topological constraint that the total number of lines of force including those negative lines expelled from the disc must be constant. Indeed, a simple estimate of the expulsion terms making use of the numerical results in Ruzmaikin, Shukurov, $\\&$ Sokoloff (1988) shows that the negative flux expelled during one growth cycle is comparable to the positive flux in the disc at the beginning of the  e-folding.  Although the $\\alpha{\\mbox -}\\Omega$ dynamo theory is complicated, this physical intuition of flux expulsion can be considered in the absence of these complications. Further, the mechanism of expulsion need not be tied to the $\\beta$ diffusion inside the disc.  The main problem with the expulsion of flux is that this flux is loaded with matter so that it is related to the expulsion of matter against the strong gravity of the galactic disc. The most likely process to expel flux is the phenomenon of sequential supernova (SN) or superbubbles (SB), which sweep up matter into dense, radiatively cooled shells. Magnetic field, tied to the matter due to the strong flux freezing in the ISM, is also swept up and deposited in these shells. If some part of the shell leaves the Galaxy, it carries the frozen-in magnetic field with it, thus producing the flux expulsion.  But most of these superbubbles  are not powerful enough to expel matter out of the gravitational well of the disc. The only possibility for flux expulsion seems to be: as the  bubble expands, the field lines in the shells of SBs are not horizontal but form arcs, along which matter can slide down, lowering the amount of matter on the top of the lines and allowing some flux to escape.  There is further difficulty with the mechanism involving the SBs which is relevant to its application to the $\\alpha{\\mbox -} \\Omega$-dynamo. The $\\alpha{\\mbox -} \\Omega$-dynamo assumes {\\it small-scale} turbulence while the cavities produced by SB explosions may  reach $\\sim 500$ pc or larger, which is  greater than some of the length scales of the galactic disk. For the dynamo theory in its conventional form to be applicable it is important that turbulence be small scale, because it involves the expansion of  the turbulent electromotive force ${\\cal E}$, which describes the effect of turbulent motions on the mean (or ensemble- averaged) magnetic field, in terms of the mean magnetic field itself and its spatial derivatives: \\begin{equation} {\\cal E}_i=\\alpha_{ij}<B_j>+\\beta_{ijk}\\frac{\\partial <B_j>}{\\partial x_k} \\end{equation} (Moffat 1978)  If the scale of the turbulence is too large, then the expansion is invalid  and usual $\\alpha{\\mbox -}\\Omega$-dynamo theory must be modified. That's why, for example, direct application of  $\\alpha$ and $\\beta$ tensors calculated by Ferri$\\grave {\\rm e}$re (1995, 1998) for SBs and SNs to the $\\alpha{\\mbox -}\\Omega$-dynamo theory can lead to an overly optimistic estimates of the rate of flux escape.  Although $\\alpha{\\mbox -}\\Omega$-dynamo theory is not strictly applicable to the case of SBs, the actual operation of them in amplifying the field is clear from the work of Ferri$\\grave {\\rm e}$re (1991, 1995, 1998). It is also clear that the rapid escape of the flux from the disk is essential. In this paper we show that because of the deep gravitational well of the Galaxy it is difficult for the matter and field lines to escape, and consequently for the mean field to grow. In Ferri$\\grave {\\rm e}$re's works she finds the lines of force rising with the SB but does not follow them long enough to see that they must fall back into the disk and inhibit the growth of the field.  In this note we quantitatively examine the dynamics of the rising field lines carried by SBs and show that even with sliding the matter and flux are unlikely to escape. Thus, the requirement of the escape of flux provides a strong constraint for the $\\alpha{\\mbox -}\\Omega$ dynamo to overcome if it is to amplify the galactic magnetic field. ",
        "conclusions": "In this paper we consider the effect of the downsliding of the matter which takes place in the expanding superbubbles for application to the expulsion of the magnetic flux from the Galaxy. This expulsion is an important ingredient of the  $\\alpha{\\mbox -}\\Omega$ dynamo theory. However, it is shown that even the inclusion of the sliding into the calculations of the kinematics of the superbubbles, does not enable matter and frozen in flux to leave the Galaxy in SBs. One must note, that the impossibility of the flux escape from the Galaxy weakens significantly the Parker's (1971) argument against the primordial magnetic field, because there is actually no mechanism  to expel it.   Some authors (Korpi et al 1999a, Korpi et al 1999b) have considered dynamics of the superbubbles in the gravitational field of the Galaxy in greater detail and they also find that the matter does not reach terminal velocities larger than the escape velocity from the Galaxy. In their simulations they do observe the development of the SB and its blowout from the disk, but at the height of several kiloparsecs the velocity of the matter is too small for the matter and the field to leave the Galaxy, which agrees with our conclusions. However, they do not comment on its relation to the dynamo.  Other authors (e.g. Hanasz \\& Lesch 1998, Moss et al 1999), more directly concerned with dynamos, do not investigate the dynamics of escape but merely assume that once the magnetic field lines reach the boundary of the disk, they are advected away by some mechanism leading to a vacuum boundary conditions. They do mention magnetic buoyancy as a possible escape mechanism but buoyancy is unlikely to be important outside the disk and clearly plays no role in the SB escape mechanism, especially when the magnetic fields are first amplified from weak seed fields.  In this paper an analytical formalism for the consideration of sliding was built on reasonable assumptions, which enabled to include the back-reaction of lowering the density of the matter on the top of the SB, on the expansion of the shell in the radial direction. All the SB, depending on their luminosity and conditions in the surrounding ISM were demonstrated to fall into two classes: low luminosity SBs which are  stopped by the Galactic gravitational field and fall back, and powerful SBs, which are possibly able to reach low density regions at high altitude. Even without sliding these powerful SBs might be able to expel matter from the Galaxy. However, before this happens the shells of these high luminosity  SBs fragment into separate blobs due to the Rayleigh-Taylor instability, which develop when the shells start accelerating at the high altitudes, where the density of the matter and the outer pressure are very small. These separate fragments of the shells continue to move ballistically with velocities too small to leave the Galaxy.    It was shown that our conclusion about the impossibility of flux to escape in SB explosions does not depend essentially upon the details of the inner pressure behavior in the SB or the shape of the SB top. The inclusion of the CR pressure gradient in the shell also does not influence the downsliding because of the high uniformity of CR density in the shell caused by the very large diffusion velocity of CR in the shell.  It should be remarked that throughout the paper we have assumed that magnetic field lines are tied to the gas. If the ionization is the shell is low, then ambipolar diffusion (under the influence of the CR pressure gradient perpendicular to the shell) could allow some small slippage of the field lines through the neutral component and, thus, some small amount of escape. We do not discuss this possibility in this paper.  All this lead us to the conclusion that if there does exist some mechanism responsible for the flux expulsions from the Galaxy needed by the $\\alpha{\\mbox -}\\Omega$ dynamo this mechanism is not superbubbles. The question of whether such a mechanism exist remain an open one."
    },
    "0004/astro-ph0004185_arXiv.txt": {
        "abstract": "We analyze cooling of neutron stars, assuming the presence of localized  protons in the densest region of their cores.  Choosing a single threshold density for proton localization and adjusting neutron star mass, we reproduce the observational data on effective surface temperatures of Vela and PSR 0656+14, with or without an accreted hydrogen envelope. However, the presence of a tiny hydrogen envelope is mandatory, in this model, for reproducing the Geminga data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004070_arXiv.txt": {
        "abstract": "We study the stochastic coagulation equation using simplified models and efficient Monte Carlo simulations.  It is known that (i) runaway growth occurs if the two-body coalescence kernel rises faster than linearly in the mass of the heavier particle; and (ii) for such kernels, runaway is instantaneous in the limit that the number of particles tends to infinity at fixed collision time per particle. Superlinear kernels arise in astrophysical systems where gravitational focusing is important, such as the coalescence of planetesimals to form planets or of stars to form supermassive black holes.  We find that the time required for runaway decreases as a power of the logarithm of the the initial number of particles.  Astrophysical implications are briefly discussed.  {\\it Key Words}: planetesimals, planetary formation; collisional physics; methods, numerical. ",
        "introduction": "A frequently-encountered process in many fields of science is the random coalescence of small bodies into larger ones, conserving total mass.  Astrophysical examples include the coalescence of planetesimals into planets \\citep{Safronov} and of stars into black holes \\citep{Lee,QS}. When the number of bodies is large, coalescence is often modeled by Smoluchowski's equation \\citep{Smo}, also known as the \\emph{statistical equation}: \\be\\label{eq:Smo} \\frac{dc_i}{d\\tau} = \\frac{1}{2}\\sum\\limits_{j=1}^{i-1} K(i-j,j)c_{i-j}c_j -c_i\\sum\\limits_{j=1}^{\\infty} K(i,j) c_j, \\ee Here the \\emph{concentration} $c_i(\\tau)$ is the number of bodies per unit volume of mass $m_i\\propto i$ at time $\\tau$, and the functional form of the \\emph{coagulation kernel} $K(i,j)=K(j,i)$ is chosen to approximate the mass dependence of the two-body collision rate.  The statistical equation is intended to model coalescence on average, smoothing over fluctuations.  It is not expected to be accurate when the initial number of particles, $N$, is small. More interestingly, eq.~(\\ref{eq:Smo}) can fail even for $N\\gg1$ if $K(i,j)$ increases sufficiently rapidly with $i$ and $j$. Coalescence should conserve mass, so that the mass density \\bd \\rho(\\tau)\\define \\sum\\limits_{k=1}^\\infty k\\, c_k(\\tau) \\ed is constant. It is easy to show from eq.~(\\ref{eq:Smo}) that $\\dot\\rho=0$ provided that the relevant summations converge and can be interchanged. But in the analytic solution of the case $K(i,j)=ij$ starting from the \\emph{monodisperse} initial conditions \\be \\label{monodisperse} c_i(0)=\\delta_{i1},\\qquad\\rho(0)=\\frac{N}{V}=1, \\ee $\\rho(\\tau)$ begins to decrease  after $\\tau=1$ \\citep{Trubnikov}. This is usually interpreted to mean that a macroscopic \\emph{runaway particle} has formed---also known as a \\emph{gel}, because of applications in physical chemistry. It is believed that for $K(i,j)\\sim (ij)^\\nu$ with $\\nu>1$, gelation begins immediately in the statistical equation \\citep[and references therein]{Jeon,MHL}. The model~(\\ref{eq:Smo}) can be extended to include the gel/runaway particle explicitly \\citep{Flory}. Monte Carlo simulations for large but finite $N$ show, however, that such models do not accurately predict the time dependence of the gel mass even on average if $\\nu>1$ \\citep{Spouge}.  Coagulation is more correctly described by the \\emph{stochastic} equation for the joint probability $f(n_1,n_2,\\ldots;t)$ for the occupation numbers $\\{n_1,n_2,\\ldots\\}$ of the mass bins: \\beq \\frac{\\partial}{\\partial t} f({\\bf n}; t)&=&\\sum\\limits_{i=1}^N \\, \\sum\\limits_{j=i+1}^N K(i,j)\\,(n_i+1)(n_j+1) f(...,n_i+1,...,n_j+1,...,n_{i+j}-1,...; t)\\nonumber\\\\ &+&\\sum\\limits_{i=1}^N \\frac{1}{2} K(i,i)\\,(n_i+2)(n_i+1) f(...,n_i+2,...,n_{2i}-1,...; t)\\nonumber\\\\ &-&\\sum\\limits_{i=1}^N \\,\\sum\\limits_{j=i+1}^N K(i,j)\\,n_i n_j f({\\bf n}; t)~-~ \\sum\\limits_{i=1}^N \\frac{1}{2} K(i,i)\\,n_i(n_i-1) f({\\bf n}; t), \\label{STOCHASTIC_EQ} \\eeq The statistical equation results from taking first moments, \\be\\label{avgn} c_i(t)\\define V^{-1}\\,\\overline{n}_i(t)\\define V^{-1}\\,\\int n_i f(n_1,\\ldots,n_{i-1},n_{i},n_{i+1},\\ldots; t) \\,dn_1\\ldots dn_{i-1} dn_{i+1}\\ldots, \\ee in the limit that $N,V\\to\\infty$ at fixed $\\rho(0)$. In order to get a closed set of equations, one assumes that $\\overline{n_i n_j}=\\overline{n}_i\\times\\overline{n}_j$.  This is justified if the occupation numbers are approximately uncorrelated. But if runaway should occur, then the occupation numbers of all the low-mass bins are correlated with bins $i\\sim N$ traversed by the runaway particle.  Like the statistical equation, the stochastic equation has analytic solutions for the kernels $K(i,j)\\propto~\\mbox{constant}$, $i+j$, and $i\\times j$ \\citep{Lushnikov, TN}. For the first two kernels, the predictions of the two equations are compatible in the following sense: after monodisperse initial conditions, $c_i(\\tau)$ computed from eq.~(\\ref{eq:Smo}) agrees with $c_i(t)$ computed from eqs.~(\\ref{STOCHASTIC_EQ})-(\\ref{avgn}) in the limit $N\\to\\infty$ if $t$ and $\\tau$ are related by the initial collision time per particle, \\be t_{\\rm coll}\\define\\left[(N-1)\\,K(1,1)\\right]^{-1},\\qquad \\tau= t/t_{\\rm coll}. \\qquad \\label{DIMENSIONAL_TIME} \\ee For these kernels, there is no runaway, and the statistical equation conserves mass. Even for $K(i,j)=ij$, a similar agreement is found between the statistical and stochastic equations at $\\tau<1$ and mass bins $i\\ll\\sqrt{N}$ \\citep{TN}.  The stochastic equation confirms that runaway begins at $\\tau>1$; more precisely, \\cite{Lushnikov} shows that the quantity \\bd \\kappa(\\tau)\\define \\lim_{N\\to\\infty}\\left[N^{-2} \\sum\\limits_{i=1}^N\\, i^2\\,\\overline{n}_i(\\tau)\\right], \\ed which can be interpreted as the runaway mass fraction on average, becomes nonzero at $\\tau=1$ and $\\approx 1-2e^{-\\tau}$ at $\\tau\\gg 1$.  Kernels such that $K(i,j)\\propto i^\\nu$ when $i\\gg j$ are sometimes called ``unphysical'' if $\\nu>1$ because if collision rate were proportional to surface area, then surface area would have to increase more rapidly than mass.  But in astrophysics (and elsewhere), gravitational focusing and mass stratification may conspire to produce $\\nu>1$ (\\S5).  Thus a better term for such kernels would be \\emph{superlinear}.  The occurrence and timescale of runaway are of great interest since the runaway particle may represent a planet or supermassive black hole, for example.  Unfortunately, few general results are known for the superlinear regime.  \\cite{Spouge} conjectured on the basis of Monte Carlo simulations of stochastic coagulation that {\\it runaway is instantaneous for $\\nu>1$ in the limit $N\\to\\infty$}: in other words, the entire mass is consumed by a single particle after an infinitesimal multiple of the single-particle collision time $t_{\\rm coll}$.  A similar conjecture had previously been made by \\cite{DLP} for the special case $\\nu=3$, also on the basis of Monte-Carlo simulations.  Recently, \\cite{Jeon} has supplied a proof.  Clearly, runaway/gelation is not instantaneous for finite $N$. To the best of our knowledge, there have been no general and quantitative statements about the scaling of the runaway time with $N$, although the question is an important one since $N$ is never truly infinite in practical applications. The main result of the present paper will be the conjecture, supported by Monte Carlo simulations, that $\\tau_{\\rm runaway}$ varies as a negative power of $\\log N$ for $\\nu>1$.  Even before performing our simulations, we were lead to this conjecture on the basis of the highly simplified model presented in \\S 2.  Our Monte Carlo algorithm for stochastic coagulation is described in \\S3 and tested against the analytic solutions of the stochastic equation cited above.  Simulations for $\\nu>1$ are reported in \\S4.  Finally, in \\S5, we discuss applications of our conjecture to cases of astrophysical importance. ",
        "conclusions": "\\label{DISCUSSION}  Our Monte Carlo simulations indicate that runaway is instantaneous in the limit $N\\to\\infty$ if the coalescence kernel is superlinear, in agreement with previous authors \\citep{Jeon,Spouge}. But the simulations, as well as the heuristic model of \\S2, suggest a more quantitative new result: namely, that the limit $\\tau_{\\rm runaway}\\to 0$ is approached only logarithmically in $N$ for power-law kernels.  We consider briefly two astrophysical applications: coalescence of planetesimals into planets in a protostellar disk, and black hole formation through stellar mergers in galactic nuclei.  In both cases, the two-body collision rate per unit volume is $\\nu_i\\nu_j\\langle\\sigma v\\rangle_{ij}$, where $\\nu_i$ is the number of particles of mass $m_i$ per unit volume, and the rate coefficient is approximately (Lee 2000) \\be\\label{RATE0} \\langle\\sigma v\\rangle_{ij} = \\pi(R_i+R_j)^2\\left[ v_{ij} + \\frac{2G(m_i+m_j)}{(R_i+R_j)}\\,\\frac{1}{v_{ij}}\\right]. \\ee Here $R_i$ is the radius corresponding to mass $m_i$, and $v_{ij}$ is the root-mean-square relative velocity of the particles. The prospects for runaway are usually more favorable when the second (\\emph{gravitational focusing}) term in the square brackets above is dominant. Henceforth, we neglect the first term.  A power-law mass-radius relation \\be\\label{MR} R\\propto m^\\alpha \\ee describes planetesimals if  $\\alpha\\approx 1/3$, and stars if $\\alpha\\approx 1$. For a maxwellian velocity distribution with equipartition of kinetic energies among mass groups, \\bd v_{ij}= v_1\\left(\\frac{i+j}{ij}\\right)^{1/2}, \\ed where $v_1$ is the velocity dispersion of particles of mass $m_1$ and $m_i=i\\times m_1$. Then \\be\\label{RATE1} \\langle\\sigma v\\rangle_{ij} \\approx \\frac{2\\pi Gm_1 R_1}{v_1}\\, \\left(i^\\alpha+j^\\alpha\\right)(i+j)^{1/2}(ij)^{1/2} \\ee Since this is proportional to $j^{\\alpha+1}$ for $j\\gg i$, runaway is possible if $\\alpha>0$.  Mass stratification somewhat intensifies the tendency towards runaway. A large-scale gravitational potential will concentrate the heavier particles into a smaller volume than the lighter ones. Thus in a disk of planetesimals with orbital angular velocity $\\Omega$, the distribution of mass group $i$ with height $z$ above the midplane is \\bd \\nu_i(z)\\approx\\nu_i(0)\\exp\\left(-z^2/2H_i^2\\right),\\qquad H_i\\equiv v_i/\\Omega\\propto i^{-1/2}. \\ed Hence upon integrating with respect to $z$, \\be\\label{STRATP} \\int\\limits_{-\\infty}^\\infty\\nu_i(z)\\nu_j(z)\\,dz= \\left(2\\pi H_1^2\\right)^{-1/2}\\,\\hat n_i\\hat n_j\\, \\left(\\frac{ij}{i+j}\\right)^{1/2}. \\ee Here $\\hat n_i$ is the number of particles of mass $m_i$ per unit area. The mass-dependent factor in this expression multiplies the rate coefficient (\\ref{RATE1}) to yield an effective collision kernel ($\\alpha\\approx 1/3$) \\be\\label{KPLANET} K(i,j)= K(1,1)\\,\\times\\,ij\\, (i^{1/3}+j^{1/3})/2, \\ee which has the same $j^{4/3}$ scaling for $j\\gg i$ as the rate coefficient (\\ref{RATE1}) but a stronger dependence on the mass of the lighter particle: $i$ versus $i^{1/2}$. This is likely to cause more rapid consumption of intermediate-mass objects, so we expect evolution as if for a multiplicative kernel [$K(i,j)\\propto(ij)^{4/3}$] rather than an additive one [$K(i,j)\\propto(i+j)^{4/3}$].  The situation in galactic nuclei is more complicated if the heavier stars concentrate into a selfgravitating population.  Setting aside this possibility (which would make runaway still more rapid) and assuming a harmonic mean potential, we get in place of eq.~(\\ref{STRATP}) \\bd \\int\\limits_0^\\infty\\nu_i(r)\\nu_j(r)\\, 4\\pi r^2dr= \\frac{ n_i n_j}{2(2\\pi)^{5/2}H_1^3}\\, \\left(\\frac{ij}{i+j}\\right)^{3/2}, \\ed since the geometry is spherical rather than planar.  Here $n_i$ is the total number of particles of type $i$.  So in this case, with $\\alpha=1$, \\be\\label{KSTAR} K(i,j)= K(1,1)\\,\\times\\,(ij)^2. \\ee This is precisely our multiplicative $\\nu=2$ kernel.  As many as $\\sim 10^{11}$ asteroid-sized planetesimals are required to make up the mass of a terrestrial planet. The smallest coalescence exponent for which our simulations give reliable averages for $\\tau_{\\rm runaway}$ is $\\nu=3/2$, and the largest $N$ we have used is $10^9$.  Nevertheless, by extrapolating the curves in Fig.~(\\ref{RUNAWAY_TIME}), we estimate $t_{\\rm runaway}\\approx 0.4\\,t_{\\rm coll}$  at $N=10^{11}$, versus $t_{\\rm runaway}\\approx 2\\,t_{\\rm coll}$ at $N=10^2$.  The dependence on $\\log N$ is somewhat stronger than one would expect from the asymptotic scaling (\\ref{ESTIMATED_RUNAWAY_TIME}) because of the curvature in Fig.~(\\ref{RUNAWAY_TIME}).  Black holes associated with quasars in the nuclei of galaxies are thought to have masses of order $10^{8-9}\\,M_\\odot$.  The prevailing wisdom is that the black hole forms by accretion from a gas disk, but formation by stellar coalescence has long been a popular alternative model \\citep[and references therein]{Rees}.  Referring again to Fig.~(\\ref{RUNAWAY_TIME}), but now for $\\nu=2$, we find $t_{\\rm runaway}\\approx .09\\,t_{\\rm coll}$ at $N=10^9$ versus $t_{\\rm runaway}\\approx .35\\,t_{\\rm coll}$ at $N=10^2$, roughly in agreement with the predicted asymptotic scaling $(\\log N)^{1-\\nu}$."
    },
    "0004/astro-ph0004246_arXiv.txt": {
        "abstract": "We examine  the excitation of  oscillations in the magnetic network of the  Sun through the footpoint   motion of photospheric magnetic flux tubes located in intergranular lanes.   The motion is derived from a  time series  of high-resolution  G band  and continuum  filtergrams using an object-tracking technique.  We model the response of the flux tube  to the footpoint motion in  terms of the Klein-Gordon equation, which   is  solved analytically as     an  initial value problem   for transverse  (kink) waves.  We compute  the  wave energy flux in upward propagating transverse waves. In general we find that the injection of energy into  the chromosphere  occurs in  short-duration pulses, which would lead to a  time variability in chromospheric emission that is incompatible with observations.  Therefore, we consider the effects of turbulent convective flows on flux tubes in intergranular lanes. The turbulent flows are simulated by adding high-frequency motions  (periods   5-50 s) with an amplitude of 1~km~ s$^{-1}$. The  latter are simulated by adding random velocity fluctuations to the  observationally determined  velocities.  In  this case we find that the energy flux is much less intermittent and can in principle carry adequate energy for chromospheric heating. ",
        "introduction": "It is now well established  that the solar photosphere  is threaded  with strong (kilogauss) magnetic fields in the form of vertical  flux tubes located in the network, where they are  observed as bright points.  Observations have revealed  that  these  network bright points  (NBPs)  are  in a highly dynamical state due  to the buffeting action  of granules (e.g. Muller 1983; Muller et  al. 1994; Berger \\& Title 1996; Berger et al. 1998; van Ballegooijen  et al. 1998).  The jostling of magnetic elements by granules can excite oscillations in  flux tubes, which  can in principle  be an important mechanism for chromospheric and coronal heating, provided that the motions are rapid enough  (Choudhuri, Auffret  \\& Priest 1993;    Steiner, Kn\\\"olker  \\& Sch\\\"ussler 1994; Hasan \\& Kalkofen 1999,  hereafter  Paper  I). As transverse  waves, excited  in the photosphere,  travel upwards their velocity amplitude  increases. In chromospheric layers, the  velocity amplitude becomes comparable to the tube speed  for kink waves, leading to  an  efficient coupling with  longitudinal waves (Kalkofen 1997).  The latter  can dissipate  by forming   shocks (e.g. Zhugzhda, Bromm \\& Ulmschneider 1995).  In Paper I it was argued that  network oscillations can be efficiently excited through the buffeting action on flux tubes by external granules with intermittent motions (of about 2 km s$^{-1}$) occurring on a  time scale of less than about half the cutoff period of kink waves. Typically, a period of 1 min. is consistent with the general result. We assumed an analytic form for the external motions. In the   present work, we use observational data obtained  at the Swedish Solar Observatory   at La Palma during 1995 (Berger et al. 1998) to infer the horizontal footpoint motions of magnetic elements in the photosphere.  We model the response of flux tubes to these motions and calculate the energy  flux in upward-propagating transverse (kink) waves. We also simulate the effect  of hypothetical motions on time  scales of  a few seconds, which  could significantly  enhance the energy  flux in kink waves.  Finally, we examine some implications of our results.  \\section {MEASUREMENT OF BRIGHT POINT MOTION} The data were collected on October 5, 1995  between 10:57 and 12:08 UT, simultaneously  in the  G band  (4305  \\AA) and the nearby continuum (4686 \\AA), using the phase-diversity method (see L\\\"ofdahl et al. 1998; Berger et al. 1998).  After filtering out  $p-$mode oscillations, a time  series of 180 images was obtained,  with  high  spatial  resolution and with an interval between frames of 23.5~s. Following Berger et al., we  subtract the  images in  the G band and   4686~\\AA, which  clearly highlights the network bright points.  The  motion of   the  bright  points was  followed    using a tracking technique with ``corks'' as tracers of bright points (van Ballegooijen et al.  1998, hereafter  Paper II).  The  corks, each with a radius of 60  km (the pixel size of  the current data set),  were advected by an artificial flow field that is   proportional  to the gradient of   the intensity in  the magnetic image, and  their positions were followed in time (for details see Paper II). In Figure 1 we  show the cork positions superposed on a typical frame of a  magnetic difference image, where $x$ and $y$ denote the horizontal coordinates on the Sun.  Figures 2a and 2b show the $x$ and $y$ components of the velocity as a function of time for  two representative corks, whose initial locations are indicated by arrows in Figure 1.  \\section {KINK WAVE EXCITATION DUE TO FOOTPOINT MOTION}  \\def\\k{\\kappa}       \\def\\p{\\perp}   \\def\\b{\\beta}      \\def\\a{\\alpha} \\def\\g{\\gamma}  \\def\\V{{\\cal  V}} \\def\\Kv{{K$_{\\rm  2v}$}} \\def\\e{{\\rm e}}  \\def\\d{{\\rm  d}}   \\def\\ratio#1#2{{{#1}\\over{#2}}}  \\def\\e#1{{\\rm e}^{#1}}      \\def\\ex#1{{e}^{#1}}    \\def\\ei#1{{\\rm  e}^{{\\rm   i}#1}} \\def\\inv#1{\\ratio{1}{#1}}          \\def\\first#1#2{{{\\d#1}\\over{\\d#2}}} \\def\\firstp#1{{\\partial\\over{\\partial#1}}} \\def\\firstpp#1#2{{{\\partial#1}\\over{\\partial#2}}}  Let us assume that  the motion  of the G  band  bright points can  be taken as a proxy for the footpoint  motion of flux tubes at the base of the photosphere ($z=0$). Furthermore, for  simplicity, we assume that flux tubes are isothermal   and ``thin'' compared to the pressure scale height $H$.  It  is convenient to   work  in terms  of  the  ``reduced'' transverse displacement $Q_\\perp(z,t)=\\xi_\\perp  \\ex{-z/4H}$,   where $\\xi_\\perp$  is  the Lagrangian  displacement.  Once the  footpoint motion is  specified, the velocity ($\\dot Q_\\perp$) at any height $z$ and time $t$ can  be determined by the following expression (Paper I): \\begin{equation} \\dot     Q_\\perp(z,t)=  \\dot    Q_\\perp(0,t-z/c_\\k)   -k_\\k z\\int_0^{t-z/c_\\k}              \\dot        Q_\\perp(0,t_0) \\ratio{J_1(\\omega_\\k\\zeta_\\k)}{\\zeta_\\k}\\ \\d t_0\\ , \\end{equation} where $J_1$ is a Bessel function, $\\zeta_\\k=\\sqrt{(t-t_0)^2 -(z/c_\\k)^2},$ $c_\\k$ and $\\omega_\\k$ are, respectively, the propagation speed and cutoff frequency for kink waves, given by:  $$ c_\\k^2=\\ratio{2}{\\gamma}\\,\\, \\ratio{c_s^2}{1+2\\b}\\,,$$  $$\\omega_\\k^2  = {g\\over 8 H}{1\\over  1+2\\beta};$$  $c_s$ is the sound speed, $k_\\k=\\omega_\\k/c_\\k$,  $\\gamma=5/3$, $\\beta=8\\pi p/B^2$,  $p(z)$ is the gas  pressure inside the  tube and  $B(z)$ is  the  magnitude of  the vertical component of the magnetic field on the tube axis.  The energy flux ($F_\\k$) in a single flux tube is given by (Paper I): \\begin{equation} {F}_{\\k} = - \\ratio{2p_{\\e,0}}{\\b+1}\\dot Q_\\perp  \\Biggl(   \\firstpp{Q_\\perp}{z} +\\inv{4H} Q_\\perp \\Biggr)\\ex{-z/2H}, \\end{equation} where $p_{e,0}$ denotes the external gas pressure at $z=0$.  Figures 3a and 3b show the vertical energy flux in transverse waves versus time at a height $z=750$~km for two observed magnetic elements located initially at the coordinates shown above the figure. The left and right panels correspond to footpoint motions with the $x $ and $y $ components of the velocity, respectively.  We find that the injection of energy into the chromosphere takes place in brief and intermittent bursts, lasting typically 30 s, separated by longer periods (longer than the time scale for radiative losses in the chromosphere) with lower energy flux. The peak energy flux into the chromosphere is as high as  $\\sim 10^9$ erg cm$^{-2}$ s$^{-1}$ in a single flux tube, although the time-averaged flux is $\\sim 10^8$ erg cm$^{-2}$ s$^{-1}$.  This scenario for heating the magnetic network would produce strongly intermittent chromospheric emission consisting of brief intense flashes superimposed on generally a very low background.  However, in dense network regions the observed chromospheric emission is stably present and exhibits relatively low amplitude, long-period variations (e.g. Dam\\'{e}, Gouttebroze and Malherbe 1984; von Uexk\\\"{u}ll et al. 1989; Lites, Rutten and Kalkofen 1993).  Therefore, the above scenario cannot readily explain the observed persistence of emission from dense network regions.  A possible remedy is to consider the effect of high frequency motions.  \\section {SIMULATING HIGH-FREQUENCY MOTIONS}  The observations of G band bright points  were taken with  a time cadence of 23.5~s.     However, the possibility  of  random  footpoint motions occurring on a shorter  time scale cannot be  ruled out.  The solar convection has very high Reynolds number and is expected to be highly turbulent, involving a wide range of length and time scales. Indeed, spectroscopic observations of the solar granulation have shown enhanced broadening of spectral lines in intergranular lanes, indicating enhanced turbulence within these region (e.g. Nesis et al. 1992, 1994). Magnetic flux tubes interact with these turbulent flows within the lanes. We speculate that these interactions produce transverse displacements of magnetic field lines on length scales much less than the flux tube diameter ($\\sim 100$ km), i.e., the magnetic field inside a flux tube has fine structure on transverse scales $\\ell \\ll 100$ km. These perturbations likely generate Alfv\\'{e}n waves which propagate upward along the field lines and dissipate their energy higher up in chromosphere.  The velocity at small spatial scale can be estimated as follows.  Assuming the turbulence has a Kolmogoroff spectrum, the velocity of eddies on length scale $\\ell$ is given by $v( \\ell ) \\sim v_0 ( \\ell / \\ell_0 )^{1/3}$, where $\\ell_0$ is the outer scale of the turbulence, and $v_0$ is the velocity of the largest eddies. We assume that $\\ell_0$ is given by the width of an intergranular lane as determined from 3D simulations of the solar granulation (e.g., Stein and Nordlund 1998), and that $v_0$ is the convective velocity predicted by such simulations: $\\ell_0 \\sim$ 100 km and $v_0 \\sim$ 2.5 km s$^{-1}$. This turbulence model predicts that $v\\sim$~1~km~s$^{-1}$ on a length scale $\\ell \\sim$ 10 km, i.e., there is significant power in motions at this scale.  The time scale of such motions is $\\ell / v \\sim$ 10 s, which is significantly shorter than the time scale for granular buffeting.  In the following we simulate the effect of such high-frequency motions on the response of the flux tube (similar to the study of longitudinal wave excitation in flux tubes by Ulmschneider and Musielak 1998). For simplicity we ignore the small transverse scale of the motions and we describe the \"turbulence\" in terms of footpoint displacements of the flux tube.  We do this by superposing random motions, with a  rms amplitude of   1 km s$^{-1}$  and a zero mean  value, on the observationally determined velocity. The random motions vary on a time scale of 2.35 s.   Such high-frequency  motions (if present) would not have been detected  in the   present observations.  Hence, it  is not possible  {\\it a priori} to   rule them   out.   Results from such simulations are shown in Figures 4a and 4b for the $x$ component of the velocity as a function of time for  two representative corks whose initial locations are indicated above each panel, and in  Figures 5a and 5b for the corresponding vertical energy wave flux.   We note that the energy flux in upward-propagating transverse waves is much  larger    and  less intermittent than in  the case without high-frequency footpoint motions.  The energy flux  in the form of  high-frequency transverse waves appears adequate to continuously heat the chromosphere.  Taking a filling factor of 10\\% at a height $z=750$~km above the photosphere,  the predicted flux is approximately \\hbox{$2\\times 10^7$} erg cm$^{-2}$ s$^{-1}$,   which  is sufficient to balance the   observed radiative  loss of   the  chromospheric network ($\\sim 10^7$ erg cm$^{-2}$ s$^{-1}$, see Model F$^\\prime$ in Avrett 1985) ",
        "conclusions": "The energy flux in  upward propagating transverse waves has been shown for a few  representative magnetic elements. We should point out that we carried out similar calculations for several corks,  before selecting two as being typical of the large sample that we examined.  We  find that the injection of energy into the  chromosphere  occurs in  brief and intermittent bursts (lasting typically  30-60 s), separated  by longer periods (longer than the typical time scale  for radiative loss) with low energy flux.   The peak energy  flux into the chromosphere  is typically $10^{9}$ erg cm$^{-2}$ s$^{-1}$ in a single flux tube.  Observationally, such a scenario  for  heating of the magnetic network on a 1 arc sec scale would lead to high variability in Ca II K emission, contrary to observations.  However, the observations of G band bright  points used here do not have the time resolution and sensitivity necessary to detect motions on shorter time scales. We simulate such high-frequency motions by adding random velocity fluctuations with a zero  mean value and  an rms amplitude of 1~km~s$^{-1}$.  We  conclude that  for   transverse waves  to provide a  viable mechanism  for  {\\it sustained}   chromospheric   heating, the  main contribution  to the  heating must come   from high-frequency motions (periods 5-50 s).  The proposed high-frequency motions have very small spatial scale (10 - 100 km) and the transverse displacements involved are also very small (for example, a motion with period 50 s and velocity amplitude of 1~km~s$^{-1}$ has a displacement amplitude of only 8 km). Such length scales are well below the angular resolution of present-day solar telescopes.  Therefore, it is unlikely that such high-frequency motions will ever be detected in studies of the proper motions of photospheric flux tubes. Spectroscopic methods, such as measurements of spectral line broadening by unresolved turbulent flows, offer greater promise but will still require very high angular resolution."
    },
    "0004/astro-ph0004136_arXiv.txt": {
        "abstract": "We have made HI 21cm-line absorption measurements using the GMRT towards 15 directions in the Galaxy which are known to have high random velocity clouds as seen in the optical absorption lines of CaII and NaI. For the first time, in 6 out of these 15 directions we detect HI absorption features corresponding to the high random velocity optical absorption lines. The mean optical depth of these detections is $\\sim$ 0.08. ",
        "introduction": "Our picture of the interstellar medium, its structure and motion emerged from the optical absorption studies in the lines of CaII and NaI towards bright O \\& B stars. Soon after the 21cm-line from atomic hydrogen was discovered, attempts were made to detect the HI line from these clouds through emission measurements. Interestingly, although the HI emission velocity is clearly correlated with that of the optical absorption line at low random velocities, no HI was detected corresponding to the higher random velocity ($|v|$ $\\ge$ 10 km sec$^{-1}$) optical absorption lines. Recently, an HI absorption study was carried out (Rajagopal et al 1998a) using the VLA towards radio sources in the sky which are aligned very close ($\\sim$ a few arc mins) to the star. No HI absorption corresponding to the high random velocity optical absorption lines was detected down to an optical depth of 0.1. ",
        "conclusions": ""
    },
    "0004/astro-ph0004300_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cyclotron lines in accreting X-ray pulsar spectra result from the resonant scattering of X-rays by electrons in Landau orbits on the intense (\\aprx $10^{12}$\\,G) magnetic fields near the neutron star poles.  For this reason they are known as cyclotron resonance scattering features (CRSFs).  Because Landau transition energies are proportional to field strength ($\\rm E_{cyc}=12$\\,keV approximately corresponds to $B = 10^{12}$\\,G), CRSF energies give us our most direct measures of neutron star magnetic fields.  Other line properties, such as depths, widths, and presence of multiple harmonics, are strongly dependent on the details of the geometry and environment at the base of the accretion column.  CRSFs therefore give us a sensitive (if difficult to interpret) diagnostic of the accretion region.   In this paper, we summarize the \\rxte\\ measurements of CRSFs in 8 accreting pulsars.  The wide bandpass and modest resolution of the \\rxte\\ instruments make them ideal for measuring these generally broad features.  In particular, the high energy response of HEXTE provides a window for discovery of new lines not detectable with proportional counters such as \\ginga\\ or the PCA alone.  Some highlights of this work are: \\begin{list}{$\\bullet$}% {\\setlength{\\topsep}{0.in}% \\setlength{\\parsep}{0.in}% \\setlength{\\partopsep}{0.in}% \\setlength{\\itemsep}{0.5ex}}% \\item{New lines in the well known pulsars Cen~X-3 and 4U~1626-67 \\cite{hei99b,San98,Orl98}. } \\item{The discovery of more than two cyclotron line harmonics in 4U~0115+63 \\cite{Hei99}}. \\item{The \\rxte\\ picture of the correlation between cyclotron line energy and width.} \\end{list} ",
        "conclusions": "Table~\\ref{tab:fits} summarizes the cyclotron line parameters in eight accreting pulsars.  In general, pulsar spectra (including line parameters) vary with pulse phase.  For this reason, cyclotron lines are often best measured in spectra from limited pulse phase ranges. This is called ``pulse phase spectroscopy''.  Table~\\ref{tab:fits} indicates the phase relative to the pulse profile for which the given parameters apply. \\begin{table} \\caption{\\label{tab:fits}Summary of \\rxte\\ Cyclotron Line Measurements. 4U~0115+63 and Vela X-1 require multiple CRSFs in their spectra.  The surface B-fields assume a gravitational redshift of 0.3 for the emitting region.} \\label{t_crsf} \\begin{tabular}{lccccc}% Source  \t&    Energy   &    Sigma   &   Optical   & Pulse Phase  & Surface B Field\\\\ &  (keV)      &   (keV)    &   Depth     & & ($ 10^{12}$\\,G)  \\\\\\hline 4U 0115+63\t& $\\rm 12.40^{+0.65}_{-0.35}$\t&$\\rm 3.3^{+0.1.9}_{-0.4}$&$\\rm 0.72^{+0.10}_{-0.17} $& $\\rm2^{nd}$ Fall & 1.4\\\\ & $\\rm 21.45^{+0.25}_{-0.38}$\t&$\\rm 4.5^{+0.7}_{-0.9} $ &$\\rm 1.24^{+0.04}_{-0.06} $& $\\rm2^{nd}$ Fall & \\\\ & $\\rm 33.56^{+0.70}_{-0.90}$\t&$\\rm 3.8^{+1.5}_{-0.9}$& $\\rm 1.01^{+0.13}_{-0.14} $& $\\rm2^{nd}$ Fall & \\\\ 4U 1907+09\t& $\\rm 19.7 \\pm 0.1$\t\t& $\\rm2.6 \\pm 0.1$\t& $\\rm0.87 \\pm 0.05$\t& Average  & 2.2\\\\ 4U 1538-52\t& $\\rm 21.3\\pm0.3 $\t\t& $\\rm3.05 \\pm0.20 $\t& $\\rm0.86 \\pm 0.07$\t& Average  & 2.3\\\\ Vela X-1\t& $\\rm 23.7^{+0.4}_{-0.3}$\t& 5 (fixed)\t\t& $0.29^{+0.03}_{-0.04}$& Main Pulse & 2.3\\\\ & $\\rm 59.7\\pm3.7$\t\t& $12.6\\pm0.8$\t\t& $1.41^{+0.67}_{-0.61}$& Main Pulse & \\\\ Cen X-3\t\t& $\\rm 31.8\\pm 0.3$\t\t& $7.5\\pm 0.9$\t\t& $0.77^{+0.16}_{-0.11}$ & Peak & 3.5\\\\ GX 301-2\t& $\\rm 42.9^{+0.9}_{-2.6}$\t& $10.0^{+1.9}_{-2.3}$  & $0.8^{+0.7}_{-0.3}$\t&  Average & 4.8 \\\\ 4U 1626-67\t& $\\rm 39.3^{+1.9}_{-1.1}$\t& $6.4\\pm 0.8$ & $2.3^{+0.6}_{-0.4}$ & Peak & 4.4\\\\ Her X-1\t\t& $\\rm 41.0\\pm1.0$\t\t& $9.8\\pm0.5$ & $1.84\\pm0.05$ & Average& 4.6 \\\\ \\hline \\end{tabular} \\end{table}  \\noindent{\\bf New Lines in Cen~X-3 and 4U~1626-67\\hspace{1.5em}} Cen~X-3 and 4U~1626-67 are two of the earliest known accreting pulsars; however, it is only recently that CRSFs were discovered in their spectra \\cite{hei99b,Ddf99}.  Figure~\\ref{fig:cen} shows folded light curves, phase resolved count spectra, and fits to the peak phase spectra. In Cen X-3, the CRSF energy moves by about 20\\% with phase. This variation has been modeled as resulting from an offset of the magnetic dipole moment from the center of the neutron star \\cite{Bur99}.  In both sources, the CRSF appears strongest in the ``peak'' and ``fall'' spectra. This may also be related to an offset dipole.  \\noindent{\\bf 4U~0115+63: Multiple Cyclotron Line Harmonics\\hspace{1.5em}} 4U~0115+63 is a transient accreting pulsar in an eccentric orbit with a massive star, a so-called ``Be X-ray binary''.  It was also the first pulsar to show two cyclotron line harmonics \\cite{wsh83}. 4U~0115+63 underwent a two month long outburst in 1999 February--April.  Observations made with \\rxte\\ and \\sax\\ reveal for the first time more than two harmonics in a single pulsar \\cite{hei99c,san99a} (see Fig.~\\ref{fig:allsix}). \\begin{figure} \\centerline{\\includegraphics[width=2.5in,bb=121 133 529 517]{4u0115.ps} \\hspace{0.25in}\\includegraphics[width=2.5in,bb=128 133 529 517]{4u1907.ps}} \\vspace{0.2in} \\centerline{\\includegraphics[width=2.5in,bb=128 134 529 517]{4u1538.ps} \\hspace{0.25in}\\includegraphics[width=2.5in,bb=128 134 529 517]{velax1.ps}} \\vspace{0.2in} \\centerline{\\includegraphics[width=2.5in,bb=128 134 529 517]{gx301.ps} \\hspace{0.25in}\\includegraphics[width=2.5in,bb=128 134 529 517]{herx1.ps}} \\caption{\\label{fig:allsix} \\rxte\\ spectra of six accreting pulsars. Top panels show inferred incident spectra (solid lines) and PCA and HEXTE counts spectra. Middle panels show the ratio of the best fit model, which includes a CRSF, to the data.  Bottom panels show this ratio for a model which has no CRSF.  In the case of 4U~0115+63, which has three lines in this spectrum, we show residuals to models with three and two lines included. For line parameters, see Table~\\ref{tab:fits}.  } \\end{figure}  \\noindent{\\bf Correlation between Line Energy and Width\\hspace{1.5em}} The width of a CRSF depends on the temperature of the emitting region (kT), the line energy ($\\omega_B$) and the viewing angle ($\\theta$) with respect to the magnetic field as: $\\Delta \\omega \\propto \\omega_B (kT/mc^2)^{1/2} cos \\theta$ \\protect\\cite{Mes85}.  This predicts that a correlation between line energy and width should exist, provided that viewing angles and plasma temperatures do not vary too greatly from source to source. These \\rxte\\ results, as well as measurements by \\sax \\protect\\cite{Ddf99}, show a strong correlation, which supports this picture. \\begin{figure} \\centerline{\\includegraphics[angle=90,width=4in]{width_vs_e.ps}} \\caption{ CRSF width versus line energy for all the eight pulsars discussed here. A positive correlation between energy and with is apparent.  Note that for the three cyclotron lines in 4U~0115+63, the correlation does not appear to apply. } \\end{figure}  {"
    },
    "0004/hep-ph0004259_arXiv.txt": {
        "abstract": "We investigate a new mechanism for the cosmological QCD phase transition: inhomogeneous nucleation. The primordial temperature fluctuations, measured to be $\\delta T/T \\sim 10^{-5}$, are larger than the tiny temperature interval, in which bubbles would form in the standard picture of homogeneous nucleation. Thus the bubbles nucleate at cold spots. We find the typical distance between bubble centers to be a few meters. This exceeds the estimates from homogeneous nucleation by two orders of magnitude. The resulting baryon inhomogeneities may affect primordial nucleosynthesis. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004314_arXiv.txt": {
        "abstract": "Understanding how  the physical properties of galaxies  (e.g.  their spectral type or  age) evolve as a function  of redshift relies on  having an accurate representation of  galaxy spectral energy  distributions.  While it  has been known for some  time that galaxy spectra can be  reconstructed from a handful of orthogonal  basis templates, the  underlying basis is  poorly constrained. The limiting factor has been the  lack of large samples of galaxies (covering a wide  range in spectral type) with  high signal-to-noise spectrophotometric observations. To alleviate this problem we introduce here a new technique for reconstructing galaxy spectral energy  distributions directly from samples of galaxies  with  broadband   photometric  data  and  spectroscopic  redshifts. Exploiting the  statistical approach  of the Karhunen-Lo\\`eve  expansion, our iterative training procedure increasingly improves the eigenbasis, so that it provides better agreement with the photometry.  We demonstrate the utility of this approach by applying these improved spectral energy distributions to the estimation of photometric  redshifts for the HDF sample  of galaxies. We find that  in a  small  number of  iterations  the dispersion  in the  photometric redshifts estimator  (a comparison between predicted  and measured redshifts) can decrease by up to a factor of 2. ",
        "introduction": "\\label{intro}  With the introduction of large-format  CCDs, the study of galaxy evolution is undergoing a  renaissance. Statistically significant volumes  of the Universe can now be surveyed at high and low redshift with relatively small amounts of observing time.  In fact the rate  at which we can undertake deep, multicolor photometric surveys is now over 2 orders of magnitude larger than the rate at which we  can followup these observations spectroscopically.   As an example, the  Sloan Digital  Sky  Survey will  spend  85\\% of  its  five year  mission measuring the  redshifts of  one million objects  with a  unique, multi-fiber spectrograph, while in the remaining 15\\%  of the time it will gather 5 color photometry for 100 million objects.  This  rapid  growth in  multicolor  surveys has  led  to  the development  of techniques for  exploiting the  information contained within  the photometric data themselves (without  relying on followup spectroscopy). One  of the most successful of  these techniques has  been the estimation of  galaxy redshifts directly from  multicolor data (photometric  redshifts) \\citep{baum62}.  (For an  overview of  recent work  in this  field, see  \\citet{weymann99}.)  These techniques can be  broken down into two basic classes.   One class of methods use a training set of  galaxies with known photometry {\\it and} spectroscopic redshifts to derive an empirical  relation between the colors of galaxies and their redshifts \\citep{connolly95a, wang98}.  After calibration, the redshift estimator   can  predict   redshifts  for   objects  with   only  photometric observations. The advantage of this  method lies in its empirical nature; all the effects of dust and galaxy evolution that are present in the training set are implicit  within the  derived correlation.  The  drawback is that  it can only robustly predict redshifts for objects  that are similar to those in the original  training   set.   In  other   words  the  estimator  is   good  for interpolation, but not  for extrapolation to galaxies that  have redshifts or spectral types far from the range covered by the training set.  The  second   class  of  techniques,  template   based  photometric  redshift estimators (or spectral energy  distribution fitting), use model or empirical galaxy  spectral  energy  distributions  \\citep{koo85,  connolly95b,  gwyn96, soto99, benitez99}.   By minimizing the  difference between the  observed and synthetic  colors one  can find  the  most likely  type and  redshift of  the galaxy. This method is simple to  implement, does not require a training set, and  has  no  extrapolation  problems.   Its  main  limitation  is  that  the underlying spectral energy distributions (SEDs) of galaxies within the sample must  be well known.   Comparisons of  the colors  of galaxies  with spectral synthesis models \\citep{bc93} have shown that the modeling of the ultraviolet spectral  interval  for  galaxies  is  uncertain.  Whether  this  is  due  to uncertainties in the modeling  of the stars or due to the  effects of dust is as  yet  unclear.   Consequently,  photometric-redshift  estimates  are  most accurate when  we apply empirical spectral energy  distributions derived from observations  of local  galaxies, e.g.   \\citet{sawicki97}.   These empirical relations are,  however, constructed  from a small  number of  local galaxies that  have been  observed  in detail  and  there is  no  guarantee that  they represent the full distribution of galaxy types (particularly when we include the effects of evolution with redshift).  In this  paper we derive a method  that combines the advantages  of the above techniques without their drawbacks. We use a training set, not for deriving a direct  relationship between  colors and  redshifts but  instead to  build an optimal set of spectral templates. We  optimize for the shape of the spectral templates  to  give  the  best  match between  the  predicted  galaxy  colors (calculated  using  the  {\\it  spectroscopic}  redshifts)  and  the  observed colors. As a result  we derive a set of spectral templates  that are a better match  to the  SEDs of  the galaxies  in the  training set  than  the initial model/empirical templates. Our tests  show that these improved templates give a  tighter photometric  redshift relation  than  do the  original SEDs.   In  section   \\ref{inversion}  we  revisit   different  photometric  redshift estimator  techniques  and  we  address   the  question  of  how  to  recover information  beyond just redshift  from photometric  observations and  how to interpret  these results. We  introduce an  interpolation scheme  for optimal spectral  type determination.   In  section \\ref{algorithm}  we describe  our novel technique for creating eigenspectra in an iterative way, and in section \\ref{app} we show a test application to the HDF-N/NICMOS catalog, where there are around 150 spectroscopic  redshifts available.  We present the comparison of photometric redshift estimates based on the most commonly used \\citet{cww} spectra (hereafter CWW) and our empirically developed templates. ",
        "conclusions": "\\label{discussion}  Our technique has  been shown to be able to  derive rest-frame templates from photometry in an iterative way, improving step-by-step the correlation of the template based  colors to the actual measurements.   The statistical approach of the KL expansion provides a robust determination of the systematic changes in the  SEDs. After  just the initial  2-3 iterations, the  overall continuum shapes  of  the  SEDs  are  adjusted to  reproduce  the  photometry.   Higher resolution detailed features emerge in later iterations.  Consistent with our expectations, significant continuum changes occur both in the ultraviolet and infrared, where the original  templates are poorly constrained. The estimated SEDs based  on the trained eigenbases tend  to be redder in  general.  The UV part of  the spectra get  fainter and  the flux increases  in the IR  for the early type  galaxies (which  change the most).   The blue galaxies  have gone through less dramatic  modulation.  Clearly, the color spread  of galaxies is larger than one might anticipate from the CWW spectra.  The limitation on  our current application is the  relatively small number of objects with  accurate multicolor  photometry and spectroscopic  redshifts in the Hubble Deep  Field. Combining several observations into  one training set is  one way  to improve  the  robustness of  our results.   Our technique  is constructed to be able to deal  with many multicolor observations at the same time. Data of  various surveys can be incorporated  regardless of photometric systems they use, meaning not just broad but also narrow or intermediate band imaging.  This  flexibility of the algorithm  allows us to  simply merge data sets coming  from several instruments, where  the data set  up constraints on the SEDs  at different  wavelengths.  With  time we will  be able  to improve other  parts  of  the  SEDs.   New upcoming  instruments  will  provide  high signal-to-noise  ratio data in  different spectral  ranges, e.g.   the Galaxy Evolution   Explorer   (GALEX)  probing   the   far   and  near   ultraviolet (1300-3000\\AA) or the Space Infrared  Telescope Facility (SIRTF) and the Next Generation Space Telescope  (NGST) looking at the infrared  ranges.  In fact, the   algorithm  is  even   more  generic.    Eventually,  spectrophotometric observations can be naturally incorporated in the training mechanism.  There remains  one class of galaxy  for which this technique  is not optimal; high  redshift  galaxies for  which  we  cannot  easily obtain  spectroscopic redshifts can  not be optimized by  our procedure (i.e. we  lack the relevant training set).  As our technique  can seamlessly incorporate  theoretical and empirical data (weighting the samples  appropriately) if we have a reasonable understanding of  the stellar populations  we can account for  missing galaxy populations. Objects of extremely unusual  colors will not be well matched by any template  fitting procedure but  the goodness of  fit to these  data will enable us to identify these anomalous objects.   In this  paper we have presented  a new method to  reconstruct the rest-frame continuum spectra of galaxies directly from a catalog of broadband multicolor photometric  observations and  spectroscopic  redshifts.  In  addition to  an improved photometric  redshift estimation the application  of eigenspectra to this  problem enables the  classification of  galaxy spectra  to move  from a series of discrete (and somewhat ad hoc) classification types to a continuous parameterization.   With   the  help  of  the   optimal  subspace  filtering, non-physical spectral energy  distributions could be avoided. As  a result of the repair algorithm, the eigenbasis improved the capability to represent the observed photometric properties. This  improvement cannot be tested directly, since we  do not  know the real,  photometrically calibrated spectra  for the objects in our training catalog, but the photometric redshift estimation test with  the  repaired  spectra  gave  substantially  better  results  than  the estimation   using   model/empirical   templates   ($\\Delta_{\\rm   rms}=0.12$ vs. $\\Delta_{\\rm rms}=0.23$). New  photometric and spectroscopic surveys will extend the size  of the training catalog, making  possible to improve further the spectral templates and enrich them with finer details."
    },
    "0004/astro-ph0004122_arXiv.txt": {
        "abstract": "Optical data in the V-band gathered with the 8.2m ESO Very Large Telescope (VLT) at the radio interferometric position of PSR 1706-44 are presented. The pulsar is close to a bright star in projection and was not detected. The pulsar magnitude limit must be fainter than V=24.5 for a distance of $\\leq$ 2\\farcs0 from the bright star. In the outer gap model for an aligned rotor the optical flux should scale with the gamma-ray flux. For pulsars which emit pulsed gamma-rays but are not detected in the optical bands, the synchrotron cutoff frequency for the tertiary photons must be well below the optical frequencies and the magnetic and spin axes may be misaligned. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004408_arXiv.txt": {
        "abstract": "We report the discovery of three cool brown dwarfs which fall in the effective temperature gap between the latest L dwarfs currently known, with no methane absorption bands in the 1---2.5~$\\mu$m range, and the previously known methane (T) dwarfs, whose spectra are dominated by methane and water. The newly discovered objects were detected as very red objects in the Sloan Digital Sky Survey imaging data and have JHK colors between the red L dwarfs and the blue Gl~229B--like T dwarfs. They show both CO and CH$_4$ absorption in their near--infrared spectra in addition to H$_2$O, with weaker CH$_4$ absorption features in the H and K bands than those in all other methane dwarfs reported to date. Due to the presence of CH$_4$ in these bands, we propose that these objects are  early T dwarfs.  The three  form part of the brown dwarf spectral sequence and fill in the large gap in the overall spectral sequence from the hottest main sequence stars to the coolest methane dwarfs currently known. ",
        "introduction": "Brown dwarfs, gravitationally condensed objects whose masses are too low for equilibrium hydrogen burning, occupy the mass range between the lowest mass stars ($\\rm \\sim 0.07 ~ M_{\\odot}$; Burrows et al. 1997) and the giant extrasolar planets ($\\rm \\sim 0.01 ~ M_{\\odot}$; Marcy \\& Butler 1998).  The lowest mass stars and brown dwarfs only slightly cooler than M dwarfs are classified as spectral type ``L'' (Kirkpatrick et al. 1999, Mart\\'{\\i}n et al. 1997) and have T$_{eff}$ in the range $\\sim$1500---2000~K.  Determining if L--type objects are brown dwarfs is difficult, because the luminosities and effective temperatures of brown dwarfs are a function of both age and mass (e.g. Burrows et al. 1997). The first unambiguous brown dwarf, Gl~229B, was discovered as a companion to a nearby M dwarf by Nakajima et al. (1995). This object is cooler than any L dwarf and has CH$_4$ in its atmosphere implying $\\rm T_{eff}<$1300 K (e.g. Fegley \\& Lodders 1996) and a sub--stellar nature. The recent discoveries of methane dwarfs (tentatively given a ``T'' spectral clasification) in the field by Strauss et al. (1999), Burgasser et al. (1999, 2000), Cuby et al. (1999) and Tsvetanov et al. (2000) demonstrate that objects like Gl~229B can form singly.  A striking characteristic of known T dwarfs is the similarity of their spectra, which resemble those of the gaseous solar system planets and are very different from even the coolest L dwarfs. The L dwarf spectra are characterized by the disappearance of the TiO and VO bands (which are strong in M dwarf spectra), the presence of atomic alkali lines and CO bands, and increasing depth of the $\\rm H_2O$ bands as T$_{eff}$ decreases.  The T dwarfs are characterized by very deep $\\rm H_2O$ and $\\rm CH_4$ bands in the 1--3 $\\rm \\mu m$ region. Alkali metal features are present between the bands, but the 2.3$\\rm \\mu m$ CO bands are absent, due to reduced CO abundance and the overwhelming strength of the CH$_4$ absorption.  A large remaining gap in the stellar to gas giant planet spectral sequence lies between the previously observed types L and T.  In this paper, we report photometry and 0.8---2.5 $\\rm \\mu m$ spectroscopy of three very red, faint objects identified in the imaging data of the Sloan Digital Sky Survey (SDSS) and observed at the United Kingdom Infrared Telescope (UKIRT) and the Hobby-Eberly Telescope (HET). These objects have spectra intermediate between those of the previously known L and T dwarfs, showing the onset of CH$_4$ absorption at 1.6 and 2.2 $\\rm \\mu m$, while still retaining observable CO absorption at 2.3 $\\rm \\mu m$.  We propose that they represent the warm end of the ``T'' spectral sequence, yet to be defined in detail.   \\begin{deluxetable}{rrrrrrrrrr} \\footnotesize \\tablenum{1} \\tablecaption{Coordinates and Photometry} \\tablehead{ \\colhead{RA} & \\colhead{Dec} & \\colhead{$i^*$} & \\colhead{$z^*$}& \\colhead{J}& \\colhead{H} & \\colhead{K}& \\colhead{J} & \\colhead{J$-$H}  & \\colhead{H$-$K} \\nl \\multicolumn{2}{c}{(2000)} & \\multicolumn{2}{c}{(AB)} & \\multicolumn{3}{c}{(UKIRT--UFTI)} & \\multicolumn{3}{c}{(UKIRT--IRCAM)}\\nl } \\tablecolumns{7} \\startdata 05 39 51.99 & $-$00 59 02.0 & 19.04 $\\pm$0.02\\phm{)}& 16.73 $\\pm$0.01 & 13.85 $\\pm$0.03 & 13.04 $\\pm$0.03 & 12.40 $\\pm$0.03 & 13.94  & 0.97  & 0.53  \\nl 08 37 17.21 &$-$00 00 18.0 &(23.51 $\\pm$0.42) & 19.95 $\\pm$0.09 & 16.90 $\\pm$0.05 & 16.21 $\\pm$0.05 & 15.98 $\\pm$0.05 & 17.08  & 0.93  & 0.11  \\nl 10 21 09.69 & $-$03 04 20.1 &  (23.73 $\\pm$0.58)   & 19.28 $\\pm$0.05 & 15.88 $\\pm$0.03 & 15.41 $\\pm$0.03 & 15.26 $\\pm$0.05 & 16.12  & 0.74 & 0.08   \\nl 12 54 53.90 &$-$01 22 47.4 &  22.22 $\\pm$0.28\\phm{)} & 18.00 $\\pm$0.04 & 14.66 $\\pm$0.03 & 14.13 $\\pm$0.03 & 13.84 $\\pm$0.03 & 14.90 & 0.84  & 0.15 \\nl \\enddata \\tablecomments{ $i^*$ and $z^*$ are asinh magnitudes on the AB system (Lupton et al. 1999).  Zero flux corresponds to $i^*=23.89$ and $z^*=22.47$.} \\end{deluxetable} ",
        "conclusions": "The spectra in Figure 4 reveal a clear spectral sequence.  Thus SDSS 0837, SDSS 1254 and SDSS 1021 are examples of the sought after ``L/T transition'' objects. However, as suggested by Figure 4, they might more properly be regarded as examples of early T dwarfs, where the ``T'' spectral type is defined by the presence of CH$_4$ absorption in the H and K bands. The strongest short wavelength infrared band of CH$_4$ is the $\\nu$$_2$ band, centered at 3.3~$\\mu$m, and objects in which this band is present but the shorter wavelength  bands reported here are absent must exist over a narrow temperature range. However, measurements at 3.3~$\\mu$m are considerably more difficult from the ground than are H and K band measurements, and objects which show only the 3.3~$\\mu$m CH$_4$ band perhaps should be classified as L rather than T (Kirkpatrick et al. (1999) suggest that the T class be defined by the presence of CH$_4$ absorption in the K band).  As the three objects presented here are at the CO/CH$_4$ transition temperature they are likely to have T$_{eff} \\approx 1300$~K  (Fegley \\& Lodders 1996)  and therefore masses of 20---70~M$_{Jupiter}$ for ages in the range 0.3---5~Gyr (Burrows et al. 1997).  Figure 2 shows that SDSS easily identifies T dwarfs by their extremely red $i^*-z^*$ color, including the early T dwarfs identified in the present paper.  On the other hand, the J$-$K colors are similar to those of the common M dwarfs and these early T objects are thus very difficult to select on the basis of the near--infrared colors alone (although the later T dwarfs are blue and {\\it can} be selected this way).  Many more L and T dwarfs will be found in the SDSS imaging data.  \\footnotesize"
    },
    "0004/astro-ph0004252_arXiv.txt": {
        "abstract": "The 3rd EGRET Catalog of High-energy Gamma-ray Sources contains 170 unidentified sources, and there is great interest in the nature of these sources. One means of determining source class is the study of flux variability on time scales of days; pulsars are believed to be stable on these time scales while blazars are known to be highly variable. In addition, previous work has demonstrated that 3EG J0241-6103 and 3EG J1837-0606 are candidates for a new gamma-ray source class. These sources near the Galactic plane display transient behavior but cannot be associated with any known blazars. Although many instances of flaring AGN have been reported, the EGRET database has not been systematically searched for occurrences of short-timescale ($\\sim1$ day) variability. These considerations have led us to conduct a systematic search for short-term variability in EGRET data, covering all viewing periods through proposal cycle 4. Six 3EG catalog sources are reported here to display variability on short time scales; four of them are unidentified. In addition, three non-catalog variable sources are discussed. ",
        "introduction": "The Energetic Gamma-Ray Experiment Telescope (EGRET) on board the {\\it Compton Gamma-Ray Observatory} ({\\it CGRO}) began observing the high-energy gamma-ray sky in 1991. Since that time, EGRET has contributed enormously to the list of gamma-ray point sources; the third EGRET catalog (\\cite{Har99}, hereafter H99) contains 271 point sources exhibiting significant flux above 100 MeV. Besides a single solar flare, the Large Magellanic Cloud (LMC), and a probable association with a radio galaxy (Cen A), the identified sources are distributed among two established classes of high-energy gamma-ray emitters: pulsars and radio-loud blazars. 170 sources remain unassociated with known objects. There is great interest in the nature of the unidentified sources, whether they be pulsars, blazars, or representatives of new classes of gamma-ray sources.  The limited spatial resolution of EGRET often hampers source identification based solely on position; an alternate means of determining source class is the study of flux variability. In particular, pulsars are believed to be stable on short ($\\sim 1$ day) time scales (\\cite{Ram95}) while blazars are known to be strongly variable (e.g., \\cite{Mat97a}). Short-term and long-term variability studies have also led to the discovery of two candidates for a new gamma-ray source class: 3EG J1837-0606 (GRO J1838-04) (\\cite{Tav97}) and 3EG J0241-6103 (2CG 135+01) (\\cite{Kni97,Tav98}). These sources display transient behavior but cannot be associated with any known radio-loud blazars.  There are many instances of gamma-ray flares of AGN on short time scales. For example, 3C 279 (\\cite{Kni93,Weh98}), PKS 0528+134 (\\cite{Hun93}), 3C 454.3 (\\cite{Har93}), PKS 1633+382 (\\cite{Mat93}), PKS 1406-076 (\\cite{Wag95}), PKS 1622-297 (\\cite {Mat97a}), and CTA 26 (\\cite{Mat95,Hal97,Mat00a}) all show strong evidence of variability on time scales of 1-3 days. Many other AGN show variations on time scales of weeks to months (\\cite{von95,McL96,Muk97}). Although these many instances of short-term variability have been noted, no exhaustive survey of the EGRET database has been performed.  It is the purpose of this study to conduct a systematic and comprehensive survey of EGRET data from proposal cycles 1-4 in search of evidence of short-term flux variability above 100 MeV. The standard method of EGRET data analysis, maximum likelihood, was applied to short-duration (one-day and two-day) maps of gamma-ray intensity. The variability statistic $V$ was used as a guide for singling out those light curves which deviate appreciably from constant flux. For the unidentified sources that display evidence of flaring, Monte Carlo simulations were performed to determine the probability of such fluctuations assuming intrinsically nonvariable sources. Six catalog sources were found to exhibit variability on one- and two-day time scales: four are unidentified, one is a flat-spectrum radio quasars (FSRQ), and one is a BL Lac object. In addition, two transient non-catalog excesses and one non-catalog sub-threshold repeating source are discussed. The limitations of this study are then considered; a summary and concluding remarks follow. ",
        "conclusions": "The results of this work are summarized in Table 1. The first column indicates the source name by increasing Right Ascension; sources found in H99 are given the 3EG designation and non-catalog sources are given the GRO designation. The second column indicates alternate names, and the third column indicates source identification as FSRQ, BL Lac, or unidentified (UID). Column four indicates those viewing periods in which the source is found to be variable, column five gives the variability index for that viewing period, and column six indicates whether the source flared in that viewing period. The Monte Carlo probability for unidentified flaring sources is given in the seventh column.  The detections of short flares in 3EG J0222+4253 and 3EG J2006-2321 demonstrate that EGRET can detect short, intense flares in relatively weak catalog sources ($3.9\\sigma$ and $4.8\\sigma$ in the viewing periods in which the flares were seen) as well as in strong blazars. Therefore most EGRET sources, except bright AGN, are not characterized by strong variability on one- and two-day time scales. If intense flares from sources like 3EG J0222+4253 and 3EG J2006-2321 were common, this survey would have detected more of them.  The same conclusion is reached for sub-threshold sources. The detections of GRO J0927-41 and GRO J1547-39, plus the repeating sources, demonstrate that EGRET can find very weak transients from sub-threshold sources with this sort of systematic analysis. The fact that few such intense flares are seen in this study indicates that the gamma-ray sky is not filled with these kind of flares from sub-threshold sources.  The flaring behavior seen in 3EG J0222+4253 adds another puzzle to the issue of identification of this source.  The flaring is seen at energies between 100 MeV and 1 GeV, where the emission may be dominated by PSR J0218+4232. This would make this source unusual in two ways; it would be the only ms pulsar seen by EGRET and the only pulsar to show strong flaring.  The long observation of steady emission from PKS 0528+134 shows that, as seen by EGRET, blazars are not always variable even on time scales of three weeks. Therefore long episodes, free of intense flares from a given source, do not rule out the identification of that source as a blazar.  The observation of clear short-term variability in 3EG J2006-2321 indicates that this source may be a blazar, although its radio flux at 5 GHz is significantly less than other EGRET-detected blazars with similar peak gamma-ray brightness. The suggestion of variability of 3EG J1410-6147 may place it alongside 3EG J0241-6103 and 3EG J1837-0606 as a third candidate for a new gamma-ray class of variable Galactic sources.  In summary, the results of this study suggest that EGRET reached the threshold of an important new feature of gamma-ray astrophysics: the ability to measure short-term time variability.  The few examples found by EGRET show that such sources do exist and are not restricted to the blazar class.  Variability offers an effective tool for matching gamma-ray sources to objects seen at other wavelengths.  With its ability to detect much fainter sources, measure variability on time scales of hours, and detect smaller flux variations that EGRET, the proposed Gamma-ray Large Area Space Telescope (GLAST) should be able to take full advantage of gamma-ray source time variability to extend the results of the present study."
    },
    "0004/astro-ph0004064_arXiv.txt": {
        "abstract": "We explore whether supermassive non-baryonic stars (in particular boson, mini-boson and non-topological soliton stars) might be at the center of some galaxies, with special attention to the Milky Way. We analyze, from a dynamical point of view, what current observational data show, concluding that they are compatible with a single supermassive object without requiring it to be a black hole. Particularly, we show that scalar stars fit very well into these dynamical requirements. The parameters of different models of scalar stars necessary to reproduce the inferred central mass are derived, and the possible existence of boson particles with the adequate range of masses is commented. Accretion to boson stars is also analyzed, and a comparison with another non-baryonic candidate, a massive neutrino ball, which is also claimed as an alternative to the central black hole, is given. Both models are capable to explain the nature of the object in Sgr A$^*$ without invoking the presence of a singularity. One difficult issue is why the accreted materials will not finally produce, in a sufficiently long time, a black hole. We provide an answer based on stellar disruption in the case of boson stars, and comment several suggestions for its possible solution in neutrino ball scenarios. Finally, we discuss the prospects for the observational detection of these supermassive scalar objects, using the new generation of X-ray and radio interferometry satellites. ",
        "introduction": "During the last years, the possible existence of a single large mass in the Galactic Center has been favored as the upper bound on its size tightens, and stability criteria rule out complex clusters. Although it is commonly believed that this central mass is a supermassive black hole, it is not yet established, as we discuss below, on a firm observational basis.  The aim of this paper is to present an alternative model for the supermassive dark object in the center of our Galaxy, formed by self-gravitating non-baryonic matter composed by bosons. This kind of objects, so-called boson stars, are well known to physicists, but up to now, observational astrophysical consequences have hardly been explored. The main characteristics of this model are: \\begin{enumerate} \\item It is highly relativistic, with a size comparable to (but slightly larger than) the Schwarzschild radius of a black hole of equal mass. \\item It has neither an event horizon nor a singularity, and after a physical radius is reached, the mass distribution exponentially decreases. \\item The particles that form the object interact between each other only gravitationally, in such a way that there is no solid surface to which falling particles can collide. \\end{enumerate}  It is the purpose of this work to show that these features are able to produce a Galactic Center model which can be confronted with all known observational constraints, and to point the way in which such a center could be differentiated from a usual supermassive black hole.  \\subsection{Why scalar fields?}  Interesting models for dark matter use weakly interacting bosons, and primordial nucleosynthesis show that most of the mass in the universe should be non-baryonic if $\\Omega \\sim 1$. Most models of inflation make use of scalar fields. Scalar-tensor gravitation is the most interesting alternative to general relativity. Recent results from supernovae, which in principle were thought of to favor a cosmological constant \\citep{per99}, can as well be supported by a variety of models, some of them with scalar fields too \\citep{cel99}. Particle physicist expect to detect the scalar Higgs particle in the next generation of accelerators. Scalar dilatons appear in low energy unified theories, where the tensor field $g_{\\mu\\nu}$  of gravity is accompanied by one or several scalar fields, and in string effective super-gravity. The axion is a scalar with a long history as a dark matter candidate, and Goldstone bosons have also already inferred masses. Symmetry arguments, which once led to the concept of neutron stars, may force to ask whether there could be stellar structures made up of bosons instead of fermions.  In recent works, Schunck and Liddle (1997), Schunck and Torres (2000), and Capozziello, Lambiase and Torres (2000) analyzed some of the observational properties of boson stars, and found them notoriously similar to black holes. In particular, the gravitational redshift of the radiation emitted within a boson star potential, and the rotational curves of accreted particles, were studied to assess possible boson star detection.  The interest in observational properties of boson stars also led to investigate them as a possible lens in a gravitational lensing configuration \\citep{dab00}. Recent studies are putting forward the gravitational lensing phenomenon in strong field regimes \\citep{vir98,vir00,tor98a,tor98b}. This would be the case for boson stars, which are genuine relativistic objects.  This work is organized as follows. Sec. II is a brief summary of the main observational results concerning Sgr A$^*$, and the main hypotheses in order to explain it are outlined there. Sec. III analyzes what dynamical observational data show, and what kind of models can support them. Sec. IV gives the basic ingredients to theoretically construct scalar stars, shows mass and radius estimates, and study effective potentials and orbits of particles. In Sec. V we study the center of the Galaxy, show which are the stable scalar star models able to fit such a huge mass, and comment on the possible existence of boson particles with the required features. We also provide there an assessment of the disruption processes in boson star scenarios. Discussion and conclusions are given in Sections VI and VII. ",
        "conclusions": "We have shown that boson stars provide the basic necessary ingredients to fit dynamical data and observed luminosity of the center of the Galaxy. They constitute viable alternative candidates for the central supermassive object, producing a theoretical curve for the projected stellar velocity dispersion consistent with Keplerian motion, relativistic rotational velocities, and having an extremely small size.  Other singularity-free models were as well considered, as non-baryonic fermion stars (e.g. neutrino, or gravitino condensations). In this case, the object is sustained by its Fermi energy, while in the boson star case, it is the Heisenberg's Uncertainty Principle which prevents the system from collapsing to a singularity. Due to this fact, boson stars are genuine relativistic objects where a strong gravitational field regime holds.  This difference in the relativistic status of both objects is not trivial. While fermion neutrino balls are extended objects, boson stars mimic a black hole. Disruption processes can not happen in a fermion condensation, and it has the unpleasant consequence of not providing with a straightforward mechanism by which stars could be diverted from the center, and through which finally avoid the formation of a massive black hole inside the condensate.  The formation of boson stars, neutrino balls, and black holes can all be competitive processes. Then, it might well be that even if we discover that a black hole is in the center of the galaxy, others galaxies could harbor non-baryonic centers. In the case of boson stars, only after the discovery of the boson mass spectrum we shall be in position to determine a priori which galaxies could be modeled by such a center. Observations of galactic centers could then suggest the existence of boson scalars much before than their discovery in particle physicists labs."
    },
    "0004/astro-ph0004228_arXiv.txt": {
        "abstract": "The publicly available EIS-DEEP optical-NIR data for the AXAF (Chandra) Deep Field have been used to construct samples of Extremely Red Objects (EROs) using various single-band and multi-band color criteria.  In this work we define as EROs objects with colors consistent with passively evolving elliptical galaxies at z $\\geq$ 1. The EROs surface densities we derive are intermediate between previous published values, emphasizing again the need for larger survey areas to constrain the effects of possible large-scale structure.  Although various single-color selected samples can be derived, the EROs sample selected using $R-Ks > 5, I-Ks > 4, J-Ks > 1.8$ jointly is likely to contain the highest fraction of passively evolving luminous field elliptical galaxies at z$\\geq$1, or conversely, the lowest fraction of lower redshift interlopers.  The surface density of this multi-band selected EROs sample is consistent with the conclusion that little or no field elliptical volume density evolution has occurred in the redshift range 0 $>$ z $>$ 1.5.  However, extensive spectroscopic followup is necessary to confirm this conclusion. ",
        "introduction": "\\label{sec:intro}  The formation history of field elliptical galaxies (i.e. elliptical galaxies located outside of massive galaxy clusters) remains very uncertain.  Although most stellar population indicators point towards star formation at z $>$ 2, it remains unclear when and how these stars were assembled into their current dynamical structures.  Two formation scenarios with distinct observational predictions continue to be debated.  Each one allows for star formation to take place at high redshift, i.e.  for the present epoch stellar populations of field ellipticals to be mostly old and metal-rich. However, the first scenario, based on the {\\it rapid dissipative collapse} model, postulates that the stars were formed and assembled into elliptical galaxies in one rapid collapse event at high (z $>$ 2) redshift (Eggen, Lynden-Bell, \\& Sandage 1962; Larson 1974; Tinsley \\& Gunn 1976).  In this ``top-down'' scenario field ellipticals have evolved quiescently since their formation, and therefore the number of massive/luminous field elliptical per co-moving volume should be constant at all redshifts between now and their formation epoch.  On the contrary, the opposing {\\it hierarchical merger} scenario assumes that ellipticals were assembled from previously formed sub-units, and that this process has continued until the current epoch (see, e.g. Toomre 1977; White \\& Rees 1978; Blumenthal et al.  1984).  In this ``bottom-up'' scenario, the number of massive/luminous elliptical field galaxies per co-moving volume increases over time.  These seemingly distinct and easily tested observational predictions have in practice been difficult to resolve at relatively low redshift. Based on the Canada-France Redshift Survey, Lilly et al.  (1995) concluded that there has been no significant number evolution for bright, early-type galaxies since z $=$ 1. This result was questioned by Kauffmann, Charlot, \\& White (1996), but it has been confirmed by Schade \\etal (1999), who took advantage in their study also of the availability of detailed morphological information derived from HST images.  Comparing the co-moving number density of field elliptical galaxy at z $>$ 1 to the current epoch number density would be more fruitful because the longer time baseline creates a larger (and therefore easier to detect) number density difference in the hierarchical model.  Under the rapid dissipative collapse model, of course, the co-moving number density would remain constant.  The availability of large format near-IR arrays has created the opportunity to search for the luminous field ellipticals at high redshift needed to carry-out this test.  These searches are based on a very simple concept: once the 4000 \\AA\\ break is redshifted beyond the I-band, the observed optical-IR colors of a quiescently evolving elliptical become very red (R--K $>$ 5; I--K $>$ 4).  In terms of brightness, a typical current epoch field elliptical redshifted to z $>$ 1 will have an observed K of $\\approx$ 19 -- 21 (Moustakas et al. 1997).  Objects with these extreme properties (faint K-band magnitudes, extremely red optical-NIR colors) are now commonly known as Extremely Red Objects (EROs).  Of course, EROs imaging surveys only measure EROs surface density. Determining the field elliptical galaxies volume density ultimately requires a combination of morphological and spectroscopic information to eliminate Galactic and extragalactic interlopers and to determine the redshifts of high redshift field elliptical galaxy candidates. Although some EROs have been found to have spectral properties consistent with young (3 - 4 Gyr) elliptical galaxies (\\eg Cowie \\etal 1996; Dunlop \\etal 1996; Spinrad \\etal 1997; Cohen \\etal 1999), spectra and/or exact redshifts are available for only a very small fraction (10 -- 15\\%) of detected EROs.  Thus, it is possible that some EROs are in fact dust-enshrouded galaxies at lower redshift (as discussed, e.g., by Graham \\& Dey 1996; Cimatti \\etal 1997) or foreground lower main sequence and brown dwarf stars, although the latter possibility is extremely unlikely given the apparent surface density of such Galactic objects (Moustakas et et.  1997; Cohen \\etal 1999).  Despite such contamination problems, observed EROs number surface density should be able to place a stringent upper limit on luminous field elliptical galaxy number volume density at z $>$ 1.  \\begin{table*} \\caption{Published Deep EROs Survey Results} \\label{tab:surveys} \\begin{tabular}{l c c c c c c} \\hline \\hline \\hfil Reference \\hfil& Color & Area & $K_{\\rm lim}$ & N(EROs) & N(EROs)/deg$^2$ \\\\ \\hfil (1) \\hfil & (2) & (3) & (4) & (5) & (6) \\\\ \\hline Djorgovski et al. 1995  & i--K & 3  & 20 &  1 & 1200 $\\pm$ 1200  \\\\ Cowie et al. 1996       & I--K & 26 & 20 & 19 & 2611 $\\pm$  600  \\\\ Moustakas et al. 1997   & I--K & 2  & 22 & 11 & 19800$^{\\mathrm{a}}$ $\\pm$ 6000 \\\\ Moustakas et al. 1997   & I--K & 2  & 20 &  3 & 5400 $\\pm$ 3100  \\\\ Barger et al. 1999      & I--K & 62 & 20 & 16 &  929 $\\pm$  230  \\\\ Cohen et al. 1999       & R--K & 15 & 20 & 19 & 4560 $\\pm$ 1050  \\\\ Thompson et al. 1999    & R--K &154 & 20 &289 & 6750 $\\pm$  400  \\\\ This work               & I--K & 23 & 20 & 20 & 3130 $\\pm$  700  \\\\ This work               & R--K & 43 & 20 & 22 & 1842 $\\pm$  390  \\\\ This work &RIJK$^{\\mathrm{b}}$ & 23 & 20 & 10 & 1565 $\\pm$  495  \\\\ \\hline \\hline \\end{tabular}  {\\bf Columns:} (1) Reference; (2) Color Criteria (R--K $>$ 5 or I--K $>$ 4); (3) Area Surveyed (square arcminutes); (4) K-band EROs Limiting Magnitude; (5) Number of EROs Found; (6) Implied Number of EROs per square degree, with associated Poissonian uncertainty. \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Note the different limiting magnitude associated with this EROs density. \\item[$^{\\mathrm{b}}$] This is a multi-color selection criterion, discussed in Sect. 4. \\end{list}  \\end{table*}  Unfortunately, published results from such surveys do not present a consistent picture of EROs surface density, let alone inferred volume density.  As Table~\\ref{tab:surveys} illustrates, measured EROs surface densities vary by factors of 2 -- 6.  Clearly, a consensus result has not been reached yet.  However, given that most of these surveys cover very small areas, it is not impossible that the observed surface density variations are caused by large-scale structure at high redshift. Additional samples, and preferable over larger areas, are still needed.  In this first in a series of three papers, we present an analysis of the optical-NIR images obtained of the AXAF (Chandra) Deep Field as part of the ESO Imaging Survey (EIS; Rengelink et al. 1999).  These data cover an area of approximately 43 square arcmin, comparable in spatial area to the largest EROs surveys already published.  The EIS data have been used to construct EROs samples using a variety of magnitude and color criteria.  In this paper, we review the EIS dataset (Sect.~\\ref{sec:data}) and then present a number of possible EROs samples using different selection criteria (Sect.~\\ref{sec:samples}).  The sample most likely to represent field elliptical galaxies at z $>$ 1 is then discussed in Sect.~\\ref{sec:ellipticals}.  Our results are summarized in Sect.~\\ref{sec:summary}.  In the next two papers, we will present an analogous analysis of the public EIS-DEEP HDF-S dataset (da Costa et al.  1999) and of a private near-IR dataset obtained as part of two independent EIS Cluster candidate followup surveys.  The total area surveyed by this joint study will be $\\sim$ 400  square arcmin, i.e.  larger than any other published EROs imaging survey at a comparable K-band limiting magnitude. ",
        "conclusions": "\\label{sec:summary}  We have used the publicly available EIS-DEEP optical-NIR data for the AXAF (Chandra) Deep Field to construct samples of Extremely Red Objects (EROs) using various color criteria consistent with passively evolving elliptical galaxies at z $>$ 1.  Our conclusions can be summarized as follows:  \\begin{enumerate}  \\item{} The surface density of EROs is still far from being reliably determined, and there are significant differences in different surveys.  The EROs surface densities derived here lie between the values derived for other studies.  A larger, and preferably non-contiguous, area should be used to better constrain EROs surface density.  \\item{} Just one color measurement is not enough to select accurately likely $z > 1$ elliptical galaxies.  However, if obtaining more than one color is impossible or undesirable for some reason, our results suggest that $R-K$ is the most efficient color to use.  We will study this issue further in our next paper.  \\item{} We have used multi-band EROs selection criterion to construct a sample of likely $z > 1$ passively evolving elliptical galaxies.  Our derived surface density is consistent with little or no volume density of such objects since z $\\sim$ 1.5.  This result is tenuous, as well as being based on very simple arguments, and requires spectroscopic followup to confirm the nature of the multi-band selected EROs population.  \\end{enumerate}"
    },
    "0004/astro-ph0004158_arXiv.txt": {
        "abstract": "We investigate the model of the disc/corona accretion flow around  the black hole. Hot accreting advective corona is described by the two-temperature plasma in pressure equilibrium with the cold disk. Corona is powered by accretion but it also exchanges energy with the disk through the radiative interaction and conduction. The model, parameterized by the total (i.e. disk plus corona) accretion rate, $\\dot m$ and the viscosity parameter, $\\alpha$, uniquely determines the fraction of energy released in the corona as a function of radius and, in particular, the transition radius to the single-phase flow.   Self-consistent solutions with the mass exchange between phases display radial dependence of the parameters qualitatively different from the 'static' case, without the mass exchange. Corona covers the entire disk. The character of the radial dependence of the fraction of energy dissipated in the corona is qualitatively different for low and high total accretion rate.  If the total accretion rate is low, the corona becomes stronger towards the central object, and finally the disc completely evaporates, changing the accretion pattern into the single hot advection-dominated accretion flow (ADAF). For intermediate accretion rates the reverse process - condensation - becomes important, allowing possibly for a secondary disc rebuilding in the innermost part of the system. High accretion rates always prevent the transition into ADAF, and the cold disk extends down to the marginally stable orbit.  The transition radius, $r_{tr}$, between the outer, two-phase flow and the inner, single-phase, optically thin flow, is equal to $4.51 \\dot m^{-4/3} \\alpha_{0.1}^7 R_{Schw}$ for $\\dot m <6.9 \\times 10^{-2} \\alpha_{0.1}^{3.3}$ and then contracts to the marginally stable orbit in a discontinuous way above this critical value of $\\dot m$..  This model reproduces all characteristic luminosity states of accretion black hole without any additional ad hoc assumptions. In particular. the mechanism of the disk evaporation leads to a new, almost horizontal branch on the accretion flow's stability curve (i.e. the dependence of accretion rate on surface density) at the critical accretion rate. This branch, together with the upper, advection dominated branch for optically thick disks, form boundaries for the time evolution of unstable, radiation pressure dominated disk. Therefore the disk at high accretion rates, corresponding to Very High State in GBH and perhaps to Narrow Line Seyfert 1, and quasar stage may oscillate between the disk dominated state and the evaporation branch state, with only a weak contribution from the cold disk emission. The position of this branch for $\\alpha = 0.1$ with respect to the gas pressure dominated branch is consistent with the presence of only weakly variable High State in GBH and the absence of a similar state in AGN: all the quasars vary considerably if monitored in timescales of years. We also suggest a new interpretation of the Intermediate State, consistent with the presence of the strong reflected component. ",
        "introduction": "\\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig1.ps} \\caption{The crossection through the central region of accretion flow into a black hole.} \\label{fig:pict} \\end{figure}  There are strong observational evidence for the existence of the hot plasma $\\sim 10^9$K in the neighborhood of the accretion disc in active galactic nuclei (AGN) and galactic black holes (GBH). Soft disc photons are Compton upscattered by this optically thin plasma and form an approximately power law hard X-ray spectral component. The presence of the Compton reflection hump and the cold iron $K_{\\alpha}$ line in the hard spectrum is a signature of strong illumination of cold disc matter by X-ray radiation (for a review see Mushotzky, Done \\& Pounds 1993, Tanaka \\& Lewin 1995, Madejski 1999).  The location of the hot plasma  is still under debate, but one of the attractive models postulated by Liang \\& Price (1977) and Paczy\\' nski (1978) introduces hot dissipative corona above accretion disc following the example of the solar corona. via magnetohydrodynamic waves, jets or other nonthermal energy transport, and then thermalized in corona. But such ways of energy transport occurred to be inefficient, only $10 \\%$ of energy generated in the disc can be carried out towards corona. Within the frame of such geometry, observations seem to require that in some sources and/or some spectral states the innermost part of the disk is either disrupted and replaced by an optically thin flow or there is a strong coronal outflow (e.g. Esin, McClintock \\& Narayan 1997, Zdziarski, Lubi\\' nski \\& Smith 1999, Beloborodov 1999). Therefore, it is important to determine whether any of such effects are expected on the basis of the physical description of the accretion flow (Narayan 1997).  The disc/corona models were recently investigated following two  basic lines.  The first group of models does not explain the mechanism which heats up the plasma  in the disk/corona part of the flow. The fraction of energy generated in the corona $\\fcor$ is treated as a free parameter of the model (Haardt \\& Maraschi 1991, Svensson \\& Zdziarski 1994) or changes according to the artificial formula being inversely proportional to  the distance form a central object (Esin et al. 1997). By changing $\\fcor$ and the geometrical shape of the hot region (continuous, patchy or clumpy corona), such models can be fitted to the observed X-ray spectra (for review see Pountanen 1999). The main conclusion from those models is that static corona agrees with observations only when it does not cover the whole disc. However, more accurate considerations of the reflection, taking into account the optical depth of the Compton heated zone (R\\' o\\. za\\' nska \\& Czerny 1996, Nayakshin, Kazanas \\& Kallman  1999) or an outflow (Beloborodov 1999) may influence this result to some extent.  The second group of models follow Shapiro, Lightman \\& Eardley (1976) solution when hot optically thin flow exists in the inner part of the cold disc. They assume both cold and hot flow to be driven via accretion and depending on the model they take into account mechanical energy transport in the radial direction: like advection (Ichimaru 1977;  Narayan \\& Yi 1994, 1995a, 1995b; Abramowicz et al. 1995; Chen et al. 1995;  for review see Narayan, Mahadevan \\& Quataert 1998  and Kato, Fukue \\& Mineshige 1998), turbulent diffusion (Honma 1996), hydrodynamical inflow/outflow (Narayan \\& Yi 1994, 1995a; Blandford \\& Begelman 1999) and/or in the vertical direction: thermal conduction, evaporation  or coronal wind (Meyer \\&  Meyer-Hofmeister 1994 hereafter MMH94, Meyer-Hofmeister \\& Meyer 1999 Liu et. al 1999, Turolla \\& Dullemond 2000).  In both groups of models the crucial thing is the transition between the hot and the cold phase. In general picture such a transition  occurs in two directions: vertical transition when corona accretes above a cold disc, and radial transition when disk/corona system is replaced with a single optically thin flow close to the black hole (Fig.~\\ref{fig:pict}). Therefore, as was pointed by MMH94, in general picture both radial and vertical structure of the hot and cold phase should be considered.  In this paper we present the simplest disc/corona model which is the combination of two groups of models described above. The similar model was presented by Esin et al. (1997) with better description of ADAF region, and with non-local radiative interaction between the hot  flow and cold  disc. Nevertheless, they were not able to determine self consistently the transition between disc/corona flow, since they did not calculate its vertical structure. Also MMH94 stressed that thermal conduction and mass exchange between disc and corona in vertical direction are important do determine the  transition. However, in their model (Meyer-Hofmeister \\& Meyer 1999 Liu et. al 1999) the radiative interaction between phases was neglected, an assumption of a single-temperature plasma was used, and the conductive flux on the bottom of the hot phase was artificially taken from the solar corona.  Our model takes into account radiative coupling between phases, thermal conduction and evaporation in vertical direction, and radial energy transport via advection. The mass exchange rate is calculated self-consistently from the physical conditions in the transition zone. Hot corona is treated as two temperature plasma and, similarly to the disc, it is powered by accretion (Witt, Czerny \\& \\.Zycki 1997; hereafter WC\\. Z97, Esin et al. 1997).  On each radius the vertical structure is calculated depending on comparison  of accretion and vertical evaporation timescales. When accretion time scale is long enough, the mass exchange can be neglected and the 'static' equilibrium is achieved. In such a case, for vertical transition we use the condition based on considerations of Krolik (1998), Dullemond (1999), and R\\' o\\. za\\' nska \\& Czerny (2000) which is relevant to the dissipative corona. When accretion time scale is short  mass exchange cannot be neglected, and we determine the evaporation rate through the disc/corona transition.  Vertical properties are used to compute the radial trends solving continuity equation. The aim of this paper is to determine the radial dependence of the fraction of energy liberated in the corona for a given accretion rate, viscosity and a mass of the black hole and to check when the disk evaporation leads to a transition from disc/corona flow to the single optically thin flow. The radial transition occurs at the radius when all the mass starts  to be carried by the corona.  Observational consequences of our model occur in the dependence of the transition radius on accretion rate, which can be seen in different shapes of the reflected component (Done \\& \\. Zycki 1999). Also the evaporation leads to the new stable branch on the accretion flow's stability curve ( i.e. $\\dot m ( \\Sigma)$), allowing for new evolution cycle to exist. If accretion rate is too low optically thick disk can evaporate completely in the innermost parts and the position of transition radius moves outward. This can explain the different luminosity states in AGNs and GBHs  In Sec \\ref{sec:mod} we present the model. Calculation are shown in Sec. \\ref{sec:res}, and we discuss our results in Sec. \\ref{sec:dis}. Evaporation rate is determined in Appendix A. ",
        "conclusions": "\\label{sec:dis}  \\subsection{Comparison to other models}  \\subsubsection{Underlying assumptions}  Our approach to modeling the disk/corona structure follows the MMH94 attitude. We also model the energy dissipation within a corona as due to the accretion flow, we take the transition between the disk and the corona which allows to determine the fraction of energy generated in the corona as a function of radius as well as the transition radius to the optically thin flow as a function of the mass of the black hole, accretion rate and the viscosity parameter in the corona.  There are, however, major important differences between our model and MMH94: (i) our corona is basically a two-temperature medium with energy dissipation due to ion collisions and the Coulomb transfer of energy from ions to electrons, as in SLE and ADAF models (ii) apart from bremsstrahlung we allow also for Compton cooling (iii) we determine the conductive flux at the basis of the corona for the adopted heating/cooling mechanism instead of using the scaled prediction from solar corona. This last modification is responsible for the dependence of the accretion rate in the corona on the viscosity parameter in the case of our model and the lack of such dependence in MMH94 and subsequent papers (Liu, Meyer \\& Meyer-Hofmeister 1995, Liu et al. 1999, Meyer, Liu \\& Meyer-Hofmeister 2000). The predicted dependence of the coronal accretion rate on the radius is qualitatively similar in both papers in the case of low accretion rate but not in the case of large accretion rate when the condensation, neglected by MMH94, is important. Our approach to conduction is therefore the same as of Dullemond (1999) but he did not consider a two-temperature plasma and Compton cooling.  Two-temperature accreting corona was discussed by WC\\. Z97 but their condition of disk/corona transition based on analogy with Compton heating medium was not accurate numerically and missed an important insight on corona formation in the innermost and outermost  parts of the flow.   \\subsubsection{Transition radius to ADAF-type flow} \\label{sect:transi}  The transition from an outer optically thick disk flow to an inner optically thin ADAF was predicted by a number of models. In some models the transition was based on ADAF principle (i.e. the transition happens whenever an ADAF solution is possible), but in other cases the necessity of the transition resulted from the model itself (MMH94 and applications, our model). It is therefore interesting to compare quantitatively those estimates as they may in principle be estimated observationally.   Condition based on the strong ADAF principle, i.e. on the absence of ADAF solutions:  \\begin{equation} \\rtr = 1.9 \\times 10^4 \\dot m^{-2} \\alpha_{0.1}^4 \\RS \\end{equation} (Abramowicz et al. 1995, Honma 1996, Kato \\& Nakamura 1998).  Approximation to the results of Esin et al. (1997) for $\\alpha = 0.25$ and $\\beta = 0.5$: \\begin{equation} \\renewcommand{\\arraystretch}{1.5} \\begin{array}{l@{\\hspace{1cm}}c} \\rtr  = \\times 10^4 \\RS  &   \\dot m  <   0.084   \\\\ \\rtr  =  30 \\RS  & \\dot m  \\approx  0.084   \\\\ \\rtr  =  3 - 30 \\RS & 0.084 <  \\dot m  < 0.092  \\\\ \\rtr  =  3 \\RS  & \\dot m  >  0.092 \\end{array} \\renewcommand{\\arraystretch}{1} \\end{equation}   Liu et al. (1999) gives:  \\begin{equation} \\rtr =18.3 \\dot m^{-1/1.17}  \\RS \\end{equation}  Our results: \\begin{equation} \\renewcommand{\\arraystretch}{1.5} \\begin{array}{l@{\\hspace{0.5cm}}l} \\rtr  =  4.51 \\dot m^{-1.33} \\alpha_{0.1}^7 \\RS & \\dot m < 6.92 \\times 10^{-2} \\alpha_{0.1}^{3.3}  \\\\ \\rtr  =  3 \\RS  &\\dot m >  6.92\\times 10^{-2}\\alpha_{0.1}^{3.3} \\end{array} \\renewcommand{\\arraystretch}{1} \\label{por:tran} \\end{equation}   All those conditions are roughly in agreement with observational requirement that transition between the high and low state happens at  luminosity $10^{37} $ erg s$^{-1}$ (Tanaka 1999). The accurate dependence of the transition radius on the accretion rate can in principle be determined from the shape of the reflected component (Done \\& \\. Zycki 1999) in X-ray novae but the errors are considerable and the presence of the hot Compton heated skin on the disk surface further complicates the interpretation of the data. However, high precision spectroscopy from new X-ray satellites should give stronger constraints on the available models.   \\begin{figure} \\epsfxsize=8.8cm \\epsfbox{fig17.eps} \\caption{The relation between the accretion rate and the surface density of the disk/corona system at $10 \\RS$ for the viscosity parameter $\\alpha = 0.1$ in case of irreversible  ADAF transition for an AGN ($M= 10^8_{\\odot}$, continuous line) and a GBH ($M= 10_{\\odot}$, dashed line). Stable solutions (positive slope of the curve) are expected for intermediate accretion rates in the case of a small mass of the black hole but this region is very narrow for massive black holes.} \\label{fig:mdotsig2} \\end{figure}  \\subsection{Observational consequences}  \\subsubsection{Luminosity states}  The transition from optically thick flow to ADAF, i.e. the accretion rate at the evaporation branch marks the transition between the hard state and soft state in accreting binaries, according to a number of suggestions (e.g. Narayan 1996, Esin et al. 1997, Zdziarski, Lubi\\' nski \\& Smith 1999). When the accretion rate is high enough and two-phase flow exists up to the marginally stable orbit then the soft flux is always dominant, cooling of corona is efficient and produces soft X-ray spectrum. Such a solution is stable in GBH for moderate accretion rates and may correspond to the soft (high) state. If the accretion rate is very high the disk is unstable and evolving between the evaporation branch and advection-dominated optically thick branch - such a parameter range may be appropriate for a Very High State, with its complex spectral behaviour (e.g. Takizawa et al. 1997, Revnivtsev et al. 2000). On the other hand, when the accretion rate is very low, the innermost part of the disk is an optically thin ADAF, dominating the spectrum, so such a case corresponds to the Low (or Hard) State.  In our model, the position of horizontal evaporation branch does not depends on the mass of the black hole (Fig. \\ref{fig:mdotsig2}). Nevertheless the turning point between the unstable radiation pressure branch and the stable  evaporation branch is different for AGNs and GBHs and the possibility of formation of the High State in AGN is strongly suppressed.  Since most of the energy generates at about 10 $\\RS$ we can roughly identify the accretion rate ranges for various luminosity states on the basis of Fig~\\ref{fig:mdotsig2}. Assuming $\\alpha=0.1$ we obtain: \\begin{equation} \\renewcommand{\\arraystretch}{1.5} \\begin{array}{l@{\\hspace{1cm}}cc} State & GBH  &  AGN   \\\\ Very~ High ~State  & \\dot m  > 0.1  &  \\dot m  > 0.07  \\\\ High ~State &   0.07 <  \\dot m  < 0.1  & \\\\ Low ~State   & \\dot m  <  0.07 & \\dot m  <  0.07 \\end{array} \\renewcommand{\\arraystretch}{1} \\end{equation}  However, we stress that those limits depend significantly on the adopted value of the viscosity parameter $\\alpha$.  \\subsubsection{The nature of the Intermediate State in GBH}  The transition of X-ray novae from soft state with the dominant disk emission to a hard state with the dominant thin plasma (e.g. ADAF) contribution proceeds through the Intermediate State. The analysis of the Nova Musca revealed that in the intermediate state the reflection component is seen with an amplitude of order of 1 while the kinematical determination of the inner disk radius from Doppler shifts of the reflected component show disk disruption (\\. Zycki, Done \\& Smith 1998). We can interpret such a state within the frame of our model as corresponding to an accretion rate for which a strong corona forms at intermediate radii. In this case the reflection is strongly biased towards larger radii instead of towards innermost disk radius. Also very high values of the ionization parameter may be consistent with the irradiating source being directly above an accretion disk. Our interpretation is therefore different from that of Esin et al. (1997) since in their approach the intermediate state corresponds to a disrupted disk with an innermost ADAF flow instead of the two-phase disk/corona flow.  \\subsubsection{Variability}  The lack of High State in AGN, as predicted by our model for the viscosity parameter $\\alpha = 0.1$ may offer an explanation why the observed variability properties of AGN and GBH are not the same. GBH, when in their High State, dominated by disk emission, are only weakly variable in X-rays in the timescales of about a second, with rms amplitude of order of a percent (e.g. Takizawa et al. 1997). Spectrally, such states are basically similar to quasar spectra. However, quasars are strongly variable, with rms amplitude in the optical band dominated by the disk emission is of order of 15 per cent (e.g. Giveon et al. 1999) in the timescales of years, so the simple scaling with the mass of the black hole does not seem to hold. It may mean that quasars are actually counterparts of the GBH in their Very High State. Since the spectra of GBH in Very High State are quite complex and not always dominated by the disk emission we can also expect such weak disk spectral states among the quasar/Seyfert class. Well studies source MCG-6-15-30 (Tanaka et al. 1995 and subsequent papers) may be an example, due to its exceptionally low mass and consequently high luminosity to the Eddington luminosity ratio (Nowak \\& Chiang 1999).  \\subsubsection{Broad Line Region and LINERs}  Nicastro (2000) proposed a very interesting model of the Broad Line Region based on the disk/corona model of WC\\. Z97. The basic idea is that the strong outflow develops at radii where the corona strength reaches its maximum and cool clouds forming in this disk wind are responsible for the emission of broad lines. Predictions based on our new model does not alter presented estimates significantly. The radius of the maximum corona strength in the new model can be determined from the condition $\\mdotz = 0$ (see Eq.~\\ref{eq:mdotz_loc_gen}), using the result with negligible evaporation to determine the local coronal accretion rate (Eq.~\\ref{eq:glob_evap}) which gives the dependence \\begin{equation} R_{\\rm max}=1.02 \\times 10^3 \\dot m^{16/23} \\alpha_{0.1}^{37/115}, \\end{equation} smaller by a factor of few than the WC\\. Z97 formula (see Eq. 3 of Nicastro 2000). On the other hand the new corona efficiency is now roughly symmetric around this point so the weighted formula for the outflow used by Nicastro would not decrease this radius, as it did in the case of sharp edge corona of WC\\. Z97 so the net effect will be roughly the same as before.  It was suggested by Lasota et al. (1996), that LINERs may represent the systems with ADAFs. The transition from objects with broad emission lines to LINERs happens, according to Nicastro (2000), due to the switch off the outflow if the accretion rate is too low for the existence of the radiation pressure dominated zone. The most efficient part of the corona, however, is located in the gas pressure dominated region so this connection is not clear. On the other hand, for very small accretion rate, within the frame of our new model, there is a transition from secondary disk rebuilding phase to a definite ADAF flow in the innermost part (see Eq.~\\ref{eq:mdotmin}). Such a change may be important from the point of view of the radiative cooling, the efficiency of advection and, consequently, the mass outflow. The mass loss through the wind as a function of radius is not calculated yet in the present model but this can be done in the future, following the method of e.g. Meyer \\& Hoffmeister (1994) and WC\\. Z97."
    },
    "0004/astro-ph0004191_arXiv.txt": {
        "abstract": "We present the discovery of several new regions of interaction between the ejecta and equatorial circumstellar ring of SN1987A, an interaction leading to a much expanded development of the supernova remnant. We also trace the development of the first such ``hot spot,'' discovered in 1997, back to 1995.\\ \\ Later hot spots seem to have emerged by early 1999.\\ \\ We discuss mechanisms for the long delay between the first and later spots. ",
        "introduction": "The development of SN~1987A provides an unprecedented opportunity to observe, at high spatial, spectral, and temporal resolution, the birth of a supernova remnant (SNR).\\ \\ Observations of SNR~1987A might reveal the angular and velocity distribution of the ejecta by highlighting where it impacts a presumably well-known nebular structure. While SNR~1987A has been observed for some years in X-rays and radio, only optical/near-IR data have sufficient sensitivity and angular resolution to detail the structures presented here.  The early discovery of these interaction regions is crucial. We demonstrate how they are best revealed with difference imaging, in H$\\alpha$ and [\\ion{N}{2}], and in \\ion{He}{1}~1.083$\\mu$m, for which the newly-formed spots have greatest contrast above the circumstellar equatorial ring (ER).\\ \\ We detect several new interaction sites, and trace known features back to epochs well before their actual discoveries. ",
        "conclusions": "\\mbox{HS~1-029} began within 8 years after the SN and at least 3--4 years ahead of the subsequent spots. Did this asymmetry arise in the velocity of the ejecta, in the density of the blue supergiant (BSG) wind, or in the morphology of the ER itself? The BSG wind should have been azimuthally smoothed by rotation of the progenitor, and relatively undisturbed by the ER (interior to the reverse shock $\\sim$5$\\times10^{17}$~cm from the star).\\ \\ Rayleigh-Taylor instabilities are expected in this decelerating wind, however over the $\\sim$20000~yr age of the ER \\citep{CH91} the sound speed would have allowed smoothing of these over scales of tenths of a parsec. This should be modelled over timescales for the BSG wind propagation to the ER ($\\sim$1000~yr).\\ \\ If the ER/ejecta interaction asymmetry were due to a large-scale misalignment between a bipolar ejecta pattern and the ER axis, a counter-jet should have impacted opposite \\mbox{HS 1-029}, contrary to observation. It seems likely that the asymmetry is due to structures in the ER or ejecta velocity field with fairly narrow extent in position angle, i.e.\\ thin fingers or jets. The homogeneity of the ER's shape versus the long delays between the \\mbox{HS~1-029} and later spots points to the ejecta velocity field as the major cause.  Valuable opportunities to obtain diagnostic spectra of the hot spots at very early times have been lost. If the fainter spots presented here are confirmed, then the ejecta have reached the ER over a wide range of PAs, and soon little pristine ER material will remain. Frequent ground and space-based monitoring of this rapidly developing event is required, with emphasis on difference imaging in \\ion{He}{1}~1.083$\\mu$m, H$\\alpha$, and [\\ion{N}{2}]."
    },
    "0004/astro-ph0004103_arXiv.txt": {
        "abstract": "The bipolar reflection nebula OH~231.8+4.2 (OH~0739-14) centred on the Mira variable star QX Pup has been mapped at K$_{S}$ in linear polarization using the ESO SOFI near-infrared multi-mode instrument. The large scale linear polarization features agree with previous studies, validating the polarimetric mode of the instrument. However the high spatial resolution of the polarization map reveals considerable structure, some of which finds correspondence in the surface brightness images. The core is crossed by low polarization filaments; the northern lobe reveals an axial finger and a rim of very high linear polarization (upto 80\\%) whilst the southern lobe shows less polarization contrast. There is a trend in polarization perpendicular to the bipolar axis in both lobes. The individual linear polarization features most probably arise from localised dust structures with distinct orientations leading to differing polarization through single scattering. A single scattering geometric model, consisting of two lobes of ellipsoidal section and a central dense disk, is considered which can explain the general features of the intensity and polarization images. ",
        "introduction": "The bipolar nebula OH~231.8+4.2 (IRAS OH~0739-14, sometimes called the Calabash nebula) is centred on the Mira variable QX Pup and is a source of line and continuum emission at many wavelengths. It displays strong extended optical and IR structure about 40$''$ in length, which is very highly polarized (Scarrott et al \\cite{scarr} for the optical; Kobayashi et al. \\cite{kob}, Heckert \\& Zeilik \\cite{heck}, Shure et al. \\cite{shur}, Kastner \\& Weintraub \\cite{kawe} for the IR). At the extremity of the southern bipolar lobe, there are high velocity emission line knots, similar to Herbig-Haro objects (Reipurth \\cite{reip}). The central variable star has an amplitude of $\\sim$2 mag., a period of $\\sim$700 days (Kastner et al. \\cite{kast92}) and a spectral type of M9~III (Cohen \\cite{coh81}; Cohen et al. \\cite{coh85}). The star is not directly visible in the optical on account of a waist of high extinction, but dominates the structure at longer wavelengths (Woodward et al. \\cite{wood}). Kastner et al. (\\cite{kast92}) have resolved the star from the surrounding nebula at K and L and deduce that it is a typical Mira variable. The object was first discovered on account of its OH maser emission (Turner \\cite{turn}) and mm wave CO mapping has shown a high velocity outflow oriented at $\\sim$40$^\\circ$ to the plane of the sky with deprojected velocities upto 330 kms$^{-1}$ for the southern lobe and 90 kms$^{-1}$ for the northern lobe (Alcolea et al. \\cite{alco}). It is generally considered to be a transition object between the Asymptotic Giant Branch (AGB) and the Planetary Nebula phases in which the high velocity stellar wind is blowing out the AGB dust envelope. The estimated mass of the molecular and dust envelope is estimated to be $\\sim$0.5 - 1M$_{\\odot}$, consistent with a main sequence mass $\\sim$3 M$_{\\odot}$ (Kastner et al. \\cite{kast98}).  Aside from its great astrophysical interest as a short-lived proto-planetary nebula, OH~231.8+4.2 is a useful test source for polarization measurements. From the optical to the L band it possesses high linear polarization (upto 50\\% in large apertures, Shure et al. \\cite{shur}) and so is a useful polarization calibrator. It cannot be stated with any certainty that it possesses the most important attribute for a calibrator - that the polarization is non-variable, but the high polarization makes it useful as a calibrator for the polarimetric efficiency and for position angle determination. In addition, since the polarization pattern is centro-symmetric along the bright lobes (although not in the equatorial direction) and extended over a length of about 50$''$, it provides an ideal source to confirm any instrumental depolarization effects. In consequence OH~231.8+4.2 was selected as a test source for the commissioning of the SOFI polarimetric mode. The quality of the instrument and the observations proved to be excellent and yielded the highest resolution near-IR polarization map of this source.  The instrumental set-up and the observations are detailed in Sect. 2 and the polarization results presented in Sect. 3. The spatial polarization structure is discussed in Sect. 4 together with conclusions about the geometry of the elongated dust cloud and the scattering properties of the grains. ",
        "conclusions": "A linear polarization map of OH~0231.8+4.2 has been obtained in the K$_{S}$ band with the SOFI multi-mode instrument. A wealth of structure is revealed in the polarization and two elongated features in the northern bipolar lobe have very high linear polarization. A geometrical model with bipolar lobes and an equatorial disc is proposed which accounts for some of these features. However the high polarization finger in the northern lobe, and the position angle of polarization vectors in the surrounding (disc) material, are not explained. Higher spatial resolution polarimetry of OH~231.8+4.2 is achievable from the ground with AO, and polarization maps at different wavelengths would further understanding of this fascinating target."
    },
    "0004/astro-ph0004335_arXiv.txt": {
        "abstract": "The temporal variations of the rotation rate in the solar interior are studied using frequency splittings from Global Oscillations Network Group (GONG) data obtained during the period 1995--99. We find alternating latitudinal bands of faster and slower rotation which appear to move towards the equator with time --- similar to the torsional oscillations seen at the solar surface. This  flow pattern appears to persist to a depth of about $0.1R_\\odot$ and in this region its magnitude is well  correlated with solar activity indices. We do not find any periodic or systematic changes in the rotation rate near the base of the convection zone. ",
        "introduction": "Helioseismic data allow us to probe the rotation rate in the solar interior as a function of radius and latitude (Thompson et al.~1996; Schou et al.~1998). With the accumulation of Global Oscillations Network Group (GONG) and Michelson Doppler Imager (MDI) data over the last five years, it has also  become possible to study the temporal variation in the rotation rate and other properties of the  solar interior.  It is well known that frequencies of solar oscillations vary with time (Elsworth et al.~1990; Libbrecht \\& Woodard 1990; Dziembowski et al.~1998; Bhatnagar, Jain \\& Tripathy 1999; Howe, Komm \\& Hill 1999). These variations are also known to be  correlated with solar activity. However, these frequency changes appear to result from variations in solar structure close to the surface; there is no clear evidence for structural changes in deep interior (Basu \\& Antia 2000a). Inversions for the rotation rate in solar interior on the other hand, have confirmed that the temporal variations in the rotation rate  penetrate to somewhat  deeper layers (Howe et al.~2000a; Toomre et al.~2000). Furthermore, the changing pattern of solar rotation rate in interior also agrees with the torsional oscillations observed at the solar surface (Howard \\& LaBonte 1980). While most of these studies are restricted to the outer part of the convection zone, the seat of the solar dynamo is generally believed to be near the base of the convection zone and rotation is believed to play an important role in the operation of the solar dynamo. This is also the region where the tachocline is located (Kosovichev 1996; Basu 1997). Thus one needs to look for possible changes in rotation rate in this region. Recently, Howe et al.~(2000b) have reported a 1.3 year periodicity in variation of equatorial rotation rate at a radial distance of $0.72R_\\odot$. It is not clear if the period of 1.3 years is associated with solar cycle variations. Considering the fact that we have only 4 years of data, this period is also uncomfortably  close to the orbital period of the Earth and one may expect systematic errors on this period being introduced in the data. If this periodicity is confirmed, it may have implications for the theories of solar dynamo. Hence, this needs to be checked by an independent analysis.  In this work we use oscillation frequencies obtained from GONG observations  covering a period from May 1995 to June 1999, to study possible temporal variations in the rotation rate as a function of radius and latitude. We use both 1.5d and 2d inversion using the Regularized Least Squares (RLS) technique (Antia, Basu \\& Chitre 1998) to determine the rotation rate from the observed splitting coefficients. ",
        "conclusions": "Inversions to determine the solar rotation rate with data from GONG months 1--42, which mostly cover the rising phase of solar cycle 23,  show a clear change in the rotation rate  with solar activity. The zonal flow, which we have defined as residual in the rotation rate obtained by subtracting a time-average of a three-term fit to the rotation rate, shows bands of faster and slower rotation. These bands appear to move towards the equator with time. The magnitude of this time varying component is about 2 nHz. This pattern is similar to that found by Howe et al.~(2000c) using helioseismic data from GONG and MDI projects. The amplitude of the flow pattern in our results is also similar to those in Howe et al.~(2000c) as can be seen by comparing their Fig.~3 with our Fig.~\\ref{fig:zonalcuta}.  The mean magnitude of the pattern in the outer parts of the solar convection zone appears to increase with solar activity and is reasonably well correlated with the mean 10.7 cm flux which tracks solar activity. The correlation is seen  up to a depth of about $0.1R_\\odot$ (70 Mm). This appears to be depth to which the zonal flow pattern penetrates. Below this depth the temporal variations in the rotation rate do not appear to be significant. This inferred depth of zonal flow pattern is consistent with estimate of at least 60 Mm obtained by Howe et al.~(2000c) using independent data sets and inversion techniques. Clearly, there is good agreement between different inversion results in the outer convection zone. There appears to be  some hint that the flow pattern moves upwards with time, however, the statistical significance of this is as yet unclear.   We do not find any evidence of periodic changes in the solar rotation rate near the base of the solar convection zone. We find that the temporal variations at this depth are random in nature and probably represent statistical fluctuations. There is no clear periodic signal in the residual rotation rate with period less than 2 years at any depth. The present data length of 4 years does not allow us to study periodicities with longer time period. Some of this fluctuations may also arise from systematic errors in data."
    },
    "0004/gr-qc0004002_arXiv.txt": {
        "abstract": "A rapidly rotating, axisymmetric star can be dynamically unstable to an $m=2$ ``bar\" mode that transforms the star from a disk shape to an elongated bar. The fate of such a bar--shaped star is uncertain. Some previous numerical studies indicate that the bar is short lived, lasting for only a few bar--rotation periods, while other studies suggest that the bar  is relatively long lived. This paper contains the results of a numerical simulation of a rapidly rotating $\\gamma=5/3$ fluid star. The simulation shows that the bar shape is long lived: once the bar is established, the star retains this shape for more than $10$ bar--rotation periods, through the end of the simulation. The results are consistent with the conjecture that a star will retain its bar shape indefinitely on a dynamical time scale, as long as its rotation rate exceeds the threshold for secular bar instability. The results are described in terms of a low density neutron star, but can be scaled to represent, for example, a burned--out stellar core that is prevented from complete collapse by centrifugal forces. Estimates for the gravitational--wave signal indicate that a dynamically unstable neutron star in our galaxy can be detected easily by the first generation of ground based gravitational--wave detectors. The signal for an unstable neutron star in the Virgo cluster might be seen by the planned advanced detectors. The Newtonian/quadrupole approximation is used throughout this work. ",
        "introduction": "\\label{sec:intro}  A self--gravitating, axisymmetric fluid body with a sufficiently high rotation rate can be dynamically unstable to non-axisymmetric perturbations. Typically, the fastest growing unstable mode is the $m=2$ ``bar\" mode which acts to transform the body from a disk--like shape to an elongated bar that tumbles end--over--end. This instability has been described analytically for the case of uniform density bodies\\cite{Chandra,ST}, and has been the subject of numerous numerical studies\\footnote{For a brief review, see Ref.~\\cite{NCT}}. The numerical results show that bar formation is accompanied by the ejection of mass and angular momentum, and that the ejected matter forms long spiral arms in the equatorial plane. The subsequent evolution is less certain. Some simulations indicate that the bar shape is short lived, with the star returning to a predominantly disk--like shape after a few bar--rotation periods. Other simulations predict that the bar  persists for many bar--rotation periods. In recent work, New, Centrella, and Tohline\\cite{NCT} address this issue with a series of simulations using two different codes at various resolutions, and conclude that the bar shape is persistent. In their highest resolution run the bar decayed after roughly $6$ or $7$ bar--rotation periods. This was believed to be caused by numerical errors that induced an unphysical center of mass motion. In a lower resolution run a symmetry condition was imposed that prevented any center of mass motion. In that case the star maintained its bar shape throughout the simulation.  The purpose of the present work is to simulate the long--time evolution of a rapidly rotating, self--gravitating star using Newtonian hydrodynamics and  gravity. As discussed in Sec.~\\ref{sec:numcode}, the numerical code is substantially different from the codes that have been used previously to address this problem. The initial data for this study consists of a $\\gamma = 5/3$ polytrope with stability parameter $\\beta = 0.30$. The stability parameter is defined by $\\beta = T/|W|$, where $T$ is rotational kinetic energy and $W$ is gravitational potential energy. The results here suggest that the bar shape is indeed long--lived---the star displays a prominent bar shape at the end of the simulation, which includes more than $10$ bar rotation periods.  Numerical studies of fluids with various equations of state and initial rotation profiles have shown that the {\\em dynamical} bar instability appears when the stability parameter $\\beta$ exceeds a certain critical value, typically close to $0.27$\\cite{DT,TIPD}. For the {\\em secular} instability, which arises through dissipative mechanisms, the critical value of $\\beta$ is near $0.14$. A neutron star might reach the critical value of $\\beta$ for dynamical or secular instability by accreting matter and angular momentum from a binary companion. A stellar core that has exhausted its nuclear fuel might reach a critical rotation rate as it collapses.  A star or stellar core that develops a rotating bar--like configuration will generate large amounts of gravitational radiation. Depending on the distance of the source, this radiation might be strong enough to be detected by the world--wide network of gravitational--wave detectors currently under construction\\cite{BandC}. Here, the question of persistence of the bar shape becomes very important. The detectability of a source depends on its characteristic amplitude $h_c \\approx h\\sqrt{n}$, where $h$ is the amplitude of the waves with frequency $f$ and $n$ is the number of wave cycles in a bandwidth near $f$\\cite{Thorne}. Thus, long--duration signals with large $n$ can be more easily detected than short--duration signals.  Should we expect the bar shape to persist or decay? One reason why the bar might decay is the loss of mass and angular momentum  from the ends of the bar. The accompanying drop in rotational kinetic energy could reduce the stability parameter and allow the star to return to axisymmetry. Loss of rotational kinetic energy through shock heating might also occur. For the simulation presented here $\\beta$ has an initial value of $0.30$, large enough to dynamically trigger the growth of the bar mode. During the initial period of bar formation, mass and angular momentum are shed from the ends of the bar and long, spiral arms are formed. The stability parameter rapidly drops below $0.27$ and eventually settles to a value of about $0.24$. However, these losses of mass and angular momentum, with the accompanying drop in $\\beta$, are insufficient to completely rob the star of its bar shape.  The results here are consistent with the conjecture that the star will retain its bar shape indefinitely on a dynamical time scale, as long as $\\beta$ stays above the critical value for secular instability\\cite{IDP}. If this is correct, the bar should still decay due to secular processes. Imamura, Durisen, and Pickett have argued that the bar continues to shed small amounts of angular momentum to the spiral arms, and these loses cause the bar to decay\\cite{IDP}. Other secular mechanisms, such as energy and angular momentum loss through gravitational radiation, could play a role as well.  It is possible that, for the equation of state and angular velocity profile used here, the star is able to retain its bar shape simply because the critical value of $\\beta$ for dynamical instability is $0.24$ or less. This possibility was tested by conducting a second simulation using the same equation of state and scaling the angular velocity by a factor of $3/4$. The resulting model has an initial value of $\\beta$ equal to $0.25$. The bar mode showed no signs of growth in this simulation. Thus, for the models tested here, the bar mode does not spontaneously grow unless $\\beta$ exceeds a critical value greater than $0.25$. However, once the bar shape is established, it can persist for many bar--rotation periods with $\\beta$ equal to $0.24$ or less.  In the present numerical calculation there are two sources of error in angular momentum. First, some of the angular momentum that is lost off the edge of the computational grid might, ideally, fall back onto the star. Second, numerical errors at each time step introduce a purely artificial loss of angular momentum. Both of these errors act to decrease the angular momentum of the star. One expects such losses to cause the bar shape to deteriorate. What is seen, both in the gravitational wave analysis and in the Fourier analysis of the density, is a relatively persistent signal with a modest decline in amplitude over the duration of the simulation.  The physical model and initial data that form the basis of the simulation are described in Sec.~II. Section III contains a discussion of the numerical code. The results of the simulation are described qualitatively and quantitatively in Sec.~IV. Conclusions and results are summarized in Sec.~V. Details of the initial data code are given in the Appendix. ",
        "conclusions": ""
    },
    "0004/astro-ph0004045_arXiv.txt": {
        "abstract": "The orbital parameters of the observed extrasolar planets differ strongly from those of our own solar system. The differences include planets with high masses, small semi-major axis and large eccentricities. We performed numerical computations of embedded planets in disks and follow their mass growth and orbital evolution over several thousand periods.  We find that planets do migrate inwards on timescales of about $10^5$ years on nearly circular orbits, during which they may grow up to about 5 Jupiter masses. The interaction of the disk with several planets may halt the migration process and lead to a system similar to the solar planetary system. ",
        "introduction": "The evolution of protoplanets can only be studied simultaneously with that of the protostellar disk in which the planets are still embedded early in their lifetime. The disk is assumed to be flat and non-self gravitating, and is modelled by the planar two-dimensional ($r-\\varphi$) Navier-Stokes equations. The mutual gravitational interaction of the planets and the star, and the gravitational torques of the disk acting on the planets and the central star are included. The time dependent hydrodynamic equations are integrated by using a finite difference scheme (Kley 1998, 1999). The computations cover the radial range from 1.2 to 20~AU. The initial surface density distribution is axisymmetric and follows $\\Sigma \\propto r^{-1/2}$; to speed up the computations an annular gap is imposed (Kley 1999). The material orbits the star with Keplerian speed. The planets of one Jupiter mass are initially placed at typical distances of several AU from the star. The whole system (disk, star and planets) is then evolved for several thousand orbital periods. ",
        "conclusions": "By initially considering one planet on a fixed circular orbit, it was shown that even though the planet opens up an annular gap in disk it is nevertheless able to accrete more mass from its surroundings. As more massive planets induce a wider and deeper gap the mass accumulation essentially terminates, which puts the upper limit to the mass of the planet mass at about 5 to 10 $M_{Jup}$, in good agreement with the observations.  By considering the evolution of a system consisting of two Jupiter-type planets, we show (Kley 2000) that by mutual gravitational interaction the inward motion of the inner planet may come to a halt. As the outer planet continues to move inward and approaches the inner, the gravitational interaction between the two planets increases. The resulting configuration may be unstable, and lead to systems similar to $\\upsilon$~And.  If the protoplanetary nebula has dissipated before the planets come very close to each other, one is left with a system of massive planets at several AU distance, similar to our own Solar System."
    },
    "0004/astro-ph0004273_arXiv.txt": {
        "abstract": "We report on a sensitive survey for radio pulsar wind nebulae (PWN) towards 27 energetic and/or high velocity pulsars. Observations were carried out at 1.4~GHz using the Very Large Array and the Australia Telescope Compact Array, and utilised pulsar-gating to search for off-pulse emission.  These observing parameters resulted in a considerably more sensitive search than previous surveys, and could detect PWN over a much wider range of spatial scales (and hence ambient densities and pulsar velocities).  However, no emission clearly corresponding to a PWN was discovered. Based on these non-detections we argue that the young and energetic pulsars in our sample have winds typical of young pulsars, but produce unobservable PWN because they reside in low density ($n \\sim 0.003$~cm$^{-3}$) regions of the ISM. However, non-detections of PWN around older and less energetic pulsars can only be explained if the radio luminosity of their winds is less than $10^{-5}$ of their spin-down luminosity, implying an efficiency at least an order of magnitude smaller than that seen for young pulsars. ",
        "introduction": "\\label{sec_intro}   Almost all radio pulsars have rotational periods which are steadily increasing with time. This spin down corresponds to a loss of rotational kinetic energy $\\dot{E} \\equiv 4\\pi^2I\\dot{P}/ P^3$, where $I$ is the moment of inertia of the neutron star (assumed to be $10^{45}$~g~cm$^2$) and $P$ is its period; for the known pulsar population $\\dot{E}$ falls in the range $10^{28}-10^{39}$~erg~s$^{-1}$. The bolometric luminosity of the radio pulses themselves is in all cases a tiny fraction of the spin-down luminosity $\\dot{E}$, and it is thought that most of the pulsar's spin-down luminosity is dissipated via a magnetised wind populated by relativistic electrons and positrons \\cite{rg74,mic82,kc84}.  However it is still not well understood how this wind is produced, how it evolves as it flows away from the pulsar, what it is composed of, how its properties depend on those of the pulsar itself, or how it changes as the pulsar ages.  Particles in the wind move along magnetic field lines as they stream away from the pulsar magnetosphere, and produce no observable emission. At some distance from the pulsar, the pressure of the wind is eventually balanced by an external pressure, and the resulting shock randomises the pitch angles of the relativistic particles.  These particles consequently gyrate in the local magnetic field and produce synchrotron emission.  The properties of the resulting {\\em pulsar wind nebula}\\ (PWN) can then be used to determine various parameters of the pulsar wind which produces it. At radio wavelengths, PWN are characterised by an amorphous or filled-centre morphology, a moderate degree of linear polarization ($\\sim$20\\%), and relatively flat spectra ($\\alpha \\sim -0.3$, $S_\\nu \\propto \\nu^\\alpha$) \\cite{wp78}.  Various types of PWN are produced, depending on the source of confinement of the wind. Young pulsars are often still located inside their associated supernova remnants (SNRs), and the hot gas produced by the SNR blast-wave provides the confining pressure.  These PWN, also known as ``plerions'', are typified by the Crab Nebula. If the SNR has dissipated, the confining pressure is then that of the ambient interstellar medium (ISM). This results in much larger ``ghost remnants'', which have been proposed but not observed \\cite{bopr73,ccgm83}. In cases where a pulsar has a high space velocity, the ram pressure resulting from its motion can dominate the ambient gas pressure, resulting in a bow-shock PWN (e.g.\\ Frail \\& Kulkarni 1991\\nocite{fk91}).   While PWN can tell us much about pulsar winds, the number of sources we have to study is small -- at radio wavelengths, fewer than 10 pulsars have observable PWN. All of these pulsars are very young and have high values of $\\dot{E}$. While $\\sim$20\\% of SNRs have a ``plerionic'' component, in most cases the associated pulsar has not been detected, and without knowing the pulsar parameters it is difficult to constrain the properties of the PWN and the corresponding wind. Thus it is of considerable interest to target candidate pulsars with the intention of either finding new PWN, or determining upper limits on such emission.  Many searches for radio PWN around energetic or fast-moving pulsars have been carried out, at varying resolutions and surface-brightness sensitivities, but usually with no success  (e.g. Sch\\\"{o}nhardt 1974\\nocite{sch74}; Weiler, Goss \\& Schwarz 1974\\nocite{wgs74}; Cohen \\etal\\ 1983\\nocite{ccgm83}).  The most recent and comprehensive of these searches, and the only one to specifically target young, energetic or high velocity pulsars, was the recent survey of Frail \\& Scharringhausen \\shortcite{fs97}, hereafter FS97. FS97 imaged regions around 35 pulsars with the Very Large Array (VLA) at 8.4~GHz, and found no nebular emission associated with any of their targets. Their stringent upper limits allowed them to conclude that most pulsars put less than $10^{-6}$ of $\\dot{E}$ into radio emission from a PWN, $\\sim100$ times less than observed for the young, high $\\dot{E}$, pulsars which power detected radio nebulae. Based on this result, FS97 concluded that pulsar winds change in some way as pulsars age and slow down, such that they are no longer efficient at producing radio emission.  However, despite these apparently constraining limits, in hindsight the observing parameters for this search were probably not ideal for looking for PWN.  First, FS97 argued that for ambient densities of 1~cm$^{-3}$ and pulsar velocities of 150~\\kms, PWN around almost all their sources were likely to be unresolved, even at their high spatial resolution of $0\\farcs8$. However, other choices of ambient density and pulsar velocity can produce PWN with much larger angular extents, resulting in a flux density limit much poorer than estimated by FS97.  Furthermore, the maximum scale to which FS97 were sensitive was only $20''$; they could not detect PWN larger than this at any flux density. Secondly, in 40\\% of their sample FS97 detected a point source at the position of the pulsar, but had no way of distinguishing between the compact PWN they were looking for and the pulsars themselves. By extrapolating pulsar flux densities from much lower frequencies, FS97 argued that the flux densities they were detecting at 8.4~GHz were consistent with the expected pulsed fluxes, and hence concluded that they were not detecting any PWN in these data.  Finally, a typical PWN (spectral index $\\alpha = -0.3$) has a significantly lower flux density at their observing frequency of 8.4~GHz than at lower frequencies.  Motivated by these points, we have undertaken an extensive survey for PWN at both northern and southern declinations, with observing parameters chosen to give much greater sensitivity to PWN.  First, we have observed at 1.4~GHz, at which frequency PWN can be expected to be $\\sim$70\\% brighter than at 8.4~GHz. Secondly, we have observed at a reduced spatial resolution of $\\sim$12$''$, to give better sensitivity to extended structure. Thirdly, our observations are in telescope configurations whose shortest spacings correspond to a spatial scale of many arcmin. Finally, and most importantly, all our observations have employed pulsar-gating, in which images are made from data taken only when the pulsar is off; we are thus sensitive to compact and unresolved PWN, which might otherwise be masked by the pulsars themselves.  Initial results from this survey, using the Australia Telescope Compact Array (ATCA), have been presented in two previous papers. Gaensler \\etal\\ \\shortcite{gsfj98}, hereafter GSFJ98, reported the discovery of a faint PWN associated with PSR~B0906--49, while Stappers, Gaensler \\& Johnston \\shortcite{sgj99}, hereafter SGJ99, presented non-detections of PWN towards four pulsars. We here report on the remainder of this survey, consisting of 1.4~GHz pulsar-gated observations of 27 more pulsars, using the VLA and ATCA. In \\S\\ref{sec_obs}, we describe our observations and analysis, while in \\S\\ref{sec_results} we present non-detections of PWN towards these sources. In \\S\\ref{sec_discuss} we quantify the improvement in sensitivity of the current survey, and discuss the constraints we can put on the radio luminosities of pulsar winds from our data. ",
        "conclusions": "\\label{sec_discuss}  As discussed in \\S\\ref{sec_intro}, the parameters of the current survey were chosen to improve on the sensitivity of previous surveys, in particular that carried out by FS97.  These factors are summarised in Fig~\\ref{fig_sens}, where the sensitivity of FS97's search is compared to that presented here. Only for PWN with radii between $0\\farcs4$ and $\\sim1''$ is FS97's search more sensitive than ours. Between $1''$ and $10''$, the current results are considerably (up to 100 times) more sensitive, while at scales smaller than $0\\farcs4$ and larger than $10''$, our data probe a parameter space to which FS97 were not sensitive at all.  \\begin{figure*} \\begin{minipage}{160mm} \\centerline{\\psfig{file=fig_sens.eps,width=16cm,angle=270}} \\caption{Comparison of 5$\\sigma$ sensitivities of the 8.4~GHz PWN search of FS97 with the 1.4~GHz observations reported here. The observations of FS97 were not gated, and so generally are not sensitive to PWN on scales smaller than the resolution limit. For our VLA (ATCA) observations, we have adopted a typical spatial resolution of $12''$ ($12''$), a maximum spatial scale of $15'$ ($4\\farcm5$), and a 1$\\sigma$ sensitivity of 0.5 (0.15) mJy~beam$^{-1}$. A spectral index $\\alpha = -0.3$ ($S_\\nu \\propto \\nu^{\\alpha}$) has been assumed in converting 8.4~GHz results to 1.4~GHz.} \\label{fig_sens} \\end{minipage} \\end{figure*}  The results reported in \\S\\ref{sec_results} showed that most of the sources in our sample had no detectable PWN associated with them. The exceptions were PSRs B0136+57, B1634--45 and B1706--16, for which unresolved off-pulse sources were detected at the pulsar position. From the current data, we are unable to conclusively determine whether this emission corresponds to emission from the pulsar or from a compact PWN. While we plan to investigate these sources further, for the purposes of the present discussion we assume these observations to be non-detections, but with a sensitivity corresponding to the flux density of the off-pulse source (rather than to the noise in the surrounding area of the image).  To quantify the significance of our non-detections, we follow FS97 in characterising a radio PWN's integrated luminosity, $L_R$, by \\begin{equation} L_R = \\epsilon \\dot{E}~{\\rm erg~s}^{-1}, \\label{eqn_l_r} \\end{equation} where  $\\dot{E}$~erg~s$^{-1}$ is the associated pulsar's spin-down luminosity, and $\\epsilon$ is the fraction of $\\dot{E}$ which goes into radio emission.  Assuming a typical PWN spectral index $\\alpha = -0.3$ and integrating from 10~MHz to 100~GHz, the corresponding 1.4~GHz flux density is  \\begin{equation} S_{1.4} = 2.1\\times10^{5}~\\frac{\\epsilon \\dot{E}_{34}}{d^2}~{\\rm mJy,} \\label{eqn_s20} \\end{equation} where $d$~kpc is the distance to the pulsar and $\\dot{E} = 10^{34} \\dot{E}_{34}$~erg~s$^{-1}$. We generally use distances from the pulsar catalogue \\cite{tmlc95}, derived either from a pulsar's dispersion measure \\cite{tc93} or from its kinematic distance based on H~{\\sc i} measurements (e.g. Frail \\& Weisberg 1990\\nocite{fw90}).  If we do not detect a PWN in our observations, we can potentially put an upper limit on $S_{1.4}$, and hence on $L_R$ and $\\epsilon$. However, as demonstrated in Fig~\\ref{fig_sens}, these limits depend on the angular size we expect for the PWN.  As previously discussed by FS97 and SGJ99, PWN can in general be divided into two distinct classes (assuming that the pulsar is not inside a SNR, which appears to be the case for all the sources in our sample): those which are confined by the gas pressure of the ambient ISM (``static PWN''), and those confined by ram-pressure resulting from motion of the pulsar through the ISM (``bow-shock PWN'').  Let us first consider the case of a static PWN. The bubble in the ISM driven by the pulsar wind will expand supersonically into the ambient medium, producing a PWN of radius \\cite{aro83} \\begin{equation} R_{\\rm static} = 0.14 \\times \\left( \\frac{\\dot{E}_{34} t^3_3}{n} \\right)^{1/5}~{\\rm pc}, \\label{eqn_rstatic} \\end{equation} where $n$~cm$^{-3}$ is the density of the ambient medium (assumed to be pure hydrogen), and $t_3$~kyr is the period for which the pulsar has been interacting with the ISM. In further discussion we assume that this age is given by the pulsar's characteristic age, $\\tau_c \\equiv P/2\\dot{P}$.  However a PWN is only static while $\\dot{R}_{\\rm static} > V_{\\rm PSR}$, where $V_{\\rm PSR}$~\\kms\\ is the pulsar's space velocity. Re-arranging Equation~(3) of FS97 (and correcting for a missing factor of $4\\pi$ in their results), we find that a PWN will be static for velocities and ambient densities for which: \\begin{equation} n V_{\\rm PSR}^5 \\la 4 \\times10^{9}~\\dot{E}_{34}/t_3^2. \\label{eqn_nv5} \\end{equation} When this condition is not met, the pulsar has ``overtaken'' its own static PWN, and a bow-shock PWN results. The resulting PWN is much smaller than the static PWN, and has a size determined by a balance between the pressure of the relativistic pulsar wind and the ram-pressure resulting from the pulsar's motion, \\begin{equation} R_{\\mbox{\\scriptsize bow-shock}} = 0.63 \\left(\\frac{\\dot{E}_{34}}{nV_{\\rm PSR}^2}\\right)^{1/2}~{\\rm pc}, \\label{eqn_rbow} \\end{equation} where we have assumed that all the spin-down luminosity of the pulsar goes into the wind, and that the radius of a bow-shock PWN is 1.5 times the radius at which the ram and wind pressures balance \\cite{vm88b}.  Thus for a given $n$ and $V_{\\rm PSR}$, we can use Equation~(\\ref{eqn_nv5}) to determine whether a PWN has a static or a bow-shock morphology, and from this use either Equation~(\\ref{eqn_rstatic}) or (\\ref{eqn_rbow}) to determine its radius, and hence its angular extent, $\\theta_{\\rm PWN}$.  For a pulsar-gated observation with a 1.4~GHz RMS sensitivity of $\\sigma$~mJy~beam$^{-1}$ at a resolution $\\theta_\\alpha '' \\times \\theta_\\delta''$, suppose no off-pulse source is detected at the 5$\\sigma$ level.  There are three possible reasons for this non-detection: the PWN is unresolved and below the point-source sensitivity of the observations, the PWN is resolved and below the surface brightness sensitivity limit, or the PWN is larger than the maximum spatial scale to which we are sensitive. We now consider what limits on $\\epsilon$ can be derived  for each situation.  If $\\theta_{\\rm PWN}$ is smaller than the resolution limit, then $S_{1.4} < 5\\sigma$ and from Equation~(\\ref{eqn_s20}) the corresponding upper limit on $\\epsilon$ is \\begin{equation} \\epsilon < \\epsilon_0 = 2.4\\times10^{-5}~\\sigma \\left( \\dot{E}_{34}/d^2 \\right)^{-1}. \\label{eqn_eps_limit} \\end{equation} Note that if an off-pulse point source of 1.4~GHz flux density $S_{\\rm pt}$ is detected at the position of the pulsar (as was the case for three of the pulsars in our sample), then $5\\sigma$ should be replaced by $S_{\\rm pt}$ in this expression.  If the PWN is extended, but smaller than the largest spatial scale to which our interferometer is sensitive, then \\begin{equation} \\epsilon < \\frac{\\epsilon_0 \\theta_{\\rm PWN}^2}{\\theta_\\alpha \\theta_\\delta}, \\label{eqn_eps_limit2} \\end{equation} where we have assumed a Gaussian profile for the PWN. Finally, if the PWN is larger than our observations can detect, we can put no limit on $\\epsilon$.  FS97 have argued, on the basis of their non-detections, that the lack of observable radio PWN around most pulsars implies $\\epsilon \\la 10^{-6}$, two orders of magnitude less than for pulsars with detected radio PWN. However, in making this calculation they assumed that the ambient density was $n=1$~cm$^{-3}$. Through Equation~(\\ref{eqn_nv5}), when combined with a reasonable velocity, this implied that all the pulsars they observed were powering bow-shock PWN, and the consequent high ram pressure ensured that these sources would largely be unresolved by their $0\\farcs8$ resolution.  However their assumed ambient density is not representative of our current best understanding of the ISM. While the relative filling fractions are uncertain, available evidence supports a multi-phase ISM of which 90\\% by volume is a combination of a warm medium of density $n = 0.3$~cm$^{-3}$ and a hot ionised component of density $n = 0.003$~cm$^{-3}$ (see Ferri\\`{e}re 1998a\\nocite{fer98a} for a recent discussion and overview).  In the hot low-density component, we still expect most PWN to be bow shocks, but the ram pressure is greatly reduced and PWN will consequently be much more extended. Thus many PWN in this low density medium will be larger than the resolution limit of FS97, and through Equation~(\\ref{eqn_eps_limit2}), the limit on $\\epsilon$ much less stringent than claimed.   In order to better consider the limits on $\\epsilon$, we therefore carry out the following calculation for 27 of the pulsars in Table~\\ref{tab_obs} (PSR~B1706--44 is excluded as it has a candidate PWN), and additionally for the 4 pulsars discussed by SGJ99. For each pulsar in our sample, we adopt possible densities of $n=0.3$~cm$^{-3}$ and $n=0.003$~cm$^{-3}$.  Ten of the pulsars in our sample have measured proper motions, and we use the corresponding 3D space velocities determined by Cordes \\& Chernoff \\shortcite{cc98}; one other pulsar (PSR~B1055--52) has had a scintillation velocity determined for it \\cite{jnk98}. For the remaining sources we set $V_{\\rm PSR} = 380$~\\kms, corresponding to the mean pulsar velocity of the distribution of Cordes \\& Chernoff \\shortcite{cc98}.  Using Equation~(\\ref{eqn_nv5}) we determine whether each corresponding PWN is bow-shock or static, and then consequently determine $\\theta_{\\rm PWN}$. Upper limits on $\\epsilon$ are then determined from Equations~(\\ref{eqn_eps_limit}) and (\\ref{eqn_eps_limit2}), and the results given in Table~\\ref{tab_limits}. For the 16 sources in our sample which were also observed by FS97, we can make a similar calculation based on their results (converting their data to 1.4~GHz assuming a spectral index $\\alpha = -0.3$); these revised values of $\\epsilon_{\\rm max}$ are also listed in Table~\\ref{tab_limits}. Note that many PWN are unresolved for either choice of ambient density, and so have the same value of $\\epsilon_{\\rm max}$ in both cases.  \\begin{table*} \\begin{minipage}{180mm} \\caption{Upper limits on $\\epsilon = L_R/\\dot{E}$ for PWN non-detections, including four non-detections from SGJ99. Pulsars are sorted by decreasing $\\dot{E}$; $\\theta$ is the predicted angular size of a PWN for a given density. For upper limits on $\\epsilon$ derived from the results of FS97, U indicates that the predicted size of the PWN is too large to have been detected by their observations, while ``$\\ldots$'' indicates that a particular pulsar was not part of their sample.} \\label{tab_limits} \\begin{tabular}{lcccccccccccc} Pulsar & $\\dot{E}_{34}$ & $t_3$ & $d$ & $V_{\\rm PSR}$ &  & \\multicolumn{3}{c}{$n=0.3$~cm$^{-3}$} & & \\multicolumn{3}{c}{$n=0.003$~cm$^{-3}$}  \\\\ & (erg~s$^{-1}$)  & (kyr) & (kpc) & (\\kms) & & $\\theta$ ($''$) & \\multicolumn{2}{c}{$\\log_{10} \\epsilon_{\\rm max}$} & & $\\theta$ ($''$) & \\multicolumn{2}{c}{$\\log_{10} \\epsilon_{\\rm max}$} \\\\ &  & & & & & & (this paper) & (FS97) & & & (this paper) & (FS97) \\\\ B1823--13     &   280 &     21 &   4.1  & $\\ldots$  &  &   5.01 &    --6.3  & --5.4  &  &   50.1 &    --5.2  &     U      \\\\ J1105--6107   &   250 &     63 &   7.1  & $\\ldots$  &  &   2.76 &    --5.8  & $\\ldots$ & &    27.6 &    --5.6  &  $\\ldots$  \\\\ B1046--58     &   200 &     20 &   3.0  & $\\ldots$  &  &   5.86 &    --6.8  &  $\\ldots$ &  &   58.6 &    --5.3  &  $\\ldots$  \\\\ B1610--50     &   160 &      7 &   7.2  & $\\ldots$  &  &   2.16 &    --5.8  &  $\\ldots$ &  &   21.6 &    --5.3  &  $\\ldots$  \\\\ B1727--33     &   120 &     26 &   4.2  & $\\ldots$  &  &   3.19 &    --5.2  &    --5.3  &  &   31.9 &    --4.8  &     U      \\\\ B1930+22     &    75 &     40 &  12.1  & $\\ldots$  &   &  0.88 &    --4.6  &    --5.4  &   &   8.8 &    --4.6  &    --3.4   \\\\ B0114+58     &    22 &    275 &   2.1  & $\\ldots$  &   &  2.70 &    --5.6  &    --5.4  &   &  27.0 &    --5.1  &     U      \\\\ J1835--1106   &    18 &    127 &   3.1  & $\\ldots$  &  &   1.69 &    --5.1  &  $\\ldots$ &  &   16.9 &    --5.0  &  $\\ldots$  \\\\ J0631+1036   &    17 &     44 &   6.6  & $\\ldots$  &   &  0.78 &    --4.0  &  $\\ldots$ &   &   7.8 &    --4.0  &  $\\ldots$  \\\\ B0740--28     &    14 &    157 &   1.9  & 276       &  &   3.36 &    --5.5  &    --4.6  &  &   33.6 &    --4.9  &     U      \\\\ B1508--57     &    13 &    298 &  12.7  & $\\ldots$  &  &   0.35 &    --4.2  &  $\\ldots$ &  &    3.5 &    --4.2  &  $\\ldots$  \\\\ B1356--60     &    12 &    319 &   5.9  & $\\ldots$  &  &  0.72 &    --4.7  &  $\\ldots$ &  &    7.2 &    --4.7  &  $\\ldots$  \\\\ B1634--45     &     7.5 &    590 &    3.8  &  $\\ldots$  & &    0.88 &    --5.1  &    --5.3  & &      8.8 &    --5.1  &    --3.3   \\\\ B0611+22     &     6.3 &     89 &    4.7  &  212       &  &   1.18 &    --4.5  &    --4.9  & &     11.8 &    --4.5  &    --2.9   \\\\ J0538+2817   &     4.9 &    619 &   1.8  &  $\\ldots$  &   &  1.54 &    --5.2  &  $\\ldots$ & &     15.4 &    --5.0  &  $\\ldots$  \\\\ B1718--35     &     4.5 &    176 &   6.4  &  $\\ldots$  &  &   0.41 &    --3.5  &  $\\ldots$ & &      4.1 &    --3.5  &  $\\ldots$  \\\\ B0355+54     &     4.5 &    563 &   2.1  &  210       &   &  2.29 &    --5.2  &    --4.8  & &     22.9 &    --4.7  &     U      \\\\ B0540+23     &     4.1 &    253 &   3.5  &  348       &   &  0.77 &    --4.5  &    --5.1  & &      7.7 &    --4.5  &    --3.3   \\\\ B1754--24     &     4.0 &    285 &   3.5 &  $\\ldots$  &   &  0.71 &    --5.0  &  $\\ldots$  & &      7.1 &    --5.0  &  $\\ldots$  \\\\ B0656+14     &     3.8 &    111 &   0.8  &  331       &   &  3.63 &    --6.0  &    --5.2  &  &   36.3 &    --5.1  &     U      \\\\ B1719--37     &     3.3 &    345 &   2.5  & $\\ldots$  &   &  0.89 &    --4.6  &    --5.4  &  &    8.9 &    --4.6  &    --3.4   \\\\ B1821--19     &     3.0 &    574 &   5.2  & $\\ldots$  &   &  0.41 &    --4.4  &    --4.7  &  &    4.1 &    --4.4  &    --3.3   \\\\ B1055--52     &     3.0 &    535 &   1.5  & 440       &   &  1.21 &    --5.8  &  $\\ldots$ &  &   12.1 &    --5.7  &  $\\ldots$  \\\\ B2011+38     &     2.9 &    412 &  13.1  & $\\ldots$  &    & 0.16 &    --2.9  &    --4.0  &   &   1.6 &    --2.9  &    --3.4   \\\\ B0136+57     &     2.1 &    403 &   2.9  & 340       &    & 0.69 &    --4.3  &    --4.6  &   &   6.9 &    --4.3  &    --3.3   \\\\ B1449--64     &     1.9 &   1035 &   1.8  & 337       &   &  1.04 &    --5.5  &  $\\ldots$ &  &   10.4 &    --5.5  &  $\\ldots$  \\\\ B1730--37     &     1.5 &    355 &   3.5  & $\\ldots$  &   &  0.44 &    --3.6  &  $\\ldots$ &  &    4.4 &    --3.6  &  $\\ldots$  \\\\ B1933+16     &     0.51 &    947 &   7.9  &  996       &  &   0.04 &    --3.0  &  $\\ldots$ & &     0.4 &    --3.0  &  $\\ldots$  \\\\ B1706--16     &     0.12 &   1655 &   1.3  & 186       &  &   0.59 &    --3.5  &  $\\ldots$ & &     5.9 &    --3.5  &  $\\ldots$  \\\\ B0736--40     &     0.089 &   3805 &   2.1  & 377       & &    0.21 &    --3.8  &  $\\ldots$ &&      2.1 &    --3.8  &  $\\ldots$  \\\\ B2148+63     &     0.012 &  36640 &  13.6 &  $\\ldots$  &  &   0.01 &    --0.7  &    --1.6  & &     0.1 &    --0.7  &    --1.6   \\\\ \\end{tabular} \\end{minipage} \\end{table*}  We can compare our upper limits to values of $\\epsilon$ for known PWN. In Table~3 of FS97, $\\epsilon$ is listed for the six pulsars then known to have detected radio PWN, to which we add PWN which have since been associated with PSR~J0537--6910 ($\\epsilon = 5 \\times 10^{-4}$; Lazendic \\& Dickel 1998\\nocite{ld98}), PSR~B0906--49 ($\\epsilon = 2\\times 10^{-6}$; GSFJ98) and PSR~J1811--1926 ($\\epsilon < 2 \\times 10^{-3}$; Morsi \\& Reich 1987\\nocite{mr87}; Torii \\etal\\ 1999\\nocite{ttd+99}). For these nine sources, six measurements lie in the reasonably narrow range $(1-5)\\times10^{-4}$. Those sources lying outside this range are PSR~B0833--45, for which it is unclear just what part of the surrounding SNR is pulsar-powered, PSR~B0906--49, which GSFJ98 argue is substantially different from other PWN, and PSR~J1811--1926, where the radio PWN is faint and has poorly constrained properties.  Thus the data available suggest that a ``typical'' detectable PWN has $\\epsilon \\approx 10^{-4}$.  FS97 argued that typical non-detections corresponded to $\\epsilon_{\\rm max} \\sim 2\\times10^{-6}$, significantly less than for detected PWN. However, it can be seen from the results in Table~\\ref{tab_limits} that for $n=0.003$~cm$^{-3}$ some PWN become too large to be detected by their data, while for the remaining pulsars the limit rises to $\\epsilon_{\\rm max} \\sim 3\\times10^{-4}$.  Thus we argue that the observations of FS97 could have missed ``normal'' PWN around most of their sample if most of these sources are in low density regions, and hence their observations do not constrain pulsars which lack detectable PWN to be any different in their wind properties from those pulsars with observed PWN.  On the other hand, the current observations can potentially provide constraining limits on $\\epsilon$.  We first consider the six young and energetic pulsars in our sample (i.e. the first six pulsars in Table~\\ref{tab_limits}). These pulsars are defined approximately by $\\dot{E}_{34} \\ga 50$ and $t_3 \\la 50$, properties similar to those pulsars with detectable radio PWN. For either assumed ISM density, our data constrain these pulsars to have upper limits on $\\epsilon$ in the range $(0.002-0.2)\\times10^{-4}$, significantly less than for pulsars with observed PWN. While this seems to imply genuinely low values of $\\epsilon$, for $n=0.003$~cm$^{-3}$, conditions are such that these pulsars have only recently overtaken their static nebulae. Lowering their velocities slightly, arguing that their actual ages are less than their characteristic ages, or accepting that realistically, the transition from static to bow-shock PWN does not happen instantaneously, it seems likely that these pulsars are still producing static nebulae, whose extents are much larger than for bow-shocks. In this case, the corresponding limits become $\\epsilon_{\\rm max} \\gg 10^{-4}$, and are not constraining.  We thus argue that our non-detections of PWN can be explained even if all young and energetic pulsars have similar wind properties. The difference between detectable and non-detectable PWN seems to be that detected PWN are either in dense regions of the ISM or in SNRs, in which there is sufficient external pressure to confine the pulsar wind and produce an observable PWN. However, pulsars with no PWN are in the low density phase of the ISM and so produce unobservable ``ghost remnants'' \\cite{bopr73}.  With the exception of PSR~B0906--49 (GSFJ98\\nocite{gsfj98}), pulsars with observed PWN also have associated SNRs, while all those young pulsars without PWN also have no associated SNRs. We indeed expect SNRs to be undetectable in the hot component of the ISM \\cite{ksbg80,gj95c}, consistent with our conclusion above that it is a low ambient density which causes a PWN around a young pulsar to be undetectable.  A notable exception is PSR~B1757--24, which is associated with both a SNR and a radio PWN despite having $n=0.003$~cm$^{-3}$ \\cite{fk91,mkj+91}. In this case, the pulsar is inferred to have a transverse space velocity $\\sim$1500~\\kms\\ \\cite{fkw94}; this not only supplies the necessary ram pressure to produce an observable PWN, but has caused the pulsar to overtake the shell of the associated SNR, re-energising the remnant with its passage. If it were not for the extreme velocity of the pulsar, neither the PWN nor the SNR would be detectable, as expected in a low density region.  The relative numbers of young pulsars with detected and undetected SNRs/PWN suggests an approximate filling fraction $\\sim$50\\% for the low density component of the ISM. This is somewhat more than recent estimates of 15--20\\% \\cite{fer98b}, but can be explained by the fact that we expect young pulsars to be preferentially located in low density regions produced by the powerful winds of their progenitors and by previous supernovae in the region.  The majority of pulsars in our sample are considerably less energetic ($\\dot{E}_{34} < 50$) and older ($t_3 \\ga 100$) than pulsars around which radio PWN have been observed.  Values of $\\epsilon_{\\rm max}$ for these pulsars are plotted in Fig~\\ref{fig_histo1}, from which it can be seen that for either choice of ambient density, the distribution peaks around $\\epsilon_{\\rm max} \\sim 10^{-5}$. These pulsars have all long since overtaken their static PWN, and will have bow shock PWN for any sensible choice of $n$ and $V_{\\rm PSR}$. The resulting size of such a PWN is not a strong function of $n$ or $V_{\\rm PSR}$; to produce values of $\\epsilon_{\\rm max} > 10^{-4}$, consistent with detected PWN, requires uniformly low space velocities, $V_{\\rm PSR} \\la 150$~\\kms, for these pulsars. However, only $\\sim$5\\% of pulsars are thought to be traveling at less than 150~\\kms\\ \\cite{cc98}. While the ages and distances we have used for these pulsars have their associated uncertainties, for bow-shock PWN values of $\\epsilon_{\\rm max}$ are independent of age, and distances would have to uniformly be increased by a factor of three to shift the peak in $\\epsilon_{\\rm max}$ to a value in agreement with that seen for detected PWN.  The lower value of $\\epsilon$ we have derived for older pulsars is thus a result quite robust to the assumptions and uncertainties involved in its derivation, and we therefore argue that these pulsars have winds which are genuinely at least an order of magnitude less efficient at producing radio emission than the winds of young and energetic pulsars.  \\begin{figure*} \\begin{minipage}{160mm} \\centerline{\\psfig{file=fig_histo.eps,width=14cm,angle=270}} \\caption{Upper limits on $\\epsilon = L_R / \\dot{E}$ from pulsar-gated 1.4~GHz data in Table~\\ref{tab_limits}; pulsars with $\\dot{E}_{34} > 50$ have been excluded. The upper panel corresponds to an assumed ambient density $n=0.3$~cm$^{-3}$, while the lower panel represents $n=0.003$~cm$^{-3}$.} \\label{fig_histo1} \\label{lastpage} \\end{minipage} \\end{figure*}  FS97 consider various reasons why older pulsars might appear to have lower values of $\\epsilon$. Possibilities include: \\begin{enumerate} \\item that the PWN are resolved out by the observations;  \\item that an increasing fraction of $\\dot{E}$ goes into pulsed X-rays and $\\gamma$-rays (see Thompson \\etal\\ 1994\\nocite{tab+94});  \\item that their winds are dominated by Poynting flux rather than relativistic particles;  \\item that the injection spectrum of particles in the pulsar wind has shifted to higher energies. \\end{enumerate}  Our observations can conclusively rule out alternative (i), as we can detect PWN produced for almost all feasible values of $n$ and $V_{\\rm PSR}$. While we cannot distinguish between the remaining three possibilities, we note that of detected PWN, that associated with the oldest pulsar, PSR~B0906--49, also has the lowest value of $\\epsilon$ and the steepest spectral index (GSFJ98). Since the spectral index of a PWN is directly related to the spectrum of injected particles \\cite{ps73}, this result tentatively suggests that the efficiency of the wind in producing radio emission is related to the injection spectrum, and that alternative (iv) might then best explain the observations."
    },
    "0004/astro-ph0004386_arXiv.txt": {
        "abstract": "We have obtained images of the Trapezium Cluster ($140\\arcsec\\times140\\arcsec$; 0.3~pc~$\\times$~0.3~pc) with the Hubble Space Telescope Near-Infrared Camera and Multi-Object Spectrometer (NICMOS). Combining these data with new ground-based $K$-band spectra ($R=800$) and existing spectral types and photometry, we have constructed an H-R diagram and used it and other arguments to infer masses and ages. To allow comparison with the results of our previous studies of IC~348 and $\\rho$~Oph, we first use the models of D'Antona \\& Mazzitelli. With these models, the distributions of ages of comparable samples of stars in the Trapezium, $\\rho$~Oph, and IC~348 indicate median ages of $\\sim0.4$~Myr for the first two regions and $\\sim1$-2~Myr for the latter. The low-mass IMFs in these sites of clustered star formation are similar over a wide range of stellar densities ($\\rho$~Oph, $n=0.2$-$1\\times10^3$~pc$^{-3}$; IC~348, $n=1\\times10^3$~pc$^{-3}$; Trapezium, $n=1$-$5\\times10^4$~pc$^{-3}$) and other environmental conditions (e.g., presence or absence of OB stars). With current data, we cannot rule out modest variations in the substellar mass functions among these clusters. We then make the best estimate of the true form of the IMF in the Trapezium by using the evolutionary models of Baraffe et al.\\ and an empirically adjusted temperature scale and compare this mass function to recent results for the Pleiades and the field. All of these data are consistent with an IMF that is flat or rises slowly from the substellar regime to about 0.6~$M_{\\odot}$, and then rolls over into a power law that continues from about 1~$M_{\\odot}$ to higher masses with a slope similar to or somewhat larger than the Salpeter value of 1.35. For the Trapezium, this behavior holds from our completeness limit of $\\sim0.02$~$M_{\\odot}$ and probably, after a modest completeness correction, even from 0.01-0.02~$M_{\\odot}$. These data include $\\sim50$ likely brown dwarfs. We test the predictions of theories of the IMF against 1) the shape of the IMF, which is not log-normal, in clusters and the field, 2) the similarity of the IMFs among young clusters, 3) the lowest mass observed for brown dwarfs, and 4) the suggested connection between the stellar IMF and the mass function of pre-stellar clumps. In particular, most models do not predict the formation of the moderately large numbers of isolated objects down to 0.01~$M_{\\odot}$ that we find in the Trapezium. ",
        "introduction": "What is the true form of the stellar initial mass function (IMF)? Is it universal, or does it depend on the properties of the natal molecular cloud or of the embedded, young stellar population?  In particular, the turnover mass and the minimum mass of the IMF and their behavior with various star forming conditions can offer vital insights into the physical processes that regulate the formation of stars and brown dwarfs (Elmegreen 1999b).  Numerous techniques have been used to measure the IMF (see reviews by Scalo 1998; Elmegreen 1999b).  Stellar open clusters have played an important role, but difficulties arise in determining cluster membership and completeness at low masses, and in properly accounting for dynamical evolution and mass segregation. These problems can be alleviated in the youngest clusters ($\\leq10$~Myr) associated with star forming regions. The compact nature and thick molecular cloud of a star-forming cluster can greatly reduce contamination by foreground and background stars. Newborn substellar objects are quite luminous and should have the same spatial distribution as the stars since these regions are too young to have undergone significant dynamical evolution. Furthermore, both the initial conditions of star formation and the resulting mass function are directly observable in the youngest clusters. Stellar populations in clusters are also relevant for comparisons to the field and other regions since it is likely that clusters represent the dominant mode of star formation in the Galaxy (Lada, Strom, \\& Myers 1993).  The Orion Nebula Cluster centered on the Trapezium OB stars is the richest of any nearby cluster and has been studied extensively through proper motions (Jones \\& Walker 1988), optical images from the ground (Herbig \\& Terndrup 1986) and space (Prosser et al.\\ 1994), wide-field infrared (IR) images (Ali \\& DePoy 1995), and high-resolution ground-based IR images (McCaughrean \\& Stauffer 1994, hereafter MS; Petr et al.\\ 1998; Simon, Close, \\& Beck 1999). Hillenbrand (1997) combined new optical spectroscopy and photometry with previous data from the literature for more than 1000 stars within $18\\arcmin$ (2.5~pc) of the Trapezium OB stars and constructed a Hertzsprung-Russell (H-R) diagram for the Orion cluster. With the theoretical evolutionary models of D'Antona \\& Mazzitelli (1994) (hereafter DM94), she inferred an average age of $<1$~Myr and an IMF that peaked at 0.2~$M_{\\odot}$ and fell rapidly to lower masses. Hillenbrand \\& Carpenter (2000) have recently extended this work to substellar masses through deep ground-based $H$ and $K$ imaging of the $5\\farcm1\\times5\\farcm1$ area surrounding the Trapezium OB stars.  Within the central $\\sim5$~arcmin$^2$ of the Orion Nebula Cluster lies the Trapezium Cluster, where stellar densities reach a peak of $\\sim5\\times10^4$~pc$^{-3}$ (MS). Several characteristics make the Trapezium Cluster a unique region for a study of the IMF. The cluster is rich (300 stars) and nearby (450~pc) and a majority of its members have minimal extinction ($A_V<5$) because of the cluster's location on the front of the molecular cloud and within the cavity created by the O star $\\theta^1$~Ori~C. In addition, the obscuration of the molecular cloud and the compactness of the cluster minimize contamination from background and foreground stars. Because of the abrupt change in reddening from the Trapezium members to the field stars behind the cloud, virtually all sources with $A_V<10$ should be cluster members, a crucial property in reliably identifying the substellar population. The special viewing geometry for the Trapezium Cluster overcomes many of the common limitations in correcting cluster measurements for the effects of obscuration. Consequently, it should be possible to construct an IMF from well below the hydrogen burning mass limit to about 50~$M_{\\odot}$.  In \\S~\\ref{sec:trap} of this paper, we describe a new study of the Trapezium Cluster.  Previous observations of the low-mass population in the Trapezium have been hindered by crowding and bright nebulosity. To overcome these obstacles, we have obtained sensitive ($H\\sim17$) high-resolution ($0\\farcs2$) images of the Trapezium Cluster ($140\\arcsec\\times140\\arcsec$) with the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) aboard the {\\it Hubble Space Telescope (HST)}. We have also measured $K$-band spectra for $\\sim100$ sources in this region. We use these measurements and data from the literature (MS; Prosser et al.\\ 1994; Hillenbrand 1997) to construct an H-R diagram for the cluster and infer individual masses and ages from theoretical evolutionary models (D'Antona \\& Mazzitelli 1997, hereafter DM97; Baraffe et al.\\ 1998, hereafter B98). The resulting star formation history is used to estimate masses for the faint sources that lack spectral types, which are combined with the masses of stars on the H-R diagram to produce a cluster IMF that reaches down to 0.01~$M_{\\odot}$.  In \\S~\\ref{sec:compare}, we compare the Trapezium IMF with the similarly derived mass functions for the star-forming clusters IC~348 (Luhman et al.\\ 1998, hereafter LRLL; Luhman 1999) and $\\rho$~Oph (Luhman \\& Rieke 1999, hereafter LR99), showing them all to be similar. Interpretation of young cluster observations is limited by the accuracy of the theoretical models for the evolution of young stars and brown dwarfs. It is now possible to mitigate this problem by testing evolutionary tracks and temperature scales against young multiple star systems that contain coeval stars of differing mass. If we use a combination of tracks (B98) and temperature scale that is closely consistent with such observational tests (see Luhman 1999), the IMF for the Trapezium and the other clusters is similar to recently determined mass functions for the Pleiades and M35 open clusters and the field. The shape of the low-mass IMF, its approximate invariance across two orders of magnitude in stellar density, and the presence of moderately large numbers of very low-mass brown dwarfs are used to test the predictions of various theoretical models for the origin of the IMF. ",
        "conclusions": "The shape of the low-mass IMF ($<1$~$M_{\\odot}$) has remained uncertain because of several issues, the most important being incompleteness for faint low-mass stars and brown dwarfs. In observations summarized in this work, we have made reliable measurements of the IMF down to well below the hydrogen burning limit and have presented a robust comparison of the low-mass IMFs in local regions of clustered star formation. Our results are as follows:  \\begin{enumerate}  \\item Despite a range of two orders of magnitude in the density of star formation, the {\\it stellar} IMFs in the core of IC~348 and the cloud core of $\\rho$~Oph are the same as in the Trapezium within the uncertainties. The IMFs in the {\\it brown dwarf} regime are also roughly similar, although more observations are required for a definitive comparison of the substellar mass functions.  \\item Data for the Trapezium, the Pleiades and M35 open clusters (Bouvier et al.\\ 1998; Barrado y Navasc\\'ues et al. 2000), and the field (Reid et al.\\ 1999) are consistent with the same IMF; this mass function is flat or slowly rising from the brown dwarf regime to 0.6-1.0~$M_\\odot$, where it rolls over to a power law with a slope of $\\sim1.7$, similar to or slightly steeper than the Salpeter value of 1.35.  The similarity of these mass functions is consistent with the suggestion that members of the field have formed predominantly in clusters.  \\item Whereas the IMF that characterizes young clusters and the field rolls over near 0.8~$M_\\odot$ and is not log-normal, recent studies have found that the IMFs of globular clusters peak near 0.3~$M_\\odot$ and can be described by a log-normal function (Paresce \\& De Marchi 1999). As a result, there appear to be fundamental differences in the IMF between globular clusters and Galactic disk clusters.  \\item Using the high spatial resolution and sensitivity of NICMOS images, we have penetrated the bright nebulosity of the Trapezium and identified a population of $\\sim50$ likely brown dwarfs, where the least massive candidate is $\\sim0.01$~$M_{\\odot}$ if the median age of the Trapezium is assumed. Most theories of the IMF do not predict the formation of free-floating objects in significant numbers at such low masses. For instance, this low-mass population rules out a log-normal IMF in this cluster, contrary to some theories of star formation.  \\end{enumerate}"
    },
    "0004/astro-ph0004392_arXiv.txt": {
        "abstract": "The gravitational lensing, as well as the velocity field and the cosmological light-cone warp, changes the observed correlation function of high-redshift objects. We present an analytical expression of 3D correlation function, simultaneously including those three effects. When two objects are separated over several hundreds Mpc along a line of sight, the observed correlation function is dominated by the effect of gravitational lensing rather than the intrinsic clustering. For a canonical lambda-CDM model, the lensing signals in the galaxy-galaxy and galaxy-QSO correlations are beyond noise levels in large-scale redshift surveys like the Sloan Digital Sky Survey. ",
        "introduction": "\\label{sec1}  The correlation function is one of the fundamental quantities in searching into the physical origin of the universe. In large-scale redshift surveys, redshifts and spherical positions on the sky of luminous objects are used for estimating the spatial distribution of mass, but the former is distorted by inhomogeneity of the universe.  Two intrinsic distortion effect on the correlation function originate in the velocity field \\citep{kai87} and the cosmological warp \\citep{bal96,mat96,mat00}. The former comes from the fact that an observed redshift corresponds to the recession velocity which is composed not only of the expansion of the universe but also of peculiar velocities. The latter distortion is brought about by the nonlinear mapping of the objects from the expanding curved space on a light cone onto a flat redshift space \\citep{alc79}.  In addition to those intrinsic distortions, there are secondary distortions which are due to perturbations of light rays. While the redshift is altered by the Sachs-Wolfe effect \\citep{sac67}, the spherical position is recast by the gravitational lensing \\citep{sch92,mel99,bar99}. Although the Sachs-Wolfe effect is not so important, the gravitational lensing can affect the observable correlation function in forthcoming redshift surveys as this effect is efficient for high-redshift objects \\citep{gun67}.  The effect of the gravitational lensing on the angular functions $w(\\theta)$ have been intensively investigated so far \\citep{bar99,moe98,kai92,vil96}. Recently, among others, cross correlations of galaxies at different redshifts \\citep{moe98} are successfully applied to the commissioning data from the Sloan Digital Sky Survey \\citep{jai00}, to single out the weak lensing effect.  As for the 3D correlation function in redshift space, a qualitative treatment of the lensing effect is given in \\citet{sut99} to estimate the upper limit of the effect, using the phenomenological Dyer-Roeder distance \\citep{dye73}. It is still not clear whether or not the lensing is actually efficient where the intrinsic correlation function is negligible on scales comparable to 100Mpc. The main purpose of this {\\em letter} is to give a quantitative treatment of this issue, consistently including velocity and cosmological distortions, and consequently to show the weak lensing actually has detectable effects on 3D correlation function in redshift space when large-scale redshift surveys like Sloan Digital Sky Survey (SDSS) are available. ",
        "conclusions": "\\label{sec4}  The absolute value of the intrinsic clustering component $\\xi_{\\rm MS}$ is a decreasing function of the line-of-sight separation, $z_2 - z_1$, except the vicinity of zero crossings. On the other hand, the lens-lens component $\\xi_{\\rm ll}$ is almost independent on the separation and the density-lens component $\\xi_{\\rm rl}$ is an increasing function. Those behaviors are understood by the fact that the weak lensing is efficient between the object and the observer. Thus, intrinsic clustering component dominates for small separations, while lens-lens and/or density-lens components dominates for large separations along the direction of line of sight.  The lower panels in Figure~\\ref{fig1} show the region where lensing contribution dominates in the case of the SDSS magnitude limits to illustrate the typical magnitude of correlations. Are those lensing signals detectable? The statistical uncertainty in estimating the correlation function is given by $(\\delta\\xi)^2 = \\Omega/(\\delta\\Omega N_1 N_2)$ \\citep{pee80}, where $N_1$ and $N_2$ are numbers of object in the bin used for redshifts $z_1$ and $z_2$, respectively, $\\Omega$ is the solid angle subtended by the survey area, and $\\delta\\Omega$ is the fraction in the bin used for angle $\\theta$. To increase the signal to noise ratio, it is desirable to use large bins for $\\theta$. To be specific, we consider a bin $[1',10']$, and theoretical curves are integrated accordingly, so that $\\delta\\Omega \\sim 100 \\pi\\ [\\mbox{arcmin}^2]$. The effective scale of this bin is given by $\\theta_{\\rm eff} = 10'/\\sqrt{2}$ for $\\xi \\propto \\theta^{-1}$. In the SDSS, $\\Omega \\sim \\pi\\ \\mbox{[str]} \\sim 1.2 \\pi \\times 10^7\\ [\\mbox{arcmin}^2]$, and the estimated numbers of galaxies and quasars are $N_{\\rm G} = 10^6$ and $N_{\\rm Q} = 1.7 \\times 10^5$, respectively. Assuming we take sufficiently large bins of redshifts (this choice is similar to considering angular correlation functions), the consequent estimates of the statistical error are given by $5.0 \\times 10^{-4}$ for G-G, $8.5\\times 10^{-4}$ for G-Q, and $2.9 \\times 10^{-3}$ for Q-Q correlations, which are plotted in lower panels. The S/N ratios turn out to be about $10$, $1.3$, and $0.22$ for G-G, G-Q, and Q-Q correlations, respectively.  Therefore, the weak lensing in 3D correlation function of galaxies in the SDSS is definitely detectable, and the detection of galaxy-QSO cross-correlation is marginal, while the quasar correlation by lensing is below the noise level in the SDSS. In order to detect the QSO-QSO lensing effect, the sample should be at least 5 times larger than the SDSS, or parameters $\\sigma_8$, $\\Omega_0$, $b$, $\\alpha$ should be larger than assumed values.  In summary, we have obtained a theoretical prediction of correlation function in redshift space, taking into account the effect of weak lensing, together with velocity distortions and cosmological distortions on a light-cone. Each effect contributes differently to the correlation function, and is realistically detectable. Our result provides a fundamental link between theoretical models and the observed correlation function in the 3D redshift survey data. Besides the determination of the power spectrum itself, various cosmological parameters, especially the bias parameter, can be estimated by proper likelihood analyses, including KL transform of the correlation matrix \\citep{vog96,mat00a}. One may also be tempted to assume the cosmological parameters before analysing data. In which case, the error originated in choosing the wrong cosmological model is roughly given by the order of the redshift $z$ times the error of cosmological parameters, since the Alcock-Paczy\\'nski effect is roughly proportional to $z$ up to $z=1$-$2$."
    },
    "0004/astro-ph0004267_arXiv.txt": {
        "abstract": "We consider the X-ray emission due to bremsstrahlung processes from spherically-symmetric low radiative-efficiency hot accretion flows around supermassive and galactic black holes. We calculate surface brightness profiles and Michelson visibility functions for a range of density profiles, $\\rho \\sim r^{-3/2+p}$, with $0 < p < 1$, to allow for the presence of outflows. We find that although the 1~keV emitting region in these flows can always extend up to $10^6$ Schwarzschild radii ($R_{\\rm{S}}$), their surface brightness profiles and visibility functions are strongly affected by the specific density profile. The advection-dominated solutions with no outflows ($p=0$) lead to centrally peaked profiles with characteristic sizes of only a few tens of $R_{\\rm{S}}$. Solutions with strong outflows ($p \\sim 1$) lead to flat intensity profiles with significantly larger characteristic sizes of up to $10^6 R_{\\rm{S}}$. This implies that low luminosity galactic nuclei, such as M87, may appear as extended X-ray sources when observed with current X-ray imaging instruments.  We show that X-ray brightness profiles and their associated visibility functions may be powerful probes for determining the relevant mode of accretion and, in turn, the properties of hot accretion flows. We discuss the implications of our results for observations with the {\\em Chandra} X-ray Observatory and the proposed X-ray interferometer {\\em MAXIM}. ",
        "introduction": "Imaging X-ray emitting regions around black holes addresses questions related both to black hole physics and to accretion phenomena. Depending on the properties and the extent of the X-ray emission as well as its proximity to the black hole, the X-ray images may have imprinted on them the accretion geometry, its physical properties, and even signatures of the space-time around the black hole. In standard models, the X-ray emission from active galactic nuclei (AGN) is assumed to arise from a hot corona above a geometrically thin accretion disk. Thermal Comptonization of disk blackbody photons in the corona leads to the production of a hard X-ray continuum, some of which is reprocessed by the underlying disk (e.g., Haardt \\& Maraschi 1991, 1993). From the properties of the hard X-ray continuum and of the reflected component (e.g., Nandra et al.~1997) it is now established that the the majority of X-ray emission from AGN originates from regions close to the black holes, corresponding to $\\lesssim 3-20 R_{\\rm{S}}$, where $R_{\\rm{S}}$ is the Schwarzschild radius of the black hole.  In recent years is has become apparent that standard thin disk accretion with high radiative efficiency may not be common in nearby galaxies.  The centers of most nearby galaxies also harbor supermassive black holes (e.g., Magorrian et al. 1998) but, unlike the more distant AGN, display little or no activity. It has been shown that the nuclei of nearby giant elliptical galaxies, which host the largest black holes known with masses of $10^8-10^{10} M_{\\odot}$ (e.g., Magorrian et al. 1998), the center of the Milky Way with its $2.5 \\times 10^6 M_\\odot$ black hole (e.g., Eckart \\& Genzel 1997), and even some of the galactic black hole sources like A0620 in quiescence are remarkably underluminous for their expected mass accretion rates. The relative quiescence and spectral characteristics of these sources can be explained if the central black holes accrete via low radiative-efficiency accretion flows or ADAFs (Rees et al. 1982; Fabian \\& Rees 1995; Narayan, Yi, \\& Mahadevan 1995; Narayan, McClintock, \\& Yi 1996; Di Matteo et al. 2000; for a review see, e.g., Narayan, Mahadevan \\& Quataert 1998 and references therein).  In this class of accretion solutions the flows are two-temperature with the electron temperature ranging from $\\approx 1~\\rm{MeV}$ in the inner regions to about 1 keV in the outer regions. At these temperatures, and when the gas is optically thin to electron scattering so that Comptonization is unimportant, most of the cooling occurs via thermal bremsstrahlung emission. In recent work it has been shown that the hard X-ray components observed in elliptical galaxy systems (Allen, Di Matteo \\& Fabian 2000) are indeed consistent with models of thermal bremsstrahlung emission (Di Matteo et al.~2000). In the same work, it has also been emphasized that bremsstrahlung emission may always be the dominant cooling mechanism because hot accretion flows are likely to drive strong mass loss (e.g. Igumenshchev \\& Abramowicz 1999; Blandford \\& Begelman 1999) so that their densities and scattering optical depths always remain low and Comptonization is never important. The same also applies to the case in which convection becomes important in hot flows, as discussed in recent work by Narayan, Igumenshchev \\& Abramowicz (2000), Quataert \\& Gruzinov (2000), and Ball, Narayan, \\& Quataert (2000).  The X-ray emission from the nuclei of nearby giant elliptical galaxies is therefore expected to have significantly different properties than that of AGN, reflecting the physical conditions of hot, extended accretion flows. In this paper we go beyond the spectral calculations carried out in earlier work and provide a set of diagnostics of the bremsstrahlung X-ray emitting region of hot accretion flows. This includes calculating the expected surface brightness profiles relevant for determining the sizes and density profiles of the X-ray emitting regions of these flows. The calculations of these quantities may have a significant impact for explaining the ongoing and upcoming X-ray observations of the cores of galactic nuclei. Conversely, the measurement of these quantities may reveal the relevant mode of accretion and the detailed properties of the accretion flows, improving our understanding of physical processes, such as convection and generation of outflows.  Here we discuss the implications of our results for observations with the {\\em Chandra} X-ray Observatory and the proposed future space-based X-ray interferometer {\\em MAXIM}.  In \\S 2, we review bremsstrahlung emission from hot flows and introduce the X-ray diagnostics that we study in this paper. In \\S 3, we present the surface brightness profiles and visibility curves for different black hole masses and accretion-flow density profiles. Finally in \\S 4, we discuss the implications of our results for current and possible upcoming X-ray missions. ",
        "conclusions": "High resolution X-ray imaging and interferometric studies, which measure directly or indirectly the brightness profiles and thus the sizes and density profiles of accretion flows around black holes, can be powerful probes for studying the relevant mode of accretion and the detailed properties of accretion flows. We have shown that the 1~keV emitting regions of hot accretion flows can be large and extend up to $10^6 R_{\\rm{S}}$ irrespective of the specifics of the density profile, provided that the ADAF itself extends out to this radius (see below for further discussion). However, the brightness profiles and visibility functions are strongly affected by the specific density profile. In particular, the flat density profiles in the presence of strong outflows (p=1) give rise to very flat surface brightness profiles with characteristic sizes of $10^6 R_{\\rm{S}}$. In contrast, the purely inflow advection dominated solutions have centrally peaked profiles with characteristic sizes of only a few tens of $R_{\\rm{S}}$.  \\vbox{ \\centerline{ \\psfig{file=Fig4.ps,width=9.0truecm} } \\figcaption[]{ \\footnotesize The resolvable scale in units of $R_{\\rm{S}}$ of each source plotted against the known distance to the sources.  If the black holes are accreting via a hot flow, these resolvable sizes represent the largest size the accretion flow can have to be resolved by a $\\it{minimum}$ baseline of 1m. }} \\vspace*{0.5cm}  If we are correct in surmising that hot accretion flows may be common in nearby galactic nuclei and that such flows may drive strong outflows (e.g., Quataert \\& Narayan 1999; Di Matteo et al. 2000) then the above results have strong implications for X-ray observations of these nuclei. An important consequence of our results is that a hot accretion flow with a flat density profile around a $\\sim 10^9 M_\\odot$ black hole at Virgo cluster distances, such as M87, may look like an extended X-ray source rather than a point source when imaged by the $\\it{Chandra}$ X-ray observatory with its 0.5 arcsec resolution mirrors. At the distance of M87 and for $m \\sim 3\\times 10^9$ (Macchetto et al. 1997), 0.5 arcsec corresponds to $10^5 R_{\\rm{S}}$ which is about an order of magnitude smaller than the extent of the regions of these flows emitting soft X-rays (at $\\lesssim 2$ keV; where most of the photon counts would be at these fluxes). Conversely, the detection of an X-ray point source in elliptical nuclei in Virgo must imply that either the physical sizes of the hot accretion flows are smaller than $10^5 R_{\\rm{S}}$ or that standard thin disk accretion with its associated corona applies. Regarding the former possibility, we stress that in our calculation we have assumed that the hot flows extend to their maximum possible X-ray emitting radius, i.e., $R \\sim 10^6 R_{\\rm{S}}$. This radius is consistent with the typical accretion radius (the radius at which the black hole gravitational potential dominates that of the galaxy) of elliptical galaxies ($\\sim 0.1$ kpc; Di Matteo et al. 2000) but does not need to coincide with the radius at which the angular momentum dominated accretion flow, i.e. an ADAF, is formed. There may be physical processes such as cooling at large radii that may cause the ADAF flow to start at smaller radii, or there may be transitions to other accretion solutions in the outer regions (see, e.g., Igumenshchev, Abramowicz, \\& Novikov 1998; Menou et al. 1999). Here, we assumed a direct transition from the hot interstellar medium (ISM) of the elliptical galaxies to an ADAF.  We note also that this hot ISM of the elliptical galaxies, with typical temperature of a few keV, may also contribute significantly to the emission in the soft X-ray band, which would further blend the images and may make the detection of the accretion flow component more difficult. However, in the hard band ($> 5$~keV), the sources are intrinsically smaller (c.f., Fig.~3), and may look more point-like provided that there are enough photon counts. This would have the additional advantage of having less background from galactic emission. Finally, we stress that {\\em Chandra} X-ray spectroscopy of the central regions of these galaxies will also complement the information provided by the imaging instruments and will be crucial for determining the presence of ADAFs in these systems. With its 1 MeV temperature, the bremsstrahlung emission from the central regions of a hot accretion flow will be much harder than that of the galaxy ISM and will have a different spectral slope than that expected from a typical AGN with low luminosity.  Variability studies are also important for determining the sizes of accretion flows and are complementary to imaging studies. Rapid variation in the X-ray flux may help constrain the size of the X-ray emitting region. For example, {\\it ROSAT} HRI observations (at $E \\sim 1$ keV) of the ``core'' of M87 show $\\approx 20 \\%$ variability on timescales of $\\approx 6$ months to a year (Harris, Biretta, \\& Junor 1997).  This is not easy to reconcile with a $p > 0.5$ bremsstrahlung model for the X-ray emission. However, identifying the possible contributions to the variability from jets associated with the accretion flows relies on high resolution imaging studies and therefore the two methods can be used as complementary probes.  Because of the large inferred black hole masses, many of the nearby supermassive black holes have also been considered some of the best targets for observations with the planned X-ray interferometry mission {\\em MAXIM}.  Owing to the short X-ray wavelengths, remarkably high spatial resolutions can be achieved with relatively short baselines with the ultimate aim of imaging of black holes on even horizon scales. (See http://maxim.gsfc.nasa.gov/ for a detailed discussion of the proposed capabilities of {\\em MAXIM} and Cash et al.~2000). Already for the first stages of such a mission, a target resolution of $0.1$ mas has been planned corresponding to a baseline of $\\sim 1~m$ at 1~keV. This should probe up to a few tens of $R_{\\rm{S}}$ for the supermassive black holes in nearby galaxies (Fig~4).  However, if indeed large black holes in nearby galaxies accrete via hot, radiatively inefficient accretion flows, e.g, M87, one important consequence of our results is that characteristic X-ray emission scale would be large and there would be little power on scales corresponding to few tens of $R_{\\rm{S}}$ or smaller. Therefore baselines of $\\sim 1$~m may not be appropriate to image black holes in the Virgo ellipticals (c.f., \\S 3 and Fig.~2) and a range of smaller baselines would be necessary to study the visibility functions of these extended flows. Thus, to study properties of accretion flows, the minimum baseline of the interferometer is an important consideration. Specifically, for studying the spatial extent of hot accretion flows and the relevance of outflows, it is important to bridge a possible gap in resolution between that of the $\\it{Chandra}$ X-ray observatory and the minimum resolution of an X-ray interferometer.  Alternatively, with a baseline of $\\sim 1$~m, galactic black hole candidates may be better targets to study hot accretion flows (Fig.~2). $\\it{Chandra}$ observations will be crucial for selecting the appropriate targets for an X-ray interferometry mission and giving us clues on the relevant mode of accretion in the nearby black holes."
    },
    "0004/gr-qc0004016_arXiv.txt": {
        "abstract": " ",
        "introduction": "The primary purpose of all physical theory is rooted in reality, and most relativists pretend to be physicists. We may often be members of departments of mathematics and our work oriented towards the mathematical aspects of Einstein's theory, but even those of us who hold a permanent position on ``scri'', are primarily looking there for gravitational waves. Of course, the builder of this theory and its field equations was {\\it the} physicist. J\\\"urgen Ehlers has always been very much interested in the conceptual and axiomatic foundations of physical theories and their rigorous, mathematically elegant formulation; but he has also developed and emphasized the importance of such areas of relativity as kinetic theory, the mechanics of continuous media, thermodynamics and, more recently, gravitational lensing. Feynman expressed his view on the relation of physics to mathematics as follows \\cite{FEM}:  ``The physicist is always interested in the special case; he is never interested in the general case. He is talking about something; he is not talking abstractly about anything. He wants to discuss the gravity law in three dimensions; he never wants the arbitrary force case in $n$ dimensions. So a certain amount of reducing is necessary, because the mathematicians have prepared these things for a wide range of problems. This is very useful, and later on it always turns out that the poor physicist has to come back and say, `Excuse me, when you wanted to tell me about four dimensions...' '' Of course, this is Feynman, and from 1965...  However, physicists are still rightly impressed by special explicit formulae. Explicit solutions enable us to discriminate more easily between a ``physical'' and ``pathological'' feature. Where are there singularities? What is their character? How do test particles and fields behave in given background spacetimes? What are their global structures? Is a solution stable and, in some sense, generic? Clearly, such questions have been asked not only within general relativity.  By studying a {\\it special} explicit solution one acquires an intuition which, in turn, stimulates further questions relevant to more {\\it general} situations. Consider, for example, charged black holes as described by the Reissner-Nordstr\\\"om solution. We have learned that in their interior a Cauchy horizon exists and that the singularities are timelike. We shall discuss this in greater detail in Section 3.1. The singularities can be seen by, and thus exert an influence on, an observer travelling in their neighborhood. However, will this violation of the (strong) cosmic censorship persist when the black hole is perturbed by weak (``linear'') or even strong (``nonlinear'') perturbations? We shall see that, remarkably, this question can also be studied by explicit exact special model solutions. Still more surprisingly, perhaps, a similar question can be addressed and analyzed by means of explicit solutions describing completely diverse situations -- the collisions of plane waves. As we shall note in Section 8.3, such collisions may develop Cauchy horizons and subsequent timelike singularities. The theory of black holes and the theory of colliding waves have intriguing structural similarities which, first of all, stem from the circumstance that in both cases there exist two symmetries, i.e. two Killing fields. What, however, about more general situations? This is a natural question inspired by the explicit solutions. Then ``the poor physicists have to come back'' to a mathematician, or today alternatively, to a numerical relativist, and hope that somehow they will firmly learn whether the cosmic censorship is the ``truth'', or that it has been a very inspirational, but in general false conjecture. However, even {\\it after} the formulation of a conjecture about a general situation inspired by particular exact solutions, {\\it newly} discovered exact solutions can play an important role in verifying, clarifying, modifying, or ruling out the conjecture. And also ``old'' solutions may turn out to act as asymptotic states of general classes of models, and so become still more significant.  Exact explicit solutions have played a crucial role in the development of many areas of physics and astrophysics. Later on in this Introduction we will take note of some general features which are specific to the solutions of Einstein's equations. Before that, however, for illustration and comparison we shall indicate briefly with a few examples what influence exact explicit solutions have had in other physical theories. Our next introductory excursion, in Section 1.2, describes in some detail the (especially early) history of Einstein's route to the gravitational field equations for which his short stay in Prague was of great significance. The role of Ernst Mach (who spent 28 years in Prague before Einstein) in the construction of the first modern cosmological model, the Einstein static universe, is also touched upon. Section 1.3 is devoted to a few remarks on some old and new impacts of the other simplest ``cosmological'' solutions of Einstein's equations -- the Minkowski, the de Sitter, and the anti de Sitter spacetimes. Some specific features of solutions in Einstein's theory, such as the observability and interpretation of metrics, the role of general covariance, the problem of the equivalence of two metrics, and of geometrical characterization of solutions are mentioned in Section 1.4. Finally, in the last (sub)sections of the ``Introduction'' we give some reasons why we consider our choice of solutions to be ``a natural selection'', and we briefly outline the main body of the article. \\newpage  \\subsection{A word on the role of explicit solutions in other parts of physics and astrophysics}  Even in a linear theory like Maxwell's electrodynamics one needs a good sample, a useful kit, of exact fields like the homogeneous field, the Coulomb field, the dipole, the quadrupole and other simple solutions, in order to gain a physical intuition and understanding of the theory. Similarly, of course, with the linearized theory of gravity. Going over to the Schr\\\"odinger equation of standard quantum mechanics, again a linear theory, consider what we have learned from simple, explicitly soluble problems like the linear and the three-dimensional harmonic oscillator, or particles in potential wells of various shapes. We have acquired, for example, a transparent insight into such basic quantum phenomena as the existence of minimum energy states whose energy is not zero, and their associated wave functions which have a certain spatial extent, in contrast to classical mechanics. The three-dimensional problems have taught us, among other things, about the degeneracy of the energy levels. The case of the harmonic oscillator is, of course, very exceptional since Hamiltonians of the same type appear in all problems involving quantized oscillations. One encounters them in quantum electrodynamics, quantum field theory, and likewise in the theory of molecular and crystalline vibrations. It is thus perhaps not so surprising that the Hamiltonian and the wave functions of the harmonic oscillator arise even in the minisuperspace models associated with the Hartle-Hawking no-boundary proposal for the wave function of the universe \\cite{HHT}, and in the minisuperspace model of homogeneous spherically symmetric dust filled universes \\cite{KP}.  In {\\it nonlinear} problems explicit solutions play still a greater role since to gain an intuition of nonlinear phenomena is hard. Landau and Lifshitz in their Fluid Mechanics (Volume 6 of their course) devote a whole section to the exact solutions of the nonlinear Navier-Stokes equations for a viscous fluid (including Landau's own solution for a jet emerging from the end of a narrow tube into an infinite space filled with fluid).  Although Poisson's equation for the gravitational potential in the classical theory of gravity is linear, the combined system of equations describing both the field and its {\\it fluid} sources (not rigid bodies, these are simple!) characterized by Euler's equations and an equation of state are nonlinear. In classical astrophysical fluid dynamics perhaps the most distinct and fortunate example of the role of explicit solutions is given by the exact descriptions of ellipsoidal, uniform density masses of self-gravitating fluids. These ``ellipsoidal figures of equilibrium'' \\cite{CHDR} include the familiar Maclaurin spheroids and triaxial Jacobi ellipsoids, which are characterized by rigid rotation, and a wider class discovered by Dedekind and Riemann, in which a motion of uniform vorticity exists, even in a frame in which the ellipsoidal surface is at rest. The solutions representing the rotating ellipsoids did not only play an inspirational role in developing basic concepts of the theory of rigidly rotating stars, but quite unexpectedly in the study of inviscid, differentially rotating polytropes. These closely resemble Maclaurin spheroids, although they do not maintain rigid rotation. As noted in the well-known monograph on rotating stars \\cite{TAS}, ``the classical work on uniformly rotating, homogeneous spheroids has a range of validity much greater than was usually anticipated''. It also influenced galactic dynamics \\cite{BT}: the existence of Jacobi ellipsoids suggested that a rapidly rotating galaxy may not remain axisymmetric, and the Riemann ellipsoids demonstrated that there is a distinction between the rate at which the matter in a triaxial rotating body streams and the rate at which the figure of the body rotates. Since rotating incompressible ellipsoids adequately illustrate the general feature of rotating axisymmetric bodies, they are also used in the studies of double stars whose components are close to each other. The disturbances caused by a neighbouring component are treated as first order perturbations. Relativistic effects on the rotating incompressible ellipsoids have been investigated in the post-Newtonian approximation by various authors, recently with a motivation to understand the coalescence of binary neutron stars near their innermost stable circular orbit (see \\cite{TI} for the latest work and a number of references).  \\vskip -0.4pt As for the last subject, which has a more direct connection with exact explicit solutions of Einstein's equations, we want to say a few words about integrable systems and their soliton solutions. Soliton theory has been one of the most interesting developments in the past decades both in physics and mathematics, and gravity has played a role both in its birth and recent developments. It has been known from the end of the last century that the celebrated Korteweg-de Vries nonlinear evolution equation, which governs one dimensional surface gravity waves propagating in a shallow channel of water, admits solitary wave solutions. However, it was not until Zabusky and Kruskal (the Kruskal of section 2.4 below) did extensive {\\it numerical} studies of this equation in 1965 that the remarkable properties of the solitary waves were discovered: the nonlinear solitary waves, named solitons by Zabusky and Kruskal, can interact and then continue, preserving their shapes and velocities. This discovery has stimulated extensive studies of other nonlinear equations, the inverse scattering methods of their solution, the proof of the existence of an infinite number of conservation laws associated with such equations, and  the construction of explicit solutions (see \\cite{ABL} for a recent comprehensive treatise). Various other nonlinear equations, similar to the sine-Gordon equation or the nonlinear Schr\\\"odinger equation, arising for example in plasma physics, solid state physics, and nonlinear optics, have also been successfully tackled by these methods. At the end of the 1970s several authors discovered that Einstein's vacuum equations for axisymmetric stationary systems can be solved by means of the inverse scattering methods, and it soon became clear that one can employ them also in situations when both Killing vectors are spacelike (producing, for example, soliton-type cosmological gravitational waves). Dieter Maison, one of the pioneers in applying these techniques in general relativity, describes the subject thoroughly in this volume. We shall briefly meet the soliton methods when we discuss the uniformly rotating disk solution of Neugebauer and Meinel (Section 6.3), colliding plane waves (Section 8.3), and inhomogeneous cosmological models (Section 12.2). Our aim, however, is to understand the meaning of solutions, rather than generation techniques of finding them. From this viewpoint it is perhaps first worth noting the interplay between numerical and analytic studies of the soliton solutions -- hopefully, a good example of an interaction for numerical and mathematical relativists. However, the explicit solutions of integrable models have played important roles in various other contexts. The most interesting {\\it multi}-dimensional integrable equations are the four-dimensional self-dual Yang-Mills equations arising in field theory. Their solutions, discovered by R. Ward using twistor theory, on one hand stimulated Donaldson's most remarkable work on inequivalent differential structures on four-manifolds. On the other hand, Ward indicated that many of the known integrable systems can be obtained by dimensional reduction from the self-dual Yang-Mills equations. Very recently this view has been substantiated in the monograph by Mason and Woodhouse \\cite{MAW}. The words by which these authors finely express the significance of exact solutions in integrable systems can be equally well used for solutions of Einstein's equations: ``they combine tractability with nonlinearity, so they make it possible to explore nonlinear phenomena while working with explicit solutions''.\\footnote{In 1998, in the discussion after his Prague lecture on the present role of physics in mathematics, Prof. Michael Atiyah expressed a similar view that even with more powerful supercomputers and with a growing body of general mathematical results on the existence and uniqueness of solutions of differential equations, the exact, explicit solutions of nonlinear equations will not cease to play a significant role. (As it is well known, Sir Michael Atiyah has made fundamental contributions to various branches of mathematics and mathematical physics, among others, to the theory of solitons, instantons, and to the twistor theory of Sir Roger Penrose, with whom he has been interacting ``under the same roof'' in Oxford for 17 years \\cite{ATI}.)}  \\subsection{Einstein's field equations}  Since J\\\"urgen Ehlers has always been, among other things, interested in the history of science, he will hopefully tolerate a few remarks on the early history of Einstein's equations to which not much attention has been paid in the literature. It was during his stay in  Prague in 1911 and 1912 that Einstein's intensive interest in quantum theory diminished, and his systematic effort in constructing a relativistic theory of gravitation began. In his first ``Prague theory of gravity'' he assumed that gravity can be described by a single function -- the local velocity of light. This assumption led to insurmountable difficulties. However, Einstein learned much in Prague on his way to general relativity \\cite{GEB}: he understood the local significance of the principle of equivalence; he realized that the equations describing the gravitational field must be nonlinear and have a form invariant with respect to a larger group of transformations than the Lorentz group; and he found that ``{\\it spacetime coordinates lose their simple physical meaning}'', i.e. they do not determine directly the distances between spacetime points.\\footnote{At that time Einstein's view on the future theory of gravity are best summarized in his reply to M. Abraham \\cite{EAB}, written just before departure from Prague.} In his ``Autobiographical Notes'' Einstein says: ``Why were seven years ... required for the construction of general relativity? The main reason lies in the fact that it is not easy to free oneself from the idea that coordinates must have an immediate metrical meaning''... Either from Georg Pick while still in Prague, or from Marcel Grossmann during the autumn of 1912 after his return to Zurich (cf. \\cite{GEB}), Einstein learned that an appropriate mathematical formalism for his new theory of gravity was available in the work of Riemann, Ricci, and Levi-Civita. Several months after his departure from Prague and his collaboration with Grossmann, Einstein had general relativity almost in hand. Their work \\cite{EG} was already based on the generally invariant line element \\renewcommand{\\theequation}{\\Roman{equation}} \\begin{eqnarray} ds^2 = g_{\\mu \\nu} dx^{\\mu} dx^{\\nu} \\end{eqnarray} in which the spacetime metric tensor $g_{\\mu \\nu}(x^\\rho), \\mu , \\nu , \\rho = 0,1,2,3,$ plays a dual role: on the one hand it determines the spacetime geometry, on the other it represents the (ten components of the) gravitational potential and is thus a dynamical variable. The disparity between geometry and physics, criticized notably by Ernst Mach,\\footnote{ Mach spent 28 years as Professor of Experimental Physics in Prague, until 1895, when he took the History and Theory of Inductive Natural Sciences chair in Vienna.} had thus been removed. When searching for the field equations for the metric tensor, Einstein and Grossmann {\\it had} already realized that a natural candidate for generally covariant field equations would be the equations relating -- in present-day terminology -- the Ricci tensor and the energy-momentum tensor of matter. However, they erroneously concluded that such equations would not yield the Poisson equation of Newton's theory of gravitation as a first approximation for weak gravitational fields (see both \\S 5 in the ``Physical part'' in \\cite{EG} written by Einstein and \\S 4, below equation (\\ref{Equ46}), in the ``Mathematical part'' by M. Grossmann). Einstein then rejected the general covariance. In a subsequent paper with Grossmann \\cite{EG1}, they supported this mis-step by a well-known ``hole'' meta-argument and obtained (in today's terminology) four gauge conditions such that the field equations were covariant only with respect to transformations of coordinates permitted by the gauge conditions. We refer to, for example, \\cite{Pai} for more detailed information on the further developments leading to the final version of the field equations. Let us only summarize that in late 1915 Einstein first readopted the generally covariant field equations from 1913, in which the Ricci tensor $R_{\\mu \\nu}$ was, up to the gravitational coupling constant, equal to the energy-momentum tensor $T_{\\mu \\nu}$ (paper submitted to the Prussian Academy on November 4). From his {\\it vacuum field equations} \\begin{eqnarray} R_{\\mu \\nu} = 0, % \\end{eqnarray} where $R_{\\mu \\nu}$ depends nonlinearly on $g_{\\alpha \\beta}$ and its first derivatives, and linearly on its second derivatives, he was able to explain the anomalous part of the perihelion precession of Mercury -- in the note presented to the Academy on November 18. And finally, in the paper \\cite{EFE} submitted on November 25 (published on December 2, 1915), the final version of the {\\it gravitational field equations}, or {\\it Einstein's field equations} appeared:\\footnote{ David Hilbert submitted his paper on these field equations five days before Einstein, though it was published only on March 31, 1916. Recent analysis \\cite{CRS} of archival materials has revealed that Hilbert made significant changes in the proofs. The originally submitted version of his paper contained the theory which is not generally covariant, and the paper did not include equations (III).} \\begin{equation} R_{\\mu \\nu} - \\frac{1}{2} g_{\\mu \\nu} R = \\frac{8\\pi G}{c^4} T_{\\mu \\nu}, % \\end{equation} where the scalar curvature $R = g^{\\mu \\nu} R_{\\mu \\nu}$. Newton's gravitational constant $G$ and the velocity of light $c$ are the (only) fundamental constants appearing in the theory. If not stated otherwise, in this article we use the geometrized units in which $G=c=1$, and the same conventions as in \\cite{MTW} and \\cite{Wa}.  Now it is well known that Einstein further generalized his field equations by adding a cosmological term $+\\Lambda g_{\\mu \\nu}$ on the left side of the field equations (III). The cosmological constant $\\Lambda$ appeared first in Einstein's work ``Cosmological considerations in the General Theory of Relativity'' \\cite{ECO} submitted on February 8, 1917 and published on February 15, 1917, which contained the {\\it closed static model of the Universe} (the Einstein static universe) -- an exact solution of equations (III) with $\\Lambda > 0$ and an energy-momentum tensor of incoherent matter (``dust''). This solution marked the birth of modern cosmology.  We do not wish to embark upon the question of the role that Mach's principle played in Einstein's thinking when constructing general relativity, or upon the intriguing issues relating to aspects of Mach's principle in present-day relativity and cosmology\\footnote{It was primarily Einstein's recognition of the role of Mach's ideas in his route towards general relativity, and in his christening them by the name ``Mach's principle'' (though Schlick used this term in a vague sense three years before Einstein), that makes Mach's Principle influential even today. After the 1988 Prague conference on Ernst Mach and his influence on the development of physics \\cite{PRG}, the 1993 conference devoted exclusively to Mach's principle was held in T\\\"ubingen, from which a remarkably thorough volume was prepared \\cite{TUB}, covering all aspects of Mach's principle and recording carefully all discussion. The clarity of ideas and insights of J\\\"urgen Ehlers contributed much to both conferences and their proceedings. For a brief more recent survey of various aspects of Mach's principle in general relativity, see the introductory section in the work \\cite{DJJ}, in which Mach's principle is analyzed in the context of perturbed Robertson-Walker universes. Most recently, Mach's principle seems to enter even into M theory \\cite{HORA}.} --  a problem which in any event would far exceed the scope of this article. Although it would not be inappropriate to include it here since exact solutions (such as G\\\"odel's universe or Ozsv\\'ath's and Sch\\\"ucking's closed model) have played a prominent role in this context. However, it should be at least stated that Einstein originally invented the idea of a closed space in order to eliminate boundary conditions at spatial infinity. The boundary conditions ``flat at infinity'' bring with them an inertial frame unrelated to the mass-energy content of the space, and Einstein, in accordance with Mach's views, believed that merely mass-energy can influence inertia. Field equations (III) are not inconsistent with this idea, but they admit as the simplest solution an empty flat {\\it Minkowski space} $(T_{\\mu \\nu}=0,\\,\\,\\, g_{\\mu \\nu}= \\eta _{\\mu \\nu}=$ diag $(-1, +1, +1, +1))$, so some restrictive boundary conditions are essential if the idea is to be maintained. Hence, Einstein introduced the cosmological constant $\\Lambda$, hoping that with this space will always be closed, and the boundary conditions eliminated. But it was also in 1917 when {\\it de Sitter} discovered the {\\it solution} \\cite{dS} of the vacuum field equations (II) with added cosmological term $(\\Lambda > 0)$ which demonstrated that a nonvanishing $\\Lambda$ does not necessarily imply a nonvanishing mass-energy content of the universe.  \\subsection{``Just so'' notes on the simplest solutions: the Minkowski, de Sitter, and anti de Sitter spacetimes}  Our brief intermezzo on the cosmological constant brought up three explicit simple exact solutions of Einstein's field equations -- the Minkowski, Einstein, and de Sitter spacetimes. To these also belongs the anti de Sitter spacetime, corresponding to a negative $\\Lambda$. The de Sitter spacetime has the topology $R^1\\times S^3$ (with $R^1$ corresponding to the time) and is best represented geometrically as the 4-dimensional hyperboloid $-v^2+w^2+x^2+y^2+z^2=(3/\\Lambda )$ in 5-dimensional flat space with metric $ds^2=-dv^2+dw^2+dx^2+dy^2+dz^2$. The anti de Sitter spacetime has the topology $S^1\\times R^3$, and can be visualized as the 4-dimensional hyperboloid $-U^2-V^2+X^2+Y^2+Z^2=(-3/\\Lambda)$, $\\Lambda < 0$, in flat 5-dimensional space with metric $ds^2 = - dU^2 - dV^2 + dX^2 + dY^2 + dZ^2$. As is usual (cf. e.g. \\cite{HE,Pen}), we mean by ``anti de Sitter spacetime'' the universal covering space which contains no closed timelike lines; this is obtained by unwrapping the circle $S^1$.  These spacetimes will not be discussed in the following sections. Ocassionally, for instance, in Sections 5 and 10, we shall consider spacetimes which become asymptotically de Sitter. However, since these solutions have played a crucial role in many issues in general relativity and cosmology, and most recently, they have become important prerequisites on the stage of the theoretical physics of the ``new age'', including string theory and string cosmology, we shall make a few comments on these solutions here, and give some references to recent literature.  The basic geometrical properties of these spaces are analyzed in the Battelle Recontres lectures by Penrose \\cite{Pen}, and in the monograph by Hawking and Ellis \\cite{HE}, where also references to older literature can be found. The important role of the de Sitter solution in the theory of the expanding universe is finely described in the book by Peebles \\cite{PEE}, and in much greater detail in the proceedings of the Bologna 1988 meeting on the history of modern cosmology \\cite{BLG}.  \\vskip -0.2pt The Minkowski, de Sitter and anti de Sitter spacetimes are the simplest solutions in the sense that their {\\it metrics} are {\\it of constant} (zero, positive, and negative) {\\it curvature}. They admit the same number (ten) of independent Killing vectors, but the interpretations of corresponding symmetries differ for each spacetime. Together with the Einstein static universe, they all are conformally flat, and can be represented as portions of the Einstein static universe \\cite{HE,Pen}. However, their {\\it conformal structure} is globally different. In Minkowski spacetime one can go to infinity along timelike geodesics and arrive to the future (or past) {\\it timelike infinity} $i^+$ (or $i^-$); along null geodesics one reaches the future (past) {\\it null infinity} ${\\cal{J}}^+ ({\\cal{J}}^-)$; and spacelike geodesics lead to {\\it spatial infinity} $i^0$. Minkowski spacetime can be compactified and mapped onto a finite region by an appropriate conformal rescaling of the metric. One thus obtains the well-known Penrose diagram in which the three types of infinities are mapped onto the boundaries of the compactified spacetime -- see for example the boundaries on the ``right side'' in the Penrose diagram of the Schwarzschild-Kruskal spacetime in Fig. 3, Section 2.4, or the Penrose compactified diagram of boost-rotation symmetric spacetimes in Fig. 13,  Section 11. (The details of the conformal rescaling of the metric and resulting diagrams are given in \\cite{HE,Pen} and in standard textbooks, for example \\cite{MTW,Wa,DIN}.) In the de Sitter spacetime there are only past and future conformal infinities $\\cal J^-, \\cal J^+$, both being {\\it spacelike} (cf. the Penrose diagram of the ``cosmological'' Robinson-Trautman solutions in Fig. 11, Section 10); the conformal infinity in anti de Sitter spacetime is {\\it timelike}.  \\vskip -0.2pt These three spacetimes of constant curvature offer many basic insights which have played a most important role elsewhere in relativity. To give just a few examples (see e.g. \\cite{HE,Pen}): both the particle (cosmological) horizons and the event horizons for geodesic observers are well illustrated in the de Sitter spacetime; the Cauchy horizons in the anti de Sitter space; and the simplest acceleration horizons in Minkowski space (hypersurfaces $t^2 = z^2$ in Fig. 12, Section 11). With the de Sitter spacetime one learns (by considering different cuts through the 4-dimensional hyperboloid) that the concept of an ``open'' or ``closed'' universe depends upon the choice of a spacelike slice through the spacetime. There is perhaps no simpler way to understand that Einstein's field equations are of local nature, and that the spacetime topology is thus not given a priori, than by considering the following construction in Minkowski spacetime. Take the region given in the usual coordinates by $|x| \\leq 1, |y| \\leq 1, |z| \\leq 1$, remove the rest and identify pairs of boundary points of the form $(t, 1, y, z)$ and $(t, -1, y, z)$, and similarly for $y$ and $z$. In this way the spatial sections are identified to obtain a 3-torus -- a flat but closed manifold.\\footnote{This very simple point was apparently unknown to Einstein in 1917, although soon after the publication of his cosmological paper, E. Freundlich and F. Klein pointed out to him that an elliptical topology (arising from the identification of antipodal points) could have been chosen instead of the spherical one considered by Einstein. Although topological questions have been followed with a great interest in recent decades, the chapter by Geroch and Horowitz in ``An Einstein Centenary Survey'' \\cite{GHO} remains the classic; for more recent texts, see for example \\cite{Jo} and references therein.}  The {\\it spacetimes of constant curvature} have been resurrected as basic {\\it arenas of new physical theories} since their first appearance. After the role of the de Sitter universe decreased with the refutation of the steady-state cosmology, it has inflated again enormously in connection with the theory of early quasi-exponential phase of expansion of the universe, due to the false-vacuum state of a hypothetical scalar (inflaton) field(s) (see e.g. \\cite{PEE}). We shall mention the de Sitter space as the asymptotic state of cosmological models with a nonvanishing $\\Lambda$ (so verifying the ``cosmic no-hair conjecture'') in Section 10 on Robinson-Trautman spacetimes. Motivated by its importance in inflationary cosmologies, several new useful papers reviewing the properties of de Sitter spacetime have appeared \\cite{HJS,GRN}; they also contain many references to older literature. For the most recent work on the quantum structure of de Sitter space, see \\cite{BOU}.  In the last two years, anti de Sitter spacetime has come to the fore in light of Maldacena's conjecture \\cite{MDC} relating string theory in (asymptotically) anti de Sitter space to a non-gravitational conformal field theory on the boundary at spatial infinity, which is timelike as mentioned above (see, e.g. \\cite{BLB}, where among others, in the Appendix various coordinate systems describing anti de Sitter spaces in arbitrary dimensions are discussed).  Amazingly, the Minkowski spacetime has recently entered the active new area of so called pre-big bang string cosmology \\cite{VEN}. String theory is here applied to the problem of the big bang. The idea is to start from a simple Minkowski space (as an ``asymptotic past triviality'') and to show that it is in an unstable false-vacuum state, which leads to a long {\\it pre}-big bangian inflationary phase. This, at later times, should provide a hot big bang. Although such a scenario has been criticized on various grounds, it has attractive features, and most importantly, can be probed through its observable relics \\cite{VEN}.  Since it is hard to forecast how the roles of these three spacetimes of constant curvature will develop in new and exciting theories in the next millennium, let us better conclude our ``just so'' notes by stating three ``stable'' results of complicated, rigorous mathematical analyses of (the classical) Einstein's equations.  In their recent treatise \\cite{CHK}, Christodoulou and Klainerman prove that any  smooth, asymptotically flat initial data set which is ``near flat, Minkowski data'' leads to a unique, smooth and geodesically complete solution of Einstein's vacuum equations with vanishing cosmological constant. This demonstrates the {\\it stability of the Minkowski space} with respect to nonlinear (vacuum) perturbations, and the existence of singularity-free, asymptotically flat radiative vacuum spacetimes. Christodoulou and Klainerman, however, are able to show only a somewhat weaker decay of the field at null infinity than is expected from the usual assumption of a sufficient smoothness at null infinity in the framework of Penrose (see e.g. \\cite{KVB} for a brief account).  Curiously enough, in the case of the vacuum Einstein equations with a {\\it nonvanishing cosmological constant}, a more complete picture has been known for some time. By using his regular conformal field equations, Friedrich \\cite{HF1} demonstrated that initial data sufficiently close to de Sitter data develop into solutions of Einstein's equations with a positive cosmological constant, which are ``asymptotically simple'' (with a smooth conformal infinity), as required in Penrose's framework. More recently, Friedrich \\cite{HF2} has shown the existence of asymptotically simple solutions to the Einstein vacuum equations with a negative cosmological constant. For the latest review of Friedrich's thorough work on asymptotics, see \\cite{HF3}.  Summarizing, thanks to these profound mathematical achievements we know that the {\\it Minkowski, de Sitter}, and {\\it anti de Sitter spacetimes} are the solutions of Einstein's field equations which are {\\it stable with respect to general, nonlinear} (though ``weak'' in a functional sense) {\\it vacuum perturbations}. A result of this type is not known for any other solution of Einstein's equations.   \\subsection{On the interpretation and characterization of metrics}  Suppose that a metric satisfying Einstein's field equations is known in some region of spacetime and  in a given coordinate (reference) system $x^{\\mu}$. A fundamental question, frequently ``forgotten'' to be addressed in modern theories which extend upon general relativity, is whether the {\\it metric tensor} $g_{\\alpha \\beta}(x^{\\mu})$ {\\it is a measurable quantity}. Classical general relativity offers (at least) three ways of giving a positive answer, depending on what objects are considered as ``primitive tools'' to perform the measurements. The first, elaborated and emphasized primarily by M{\\o}ller \\cite{MOL}, employs standard rigid rods in the measurements. However, a ``rigid rod'' is not really a simple primitive concept. The second procedure, due to Synge \\cite{SG}, accepts as the basic concepts a ``particle'' and a ``standard clock''. If $x^{\\mu}$ and $x^\\mu + dx^\\mu$ are two nearby events contained in the worldline of a clock, then the separation (the spacetime interval) between the events is equal to the interval measured by the clock. The main drawback of this approach appears to lie in the fact that it does not explain why the same  functions $g_{\\alpha \\beta}(x^{\\mu})$ describe the behavior of the clock as well as paths of free particles, as explained in more detail by Ehlers, Pirani and Schild \\cite{EPS}, in the motivation for their own axiomatic but constructive procedure for setting up the spacetime geometry. Their method, inspired by the work of Weyl and others, uses neither rods nor clocks, but instead, light rays and freely falling test particles, which are considered as basic tools for measuring the metric and determining the spacetime geometry. (For a simple description of how this can be performed, see exercise 13.7 in \\cite{MTW}; for some new developments which build upon, among others, the Ehlers-Pirani-Sachs approach, see \\cite{MAS}.) After indicating that the metric tensor is a measurable quantity let us briefly turn to the {\\it role of spacetime coordinates}.  \\vskip -0.3pt In special relativity there are infinitely many global inertial coordinate systems labelling events in the Minkowski manifold $\\RRe ^4$; they are related by elements of the Poincar\\'e group. The inertial coordinates labels $X^0, X^1, X^2, X^3$ of a given event do not thus have intrinsic meaning. However, the spacetime interval between two events, determined by the Minkowski metric $\\eta_{\\mu \\nu}$, represents an intrinsic property of spacetime. Since the Minkowski metric is so simple, the differences between inertial coordinates can have a metrical meaning (recall Einstein's reply to Abraham mentioned in Section 1.2). In principle, however, both in special and general relativity, it is the metric, the line element, which exhibits intrinsically the geometry, and gives all relevant information. As Misner \\cite{MER} puts it, if you write down for someone the Schwarzschild metric in the ``canonical'' form (equation (\\ref{Equ2}) in Section 2.2) and receive the reaction ``that [it] tells me the $g_{\\mu \\nu}$ gravitational potentials, now tell me in which $(t, r, \\theta, \\varphi)$ coordinate system they have these values?'', then there are two valid responses: (a) indicate that it is an indelicate and unnecessary question, or (b) ignore it. Clearly, the Schwarzschild metric describes the geometrical properties of the coordinates used in (\\ref{Equ2}). For example, it implies that worldlines with fixed $r, \\theta , \\varphi$ are timelike at $r>2M$, orthogonal to the lines with $t = \\const$. It determines local null cones (given by $ds^2=0$), i.e. the {\\it causal structure} of the spacetime. In addition, in Schwarzschild coordinates the metric (\\ref{Equ2}) indicates how to {\\it measure} the radial coordinate of a given event, because the proper area of the sphere going through the event is given just by the Euclidean expression $4\\pi r^2$ ($r$ is thus often called ``the curvature coordinate''). On each sphere the angular coordinates $\\theta, \\varphi$ have the same meaning as on a sphere in Euclidean space. The Schwarzschild coordinate time $t$, geometrically preferred by the timelike (for $r>2M$) Killing vector, which is just equal to $\\partial/\\partial t$, can be measured by radar signals sent out from spatial infinity ($r \\gg 2M$) where $t$ is the proper time (see e.g. \\cite{MTW}). The coordinates used in (\\ref{Equ2}) are in fact ``more unique'' than the inertial coordinates in Minkowski spacetime, because the only possible continuous transformations preserving the form (\\ref{Equ2}) are rigid rotations of a sphere, and $t \\rightarrow t + \\const$. Such a simple interpretation of coordinates is exceptional. However, the simple case of the Schwarzschild metric clearly demonstrates that all intrinsic information is contained in the line element.  It is interesting, and for some purposes useful, to consider not just one Schwarzschild metric with a given mass $M$ but the {\\it family} of such metrics for all possible $M$. In order to cover also the future event horizon let us describe the metrics by using Eddington-Finkelstein ingoing coordinates as in equation (\\ref{Equ4}), Section 2.3. This equation can be interpreted as a family of metrics with various values of $M$ given on a {\\it fixed background manifold} ${\\bar{\\cal M}}_1$, with $v \\in \\RRe, r \\in(0, \\infty)$, and $\\theta \\in [0, \\pi], \\varphi \\in [0, 2\\pi)$. Alternatively, however, we may use, for example, the Kruskal null coordinates $\\tilde U, \\tilde V$ in which the metric is given by equation (\\ref{Equ6}), Section 2.3, with $\\tilde U = V-U, \\tilde V = V + U$. We may then consider metrics on a background manifold $\\bar {\\cal {M}}_2$ given by $\\tilde{U} \\in \\RRe, \\tilde V \\in(0, \\infty),\\;\\theta \\in [0, \\pi]$, and $\\varphi \\in [0, 2 \\pi)$, which corresponds to $\\bar {\\cal {M}}_1$. However, these two background manifolds are {\\it not} the same: the transformation between the Eddington-Finkelstein coordinates and the Kruskal coordinates is not a map from $\\bar {\\cal {M}}_1$  to $\\bar {\\cal{M}}_2$ because it depends on the value of mass $M$. Therefore, the ``background manifold'' used frequently in general relativity, for example in problems of conservation of energy, or in quantum gravity, is not defined in a natural, unique manner. The above simple pedagogical observation has recently been made in connection with gauge fixing in quantum gravity by H\\'aj\\'{\\i}\\v{c}ek \\cite{HP} in order to explain the old insight by Bergmann and Komar, that the gauge group of general relativity is much larger than the diffeomorphism group of one manifold. To identify points when working with backgrounds, one usually fixes coordinates in all solution manifolds by some gauge condition, and identifies those points of all these manifolds which have the same value of the coordinates.  Returning back to a single solution $({\\cal {M}}, g_{\\alpha \\beta})$, described by a manifold ${\\cal {M}}$ and a metric $g_{\\alpha \\beta}$ in some coordinates, a notorious (local) {\\it ``equivalence problem''} often arises. A given (not necessarily global) solution has the variety of representations which equals the variety of choices of a 4-dimensional coordinate system. Transitions from one choice to another are isomorphic with the group of 4-dimensional diffeomorphisms which expresses the general covariance of the theory.\\footnote{As pointed out by Kretschmann soon after the birth of general relativity, one can always make a theory generally covariant by taking more variables and inserting them as new dynamical variables into the (enlarged) theory. Thus, standard Yang-Mills theory is covariant with respect to the transformations of Yang-Mills potentials, corresponding to a particular group, say $SU(2)$. However, the theory is usually formulated on a fixed background spacetime with a given metric. The evolution of a dynamical Yang-Mills solution is thus ``painted'' on a given spacetime. When the metric -- the gravitational field -- is incorporated as a dynamical variable in the Einstein-Yang-Mills theory, the whole spacetime metric and Yang-Mills field are ``built-up'' from given data (cf. the article by Friedrich and Rendall in this volume). The resulting theory is covariant with respect to a much larger group. The dual role of the metric, determined only up to 4-dimensional diffeomorphisms, makes the character of the solutions of Einstein's equations unique among solutions of other field theories, which do not consider spacetime as being dynamical.} Given another set of functions $g'_{\\alpha \\beta}(x'^{\\gamma})$ which satisfy Einstein's equations, how do we learn that they are not just transformed components of the metric $g_{\\alpha \\beta}(x^{\\gamma})$? In 1869 E. B. Christoffel raised a more general question: under which conditions is it possible to transform a quadratic form $g_{\\alpha \\beta}(x^\\gamma )dx^\\alpha dx^\\beta$ in $n$-dimensions into another such form $g'_{\\alpha \\beta}(x'^\\gamma)dx'^\\alpha dx'^\\beta$ by means of smooth transformation $x^\\gamma (x'^\\kappa)$? As Ehlers emphasized in his paper \\cite{ECH} on the meaning of Christoffel's equivalence problem in modern field theories, Christoffel's results apply to metrics of arbitrary signature, and can be thus used directly in general relativity. Without going into details let us say that today the solution to the equivalence problem as presented by Cartan is most commonly used. For both metrics $g_{\\alpha \\beta}$ and $g'_{\\alpha \\beta}$ one has to find a frame (four 1-forms) in which the frame metric is constant, and find the frame components of the Riemann tensor and its covariant derivatives up to -- possibly -- the 10th order. The two metrics $g_{\\alpha \\beta}$ and $g'_{\\alpha \\beta}$ are then equivalent if and only if there exist coordinate and Lorentz transformations under which one whole set of frame components goes into the other. In a practical algorithm given by Karlhede \\cite{KL}, recently summarized and used in \\cite{PRM}, the number of derivations required is reduced.  A natural first idea of how to solve the equivalence problem is to employ the scalar invariants from the Riemann tensor and its covariant derivatives. This, however, does not work. For example, in all Petrov type {\\it N} and {\\it III} nonexpanding and nontwisting solutions all these invariants vanish as shown recently (see Section 8.2), as they do in Minkowski spacetime.  However, even without regarding invariants, at present much can be learnt about an exact solution (at least locally) in geometrical terms, without reference to special coordinates. This is thanks to the progress started in the late 1950s, in which the group of Pascual Jordan in Hamburg has played the leading role, with J\\\"urgen Ehlers as one of its most active members. Ehlers' dissertation\\footnote{The English translation of the title of the dissertation reads: ``The construction and characterization of the solutions of Einstein's gravitational field equations''. In \\cite{EK} the original German title is quoted, as in our citation \\cite{EHD}, but ``of the solutions'' is erroneously omitted. This error then reemerges in the references in \\cite{Wa}.} \\cite{EHD} from 1957 is devoted to the characterization of exact solutions.  The problem of exact solutions also forms the content of his contribution to the Royaumont GR-conference \\cite{EHRO}, as well as his plenary talk in the London GR-conference \\cite{EHLO}. A detailed description of the results of the Hamburg group on invariant geometrical characterization of exact solutions by using and developing the Petrov classification of Weyl's tensors, groups of isometries, and conformal transformations are contained in the first paper \\cite{JEK} in the (today ``golden oldies'') series of articles published in the ``Abhandlungen der Akademie der Wissenschaften in Mainz''. An English version, in a somewhat shorter form, was published by Ehlers and Kundt \\cite{EK} in the ``classic'' 1962 book ``Gravitation: An Introduction to Current Research'' compiled by L. Witten. (We shall meet these references in the following sections.) In the second paper of the ``Abhandlungen'' \\cite{JEK2}, among others, algebraically special vacuum solutions are studied, using the formalism of the 2-component spinors, and in particular, geometrical properties of the congruences of null rays are analyzed in terms of their expansion, twist, and shear.  These tools became essential for the discovery by Roy Kerr in 1963 of the solution which, when compared with all other solutions of Einstein's equations found from the beginning of the renaissance of general relativity in the late 1950s until today, has played the most important role. As Chandrasekhar \\cite{CHS} eloquently expresses his wonder about the remarkable fact that all stationary and isolated black holes are {\\it exactly} described by the Kerr solution: ``This is the only instance we have of an exact description of a macroscopic object. Macroscopic objects, as we see them all around us, are governed by a variety of forces, derived from a variety of approximations to a variety of physical theories. In contrast, the only elements in the construction of black holes are our basic concepts of space and time ...'' The Kerr solution can also serve as one of finest examples in general relativity of ``the incredible fact that a discovery motivated by a search after the beautiful in mathematics should find its exact replica in Nature...'' \\cite{SCH}.  The technology developed in the classical works \\cite{EK,JEK}, and in a number of subsequent contributions, is mostly concerned with the local geometrical characterization of exact spacetime solutions. A well-known feature of the solutions of Einstein's equations, not shared by solutions in other physical theories, is that it is often very complicated to analyze their global properties, such as their extensions, completeness, or topology. If analyzed globally, almost any solution can tell us something about the basic issues in general relativity, like the nature of singularities, or cosmic censorship.  \\subsection{The choice of solutions}  Since most solutions, when properly analyzed, can be of potential interest, we are confronted with a richness of material which puts us in danger of mentioning many of them, but remaining on a general level, and just enumerating rather than enlightening. In fact, because of lack of space (and of our understanding) we shall have to adopt this attitude in many places. However, we have selected some solutions, hopefully the fittest ones, and when discussing their role, we have chosen particular topics to be analyzed in some detail, and left other issues to brief remarks and references.  Firstly, however, let us ask what do we understand by the term ``exact solution''. In the much used ``exact-solution-book'' \\cite{KSH}, the authors ``do not intend to provide a definition'', or, rather, they have decided that what they ``chose to include was, by definition, an exact solution''. A mathematical relativist-purist would perhaps consider solutions, the existence of which has been demonstrated in the works of Friedrich or Christodoulou and Klainerman, mentioned at the end of Section 1.3, as ``good'' as the Schwarzschild metric. Most recently, Penrose \\cite{ROPE} presented a strong conjecture which may lead to a general vacuum solution described in the complicated (complex) formalism of his twistor theory. Although in this article we do not mean by exact solutions those just mentioned, we also do not consider as exact solutions only those explicit solutions which can be written in terms of elementary functions on half of a page. We prefer, recalling Feynman, simple ``special cases'', but we also discuss, for example, the late-time behaviour of the Robinson-Trautman solutions for which rigorously convergent series expansions can be obtained, which provide sufficiently rich ``special information''.  \\vskip -0.15pt Concerning the selection of the solutions, the builder of general relativity and the gravitational field equations (III) himself indicates which solutions should be preferred \\cite{AI}: ``The theory avoids all internal discrepancies which we have charged against the basis of classical mechanics... But, it is similar to a building, one wing of which is made of fine marble (left part of the equation), but the other wing of which is built of low grade wood (right side of equation). The phenomenological representation of matter is, in fact, only a crude substitute for a representation which would correspond to all known properties of matter. There is no difficulty in connecting Maxwell's theory... so long as one restricts himself to space, free of ponderable matter and free of electric density...''  \\vskip -0.15pt Of course, Einstein was not aware when he was writing this of Yang-Mills-Higgs fields, or of the dilaton field, etc. However, remaining on the level of field theories with a clear classical meaning, his view has its strength and motivates us to prefer (electro)vacuum solutions. A physical interpretation of the vacuum solutions of Einstein's equations have been reviewed in papers by Bonnor \\cite{Bo}, and Bonnor, Griffiths and MacCallum \\cite{BoGM} five years ago. Our article, in particular in emphasizing and describing the role of solutions in giving rise to various concepts, conjectures, and methods of solving problems in general relativity, and in the astrophysical impacts of the solutions, is oriented quite differently, and gives more detail. However, up to some exceptions, like, for example, metrics for an infinite line-mass or plane, which are discussed in \\cite{Bo}, and new solutions which have been discovered after the reviews \\cite{Bo,BoGM} appeared as, for example, the solution describing a rigidly rotating thin disk of dust, our choice of solutions is similar to that of \\cite{Bo,BoGM}.  \\vskip -0.15pt In selecting particular topics for a more detailed discussion we will be led primarily by following overlapping aspects: (i) the ``commonly acknowledged'' significance of a solution -- we will concentrate in particular on the Schwarzschild, the Kerr, the Taub-NUT, and plane wave solutions, and (ii) the solutions and their properties that I (and my colleagues) have been directly interested in, such as the Reissner-Nordstr\\\"om metric, vacuum solutions outside rotating disks, or radiative solutions such as cylindrical waves, Robinson-Trautman solutions, and the boost-rotation symmetric solutions. Some of these have also been connected with the interests of J\\\"urgen Ehlers, and we shall indicate whenever we are aware of this fact.  \\vskip -0.15pt Vacuum cosmological solutions are discussed in less detail than they deserve. A possible excuse -- from the point of view of being a relativist, a rather unfair one -- could be that a special recent issue of Reviews of Modern Physics (Volume 71, 1999), marking the Centennial of the American Physical Society, contains discussion of the Schwarzschild, the Reissner-Nordstr\\\"om and other black hole solutions, and even remarks on the work of Bondi et al. \\cite{e} on radiative solutions, but among the cosmological solutions only the standard models are mentioned. A real reason is the author's lack of space, time, and energy. In the concluding remarks we will try to list at least the most important solutions (not only the Friedmann models!) which have not been ``selected'' and give references to the literature in which more information can be found.   \\subsection{The outline}  Since the titles of the following sections characterize the contents rather specifically, we restrict ourselves to only a few explanatory remarks. In our discussion of the Schwarschild metric, after mentioning its role in the solar system, we indicate how the Schwarzschild solution gave rise to such concepts as the event horizon, the trapped surface, and the apparent horizon. We pay more attention to the concept of a bifurcate Killing horizon, because this is usually not treated in textbooks, and in addition, J\\\"urgen Ehlers played a role in its first description in the literature. Another point which has not received much attention is Penrose's nice presentation of evidence against Lorentz-covariant field theoretical approaches to gravity, based on analysis of the causal structure of the Schwarzschild spacetime. Among various astrophysical implications of the Schwarzschild solution we especially note recent suggestions which indicate that we may have evidence of the existence of event horizons, and of a black hole in the centre of our Galaxy.  The main focus in our treatment of the Reissner-Nordstr\\\"om metric is directed to the instability of the Cauchy horizon and its relation to the cosmic censorship conjecture. We also briefly discuss extreme black holes and their role in string theory.  About the same amount of space as that given to the Schwarzschild solution is devoted to the Kerr metric. After explaining a few new concepts the metric inspired, such as locally nonrotating frames and ergoregions, we mention a number of physical processes which can take place in the Kerr background, including the Penrose energy extraction process, and the Blandford-Znajek mechanism. In the section on the astrophysical evidence for a Kerr metric, the main attention is paid to the broad iron line, the character of which, as most recent observations indicate, is best explained by assuming that it originates very close to a maximally rotating black hole. The discussion of recent results on black hole uniqueness and on multi-black hole solutions concludes our exposition of spacetimes representing black holes. In the section on axisymmetric fields and relativistic disks a brief survey of various static solutions is first given, then we concentrate on relativistic disks as sources of the Kerr metric and other stationary fields; in particular, we summarize briefly the recent work on uniformly rotating disks.  An intriguing case of Taub-NUT space is introduced by a new constructive derivation of the solution. Various pathological features of this space are then briefly listed.  Going over to radiative spacetimes, we analyze in some detail plane waves -- also in the light of the thorough study by Ehlers and Kundt \\cite{EK}. Some new developments are then noted, in particular, impulsive waves generated by boosting various ``particles'', their symmetries, and recent use of the Colombeau algebra of generalized functions in the analyses of impulsive waves. A fairly detailed discussion is devoted to various effects connected with colliding plane waves.  In our treatment of cylindrical waves we concentrate in particular on two issues: on the proof that these waves provide explicitly given spacetimes, which admit a smooth global null infinity, even for strong initial data within a $(2+1)$-dimensional framework; and on the role that cylindrical waves have played in the first construction of a midisuperspace model in quantum gravity. Various other developments concerning cylindrical waves are then summarized only telegraphically.  A short section on Robinson-Trautman solutions points out how these solutions with a nonvanishing cosmological constant can be used to give an exact demonstration of the cosmic no-hair conjecture under the presence of gravitational radiation, and also of the existence of an event horizon which is smooth but not analytic.  As the last class of radiative spacetimes we analyze the boost-rotation symmetric solutions representing uniformly accelerated objects. They play a unique role among radiative spacetimes since they are asymptotically flat, in the sense that they admit global smooth sections of null infinity. And as the only known radiative solutions describing finite sources they can provide expressions for the Bondi mass, the news function, or the radiation patterns in explicit forms. They have also been used as test-beds in numerical relativity, and as the model spacetimes describing the production of black hole pairs in strong fields.  Vacuum cosmological solutions such as the vacuum Bianchi models and Gowdy solutions are mentioned, and their significance in the development of general relativity is indicated in the last section. Special attention is paid to their role in understanding the behaviour of a general model near an initial singularity.  In the concluding remarks, several important, in particular {\\it non}-vacuum solutions, which have not been included in the main body of the paper, are at least listed, together with some relevant references. A few remarks on the possible future role of exact solutions ends the article.  Although we give over 360 references in the bibliography, we do not at all pretend to give all relevant citations. When discussing more basic facts and concepts, we quote primarily textbooks and monographs. Only when mentioning more recent developments do we refer to journals. The complete titles of all listed references will hopefully offer the reader a more complete idea of the role the explicit solutions have played on the relativistic stage and in the astrophysical sky. \\newpage ",
        "conclusions": "It is hoped that the preceding pages have helped to elucidate at least one issue: that in such a complicated nonlinear theory as general relativity, it is not possible to ask relevant questions of a general character without finding and thoroughly analyzing specific exact solutions of its field equations. The role of some of the solutions in our understanding of gravity and the universe has been so many-sided that to exhibit this role properly on even more than a hundred pages is not really feasible ...  Although we have concentrated on only (electro)vacuum solutions, there remains a number of such solutions that have also played some role in various contexts, but, owing to the absence of additional space and time, or the presence of the author's ignorance, have not been discussed. Tomimatsu-Sato solutions and their generalizations, static plane and cylindrical metrics, and some algebraically special solutions are examples.  In his review of exact solutions, Ehlers \\cite{EHLO} wrote 35 years ago that ``it seems desirable to construct material sources for vacuum solutions'', and 30 years later Bonnor \\cite{Bo}, in his review, expressed a similar view. In the above we have noted only some of the thin disk sources of static and stationary spacetimes in Section 6. To find physically reasonable material sources for many of the known vacuum solutions remains a difficult open task. In order to make solutions of Einstein's equations with the right-hand side more tractable, one is often tempted to sacrifice realism and consider materials, again using Bondi's phraseology, which are not easy to buy in the shops. Nevertheless, there are solutions representing spacetimes filled with matter which would certainly belong in a more complete discussion of the role of exact solutions.  For example, one of the simplest, the spherically symmetric Schwarzschild interior solution with an incompressible fluid as matter source, modelling ``a star of uniform density'', gives surprisingly good estimates of an upper bound on the masses of neutron stars; on a more general level, it supplies an instructive example of relativistic hydrostatics \\cite{MTW}. Many other spherical perfect fluid solutions are listed in \\cite{KSH}. The proof of a very plausible fact that any equilibrium, isolated stellar model which is nonrotating must be spherically symmetric, was finally completed in \\cite{BESI} and \\cite{LIM}. Physically more adequate spherically symmetric static solutions with collisionless matter described by the Boltzmann (Vlasov) equation have been studied \\cite{ZN} (yielding, for example, arbitrarily large central redshifts); and some of their aspects have been recently reviewed from a rigorous, mathematical point of view \\cite{RFL}. Going over to the description of matter in terms of physical fields, we should mention the first spherically symmetric regular solutions of the Einstein-Yang-Mills equations (``non-Abelian solitons'' discovered by Bartnik and McKinnon \\cite{BMK} in 1988), and non-Abelian black holes with ``hair'', which were found soon afterwards. They stimulated a remarkable activity in the search for models in which gravity is coupled with Yang-Mills, Higgs, and Skyrmion fields. Very recently these solutions have been surveyed in detail in the review by Volkov and Gal'tsov \\cite{VG}.  \\vskip -0.3pt The role of the standard FRW cosmological models on the development of relativity and cosmology can hardly be overemphasized. As for two more recent examples of this influence let us just recall that the existence of cosmological horizons in these models was one of the crucial points which inspired the birth of inflationary cosmology (see e.g. \\cite{PEE}); and the very smooth character of the initial singularity has led Penrose \\cite{ROG} to formulate his Weyl curvature hypothesis, related to a still unclear concept of gravitational entropy. Homogeneous but anisotropic Bianchi models filled with perfect fluid are extensively analyzed in \\cite{WE}. Very recent studies of Bianchi models with collisionless matter \\cite{RET} reveal how the matter content can qualitatively alter the character of the model.  \\vskip -0.3pt A number of Bianchi models approach self-similar solutions. Perfect fluid solutions admitting a homothetic vector, which in this case implies  both geometrical and physical self-similarity, have been reviewed most recently by Carr and Coley \\cite{CAC}. In their review various astrophysical and cosmological applications of such solutions are also discussed. Self-similar solutions have played a crucial role in the critical phenomena in gravitational collapse. Since their discovery by Choptuik in 1993, they have attracted much effort, which has revealed quite unexpected facts. In \\cite{CAC} these phenomena are analyzed briefly. For a more comprehensive review, see \\cite{GUN}.  \\vskip -0.3pt Self-similar, spherically symmetric solutions have been very relevant in constructing examples of the formation of naked singularities in gravitational collapse (see \\cite{CAC} for a brief summary and references). In particular, the inhomogeneous, spherically symmetric Lema\\^{\\i}tre-Bondi-Tolman universes containing dust have been employed in this context. Solutions with null dust should be mentioned as well, especially the spherically symmetric Vaidya solutions: imploding spherical null-dust models have been constructed in which naked singularities arise at their centre (see \\cite{Jo} for summary and references).  The Lema\\^{\\i}tre-Bondi-Tolman models are the most frequently analyzed inhomogeneous cosmological models which contain the standard FRW dust models as special cases (see e.g. \\cite{PEE,Jo}). In his recent book Krasi\\'{n}ski \\cite{KRA} has compiled and discussed most if not all of these exact inhomogeneous cosmological solutions found so far which can be viewed as ``exact perturbations'' of the FRW models.  Many solutions known already still wait for their role to be uncovered. The role of many others may forever remain just in their ``being''. However, even if new solutions of a ``Kerr-like significance'' will not be obtained in the near future, we believe that one should not cease in embarking upon journeys for finding them, and perhaps even more importantly, for revealing new roles of solutions already known. The roads may not be easy, but with todays equipment like Maple or Mathematica, the speed is increasing. Is there another so explicit way of how to learn more about the rich possibilities embodied in Einstein's field equations?   The most remarkable figure of Czech symbolism, Otokar B\\v rezina (1868-1929) has consoling words for those who do not meet the ``Kerr-type'' metric on the road: ``Nothing is lost in the world of the spirit; even a stone thrown away may find its place in the hands of a builder, and a house in flames may save the life of someone who has lost his way...''. \\vskip 6mm  \\noindent {\\bf  Acknowledgements}  \\vskip 4mm  \\noindent Interaction with J\\\"{u}rgen Ehlers has been important for me over the years: Thanks to my regular visits to his group, which started seven years before the hardly penetrable barrier between Prague and the West disappeared, I have been in contact with ``what is going on'' much more than I could have been at home. Many discussions with J\\\"{u}rgen, collaboration and frequent discussions with Bernd Schmidt, and with other members of Munich$\\rightarrow$Potsdam$\\rightarrow$Golm relativity group are fondly recalled and appreciated.  For helpful comments on various parts of the manuscript I am grateful to Bobby Beig, Jerry Griffiths, Petr H\\'aj\\'\\i\\v cek, Karel Kucha\\v r, Malcolm MacCallum, Alan Rendall and Bernd Schmidt. For discussions and help with references I thank also Piotr Chru\\'sciel, Andy Fabian, Joseph Katz, Jorma Louko, Donald Lynden-Bell, Reinhard Meinel, Vince Moncrief, Gernot Neugebauer, Martin Rees, Carlo Ungarelli, Marsha Weaver, and my Prague colleagues. Peter Williams kindly corrected my worst Czechisms. My many thanks  go to  Eva Kotal\\'\\i kov\\'a for her patience and skill in technical help with the long manuscript. Very special thanks to Tom\\'a\\v s Ledvinka: he prepared all the figures and provided long-standing technical help and admirable speed, without which the manuscript would certainly not have been finished in the required time and form. Support from the Albert Einstein Institute and from the grant No. GA\\v CR 202/99/0261 of the Czech Republic is gratefully acknowledged."
    },
    "0004/astro-ph0004098_arXiv.txt": {
        "abstract": " ",
        "introduction": "Magnetic fields have been observed in a large number of spiral galaxies, including the Milky Way \\cite{Kronberg}. In almost all cases, the galactic field strengths are measured to be a few times $10^{-6}$ G. This particular value has also been found at a redshift of $z=0.395$ \\cite{redshift} and between the galaxies in clusters. In spiral galaxies, one discovers that the magnetic field is aligned with the spiral arms and density waves in the disk. A plausible explanation is that it was created by a mean-field dynamo mechanism \\cite{dynamo}, in which a much smaller seed field was exponentially amplified by the turbulent motion of ionised gas in conjunction with the differential rotation of the galaxy.  For the dynamo to work, the initial seed field must be correlated on a scale of $100$ pc, corresponding to the largest turbulent eddy \\cite{dynamo}. The required strength of the seed field is subject to large uncertainties; past authors have quoted $10^{-21\\pm 2}$ G as the lower bound at the time of completed galaxy formation. This would present a problem for most particle-physics and field-theory inspired mechanisms of magnetic field generation. However, in recent work with A.-C.\\ Davis and M.J.\\ Lilley \\cite{bbounds}, I have shown that the lower bound on the dynamo seed field can be significantly relaxed if the universe is flat with a cosmological constant, as is suggested by recent supernovae observations \\cite{supernovae}. In particular, for the same dynamo parameters that give a lower bound of $10^{-20}$ G for $\\Omega_0=1$, $\\Omega_\\Lambda=0$, we obtain $10^{-30}$ G for $\\Omega_0=0.2=1-\\Omega_\\Lambda$, implying that particle-physics mechanisms could still be viable. The observation at redshift $z=0.395$ \\cite{redshift} can also be accounted for with these parameters, but requires a seed field of at least $10^{-23}$ G \\cite{bbounds}.  The magnetic field is amplified by the dynamo until its energy reaches equipartition with the kinetic energy of the ionised gas, $\\langle B^2/2\\rangle=\\langle\\rho v^2/2\\rangle$. Thus a final field of $B_0\\approx 10^{-6}$ G results for any seed field of sufficient strength.  It is difficult to explain the galactic field strength without a dynamo mechanism. To this end, one would require a strong primordial field of $10^{-3} (\\Omega_0 h^2)^{1/3}$ G  at the epoch of radiation decoupling $t_{\\rm dec}$, corresponding to a field strength $10^{-9} (\\Omega_0 h^2)^{1/3}$ G on comoving scales of $1$ Mpc. Future precision measurements of the CMB will put severe constraints on such a primordial field \\cite{CMB}. Moreover, magnetic fields on Mpc scales have been probed by observations of the Faraday rotation of polarised light from distant luminous sources, which give an upper bound of about $10^{-9}$ G \\cite{Kronberg}. The observation of micro-Gauss fields between galaxies in clusters presents an interesting dilemma. Because such regions are considerably less dense than galaxies, it is doubtful whether a dynamo could have been operative. Thus the intra-cluster magnetic fields, unless somehow ejected from galaxies, have formed directly from a primordial field stronger than $10^{-3} (\\Omega_0 h^2)^{1/3}$ G at $t_{\\rm dec}$. Such a field would certainly leave a signature in future CMB data \\cite{CMB}.  Although particle-physics inspired models typically produce weak seed fields, they tend to give precise predictions, unlike many astrophysical mechanisms, where the magnetic field strength is determined by complicated nonlinear dynamics, or solutions of general relativity with a magnetic field \\cite{Thorne}, where the field strength must be fixed by observations. With the possible exception of the last-mentioned model, there is no compelling scenario that produces a primordial field strong enough to eliminate the need for a dynamo.  Seed fields for the dynamo can be astrophysical or primordial. In the former category there is the important possibility that a seed field may arise spontaneously due to non-parallel gradients of pressure and charge density during the collapse of a protogalaxy \\cite{Kulsrud}. For the rest of this talk, however, I shall assume that the seed field is primordial. ",
        "conclusions": ""
    },
    "0004/astro-ph0004051_arXiv.txt": {
        "abstract": "A supermassive rotating toroidal black hole (TBH) is proposed as the fundamental component of quasars and other jet-producing active galactic nuclei. Rotating protogalaxies gather matter from the central gaseous region leading to the birth of massive toroidal stars, the internal nuclear reactions of which proceed very rapidly. Once the nuclear fuel is spent, gravitational collapse produces a slender TBH remnant. Transitory electron and neutron degeneracy stabilised collapse phases, although possible, are unlikely owing to the large masses involved thus these events are typically the first supernovae of the host galaxies. Given time, the TBH mass increases through continued accretion by several orders of magnitude, the event horizon swells whilst the central aperture shrinks. The difference in angular velocities between the accreting matter and the TBH induces a magnetic field that is strongest in the region of the central aperture and innermost ergoregion. Due to the presence of negative energy states when such a gravitational vortex is immersed in an electromagnetic field, circumstances are near ideal for energy extraction via non-thermal radiation including the Penrose process and superradiant scattering. This establishes a self-sustaining mechanism whereby the transport of angular momentum away from the quasar by relativistic bi-directional jets reinforces both the modulating magnetic field and the TBH/accretion disc angular velocity differential. Continued mass-capture by the TBH results in contraction of the central aperture until the TBH topology transitions to being spheroidal, extinguishing quasar behaviour. Similar mechanisms may be operating in microquasars, supernovae and sources of recurrent gamma-ray bursts when neutron density or black hole tori arise. In certain circumstances, long-term TBH stability can be maintained by a negative cosmological constant, otherwise the classical topology theorems must somehow be circumvented. Preliminary evidence is presented that Planck-scale quantum effects may be responsible. ",
        "introduction": "The commonly cited method of producing ultra-relativistic bi-directional jets as observed in quasars is the mechanism described by Blandford \\& Znajek (1978), whereby magnetic field lines thread the poles of a rotating BH as they descend towards the event horizon. Rotational energy may be extracted from the BH by this technique which is ejected in the form of radiation and matter travelling at high velocity along the BH's spin axis. Critical assessments by Ghosh \\& Abramowicz (1997) and Livio, Ogilvie \\& Pringle (1999) suggest that the role of the Blandford-Znajek mechanism has been generally overestimated and inadequately accounts for the larger double radio lobe structures. Numerical simulations indicate that the observed gamma ray energy release along quasar jets is four orders of magnitude more energetic than the Blandford-Znajek mechanism predicts. Issues that are difficult to reconcile with this model are the variability of jet dispersion angles, the finite quasar lifetime and the multiplicity of red-shifts in the very metallic absorption spectra. Any viable alternative model must simultaneously cater for all features.  Speculation concerning the fundamental processes governing quasars invariably involves discussion of compact massive central bodies, the consensus being that these are rotating spheroidal BHs of mass $10^6$--$10^9\\msun$. Profiles of stellar orbital velocities within AGN haloes lend weight to the premise that massive objects reside at the galactic nuclei. Inactive galactic nuclei (IGN) have yielded comparable velocity profiles, suggesting that the masses of AGN and IGN are similar, if not identical. The vastly differing activity levels could signify an inherent defect with current AGN models. Little theoretical progress seems to have been made over past decades towards a full explanation for these exceedingly energetic phenomena. Closer inspection and revision of existing theories may be necessary to attain a consistent understanding of quasars and AGN. The crucial test of any theory is the correspondence between predictions and observations, it is argued that existing models are struggling in this respect. The purpose of this discussion is to advocate a new model and describe how, with relatively minor theoretical embellishments, compelling explanations for AGN and anisotropic gamma-ray bursts accompanying core-collapse supernovae can be developed. Possible mechanisms responsible for originating, accelerating and collimating jets are then discussed.  The well known black hole uniqueness theorems rest on the classical theory of general relativity. It is thought that any non-distorted and asymptotically flat black hole spacetime can be represented by the Kerr-Newman set of solutions. But, while an ultimate theory remains elusive, one cannot be entirely sure of the validity of the uniqueness theorems. Hence, it is important to explore possibilities beyond those anticipated by purely classical calculations. Small departures from classical physics, e.g. of quantum mechanical origin, might lead to profound macroscopic changes, even to the extent that the topology of black hole horizons can be altered. Similarities between a rotating toroidal black hole (TBH) at the galactic centre accreting matter from its surroundings will be compared with observational evidence from quasars, Seyferts, BL Lacertae and blazars, which have long been suspected to be manifestations of the same underlying astrophysical phenomena. Attention will be paid to the formation of such a TBH, its long-term stability in our universe, jet production and its evolution with time. The possibility is examined that quasars may have been present at some stage of almost every galaxy's development in the early universe. A diagram encapsulating the life-cycles of toroidal BHs, in qualitative agreement with those of quasars is finally presented. ",
        "conclusions": "It has been qualitatively described how rotating TBHs might evolve from protogalactic gas clouds and accrete matter from the galactic centre until their inner apertures contract and highly focused relativistic jets form. The viability of the model can be tested by observing the evolution of jet collimation with red-shift as this model predicts the degree of collimation continues to increase (though jet energetics may decline at later times) until the topological transition. Such an approach could further address the issue of whether the entire AGN population or some subset thereof is accounted for by a TBH model. The direct observations of AGN and quasars in our universe suggests that TBHs are more than abstract mathematical constructs. Classical general relativity still remains to be unified satisfactorily with quantum mechanics. Evidently, TBH stability is intertwined with this issue. As direct experiments cannot be performed in intensely curved spacetimes, astrophysical observations must be our guide. The proposed stability of TBHs allows the event horizon's interior to be metaphorically unveiled, providing clues to the nature of quantum gravity and grand unification theories.  General relativity demands that the cosmological constant $\\Lambda$ be sufficiently negative in order to provide long-term TBH stability. Observational estimates of $\\Lambda$ based upon universal expansion reliant on the behaviour of general relativity in weak field environments suggest that its value is very small but probably positive. One possibility is that $\\Lambda$ is primarily a function of local spacetime curvature. Alternatively, the presence of external matter (accretion disk and galaxy) may provide the necessary curvature and non-stationarity permitting TBH stability over indefinite periods.  The TBH model, despite its controversial nature, presents a promising means of understanding the following characteristics of quasars: extreme jet energies, varying jet emergence angles, abrupt extinction, high gamma-ray radiation, the presence of heavy elements and the multiplicity of red-shifts in absorption spectra. None of these features are readily explained by spheroidal BH models. It is encouraging that the TBH model also appears to lead to plausible models for macroscopic processes within supernovae, microquasars and gamma-ray bursts. Fortunately, neutron tori -- unlike TBHs -- are exempt from topological censorship. Hence, these will certainly exist in astrophyscial circumstances, if only very briefly. Betraying their existence in distinctive and overtly energetic ways, these curiosities should be amenable to observational identification and study.  Numerical simulations are crucial if accurate comparisons with further detailed observations are to be made. Gravitational wave detectors and planned optical/X-ray interferometer technology will be sufficiently advanced in forthcoming decades to conclusively resolve the question of whether toroidal black holes truly exist."
    },
    "0004/astro-ph0004117_arXiv.txt": {
        "abstract": "We present a study of Jovian Trojan objects detected serendipitously during the course of a sky survey conducted at the University of Hawaii 2.2-meter telescope.  We used a 8192$\\times$8192 pixel charge-coupled device (CCD) mosaic to observe 20 deg$^2$ at locations spread over the L4 Lagrangian swarm and reached a limiting magnitude $V = 22.5$~mag (50\\% of maximum detection efficiency).  Ninety-three Jovian Trojans were detected with radii $2 \\leq r \\leq 20$ km (assumed albedo 0.04).  Their differential magnitude distribution has a slope of $0.40 \\pm 0.05$ corresponding to a power law size distribution index $3.0 \\pm 0.3$ ($1\\sigma$).  The total number of L4 Trojans with radii $\\geq 1$ km is of order $1.6 \\times 10^5$ and their combined mass (dominated by the largest objects) is $\\sim$$10^{-4}~M_{\\rm Earth}$. The bias-corrected mean inclination is $13.7^{\\circ} \\pm 0.5^{\\circ}$.  We also discuss the size and spatial distribution of the L4 swarm. ",
        "introduction": "The Jovian Trojans are asteroidal objects confined to two swarms in Jupiter's orbit, leading and trailing the planet by $60^{\\circ}$ of longitude (known as the L4 and L5 Trojans, respectively).  The first recognized Jovian Trojan (588 Achilles), discovered in 1906 by Max Wolf, was taken as providing observational confirmation of Lagrange's prediction of stable orbits at the triangular points.  Currently, 132 Jovian Trojans have been numbered while another 125 await permanent designations.  These objects follow loose orbits that librate around the L4 and L5 points with periods near 150 years.  Recent work has shown that Trojan orbits are destabilized by collisional ejection (for which the loss rate of bodies larger than 1 km in diameter is estimated at $\\sim$$10^{-3}$~yr$^{-1}$; Marzari et al. 1997) and, to a lesser extent, by dynamical chaos (corresponding loss rate $\\sim$$6 \\times 10^{-5}$~yr$^{-1}$; Levison, Shoemaker and Shoemaker 1997).  The implication is that the Trojans must either be the remnants of a much more substantial initial population of trapped bodies or that these objects are continually replenished from an unidentified external source.  The origin of the Trojans is a subject of much conjecture.  The principal dynamical problem concerns the nature of the dissipation needed to stabilize objects in weakly bound orbits librating about L4 and L5.  Schemes under consideration include capture of near-Jupiter planetesimals by gas drag in an early phase of the solar nebula (Peale 1993), stabilization of planetesimals near the L4 and L5 points due to the rapidly increasing mass of Jupiter in the late stages of its growth (Marzari and Scholl 1998), and collisional dissipation followed by capture of asteroidal fragments (Shoemaker et al. 1989). Physical observations provide only limited clues about the source of the Trojans.  The optical (Jewitt and Luu 1990; Fitzsimmons et al. 1994) and near-infrared (Luu, Jewitt and Cloutis 1994; Dumas, Owen and Barruci 1998) reflection spectra appear featureless, and are reminiscent of the spectra of the nuclei of short-period comets.  Like cometary nuclei, the Trojans have very low ($\\sim$4\\%) visual albedos (Cruikshank 1977; Tedesco 1989) that suggest carbonized surface compositions.  If the Trojans formed near or beyond Jupiter's orbit, temperatures were probably low enough for water to exist as solid ice (rather than vapor, as in the inner nebula).  This fact has led to the suggestion that the Trojans might possess ice-rich interiors equivalent to those of the cometary nuclei, a possibility which is not contradicted by any available observations (Jewitt 1996).  In this paper, we discuss the results of an optical survey taken in the direction of the L4 Jovian swarm.  The survey differs from most previous work on these objects in two main respects.  First, it is based on the use of a digital (CCD) detector instead of photographic plates and so has relatively high sensitivity to faint (small) Jovian Trojans.  Second, the parameters of the survey are extremely well known as a consequence of the relative ease with which digital data may be calibrated (compared to non-linear, analog photographic data). Therefore, we are able to measure the statistical properties of the L4 Trojans with greater confidence than would be possible with photographic data.  A preliminary abstract describing this work (Chen et al. 1997) is superceded by the present report. ",
        "conclusions": "\\noindent {\\it (a) \tLuminosity Function}  Photometry was performed using a circular aperture $4.7''$ in projected diameter, with sky subtraction from a contiguous annulus $3''$ wide (Table 1).  This aperture was selected by trial and error to give a stable measure of the flux while minimizing photometric noise from the background sky.  The statistical photometric uncertainty is $\\pm$0.2 mag at the faint end and less than $\\pm$0.1 mag at the bright end of the magnitude distribution.  For comparison with other work, it is useful to employ the absolute $V$ magnitude, $V(1,1,0) = V - 5\\log_{10}(R \\Delta)$, where $R$ [AU] and $\\Delta$ [AU] are the heliocentric and geocentric distances, respectively.  The correction to zero phase angle has been ignored (since the phase angles near opposition are small).  The distances $R$ and $\\Delta$ are not accurately known for the individual Trojan asteroids.  We adopt the mean distances of the numbered L4 Trojans at the epoch of observation, namely $R = 5.1 \\pm 0.2$ AU and $\\Delta = 4.1 \\pm 0.2$ AU.  Here, the quoted errors are $1\\sigma$ standard deviations and the dispersion in $R$ and $\\Delta$ results from the finite extent of the L4 swarm along the line of sight.  With these values we obtain $V(1,1,0) = V - (6.60 \\pm 0.24)$.  For reference, we further compute the Trojan radii, $r$ [km], from the relation  \\begin{equation} r = \\sqrt{\\frac{2.24 \\times 10^{16+0.4(V_{\\rm Sun} - V(1,1,0))}}{p_V}} \\end{equation}  \\noindent where $p_V$ is the geometric albedo at the $V$ wavelength and $V_{\\rm Sun} = -26.74$ is the apparent $V$ magnitude of the sun. The mean geometric albedo of the Trojans recorded by the IRAS satellite (Tedesco 1989) is $p_V = 0.040 \\pm 0.005$ (however, the quoted statistical uncertainty is probably smaller than inherent systematic uncertainties due to assumptions made in the radiometric modelling of the IRAS data).  With this $p_V$, we obtain  \\begin{equation} r_{0.04} {\\rm [km]} = 10^{0.2(24.23 - V)} \\end{equation}  \\noindent (Figure~4).  A Trojan at the 50\\% detection threshhold $V = 22.5$ has $V(1,1,0) = 15.9$ and $r_{0.04}$ = 2.2 km.  The brightest (faintest) Trojan detected in the present survey has $V = 17.7 (23.4)$ corresponding approximately to $r_{0.04} = 20.2$ km (1.5 km).  The distance variation across the diameter of the L4 swarm introduces an uncertainty to the derived radii of about $\\pm 10$\\%.  Figure~5 shows the cumulative luminosity function (CLF) computed from the present data.  The ``$\\times$'' marks in the figure show the distribution of the raw counts while the filled circles show the distribution corrected for the detection efficiency.  We have not included a correction for contamination of the Trojan sample by main-belt interlopers.  As noted above, this is a small effect whose inclusion would decrease the estimated surface densities at all magnitudes by about 10\\%.  Error bars in Figure~5 were estimated from Poisson statistics.  The luminosity function is taken to be of the form  \\begin{equation} N(V) {\\rm d} V = 10^{\\alpha(V-V_0)}~{\\rm d} V \\end{equation}  \\noindent A least-squares fit to the differential magnitude distribution gives slope parameter $\\alpha = 0.40 \\pm 0.05$.  The three lines in Figure~5 have gradients $\\alpha-1\\sigma, \\alpha$, and $\\alpha + 1\\sigma$, where $1\\sigma = 0.05$.  We assume that the radii of Trojans follow a differential power law distribution such that the number of objects having radii in the range $r \\rightarrow r+{\\rm d}r$ is $n(r)~{\\rm d}r = \\Gamma~r^{-q} {\\rm d}r$, where $\\Gamma$ and $q$ are constants. For objects all located at a single heliocentric and geocentric distance, $\\alpha$ and $q$ are related by $q = 5\\alpha + 1$ (Irwin et al. 1995).  Thus, the present data suggest $q = 3.0 \\pm 0.3$ in the radius range $2.2 \\leq r_{0.04} \\leq 20$ km.  For comparison, Shoemaker et al. 1989 estimated $\\alpha = 0.433$ for $10.25 \\leq B(1,0) \\leq 14$~mag (corresponding to $q = 3.17$ in the approximate range $4 \\leq r_{0.04} \\leq 40$ km, assuming $B-V \\sim 0.65$), but did not state their uncertainty.  We consider these determinations to be in good agreement.  The difference between $q$ measured here and the canonical $q = 3.5$ distribution produced by collisional shattering (Dohnanyi 1969) is statistically insignificant.  In any case, unmodeled effects will cause the Trojan distribution to differ from a Dohnanyi power law. For example, the velocity of ejection of collision fragments varies inversely with fragment size (as $r^{-1/2}$, Nakamura and Fujiwara 1991). Following collisional production the small Trojans should preferentially escape from the L4 region, leading to a distribution flatter than the Dohnanyi power law (i.e., $q < 3.5$).  We consider it likely that the small Trojan asteroids are collisionally produced fragments of once larger bodies (c.f. Marzari et al. 1997).  \\noindent {\\it (b) \tInclination-Frequency Distribution}  We are able to measure the position angles of the proper-motion vectors of the Trojans with an uncertainty of about $\\pm 1^{\\circ}$.  The proper-motion vectors cannot be accurately converted to orbital inclinations without a fuller knowledge of the orbits than we possess. The approximate relation between the direction of apparent motion and the true inclination is given by  \\begin{equation} \\tan\\theta_a = \\frac{\\tan(i)}{\\sqrt{R[1+\\tan^2(i)]}-1} \\end{equation}  \\noindent where $\\theta_a$ is the angle between the apparent direction of motion and the projected ecliptic, $i$ is the true orbital inclination, and $R$ is the semi-major axis of the orbit.  In deriving Eq.~(5) we have assumed that the orbital eccentricities are zero, that the proper motion is the vector difference of the intrinsic motion of the Trojans from the orbital motion of earth and that the observations are taken at opposition.  In the limit $i \\rightarrow 90^{\\circ}$, $\\tan\\theta_a \\rightarrow R^{-1/2}$ so that with $R = 5.2$ AU we find $\\theta_a \\rightarrow 24^{\\circ}$.  One object in Table 1 has $\\theta_a > 24^{\\circ}$ (\\#26).  Presumably, it is a main belt asteroid whose sky-plane velocity falls fortuitously within the Trojan domain.  The distribution of intrinsic inclinations (Figure~6a) has a mean $i = 7.4^{\\circ} \\pm 0.7^{\\circ}$, median $i = 6.2^{\\circ}$.  Trojans with large inclinations spend most of their time away from the ecliptic, leading to a bias against their detection.  The bias correction varies approximately in inverse proportion to the orbital inclination. Figure~6b shows the inclination distribution after correction by the factor $1/i$ and normalization to Figure~6a in the range $12 \\leq i \\leq 14^{\\circ}$.  The bias-corrected mean inclination is $\\bar{i} = 13.7^{\\circ} \\pm 0.5^{\\circ}$ which compares favorably with Shoemaker et al.'s 1989 best estimate for the mean inclination of the larger Trojans ($\\bar{i} = 17.7^{\\circ}$).  Observers of Trojan asteroids two decades ago suspected that ``there is a possibility that the inclinations are bimodal... with groups separated by a minimum at $i \\sim 13^{\\circ}$'' (Degewij and van Houten 1979).  There is indeed an apparent lack of Trojans in the corrected inclination distribution with $i = 14^{\\circ} \\pm 1^{\\circ}$, giving a bimodal appearance (Fig. 6b).  However, inspection of the raw data in Figure~6a shows that the local minimum is statistically insignificant.  Furthermore, the inclination distribution of numbered and un-numbered L4 Trojans (from an electronic list maintained by Brian Marsden at the Minor Planet Center) shows no evidence for bimodality (Figure~7).  Therefore, we conclude that the data provide no compelling evidence for a bimodal distribution of inclinations.  \\noindent {\\it (c) Size and Content of the L4 Trojan Swarm}  Figure~3 shows the variation of the surface density, $\\Sigma(\\theta)$~[deg$^{-2}$], of L4 Trojans with angular distance, $\\theta$, from the L4 point along the ecliptic.  The data can be fitted by a Gaussian function  \\begin{equation} \\Sigma(\\theta) d\\theta = \\Sigma_0 + \\Sigma_1\\exp \\left[\\frac{-\\theta^2}{2\\sigma_T^2} \\right] d\\theta,~~~~~~~(15^{\\circ} \\leq \\theta \\leq 60^{\\circ}) \\end{equation}  \\noindent with $\\Sigma_0 = 1.1 \\pm 0.4$~[deg$^{-2}$], $\\Sigma_1 = 27.7 \\pm 4.6$~[deg$^{-2}$] and $\\sigma_T = 11.2^{\\circ} \\pm 0.9^{\\circ}$~[deg$^{-2}$].  The apparent FWHM of the L4 swarm measured along the ecliptic is $26.4^{\\circ} \\pm 2.1 ^{\\circ}$, with measurable surface density for $\\theta \\leq 40^{\\circ}$.  The linear size corresponding to FWHM = $26.4^{\\circ}$ is $\\sim$2.4 AU.  For comparison, Holman and Wisdom (1992) found that the theoretical stable zones of the Jupiter Lagrange point have an angular radius of about $35^{\\circ}$.  The angular half width of the swarm along the ecliptic ($13.2^{\\circ} \\pm 1.0^{\\circ}$) is comparable to the bias-corrected mean inclination ($13.7^{\\circ} \\pm 0.5^{\\circ}$), meaning that we may take the projected outline of the swarm as approximately circular in the plane of the sky.  The number of L4 Trojans within angle $\\theta_{\\rm max}$ of L4 is then  \\begin{equation} N(\\theta_{\\rm max}) = \\int^{\\theta_{\\rm max}}_{0^{\\circ}} 2\\pi \\sin(\\theta) \\Sigma(\\theta){\\rm d}\\theta. \\end{equation}  \\noindent We plot solutions to Eq.~(7) in Figure~8.  The amplitude of libration about L4 ranges up to $60^{\\circ}$ (Shoemaker et al. 1989). Accordingly, we obtain the nominal population estimate $N(60^{\\circ}) = 3.4 \\times 10^4$ ($V \\leq 22.5, r_{0.04} \\geq 2.2$ km), plotted in Figure 8 as Model A.  The uncertainty on this number may be estimated in two ways.  Uniformly increasing (decreasing) the fitted parameters $\\Sigma_0$, $\\Sigma_1$ and $\\sigma_T$ by $1\\sigma$ changes $N$ by $\\pm35$\\% (c.f. Figure~8).  Systematic errors may also affect $N(\\theta_{\\rm max})$, particularly since the surface density at $\\theta \\leq 15^{\\circ}$ is not measured in our survey.  However, we find these errors to be small.  If, to consider an extreme case, we arbitrarily (and unphysically) assume that $\\Sigma(\\theta \\leq 15^{\\circ}) = 0$, we obtain (from Eq.~7) $N(60^{\\circ}) = 2.1 \\times 10^4$, still only 40\\% less than the nominal estimate. We conclude that $N(\\theta_{\\rm max})$ is uncertain to within a factor of order 2.  The cumulative luminosity function is replotted in Figure~9, including 132 numbered and 125 unnumbered Trojan asteroids (hollow circles) from orbital element catalogs maintained by the Minor Planet Center. We used $V(1,1,0) = H + 0.36$ to correct the catalog magnitudes to the $V$ band magnitudes employed here.  In making the comparison between the number of Trojans deduced from the present survey (Eq.~8) with those from the Minor Planet lists, we have corrected the former by a factor of two, to account for the unobserved L5 swarm. (Early suspicions that L4 might be more populated than L5 have not been borne out by recent data, supporting our application of a factor of two, Shoemaker et al. 1989). Curvature of the CLF at $V(1,1,0) \\geq 9.5~(r_{0.04} \\leq 42$ km) indicates observational incompleteness in the Minor Planet Trojan sample, as does the fact that the catalog asteroids are less numerous than those of the present survey by 1 to 2 orders of magnitude in the common $11 \\leq V(1,1,0) \\leq 14$ mag range (Figure~9).  Thirty seven objects have $V(1,1,0) < 9.5$.  Their effective CLF has slope $\\alpha = 0.89 \\pm 0.15$, significantly steeper than measured from the fainter objects of the 8k survey.  The implied size distribution index is $q = 5.5 \\pm 0.9~(V(1,1,0) < 9.5, r_{0.04} \\geq 42$ km).  The data of Figure~9 are adequately fitted by the following differential size distributions:  \\begin{equation} n_1 \\left(r_{0.04} \\right) {\\rm d}r_{0.04} = 1.5 \\times 10^6 \\left( \\frac{1{\\rm~km}}{r_{0.04}} \\right)^{3.0 \\pm 0.3} {\\rm d}r_{0.04}~~~~~(2.2 \\leq r_{0.04} \\leq 20{\\rm~km}) \\end{equation}  \\noindent and  \\begin{equation} n_2 \\left(r_{0.04} \\right) {\\rm d}r_{0.04} = 3.5 \\times 10^9 \\left( \\frac{1{\\rm~km}} {r_{0.04}} \\right)^{5.5 \\pm 0.9} {\\rm d}r_{0.04}~~~~~(r_{0.04} \\geq 42{\\rm~km}). \\end{equation}  \\noindent The corresponding integral distributions are  \\begin{equation} N(>r_{0.04}) = 1.6 \\times 10^5 \\left( \\frac{1{\\rm~km}}{r_{0.04}} \\right)^{2.0 \\pm 0.3}~~~~~(2.2 \\leq r_{0.04} \\leq 20{\\rm~km}) \\end{equation}  \\noindent and  \\begin{equation} N(>r_{0.04}) = 7.8 \\times 10^8 \\left( \\frac{1{\\rm~km}}{r_{0.04}} \\right)^{4.5 \\pm 0.9}~~~~~(r_{0.04} \\geq 42{\\rm~km}) \\end{equation}  \\noindent From Eq.~(10) we find the number of L4 Trojans with $r_{0.04} \\geq 5$ km is $N \\sim 6400$, to within a factor of order 2. For comparison, there were 5700 numbered and 1100 unnumbered main-belt asteroids with $r_{0.04} \\geq 5$ km as of 1999 July 29.  These estimates validate the assertion by Shoemaker et al. (1989) to the effect that the populations of the main-belt and the Trojan swarms are of the same order.  The number of L4 Trojans with $r_{0.04} \\geq 1$ km is $\\sim 1.6 \\times 10^5$, from Eq. (11).  The most straightforward explanation of the slope differences (Eq.~8 and 9 and Figure~9) is that the large objects represent a primordial population while Trojans smaller than a critical radius, $r_c$, are produced from the larger ones by collisional shattering (Shoemaker et al. 1989, Marzari et al. 1996).  By equating Eqs.~(10) and (11) we find $r_c \\sim 30$ km (corresponding to $V(1,1,0) = 10.2 \\pm 0.5$ (Figure~9)).  Binzel and Sauter (1992) reported that Trojans with $r_{0.04} > 45$ km have a larger mean lightcurve amplitude (a measure of elongated body shape) than their low albedo main-belt counterparts. They suggested that this might mark the primordial/fragment transition size, with larger bodies retaining the aspherical forms in which they were created (we note that this explanation is clearly non-unique). Inspection of their Figure~21 shows that the transition radius defined in this way is uncertain to within a factor of 2 and fully compatible with $r_c \\sim 30$ km as found here.  We conclude that two independently measured physical parameters (the size distribution and the lightcurve amplitude distribution) show evidence for a change near $r_c \\approx 30$ km to 40 km.  The total mass of Trojans is  \\begin{equation} M_T = \\int^{r_c}_0 \\frac{4}{3} \\pi \\rho n_1(r){\\rm d}r + \\int^{\\infty}_{r_c} \\frac{4}{3} \\pi \\rho n_2(r) {\\rm d}r \\end{equation}  \\noindent where $n_1$ and $n_2$ are from Eqs.~(8) and (9) and $\\rho = 2000$ kg m$^{-3}$ is the assumed bulk density.  We find $M_T \\approx 5 \\times 10^{20}$ kg $\\approx 9 \\times 10^{-5} M_{\\rm Earth}$, equivalent to a 400 km radius sphere having the same density. Dynamical calculations show that escaped Trojans would be quickly scattered into orbits indistinguishable from those of some short-period comets (Marzari et al. 1995).  Therefore it is of interest to compare $M_T$ with the mass of short-period comets delivered to the inner solar system over the past 4.5 Gyr.  The rate of supply of short period comets is $f \\sim 10^{-2}$ yr$^{-1}$ (Fernandez 1985).  The size and mass distributions of the cometary nuclei have not been adequately measured.  We assume that the nuclei follow a differential size distribution with index $q = 3$ and minimum and maximum radii $r_1 = 0.5$ and $r_2 = 30$ km.  The mass-weighted mean radius drawn from this distribution is $\\bar{r} \\approx (2r_2 r_1^2)^{1/3} = 2.5$ km.  A small number of well-measured cometary nuclei have radii $\\sim$ few km (Jewitt 1996), consistent with this estimate.  The delivered mass of short-period comets is then $M_C \\approx 4 \\pi \\rho \\bar{r}^3 fT/3$, where $T = 4.5 \\times 10^9$ yr is the age of the solar system.  We find $M_C \\approx 6 \\times 10^{21}$ kg, corresponding to $M_C \\approx 10~M_T$.  This mass, if taken at face value, makes it unlikely that the Trojan swarms could be the dominant source of the comets.  It is entirely possible, however, that a fraction of the short-period comets could be escaped Trojans. A definitive estimate of this fraction will require better knowledge of the cometary parameters ($r_1$, $r_2$, $q$, $f$, and density) than we now possess, as well as detailed understanding of the physics of collisional ejection from the Trojan swarms.  The present results were extracted from data taken for an independent (Kuiper Belt) purpose.  They serve to give an idea of the power of modern CCD arrays on a telescope of rather modest diameter.  Much more could be learned from a survey specifically targeting the Trojan asteroids and including astrometric follow-up, so that orbital elements can be determined for individual objects.  Future work should focus on a more complete digital survey of both L4 and L5 swarms in order to determine the total population and size distribution with greater confidence. Observations taken away from the ecliptic will provide a better measure of the high inclination objects.  Observations to fainter limiting magnitudes will allow us to probe the sub-kilometer population.  Carefully planned observations will produce stronger constraints on the collisional and dynamical states of the Jovian Trojans, leading ultimately to a deeper understanding of these enigmatic bodies."
    },
    "0004/astro-ph0004231_arXiv.txt": {
        "abstract": "We have obtained V and I-band photometry for 1886 stars down to $I=27$ and $V=28$ in the field of the dwarf irregular galaxy WLM, using deep STIS $CL$ and $LP$-band images taken in the parallel mode with the Hubble Space Telescope. The photometry is used to study the horizontal branch identified in WLM. The horizontal branch, extending blueward from the red giant clump, is an unambiguous signature of an old population.  We demonstrate that it is possible to reach the horizontal branch of an old population at a distance of $\\simeq 1$ Mpc using STIS, with relatively short exposure times.  {}From the VI color--magnitude diagram we obtain an accurate distance modulus $(m-M)_0 = 24.95 \\pm 0.13$ for WLM by using the V-magnitude of the horizontal branch, and by adopting $E(V-I) = 0.03$.  The implications are: 1) WLM formed stars at high redshift, and 2) the old population of WLM can be representative of a protogalactic fragment, related to those accreted to form the Milky Way halo. ",
        "introduction": "In 1997 February, the second Space Shuttle servicing mission to the Hubble Space Telescope (HST) installed the Space Telescope Imaging Spectrograph (STIS). This new spectrograph can be used in direct imaging mode providing very deep high-resolution images. The two optical wideband choices on STIS are a clear aperture ($CL$) and a longpass ($\\lambda > $5500 \\AA) filter ($LP$). Based on these two wide bands, Gregg \\& Minniti (1997) defined an effective shortpass filter (SP), thus establishing a two color photometric system. They suggested that these unusually wide bandpasses of STIS offer the possibility to reach much fainter limiting magnitudes with respect to the other HST imaging instruments like the Wide Field Planetary Camera 2 (WFPC2). STIS is expected to reach $\\sim$1.5 mag fainter than WFPC2 for a given integration time. This advantage of STIS over WFPC2 is due to three factors: (i) much broader passbands; (ii) the higher quantum efficiency and (iii) the lower readout noise of STIS's CCD.  By using spectrophotometric stellar libraries, Gregg \\& Minniti (1997) have argued that it is possible to convert the STIS magnitudes and (SP$-$LP) color to the standard photometric systems. They derived preliminary transformations to the VI Kron--Cousins wideband system. These transformations were based upon preflight specifications of the instrument and had to be empirically calibrated. Such an empirical calibration is now possible for the data collected during the STIS Parallel Survey (see, e.g., Gardner et al.\\ 1998).  The STIS Parallel Survey  data contain observations of the dwarf irregular galaxy DDO221 (= WLM) located at the edge of the Local Group (Sandage \\& Carlson 1985; van den Bergh 1994). In a photometric study of this galaxy, Minniti \\& Zijlstra (1996, 1997) presented VI photometry obtained with the NTT at La Silla Observatory. There is an overlap between he STIS field of WLM and the NTT field, enabling us to establish the empirical calibrations of the STIS photometric system. Based on this calibration, we construct a VI color--magnitude diagram (CMD) that is much deeper than those of any other Local Group galaxy studied so far with a comparable amount of exposure time (see Section~2.2), reaching V$\\sim$28 and  I$\\sim$27. The analysis of the CMD provides us with the fundamental properties of WLM's stellar populations, such as, metallicities, reddening, age and distance (Lee 1993; Minniti and Zijlstra 1996; Dolphin 2000). These properties are essential for establishing the formation history and evolution of WLM, a typical dwarf irregular galaxy that is representative of the most numerous group of galaxies in the Universe (Dekel \\& Silk 1986). The importance of WLM is stressed by the fact that it is located far from any other large spiral member of the Local Group, and therefore most probably has remained isolated over the Hubble time (Peebles 1995). The nearest galaxy to WLM is the recently discovered Cetus dSph, at a distance of $\\sim175$ kpc (Whiting et al.\\ 1999). Having remained isolated, WLM carries the unperturbed imprint of its formation and evolution in its stellar content.  The observations, data reductions, and photometry of WLM, and the resulting color--magnitude diagrams are given in Section 2. An empirical calibration of the STIS photometric system is given in section 2.3.  Fundamental parameters of WLM (metallicity, age and distance) are redetermined in Section 3.  The conclusions and implications of this work are summarized in Section 4.  \\section {The Data}  \\subsection {Ground-based Photometry}  The ground-based observations of WLM were obtained during the night of September 5, 1995, as part of a long term program monitoring the variable stars in several galaxies (Zijlstra, Minniti \\& Brewer 1997). We used the red arm (RILD mode) of the ESO Multi-Mode Instrument (EMMI) on the New Technology Telescope (NTT) at La Silla Observatory, operated remotely from the ESO headquarters in Garching, Germany. We used the 2048$\\times$2048 Tek CCD with a scale of $0.268\\arcsec$ pix$^{-1}$.  The weather was photometric, with seeing varying from $0.5\\arcsec$ to $1.0\\arcsec$.  The observations of WLM were secured with the V and Cousins I filters. They consist of pairs of frames with exposure times of 900 sec in I (airmass 1.1), and 1800 sec in V (airmass 1.1). These frames, covering an area of $9\\arcmin \\times 9\\arcmin$, are centered on the galaxy. Figure~1 shows a portion, covering only $1\\arcmin \\times 1\\arcmin$, of the NTT+EMMI images.  The reductions of the CCD frames were carried out by following standard procedures, using the package CCDRED within the IRAF environment. The pairs of images observed through each filter were combined. The final FWHM of stellar images were 0.8 and 0.7 arcsec in V and I respectively.   A total of 5 standards were observed during the night in order to calibrate our photometry. The photometric transformations to the standard system were done following the procedure described by Minniti \\& Zijlstra (1997). These transformations reproduce, with an rms of 0.016 in V and 0.031 in I, the magnitudes of the standard stars.  The photometric measurements were performed within the IRAF environment using DAOPHOT II, an improved version of the original DAOPHOT package developed by Stetson (1987).  In particular, DAOPHOT II can accommodate a spatially varying point-spread function across the field. All stars in the field of WLM that were more than 4$\\sigma$ above the background in the V and I frames were located, and their magnitudes measured by fitting a Moffat point-spread function. The resulting NTT limiting magnitudes ($4\\sigma$) were $V= 24$, and $I= 23.5$.   \\subsection {STIS Photometry}  The STIS images were acquired through the HST parallel program 7912 in September 1998. The center of the observed field was $\\alpha_{2000}=00^{\\rm h} 02^{\\rm m} 03.19^{\\rm s}$, $\\delta_{2000}=-15^\\circ 23\\arcmin 57.5\\arcsec$, at the N-E edge of WLM's disk. The $LP$-band image covers $28\\arcsec \\times 50\\arcsec$, and the $CL$-band image covers $50\\arcsec \\times 50\\arcsec$, but only the overlap region has been used. Four images with a total exposure time of 1680 seconds were combined to obtain the final $CL$-band image, and four images with a total exposure time of 2665 seconds were combined to obtain the final $LP$-band image. Cosmic rays were identified and eliminated in this process.  In the $28\\arcsec \\times 50\\arcsec$ STIS images we detect 1886 point sources, reaching more than 3 magnitudes fainter than the faintest stars seen in the NTT images. In particular, some of the single sources in the NTT images are resolved into multiple sources by STIS. Objects containing bad pixels were eliminated from our final photometry list.  Figure~2 shows the STIS $LP$-band image, with the orientation indicated in Figure~1. This parallel observation with STIS falls at the edge of the WLM disk, and is expected to include disk and halo stars.  The photometry was performed with DAOPHOT~II package in IRAF. For the zero-point calibration we followed instructions in the ``HST Data Handbook'' (1997). Using the values of the inverse sensitivity (PHOTFLAM) and HST zero-point (PHOTZPT) keywords in the headers of the images, zero-points of 26.067 and 26.518 were calculated for $LP$ and $CL$ respectively. The aperture corrections applied to the magnitudes measured through apertures of 2 pixels in diameter were $-0.575$ for $LP$ and $-0.484$ for $CL$.  The STIS CCD suffers from imperfect Charge Transfer Efficiency (CTE), as recently reported by Gilliland, Goudfrooij \\& Kimble (1999). We have used their formulation of the CTE to estimate its effect on our photometry. It turns out that, except for very faint stars that are near our detection threshold ($V \\ga 27.5$) and near the bottom of the CCD, the effect on our magnitudes is below 0.01 mag, which is negligible for our purposes.  Figure~3 shows the photometric errors as function of instrumental $CL$ and $LP$ magnitudes.  We note that since the galaxy/star ratio increases rapidly at faint magnitudes (Tyson 1988), the background galaxy contamination also has to be  taken into account. Using the number counts in the STIS field of the Hubble Deep Field South image (Gardner et al.\\ 2000), we estimate that in the field $\\sim$85 galaxies are present between $CL$ magnitudes of 26 and 28. Most of the galaxies are resolved on our STIS images and are discarded by our choice of the ``sharpness'' parameter in the DAOPHOT~II routines. Only a few fainter and very compact galaxies may be confused with stars, but these do not have a significant impact on the present work.  Figure 4 shows the $CL ~vs ~CL-LP$, and $LP ~vs ~CL-LP$ instrumental color-magnitude diagrams for a total of 1886 stars with centroids matched in $CL$ and $LP$ frames to better than 1.0 pixels (0\\farcs05), and which satisfy stringent criteria of photometric quality ($\\sigma \\leq 0.5$, $\\chi \\leq 2$, and sharpness $\\geq -1$). The sharpness criterion eliminates galaxies as well as possible star clusters at the distance of WLM (see Sect.~3.1). The filled circles in Figure~4 represent the stars in common with the NTT images (41 in total) that are used for the empirical calibration of the STIS photometry.     \\subsection {Empirical Calibration of the STIS Photometric System}  After STIS became operational in space, a significantly lower throughput of the CCD+LP system has been observed (e.g. Gardner et al. 2000). Although it is still possible to define the effective short bandpass by subtracting, from the $CL$ flux, the $LP$ flux scaled by 1.31, we use here only the $CL$ and $LP$ bandpasses for calibration. This simplifies the estimation of the uncertainties in the final photometry.  The photometry of the 41 STIS stars in common with the NTT observations is listed in Table~1. Respectively, columns 2 and 3 give the x and y positions of the stars on the NTT chip, columns 4 and 5 give V and I magnitudes of these stars, columns 6, 7, 8 and 9 give x and y positions of the stars on the STIS chip and their $CL$ and $LP$ magnitudes. For some of these stars careful inspection of the STIS images showed crowding, i.e. more than 2-3 stars with different magnitudes at the position of one NTT ``star''. These were not used for the calibration and are flagged with two asterisks in the last column of Table~1. Two more calibrators were rejected because of their marginal detection in the NTT image which carries a big magnitude error (flagged with one asterisk in Table~1).  The remaining 33 stars were used to transform $CL$ into V magnitude, and $LP$ into I magnitude. Figure~5 shows the comparison of the photometry. The relationships are approximately linear, as expected within the color range  $-0.8<V-I<2.0$ under consideration (Gregg \\& Minniti 1997). Using a least square fit the following transformation equations were obtained:  \\begin{equation} (V-I) = 1.76 (\\pm 0.13) \\times (CL-LP) + 1.18 (\\pm 0.03) \\end{equation} \\begin{equation} (V-CL) = 0.26 (\\pm 0.19) \\times (CL-LP) +  0.01 (\\pm 0.04) \\end{equation} \\begin{equation} (I-LP) = -0.51 (\\pm 0.14) \\times (CL-LP) - 1.17  (\\pm 0.03) \\end{equation}  The scatter about the mean relations is significant: $\\sigma_{\\rm rms}=0.16$, 0.24 and 0.17 for the first, second and third equation respectively.  This is largely due to the NTT photometric errors at faint magnitudes ($\\sigma_V>0.08$ for $V>22.5$, $\\sigma_I>0.10$ for $I>22$; Minniti \\& Zijlstra 1997), and the small number of acceptable calibrators in the blue part of the color range, $-0.8<V-I<-0.4$. There is also an inherent scatter because of the differences in widths and central wavelengths between $CL$ ($\\lambda_c$=5851.5\\AA) and V ($\\lambda_c$=5426\\AA) as well as between $LP$ ($\\lambda_c$=7228.5\\AA) and I ($\\lambda_c$=7985\\AA) filters. This scatter will depend on abundances, temperature and surface gravities, but because of the homogeneity of our comparison stars,almost all being red giants or supergiants, we estimate that it is smaller than the scatter induced by the NTT photometric errors. The complete photometry list, given in Table~2, will appear only in the electronic edition. The table contains x and y positions of each star on the chip, CL and LP magnitudes with their respective uncertainties as well V and I magnitudes calibrated using the calibration equations (1)--(3). The uncertainties given for V and I magnitudes come from the errors in CL and LP photometry and from the transformations.  Our calibration equations are only preliminary and valid for a limited color range ($-0.8<(V-I)<2$). Since essentially all our comparison stars are red giants or supergiants, the transformation may be different for dwarfs (Gregg \\& Minniti 1997). The complete calibration of the STIS photometry to the standard photometric system, as well as a more quantitative analysis of the influence of different physical parameters (i.e. temperature, abundances, surface gravity) on the transformation, will be published elsewhere (Goudfrooij et al. in preparation). ",
        "conclusions": "We have obtained deep STIS $CL$ and $LP$-band photometry for 1886 stars in the field of the dwarf irregular galaxy WLM, located at the edge of the Local Group.  We have calibrated the STIS photometric system empirically onto the $V$ and $I$-band system using our previous observations with the NTT telescope (Minniti \\& Zijlstra 1997). The VI color-magnitude diagram of a field located in the transition region between WLM's disk and halo shows a prominent RGB, a blue main sequence, a red clump, and an extended HB. The presented CMD was obtained with HST+STIS over less than 2 orbits, and revealed the presence of a HB in this distant Local Group galaxy. These results prove that it is now possible to reach the blue horizontal branch out to $D\\simeq 1$ Mpc using STIS with relatively short exposure times ($\\la 1$ orbit), allowing the study of the horizontal branches throughout the Local Group.  We have obtained an accurate distance modulus for WLM of $(m-M)_0=24.95\\pm 0.13$, based on the magnitude of the horizontal branch. This value agrees well with recent distance determinations using the RGB tip and Cepheids.  Our main result is, however, the detection of a horizontal branch that extends blueward from the red giant clump, an unambiguous indicator of the presence of an old population (age$ > 10^{10}$ yr) in WLM. It is likely to be associated with the halo of the galaxy, which is found to be an  extended metal-poor component by Minniti \\& Zijlstra (1996). The old age measured for the globular cluster of this galaxy by Hodge et al. (1999) favors such an interpretation.  Kinematics of individual stars is needed to compare the disk population with the halo population to check if the latter is kinematically hotter (with higher velocity dispersion and lower rotational velocity).  There are two major implications of these results:  1) WLM formed stars more than 10 Gyr ago. Various cosmological models give different redshifts for the same lookback time (Carroll, Press \\& Turner 1992). For example, according to the model of a flat universe with a plausible matter ($\\Omega_M=0.1$) and without the cosmological constant ($\\Omega_{\\Lambda}=0$), WLM formed stars at $z>4$, while according to the cosmological model of a flat universe with $\\Omega_M=0.1$ and $\\Omega_{\\Lambda}=0.9$, WLM formed stars at $z>1$. Star formation in this galaxy was not inhibited at early times by photoionization by UV sources such as QSOs and AGNs. This does not favor scenarios where dwarf galaxies form recently (e.g. Babul \\& Rees 1992).  2) The old population found in the WLM halo is similar to that of the old and metal-poor halos of the spiral galaxies in the Local Group. Having remained isolated throughout its lifetime, WLM can now be considered as a good example of a surviving protogalactic fragment, related to those that built up the Milky Way halo."
    },
    "0004/astro-ph0004007_arXiv.txt": {
        "abstract": "We present a measurement of the abundance of Fe relative to H in the solar corona using a technique which differs from previous spectroscopic and solar wind measurements. Our method combines EUV line data from the CDS spectrometer on {\\it SOHO} with thermal bremsstrahlung radio data from the VLA. The coronal Fe abundance is derived by equating the thermal bremsstrahlung radio emission calculated from the EUV Fe line data to that observed with the VLA, treating the Fe/H abundance as the sole unknown. We apply this technique to a compact cool active region and find Fe/H = 1.56 $\\times$ 10$^{-4}$, or about 4 times its value in the solar photosphere. Uncertainties in the CDS radiometric calibration, the VLA intensity measurements, the atomic parameters, and the assumptions made in the spectral analysis yield net uncertainties $\\sim$ 20\\%.  This result implies that low first ionization potential elements such as Fe are enhanced in the solar corona relative to photospheric values. ",
        "introduction": "It seems remarkable that the absolute abundances (here taken to mean abundances relative to hydrogen) of trace elements in the solar corona are still a matter for debate when many measurements for much more distant stars exist. The reason for this situation is easy to understand. Classic spectroscopic techniques for measuring relative abundances require that we compare the amount of radiation emitted in spectral lines whose atomic physics we understand and which arise from the two different elements in question. In the solar corona H and He are fully ionized, and hence produce no lines. In stars with hotter coronae, thermal bremsstrahlung produces a continuum emission which can be used instead of a spectral line as a measure of the amount of H present, but the Sun's quiescent corona is so cool that the continuum present at X--ray or EUV wavelengths is relatively weak and difficult to measure.  It has been known for some time that there is a difference in the relative abundances of elements in the photosphere and the corona organized according to the first ionization potential (FIP) of the element: high--FIP elements such as O are underabundant in the corona relative to low--FIP elements such as Fe by a factor of 4 \\citep{Mey85}. However, it has not been clear whether low--FIP elements are enhanced in the corona in an absolute sense (i.e., as a fraction of the total ion number density, dominated by H  and He) or the high--FIP elements are depleted relative to H. Two main techniques have been used to infer absolute abundances in the solar corona \\citep[]{FSH99}. One can observe solar flares, which produce plasma hot enough for the continuum to be measured, and compare the continuum with a suitable set of lines; or one can measure particles directly in the solar wind, and assume that the abundances there truly reflect the abundances in the solar corona. Unfortunately, these two techniques have resulted in conflicting measurements for, e.g., the absolute abundance of iron. The flare measurements suggest that it is close to the absolute abundance in the photosphere \\citep{VeP81} or perhaps factors of 1.5 -- 2 larger \\citep{FlS99}, while the solar wind particle measurements suggest that it is 4 times the photospheric ratio \\citep{Rea92}.  In this paper we use a third technique to measure the absolute abundance $N_{\\rm Fe}/N_{\\rm H}$ which has several desirable features. We use EUV observations of lines of Fe in order to determine the amount of Fe present in the corona, but to measure the amount of H present we use radio continuum measurements of thermal bremsstrahlung from coronal plasma. This measurement therefore pertains to the quiescent solar corona, including non-flaring weak active regions, rather than flare or solar wind plasma. Like classic spectroscopic techniques, this technique should be independent of the filling factor of plasma in the volume studied since both types of emission have the same dependence on number density. We describe the measurement technique and the observations used for this measurement in the following section. We then discuss the implications of our measurement. ",
        "conclusions": "Our measurement of an absolute abundance for Fe which is 4 times photospheric is clearly consistent with the interpretation that low--FIP elements such as Fe are enhanced relative to their photospheric abundances. Of the commonly--used solar abundance tables, Feldman's (1992) \\nocite{Fel92} is closest to our measurement with $\\log_{10}{A_{\\rm Fe}}$ = 8.10; his value  is based more on the argument that H should be treated as a high--FIP element rather than any specific measurement of $N_{\\rm Fe}/N_{\\rm H}$. Other commonly used values for $\\log_{10}{A_{\\rm Fe}}$ are the photospheric results 7.60 \\citep{All73}, 7.51 \\citep{AnG89} and 7.67 \\citep{GrA91}, the coronal results 7.59 \\citep{Mey85}, 7.83 \\citep{FlS99} and 8.50 \\citep{WMD94}, and the solar wind measurements 7.93 \\citep{Rea95} and 8.51 \\citep{Rea99}. We note in passing that spectroscopic measurements of the Fe abundance in active stars with coronae much hotter than the Sun's have tended to suggest abundances smaller than photospheric \\citep{MKW96,Dra96}, in contrast to the solar behaviour we find here.  We have reported here a single measurement of the Fe abundance. Spectroscopic studies of relative abundances suggest that they may vary considerably from flare to flare and from one region to the next \\citep{SLM84,Sab95,SLB98}, so it is clearly important to confirm our measurement. We have additional data sets which will be published in a longer paper elsewhere, but with larger uncertainties expected in the radio fluxes, and expect to carry out further studies of this type in the future."
    },
    "0004/astro-ph0004377_arXiv.txt": {
        "abstract": "In this contribution we investigate how a disk determines the physical and geometrical properties of the circumstellar environment of the T Tauri SU Aurigae. Our model of the spectral energy distribution of this star includes the central young star, a flat black disk and a diffuse envelope. We also describe the inner interaction region between the disk and the star, the magnetosphere, by analysing the accretion and wind signatures in several spectral lines observed during the MUSICOS~96 campaign. ",
        "introduction": "\\indent Young stellar objects have extremely complex circumstellar environments: material in an infalling envelope surrounding the central star falls onto a circumstellar disk, dissipating angular momentum; magnetic fields can then channel the disk material onto the star by way of accretion funnels (magnetospheric accretion), causing hot shock regions on the stellar photosphere.\\\\ \\indent Classical T Tauri Stars (CTTS) are low-mass pre-main sequence stars. These stars are active, with broad energy distributions and Balmer and metallic (permitted and forbidden) line emission. CTTS are class II objects, i.e. the systems include a central young star and an accretion disk. The presence of the disk was inferred from the IR and UV excesses in the spectral energy distributions and it can also account for some of the photometric variability and optical veiling; for some stars the disk has also been imaged directly.\\\\ \\indent Magnetospheric accretion can explain the photospheric hot spots and the He~I spectral line, as a shock region at the stellar surface results in the release of kinetic energy. Spectroscopic evidence of magnetospheric accretion are the redshifted absorption components on several spectral lines. The spectra of T Tauri stars also present blueshifted absorption features that are signatures of stellar or disk winds.\\\\ \\indent SU Aurigae is either classified as a CTTS (e.g. Bouvier et al. 1993) or as the prototype of a separate class of young stellar objects (Herbig \\& Bell 1988). The star's spectral type is G2 (Herbig 1960), earlier than most CTTS, and it is an exceptionally fast-rotator, v $\\sin i \\sim 66$~km~s$^{-1}$ (Hartmann \\& Stauffer 1989). Its photometric period has been rather difficult to determine, but it is usually considered to be close to 3~days. From our data set, we favour a somewhat shorter period of 2.5-2.8~days (Oliveira et al. 2000).\\\\ \\indent In this contribution we describe the SU~Aur system, from two different viewpoints: i) using the spectral energy distribution to characterize the larger scale circumstellar environment and ii) probing the disk-star interaction by analysing the inner magnetosphere, using results from our multi-site spectroscopic campaign on this target. ",
        "conclusions": "\\indent The enhanced activity observed in SU Aur can be explained by the presence of an accretion disk and spherical dust shell. We describe the SED of SU Aur as including the star, a flat black disk (a hot dust component to account for the near-IR excess emission) and a diffuse envelope (a cold dust component to account for the far-IR excess). The model is rather simple but it gives a very good agreement with the observations available.\\\\ \\indent Probing the inner magnetosphere, we observe a time-lagged behaviour of the variability of spectral lines with different excitation conditions that is not compatible with the standard magnetospheric models, even when the magnetic dipole is inclined with relation to the rotation axis. The behaviour we observe requires that the magnetic field lines are azimuthally distorted probably due to differential rotation in the star~+~disk system. Obviously, more stars need to be observed quasi-continuously to allow a generalization of the results presented here."
    },
    "0004/astro-ph0004188_arXiv.txt": {
        "abstract": "Recently, the detection of discrete features in the X-ray afterglow spectra of GRB970508 and GRB970828 was reported. The most natural interpretation of these features is that they are redshifted Fe K emission complexes. The identification of the line emission mechanism has drastic implications for the inferred mass of radiating material, and hence the nature of the burst site. X-ray spectroscopy provides a direct observational constraint on these properties of gamma-ray bursters. We briefly discuss how these constraints arise, in the context of an application to the spectrum of GRB970508. ",
        "introduction": "The detection of a discrete spectral feature in the X-ray afterglow spectrum of GRB970508 was recently reported by Piro {\\it et al.} (1999a,b). A similar feature in the afterglow of GRB970828 was reported by Yoshida {\\it et al.} (1999). The redshift of the host galaxy of GRB970508 was determined to be $z = 0.835$ (Metzger {\\it et al.} 1997; Bloom {\\it et al.} 1998). The apparent energy of the discrete feature in the X-ray spectrum is consistent with the redshifted energies of the $n = 2-1$ transitions in all possible charge states of Fe, {\\it i.e.}, from Fe K$\\alpha$ at 6.4 keV, to H-like Fe Ly$\\alpha$ at 6.95 keV. Therefore, none of the possible line excitation mechanisms (fluorescence, collisional excitation in hot gas, or recombination in photoionized gas or a transient collisional plasma) is currently ruled out on the basis of the measured line energy alone.  The conditions required for each of these emission mechanisms ({\\it i.e.} the presence of a given charge state of Fe) are all still compatible with the constraints on ionization parameter and gas temperature that can be derived from the fact that the line source must be located no more than about 1 light day from the ionizing source, and from the shape of the ionizing spectrum. All three mechanisms require high gas density ($n_e \\gax 10^{11}$ cm$^{-3}$) for bound Fe to exist at all this close to the burst site. The presence of a large mass of dense gas close to the burst has led to the suggestion that a merging neutron star-neutron star binary is an unlikely site for this burst, and that a scenario involving some sort of stellar collapse is more likely (Piro {\\it et al.} 1999b). In that case, the line source may be associated with debris from the stellar collapse.  Resolved X-ray spectroscopy can of course distinguish between the various line emission mechanisms, and lead to the correct characterization of the conditions in the source and the implied mass of radiating Fe. In the following, we will briefly recapitulate the constraints imposed by proximity and the energy budget, and then draw attention to an example of direct spectroscopic analysis of the X-ray spectrum: the fact that the measured centroid position, and to some extent even the shape, of the X-ray emission feature in the afterglow of GRB970508 already imply that recombination in afterglow-photoionized gas is unlikely to be the emission mechanism for this burst.  Our analysis applies to an optically thin, homogeneous medium. Complications arising from radiative transfer effects or external heat input by a relativistic shock may conspire to alter the shape of the spectrum, perhaps to the point that no definitive conclusions can be drawn from an intantaneous, low sensitivity spectrum. Such effects have been calculated by Weth et al. (1999) for the case of photoionization in the context of two specific geometries for the burst site, and by Vietri (1999) for a scenario in which the relativistic shock heats the medium shortly after the onset of the afterglow. With better data, it may hopefully be possible to reverse the argument, and infer the geometry, thermal history, and abundances from spectroscopy. ",
        "conclusions": "We find that Fe emission from a cool photoionized source in equilibrium with the afterglow cannot account for the X-ray spectrum of GRB970508. Fitting a spectral model for such a source either implies a large redshift, or indicates such high electron temperatures that the source would be closer to collisional ionization equilibrium.  This would imply that the discrete feature in the spectrum of GRB970508 arises from either Fe fluorescence or from collisional excitation, each of which implies a very different estimate for the total mass of radiating Fe ($M_{\\rm Fe} \\sim 2.6 \\times 10^{-4} M_{\\odot}$ for fluorescence, $M_{\\rm Fe} \\sim 7 \\times 10^{-2} M_{\\odot}$ for collisional excitation, which corresponds to $M_{\\rm total} \\sim 8 \\times 10^{-2} M_{\\odot}\\ (A_{\\rm Fe}/A_{{\\rm Fe},\\odot})^{-1}$ and $M_{\\rm total} \\sim 28 M_{\\odot}\\ (A_{\\rm Fe}/A_{{\\rm Fe},\\odot})^{-1}$, respectively (see, for instance, Meszaros \\& Rees [1998], Piro et al. [1999b], Lazzati et al. [1999]); here, $A_{\\rm Fe}$ is the abundance of Fe).  With higher spectral resolution and sensitivity, we may be able to distinguish spectroscopically between a fluorescent and a collisional spectrum (a recombination spectrum is readily identified from the presence of the RRC). Working from the K-shell spectrum of a single ionization stage of Fe alone, one might use the $n=2-1$ and $n=3-1$ transitions, which are resolved at $E/\\Delta E \\sim 10$. Intensity ratios do not effectively discriminate between collisional excitation and fluorescence: $I({\\rm Ly}\\beta)/I({\\rm Ly}\\alpha) = 0.10-0.15$ (Hydrogenic Fe) for electron temperatures between $3 \\times 10^7$ and $10^9$ K (Mewe, Gronenschild, \\& van den Oord 1985), while for neutral Fe $I({\\rm K}\\beta)/I({\\rm K}\\alpha) = 0.13$ (Kaastra \\& Mewe 1993). Instead, with somewhat higher resolution, one might use the fact that the ratio of the transition energies depends on the charge state: $E({\\rm Ly}\\beta)/E({\\rm Ly}\\alpha) = 1.18$, while $E({\\rm K}\\beta)/E({\\rm K}\\alpha) = 1.10$.  At higher resolution, one would resolve the $1s-2p_{1/2}$ and $1s-2p_{3/2}$ transitions ($E/\\Delta E > 500)$. The ratios of the energies are $E({\\rm Ly}\\alpha_2)/E({\\rm Ly}\\alpha_1) = 1.0030$ and $E({\\rm K}\\alpha_2)/E({\\rm K}\\alpha_1) = 1.0020$, and a resolving power of 1000 is required to distinguish between these two.  A collisional source at $kT < 15 $ keV will also show emission from He-like Fe, and the simultaneous appearance of both the H- and the He-like charge states readily indicates a high degree of ionization of the source. For temperatures in the range $3 < kT < 15$ keV, the He-like emission dominates, which could still be identified as such if the characteristic \"triplet\" structure of the $n=2-1$ lines is resolved ($E/\\Delta E > 300$). The highest resolution available at energies above 2 keV is provided by the High Energy Transmission Grating Spectrometer on {\\it Chandra} (Canizares et al. 2000). The Fe K spectrum will be resolved at $E/\\Delta E \\approx 150~(1+z)$, which would at least enable charge-state spectroscopy at moderately large redshift.  Finally, given sufficient sensitivity at low energies (and sufficiently low absorption in the source), we may of course detect emission from lower-$Z$ elements, either collisional or fluorescent, which would greatly facilitate the spectroscopic diagnosis of the excitation mechanism. With the grating spectrometers on {\\it Chandra} and {\\it XMM}, standard detailed plasma diagnostics can be invoked, at resolving powers of $E/\\Delta E \\sim 100-1500$ at $E \\lax 2$ keV, depending on photon energy (Canizares et al. 2000, Brinkman et al. 2000, Brinkman et al. 1998). Clearly, it would be very interesting to pursue more sensitive X-ray spectroscopy of gamma ray burst afterglows.  \\vskip 0.1in  We gratefully acknowledge discussions with Marten van Kerkwijk. D. A. L. was supported in part by a NASA Long Term Space Astrophysics Program grant (LTSA S-92654-F).  Work at LLNL was performed under the auspices of the U. S. Department of Energy, Contract No. W-7405-Eng-48."
    },
    "0004/astro-ph0004141_arXiv.txt": {
        "abstract": "BeppoSAX observations of the Seyfert galaxies Mrk 509 and MCG~+8-11-11 are presented. Earlier evidence of a soft excess in Mrk 509 is confirmed. This excess is found to be better represented by a power law than by a black body: with a photon slope $\\Gamma_s$ of 2.5, its extrapolation matches the flux recorded in the far UV. An ASCA observation, which appeared to exclude the presence of the excess while showing instead evidence of a warm absorber, turns out to be compatible with the coexistence of the excess seen with BeppoSAX and of the warm absorber. The hard power law of Mrk 509 is seen for the first time to be affected by a cut--off at high energies, with an e--folding energy of about 70 keV. In MCG~+8-11-11 the cut--off is found at about 170 keV, consistent within the combined errors with a previous estimate from a ASCA+OSSE/CGRO observation. In both objects the reflection component is clearly detected. In Mrk 509 its strength, together with that of the iron K line, indicates a solid angle $\\Omega$, subtended by the reprocessing gas in the accretion disk, much less than 2$\\pi$, but lacking a valid constraint on the inclination angle this evidence is not as convincing as that found with BeppoSAX in IC 4329A. In MCG~+8-11-11 the same parameters are instead consistent with $\\Omega$=2$\\pi$, but the comparison with an ASCA observation, when the flux level was about 2.5 times weaker, suggests that a substantial fraction of the angle might be associated with gas farther out than the accretion disk. ",
        "introduction": "As part of a broad band spectral survey, with the BeppoSAX satellite, of a sample of bright Seyfert 1 galaxies (2--10 keV flux greater than 1--2 mCrab), this paper follows those on NGC 4593 (Guainazzi et al. 1999) and IC 4329A (Perola et al. 1999), where the goals of the program are described in detail. In short terms, the survey aims at determining the parameters of the various spectral components typical of Seyfert 1s, and especially of the broad ones, with better confidence than could be done before, by exploiting the opportunity to cover simultaneously the band 0.1--200 keV, offered by the Narrow Field Instruments (NFI) onboard BeppoSAX.  This paper reports on the observations of two objects, namely Mrk 509 and MCG +8-11-11. Mrk 509 is a Seyfert 1.0 galaxy at $z$=0.0344 with a luminosity $L$(2--10 keV) $\\sim 3\\times10^{44}$ erg s$^{-1}$ ($H_o$= 50 km s$^{-1}$ Mpc$^{-1}$). The galactic interstellar column toward this object amounts to $N_{H,g}$=4.4$\\times10^{20}$ cm$^{-2}$ (Murphy et al. 1996).  MCG +8-11-11 is a Seyfert 1.5 galaxy at $z$=0.0205 with a luminosity $L$(2-10 keV)$\\sim10^{44}$ erg s$^{-1}$. Because of the low galactic latitude (b=10.4$^{\\circ}$) its spectrum is absorbed by a large galactic column, $N_{H,g}$=2.03$\\times10^{21}$ cm$^{-2}$ (Elvis et al. 1989).  Special reasons of interest from earlier observations of the two objects are the following. In Mrk 509 the report of a sizeable soft excess below about 1 keV in Ginga and ROSAT simultaneous observations by Pounds et al. (1994), which appeared to confirm previous claims in data from HEAO-1 (Singh et al. 1985), EXOSAT (Morini et al. 1987) and Ginga (Singh et al. 1990), was itself not confirmed in one ASCA observation by Reynolds (1997) and by George et al. (1998), who claim instead evidence of a warm absorber. In MCG +8-11-11 a high energy cut--off in the power law was detected in simultaneous ASCA and OSSE/CGRO observations at about 270 keV (Grandi et al. 1998).  Observations and data reduction are described in Sect. 2, the spectral analysis in Sect. 3. The results are commented and compared with previous ones in Sect. 4, and the conclusions are drawn in Sect. 5. ",
        "conclusions": "A particularly interesting result of the present study is the overall shape of the \"direct\" continuum in Mrk 509, which can be empirically described as a combination of two power laws. The soft power law we detect is, in the first place, a confirmation of earlier evidence of a low energy excess (see Sect. 1), and it seems to us a better description of the excess than the simple black body adopted by Pounds et al. (1994). With a slope $\\Gamma_s$ about 2.5, it appears as a continuation into the X-rays of the UV emission, and according to models sometimes invoked for this component (e. g. by Zheng et al. 1997, but see also Walter et al. 1994, Fiore et al. 1995, Laor et al. 1997 for a broader discussion of the problem), it could arise from comptonization in the innermost parts of the disk in regimes of high accretion rate (e. g. Czerny \\& Elvis 1987, Maraschi \\& Molendi 1990). The hard power law is affected by a high energy cut--off, which can be empirically described as an exponential factor with $E_f$ about 70 keV. The currently most popular models for this component involve comptonization by a hot corona (e. g. Haardt \\& Maraschi 1993, Zdziarsky et al. 1994, Ghisellini et al. 1998), in which the cut--off is related to the temperature of the corona. Our results on Mrk 509 are similar to those on NGC 4151, possibly indicating that Seyferts with unusually flat continua feature unusually low cut--off energies (Piro 1999). However it is fair to say that proper, selfconsistent models need to be used in the fits in order to recover correctly the physical parameters. As extensively discussed by Petrucci et al. (2000) in the context of a BeppoSAX observation of NGC 5548, the exponential cut--off powerlaw is indeed a quite rough approximation to realistic componization models.  Overall the spectrum of Mrk 509 requires a complex structure for the accretion disk, to explain both the soft and the hard components. This will be dealt with in a future paper, where we would like to address the problem in a comprehensive manner for all the objects in the BeppoSAX sample of bright Seyfert 1. There we shall include also a more advanced analysis of the comptonization model, underlying the continuum shape and in particular the high energy cut--off seen in MCG~+8-11-11, already discussed in Grandi et al. (1998).  This approach will also attempt to deal selfconsistently with the \"reprocessed\" spectral components, namely the amount of continuum reflection along with the energy/width/intensity of the iron line, contributed by the matter in the accretion disk. Here we limit ourselves to note that, although not conclusively yet, directly from this observation in Mrk 509 and from comparison with an ASCA observation in MCG~+8-11-11, there are indications of the accretion disk contributing to the reprocessing much less than one would expect if the solid angle $\\Omega$ were 2$\\pi$. Such a situation was also, and rather more convincingly found to hold in IC 4329A (Perola et al. 1999).  Finally, in both Mrk 509 and MCG~+8-11-11 our observations present some evidence of a warm absorber. This evidence is roughly consistent with the more significant one found with ASCA in Mrk 509 by Reynolds (1997) and by George et al. (1998), who did not however confirm in their analysis the presence of a soft excess. We have shown that the ASCA data are compatible with the coexistence of the warm absorber together with the soft excess observed by BeppoSAX. While doing that, we have noted that the interpretation of the ASCA data is complicated by a not yet perfectly understood, epoch-dependent, problem with the response in the ASCA-SIS low energy channels (T. Yaqoob, private communication): the implication (which is noted to be present, but with less of a consequence, in the ASCA observation of MCG~+8-11-11, Grandi et al. 1998) is equivalent to an additional $N_H$ of a few $10^{20}$ cm$^{-2}$, which is largely incompatible with our data.   \\bigskip \\bigskip  \\small {\\it Acknowledgements.} The BeppoSAX satellite is a joint Italian--Dutch program. We wish to thank the Scientific Data Centre for assistance. This research has made use of the TARTARUS database, which is supported by Jane Turner and Kirpal Nandra under NASA grants NAG5-7385 and NAG5-7067, and 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. This work was supported by the Italian Space Agency, and by the Ministry for University and Research (MURST) under grant COFIN98--02--32. \\normalsize"
    },
    "0004/astro-ph0004411_arXiv.txt": {
        "abstract": "We consider the effect of rapid rotation on the light curves of neutron stars with hot polar caps.  For $P \\approx 3$ms spin periods, the pulse fractions can be as much as an order of magnitude larger than with simple slowly-rotating (Schwarzschild) estimates. Doppler boosting, in particular, leads to characteristic distortion and ``soft lags'' in the pulse profiles, which are easily measurable in light curves with moderate energy resolution. With $\\sim 10^5$ photons it should also be possible to isolate the more subtle distortions of light travel time variations and frame dragging. Detailed analysis of high quality millisecond pulsar data from upcoming X-ray missions must include these effects. ",
        "introduction": "Many neutron star models have thermal surface emission concentrated at magnetic polar caps, which are heated by accretion columns, by precipitating magnetospheric particles, or by anisotropy in the thermal conductivity and heat flow from the interior.  A number of authors have computed light curves from such heated polar caps by assuming a stationary dipolar magnetic field and noting that gravitational focusing of the photons can have a substantial effect on the pulse shape (see, e.g. Riffert \\& M\\'{e}sz\\'{a}ros~1988; Zavlin, Shibanov, \\& Pavlov~1995).  It has also been suggested that resonant scattering on magnetospheric $e^\\pm$ may modify the thermal light curves \\citep{raj97,zhu97}.  Recently, pulsed X-ray emission, apparently from the stellar surface, has been detected from a number of very short period neutron stars, including the 5.75ms radio pulsar J0437-4715 \\citep{zav98} and the 2.49ms LMXB binary J1808.4-3658 \\citep{cui98,wij98}.  For such objects, existing static field and metric approximations may be inadequate.  In this paper we model the effects of rapid rotation on the the surface emission from polar caps.  The most prominent effects are associated with simple Doppler boosting at the surface \\citep{for99,str98}. We also model the perturbations associated with the varying time delays of the curved light paths and the pulse distortions resulting from frame-dragging (labeled here as ``Kerr'') effects on the space-time metric. Previously, \\cite{che89} noticed the above effects.  They modeled the neutron star as a Kerr black hole, and computed values for the star's angular momentum while assuming that the neutron star's surface is stationary in the LNRF (see below).  Because of this, they arrive at unphysical values for the angular momentum and only include Doppler shifts due to frame-dragging.  We use more realistic values of the frame-dragging angular velocity and include the dominant Doppler boosts caused by the motion of the star's surface in the LNRF. These effects are subtle, but with the increased X-ray sensitivity of {\\it AXAF}, {\\it XMM}, and possibly {\\it USA} careful observations may uncover signatures of rotational perturbations in high statistics energy-resolved light curves of the shortest-period pulsars. The effects will be even more prominent when emission arises from higher altitudes, as is likely for the non-thermal radio and high energy pulsations found in a number of millisecond pulsars (e.g.  PSR B1821-24, PSR J0218+4232).  Resonant scattering perturbations can also show much larger Doppler and Kerr effects since the flux is re-directed at several $R_{*}$ and sent back towards the star. Modeling these sources additionally requires a detailed treatment of the retarded (swept-back) $B$ field. ",
        "conclusions": "With sufficient count statistics we might hope to isolate the various contributions to the rotational perturbations of the pulsar light curves.  To gauge the observational precision required, we simulate data drawn from our model light curves and compare various scenarios.  One measure of the significance of the departures is to simulate the data with all rotational effects and to compare with model light curves ignoring either Kerr effects or time delay induced perturbations.  For example, for our $P=3$ms fiducial light curves, we find that pulse profiles depart significantly (95\\% confidence K-S statistic) from Doppler-alone models when the energy-summed total light curve contains $3\\times10^4$ counts. This does assume a simple model for the surface $T$ distribution (although gravitational focusing renders results surprisingly insensitive to the detailed cap shape).  To isolate rotational perturbations in the presence of an arbitrary cap shape and $T$ profile would require significantly more data.  Table~\\ref{tbl-1} summarizes the counts required to see Doppler and time-delay/frame-dragging induced distortions for various $P$.  Intra-band comparisons for a single pulsar are even more robust, since cap shape and $T$ effects will be correlated at the various energies. We have modeled the energy-dependent shape changes. At $P=3$ms we find that a total of $6\\times10^4$ counts are required to measure the Doppler induced differences between narrow-band lightcurves (10 bands assumed, logarithmically-spaced in energy).  For a nearby millisecond pulsar with $T_{\\rm cap, surf}=1.3 \\times 10^6$K, $N_H = 10^{19.5}$, we find the most statistically significant differences are typically between the energies (at infinity) less than 0.14~keV and greater than 0.28~keV.  These estimates are conservative in that we use only the pulse shape and phase differences.  The actual count ratios between energy channels may be used to check the surface $T$ dependence or to allow detection with even lower total counts.  We conclude that high precision models of pulsars such as J0437-4715 must include at least Doppler boosting at the surface. For example \\citet{zav98} have published ROSAT PSPC light curves with good statistics in energy bands extending to $1.1-2.0$~keV with $36\\arcdeg$ binning.  Using model atmospheres and Schwarzschild photon propagation, they have fit for $\\alpha$ and $\\zeta$, cap size and the detailed $T$ distribution.  Our simulations (Fig.~\\ref{all}) suggest that the model light curves (and the parameter values fit) ignoring rotational effects may be substantially in error. For example, fitted $M/R_*$ ratios must be considered preliminary until the effects of rotation are checked.  On the other hand the (formally independent) $R_*$ and $M$ dependencies that can be traced to Doppler and gravitational effects suggest that more precise tests of the EOS could be effected with very high quality data and detailed models.  High precision {\\it XMM} and {\\it AXAF} data, which are expected, should be suitable for such fitting.  A second example is the newly-discovered 2.49ms accretion-powered pulsar J1808-3658. This source displays strong ``soft lags'' \\citep{wij98,cui98} in the light curves. \\citet{for99} has already suggested that this pulsar, as well as a number of QPO sources, exhibits significant Doppler boosting which causes the high energy pulse to lead the softer emission.  Alternative explanations, such as Compton down-scatter in a soft corona have been suggested; however, the energy-dependent pulse shape distortions \\citep{cui98} appear very similar to those of Fig.~\\ref{energy}.  Our computations thus provide significant support for the Doppler boosting interpretation.  With more careful modeling it may possible to detect the effect of other rotational distortions to the light curves of this source.  Doppler boosting effects are very significant for these short period thermal cap emitters.  Time-delay and frame-dragging distortions, although more subtle, have characteristic shapes that may allow detection.  Sources with sharp pulse features, such as the radio and magnetospheric X-ray millisecond pulsars, provide other interesting targets. Resonant Compton scattering can also impart sharp light curve features relatively high in the magnetosphere \\citep{raj97}.  Doppler effects will, of course, be larger at these altitudes, but significant flux is ``lensed'' near the star surface in the light curves meaning that the higher-order effects (time-delay and Kerr) may also be significantly enhanced.  Modeling of such effects requires detailed treatment of the curved space, swept-back dipole field structure. We will describe such models in a future publication."
    },
    "0004/astro-ph0004282_arXiv.txt": {
        "abstract": "Using an approximate likelihood method adapted to band--power estimates, we analyze the ensemble of first generation cosmic microwave background anisotropy experiments to deduce constraints over a six--dimensional parameter space describing Inflation--generated adiabatic, scalar fluctuations.  The basic preferences of simple Inflation scenarios are consistent with the data set: flat geometries $(\\OmT \\equiv 1-\\Omk \\sim 1)$ and a scale--invariant primeval spectrum ($n\\sim 1$) are favored.  Models with significant negative curvature ($\\OmT < 0.7$) are eliminated, while constraints on postive curvature are less stringent.  Degeneracies among the parameters prevent independent determinations of the matter density $\\OmM$ and the cosmological constant $\\Lambda$, and the Hubble constant $\\Ho$ remains relatively unconstrained. We also find that the height of the first Doppler peak relative to the amplitude suggested by data at larger $l$ indicates a high baryon content ($\\Omb h^2$), almost independently of the other parameters.  Besides the overall qualitative advance expected of the next generation experiments, their improved dipole calibrations will be particularly useful for constraining the peak height. Our analysis includes a {\\em Goodness--of--Fit} statistic applicable to power estimates and which indicates that the maximum likelihood model provides an acceptable fit to the data set. ",
        "introduction": "The newest and perhaps most powerful tool in the cosmologist's toolbox are the temperature fluctuations in the cosmic microwave background (CMB).  Their very existence (e.g., Smoot et al. 1992) lends much credence to the general picture of gravitational instability forming galaxies and the observed large--scale structure.  The last two decades have witnessed the elaboration of this idea, with numerical studies of gravitational growth permitting quantitative comparison to actual survey data, and with the introduction of a physical mechanism -- namely, Inflation -- for the creation of the initial perturbations. Daring in scope, the resulting scenario would encompass the evolution of the Universe from possibly the Planck era to the present, explaining not only the origin of the required density perturbations, but also dispelling a handful of misgivings about the initial conditions of Big Bang model, such as the impressive homogeneity on large scales (Kolb \\& Turner 1990; Peebles 1993; Peacock 1999) Like any theory, it can never be proved; but it may be tested.  And like any good theory, it provides a physics, dependent on the validity of the theory, specific enough to be used as a tool to other ends:  it is now well appreciated that detailed study of the CMB fluctuations may be applied to determine the fundamental parameters of the Big Bang model itself (Bond et al. 1994; Knox 1995; Jungmann et al. 1996).  Our object with the present study is to examine what may be learned from present CMB data within the Inflationary context. It should be noted at the outset that there are other models contending to explain the origin of density perturbations, for example, defect models (Durrer 1999). Neither these nor any other alternative shall be our concern in the following, and it is important to emphasize that our results are therefore limited to the Inflationary context. Very specifically, we shall be concerned only with temperature fluctuations caused by adiabatic, coherent and passively evolving density (scalar) perturbations, dominated by cold dark matter (CDM) (Bond \\& Efstathiou 1984; Vittorio \\& Silk 1984).  Inflation can generate isocurvature modes, but we ignore these in the following, along with gravitational waves, the only non--scalar modes expected in Inflationary scenarios, and reionization.  All model predictions have been calculated using the CAMB Boltzmann code developed by Lewis, Challinor \\& Lasenby (1999) and built upon CMBFAST (Seljak \\& Zaldarriaga 1996; Zaldarriaga, Seljak \\& Bertschinger 1998) The reason for these restrictions is of course that it is impossible to explore the otherwise vast parameter space.  Even in this seemingly restrained setting, the physics of CMB temperature anisotropies is quite rich, but it may nevertheless be boiled down to two regimes: large scales where purely gravitational effects operate [Sachs--Wolfe (SW) effect; Sachs \\& Wolfe 1967], and small scales, within the horizon at recombination where causal physics plays its role.  In the latter regime, pressure of the coupled baryon--photon fluid resists the inward pull of gravity and therein establishes oscillating sound waves that will be observed in the CMB fluctuation power spectrum as a sequence of power peaks (so--called Doppler peaks; e.g., Hu \\& Sugiyama 1996).  These are damped towards the smallest scales by smoothing due to the finite thickness of the last scattering surface from which emanates the CMB. The well defined aspect of these power peaks owes to the coherent nature and passive evolution of Inflation generated perturbations; the loss of both of these in defect models has the effect of broadening or completely smearing out the peaks (see, e.g,. Durrer 1999 and references therein). The very existence of such a series of power peaks is thus a strong discriminator between models.   Besides the exact details of the physics of Inflation, the development of such perturbations depends, unsurprisingly, on the constituents of the primordial plasma in which they reside, as well as on the metric background in which they evolve.  This provides the link between the observable sky temperature fluctuations and the fundamental cosmological parameters. Actual data from the first generation of CMB experiments\\footnote{we use this term to refer to those experiments, with the exception of COBE, which were not specifically designed for map making; our detailed list is given in Table 2 and discussed below.} already indicates the existence of the first Doppler peak, a fact that leads to interesting and non--trivial conclusions, as noted by many authors and as detailed below.  In this {\\em paper} we present our constraints over a six--dimensional parameter space (see Table 1) resulting from a large compilation of first generation experiments (see Table 2).  Our conclusions are based not on a traditional $\\chi^2$ approach, but on methods developed and presented elsewhere which attempt to account for the non--Gaussian nature of power spectrum estimates as well as information at times lost in simple band--power estimates [Bartlett et al. 1999 (BDBL); Douspis et al. 2000 (DBBL); Douspis, Bartlett \\& Blanchard (DBB)]; we have tested these methods against complete likelihood analyses over subsets of the present CMB data.  Our parameter space here is spanned by the quantities listed in Table 1, namely, the Hubble constant $\\Ho$; the total energy density $\\OmT\\equiv \\OmM+\\lamo$, where $\\OmM$ is the matter density and $\\lamo\\equiv \\Lambda/3$; the vacuum energy density parameter $\\lamo$; the baryon density, $\\Omb$,  in terms of $\\etab\\equiv$ (baryon--to--photon number ratio)$\\times 10^{10}$ ($\\Omb h^2 = 0.00366\\etab$); and the primordial scalar spectral parameters $n$ and $Q$, i.e., the slope and normalization.  Our data set is listed in detail in Table 2.  While some recent analyses have covered a larger parameter space (e.g., Tegmark \\& Zaldarriaga 2000), they are based on traditional $\\chi$--squared methods. In a key work, Dodelson \\& Knox (1999) applied approximate likelihood methods similar to those employed here, although over a more restricted range of parameters.  The present work thus extends the approximate likelihood approach to a larger parameter space and includes an appropriate \\GOF\\ statistic.  On the other hand, we have not considered the effects of calibration uncertainties; the agreement between our results and those of previous authors supports the idea that these do not make a drastic difference to the final conclusions.  \\begin{figure}\\label{fig_powplot} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,totalheight=8.4cm, width=8.cm]{powplot.ps}} \\end{center} \\caption{The power plane: measured flat--band power estimates and model spectra as a function of multipole.  The solid line shows our best fit model (by approximate maximum likelihood) excluding the Python V points (see text): $(\\Ho,\\OmT,\\lamo,\\Omb h^2,n,Q)=(60\\; $km/s/Mpc,$ 1.0,0.3,0.032,1.06,16.0\\; \\mu$K). Our \\GOF\\ statistic indicates that this is an acceptable fit.  We also plot as the dotted line the best model with with a fixed $\\Omb h^2=0.011$, falling $2\\sigma$ below the prefered value: $(\\Ho,\\OmT,\\lamo,\\Omb h^2,n,Q)=(50\\; $km/s/Mpc,$ 1.1,0.8,0.011,1.00,19.0\\; \\mu$K). The dashed line is an open model with $\\OmT=0.2$ and $\\Ho=60$ km/s/Mpc, shown for illustration and clearly excluded by the data. } \\end{figure}  The {\\em paper} develops as follows: we begin in the next section with a presentation of our analysis methods; for the most part, this is a brief review of work presented in BDBL, DBBL and DBB. Our parameter constraints are then given in Section 3, followed by our conclusions in Section 4.  The basic and robust result at this stage must be considered as the conclusion that the spatial curvature is close to zero and the primordial spectral index close to its scale--invariant value of unity.  The results of our adapted statistical analysis are thus in agreement with many others (Lineweaver et al. 1997; Bartlett et al. 1998ab; Bond \\& Jaffe 1998; Efstathiou et al. 1998; Hancock et al. 1998; Lahav \\& Bridle 1998; Lineweaver \\& Barbosa 1998ab; Lineweaver 1998; Webster et al. 1998; Lasenby et al. 1998; Dodelson \\& Knox 1999; Melchiorri et al. 1999; Tegmark \\& Zaldarriaga 2000; Knox \\& Page 2000) and confirm the two fundamental predictions of the Inflationary model. ",
        "conclusions": "Our purpose in this {\\em paper} has been to analyze the ensemble of first generation CMB anisotropy experiments to see what conclusions may be drawn concerning certain fundamental cosmological parameters from the CMB data {\\em alone}. Our approach employs approximate likelihood methods that are adapted to power estimates, and which have been detailed elsewhere (BDBL, DBBL and DBB).  Our primary conclusions are that a flat geometry ($\\Omk\\sim 0$ or $\\OmT\\sim 1$) and a scale--invariant primeval spectrum ($n\\sim1$) are favored, while strongly hyperbolic models are ruled out with high significance -- in short, Inflation remains a good theory. Specifically, the best--fit model parameters are $(\\Ho,\\OmT,\\lamo,\\Omb h^2,n,Q)=(60\\; {\\rm km/s/Mpc},1.0,0.3,0.032,1.06,16.0\\; \\mu{\\rm K})$.  Our analysis includes a \\GOF\\ statistic that indicates that this model, and therefore the entire class (Inflation with adiabatic, scalar perturbations and without re--ionization) provides an acceptable description of the data. Many authors have recently explored these issues (Lineweaver et al. 1997; Bartlett et al. 1998ab; Bond \\& Jaffe 1998; Efstathiou et al. 1998; Hancock et al. 1998; Lahav \\& Bridle 1998; Lineweaver \\& Barbosa 1998ab; Lineweaver 1998; Webster et al. 1998; Lasenby et al. 1998; Dodelson \\& Knox 1999; Melchiorri et al. 1999; Tegmark \\& Zaldarriaga 2000; Knox \\& Page 2000), with various combinations of the present data set, and most would agree with these conclusions. The extensive spectral coverage (in multipole order $l$) over the first, and perhaps second Doppler peaks, and the low noise expected of the next generation instruments should qualitatively change the confidence in and precision of this kind of study.  Very little may be said about $\\Ho$ or about the relative contributions of matter, $\\OmM$, and vacuum, $\\lamo$, to the total energy density, the latter due to a well--known degeneracy when considering CMB data alone.  One must turn to other observational constraints, coming from, for example, cluster evolution (Oukbir \\& Blanchard 1992), cluster baryon fractions (White et al. 1993), SNIa Hubble diagrams (Reiss et al. 1998; Perlmutter et al. 1999), weak cosmic shear (Blandford et al. 1991; Mellier 1999), etc..., to eliminate such degeneracies (so--called ``cosmic complementarity'', Eisenstein, Hu \\& Tegmark 1999). Many of the authors listed above have included such constraints in their analysis; our own work along these lines is left to a forthcoming paper.  A surprising note is the very high baryon densities that we have found are prefered by the data set: $\\etab\\sim 8.9$ or $\\Omb h^2\\sim 0.032$.  This is even higher than the values indicated by the ``low'' D/H values found in several QSO absorption systems (Tytler et al. 2000), although within ``$2\\sigma$'' (a little more for high values if $\\Ho$; see Figure 5) the results are consistent.  It remains to be seen if this intriguing result bears the scrutiny of the next generation experiments. In particular, their ability to significantly reduce overall calibration uncertainty by using the CMB dipole will be crucial to such issues.  \\begin{table*}[h] \\begin{center} \\begin{tabular}{|c|c|c|c|} \\hline Experiment&$\\delta T \\pm \\sigma (\\mu K)$&$\\ell_{eff}$&reference\\\\ \\hline \\hline ARGO 1 &$32.9^{+4.8}_{-4.1}$&98&Ratra et al. 1999\\\\ \\hline ARGO 2 &$46.7^{+9.5}_{-12.1}$&98&Masi et al. 1996\\\\ \\hline BAM&$55.6^{+27.4}_{-9.8}$&74&Tucker et al. 1997\\\\ \\hline BOOM NA&$29.1^{+13.0}_{-11.0}$&58&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$48.8^{+9.0}_{-9.0}$&102&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$67.1^{+10.0}_{-9.0}$&153&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$71.9^{+10.0}_{-10.0}$&204&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$60.8^{+11.0}_{-12.0}$&255&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$55.4^{+14.0}_{-15.0}$&305&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$32.1^{+13.0}_{-22.0}$&403&Mauskopf et al. 1999\\\\ \\hline BOOM NA&$15.8^{+114.}_{-15.8}$&729&Mauskopf et al. 1999\\\\ \\hline CAT 1a&$51.8^{+13.6}_{-13.6}$&410&Scott et al. 1996\\\\ \\hline CAT 1b&$49.1^{+19.1}_{-13.6}$&590&Scott et al. 1996\\\\ \\hline CAT 2a&$57.3^{+11.}_{-13.6}$&422&Baker et al. 1997\\\\ \\hline CAT 2b&$0.0^{+55.0}_{-0.0}$&615&Baker et al. 1997\\\\ \\hline COBE 1&$8.5^{+16.}_{-8.5}$&2&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 2&$28.0^{+7.4}_{-10.4}$&3&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 3&$34.0^{+5.9}_{-7.2}$&4&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 4&$25.1^{+5.2}_{-6.6}$&6&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 5 &$29.4^{+3.6}_{-4.1}$&8&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 6 &$27.7^{+3.9 }_{-4.5}$&11&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 7 &$26.1^{+4.4}_{-5.3}$&14&Tegmark \\& Hamilton 1997\\\\ \\hline COBE 8 &$33.0^{+4.6}_{-5.4}$&19&Tegmark \\& Hamilton 1997\\\\ \\hline FIRS &$29.4 ^{+7.8}_{-7.7}$&10&Ganga et al. 1994\\\\ \\hline IAB   &$124.0^{+60.}_{-60.}$&203&Piccirillo et al. 1993\\\\ \\hline IAC 1 &$113.0^{+49.}_{-44.}$&33&Femenia 1998\\\\ \\hline IAC 2 &$55.0^{+27.}_{-28.}$&53&Femenia 1998\\\\ \\hline IAC 33&$43.0^{+12.5}_{-11.5}$&105 &Dicker et al. 1999\\\\ \\hline MAX GUM&$54.5^{+16.4}_{-10.9}$&138&Tanaka et al. 1996\\\\ \\hline MAX HR&$27.9^{+11.5}_{-4.7}$&130& Our analysis \\\\ \\hline MAX ID&$52.7^{+32.1}_{-10.4}$&120& Our analysis \\\\ \\hline MAX PH&$72.9^{+30.6}_{-10.7}$&131& Our analysis \\\\ \\hline MAX SH&$82.0^{+53.0}_{-15.0}$&121& Our analysis \\\\ \\hline MSAM 1&$35.0^{+15.0}_{-11.0}$&84&Wilson et al. 2000\\\\ \\hline MSAM 1&$49.0^{+10.0}_{-8.0}$&201&Wison et al. 2000\\\\ \\hline MSAM 1&$47.0^{+7.0}_{-6.0}$&407&Wilson et al. 2000\\\\ \\hline OVRO&$56.1^{+8.5}_{-6.6}$&537&Leitch et al. 2000\\\\ \\hline Pyth A&$60.0^{+9.0}_{-5.0}$&87&Platt et al. 1997\\\\ \\hline Pyth B&$66.0^{+11.}_{-9.0}$&170&Platt et al. 1997\\\\ \\hline Python V&$22.0^{+4.0}_{-5.0}$&44&Coble 1999 thesis\\\\ \\hline Python V&$24.0^{+6.0}_{-7.0}$&75&Coble 1999 thesis\\\\ \\hline Python V&$34.0^{+7.0}_{-9.0}$&106&Coble 1999 thesis\\\\ \\hline Python V&$50.0^{+9.0}_{-12.0}$&137&Coble 1999 thesis\\\\ \\hline Python V&$61.0^{+13.0}_{-17.0}$&168&Coble 1999 thesis\\\\ \\hline Python V&$77.0^{+20.0}_{-28.0}$&199&Coble 1999 thesis\\\\ \\hline Python V&$69.0^{+71.0}_{-69.0}$&261&Coble 1999 thesis\\\\ \\hline Python V&$0.0^{+87.0}_{-0.0}$&230&Coble 1999 thesis\\\\ \\hline QMap,K1&$47.0^{+6.0}_{-7.0}$&80&de Oliveira-Costa 1998\\\\ \\hline QMap,K2&$59.0^{+6.0}_{-7.0}$&126&de Oliveira-Costa 1998\\\\ \\hline Sask 1&$51.5^{+8.4}_{-5.3}$&86&Netterfield et al. 1997\\\\ \\hline Sask 2&$72.5^{+7.4}_{-6.3}$&166&Netterfield et al. 1997\\\\ \\hline Sask 3&$89.3^{+10.5}_{-8.4}$&236&Netterfield et al. 1997\\\\ \\hline Sask 4&$90.3^{+12.6}_{-10.5}$&285&Netterfield et al. 1997\\\\ \\hline Sask 5&$72.5^{+20.}_{-29.4}$&348&Netterfield et al. 1997\\\\ \\hline SP91   &$28.0^{+9.5}_{-6.7}$&58&Gundersen et al. 1995\\\\ \\hline SP94   &$36.3^{+13.6}_{-6.1}$&62&Gundersen et al. 1995\\\\ \\hline SuZIE&$0.0^{+44.0}_{-0.0}$&2340&Ganga et al. 1997\\\\ \\hline Tenerife &$34.0^{+15.5}_{-9.3}$&18&Guti\\'errez et al. 2000\\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{table*}[h] \\begin{center} \\begin{tabular}{|c|c|c|c|} \\hline Experiment&$\\delta T \\pm \\sigma (\\mu K)$&$\\ell_{eff}$&reference\\\\ \\hline \\hline \\hline Toco 97&$40.0^{+10.0}_{-9.0}$&63 &Torbet et al. 1999\\\\ \\hline Toco 97&$45.0^{+7.0}_{-6.0}$&85 &Torbet et al. 1999\\\\ \\hline Toco 97&$70.0^{+6.0}_{-6.0}$&119 &Torbet et al. 1999\\\\ \\hline Toco 97&$89.0^{+7.0}_{-7.0}$&155 &Torbet et al. 1999\\\\ \\hline Toco 97&$85.0^{+8.0}_{-8.0}$&194&Torbet et al. 1999\\\\ \\hline Toco 98&$55.0^{+18.0}_{-17.0}$&128 &Miller et al. 1999\\\\ \\hline Toco 98&$82.0^{+11.0}_{-11.0}$&152 &Miller et al. 1999\\\\ \\hline Toco 98&$83.0^{+7.0}_{-8.0}$&226 &Miller et al. 1999\\\\ \\hline Toco 98&$70.0^{+10.0}_{-11.0}$&306 &Miller et al. 1999\\\\ \\hline Toco 98&$0.0^{+67.0}_{-0.0}$&409&Miller et al. 1999\\\\ \\hline Viper&$61.6^{+31.1}_{-21.3}$&108 &Peterson et al. 2000\\\\ \\hline Viper&$77.6^{+26.8}_{-19.1}$&173 &Peterson et al. 2000\\\\ \\hline Viper&$66.0^{+24.4}_{-17.2}$&237 &Peterson et al. 2000\\\\ \\hline Viper&$80.4^{+18.0}_{-14.2}$&263 &Peterson et al. 2000\\\\ \\hline Viper&$30.6^{+13.6}_{-13.2}$&422&Peterson et al. 2000\\\\ \\hline Viper&$65.8^{+25.7}_{-24.3}$&589&Peterson et al. 2000\\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{First generation CMB data set listed in alphabetical order. This information may also be found on our \\Web\\ site ({\\tt http://webast.ast.obs-mip.fr/cosmo/CMB/}). Our final analysis excludes the Python V results (see text).} \\end{table*}"
    },
    "0004/astro-ph0004069_arXiv.txt": {
        "abstract": "We study gravitational lens time delays for a general family of lensing potentials, which includes the popular singular isothermal elliptical potential and singular isothermal elliptical density distribution but allows general angular structure.  Using a novel approach, we show that the time delay can be cast in a very simple form, depending only on the observed image positions.  Including an external shear changes the time delay proportional to the shear strength, and varying the radial profile of the potential changes the time delay approximately linearly.  These analytic results can be used to obtain simple estimates of the time delay and the Hubble constant in observed gravitational lenses.  The naive estimates for four of five time delay lenses show surprising agreement with each other and with local measurements of $\\H0$; the complicated Q~0957+561 system is the only outlier.  The agreement suggests that it is reasonable to use simple isothermal lens models to infer $\\H0$, although it is still important to check this conclusion by examining detailed models and by measuring more lensing time delays. ",
        "introduction": "The time delay in gravitational lenses is given by (e.g., Schneider, Ehlers \\& Falco 1992) \\beq \\label{eq:time} \\Delta t (x,y) = \\frac{D}{2c} (1+\\zd) \\left[ (x-\\xi)^2 + (y-\\eta)^2 - 2 \\phi(x,y) \\right], ~~D \\equiv {\\Dd \\Ds \\over \\Dds}, \\eeq where $\\zd$ is the redshift of the lens, $(\\xi, \\eta)$ is the angular source position, $(x, y)$ is the angular image position, $\\Dd$ and $\\Ds$ are angular diameter distances from the observer to the lens and source, respectively, and $\\Dds$ is the angular diameter distance from the lens to the source. The dimensionless potential $\\phi$ satisfies the two-dimensional Poisson equation \\beq \\nabla^2\\phi(x,y) = 2 \\kappa, \\quad \\kappa = {\\Sigma(x, y) \\over \\Sigma_{\\rm cr}}\\ , \\label{eq:sigcr} \\eeq where $\\Sigma$ is the projected surface mass distribution of the lens and $\\Sigma_{\\rm cr} = {c^2 \\Ds / (4 \\pi G \\Dd \\Dds)}$ is the critical surface density for lensing.  To proceed further we adopt a specific potential or density form. Individual lenses and lens statistics are usually consistent with isothermal models (e.g., Maoz \\& Rix 1993; Kochanek 1995, 1996; Grogin \\& Narayan 1996), so it is common to adopt either the isothermal elliptical potential, \\beq \\label{eq:sep} \\quad \\phi = \\frac{a_0}{2} (x^2+y^2/q^2)^{1/2}, \\eeq or the elliptical density distribution, \\beq \\label{eq:sed} \\nabla^2 \\phi = 2 \\kappa, \\quad \\kappa = \\frac{a_0}{2q} (x^2+y^2/q^2)^{-1/2}, \\eeq where $a_0$ specifies the overall angular size of the lens. For the SIED, the lens potential is given by, e.g., Kassiola \\& Kovner (1993), Kormann, Schneider \\& Bartelmann (1994), and Keeton \\& Kochanek (1998).  Note that we have chosen a coordinate system that is centered on the lens galaxy and aligned the $x$-axis along the lens galaxy's major axis.  For simplicity, we have also assumed that the isothermal potential and density distributions are singular at the origin, although we return to this issue briefly in \\S 5.  In this paper, we study a more general family of lens models. The potential is assumed to obey the relation \\beq \\label{eq:phi} \\phi = x \\phix + y \\phiy, \\eeq where $\\phix$ and $\\phiy$ are the first derivatives of the potential, which are just the components of the deflection angle in the $x$ and $y$ directions.  In polar coordinates $(r, \\theta)$, eq.~(\\ref{eq:phi}) can be written in the simple form \\beq \\label{eq:phiPolar} \\phi = r {\\partial \\phi \\over \\partial r}, \\eeq whose general solution is \\beq \\label{eq:phiSolution} \\phi = r {\\cal F}(\\theta), \\eeq where ${\\cal F}(\\theta)$ is an arbitrary function of $\\theta$. This potential corresponds to a density of the form \\beq \\label{eq:isokappa} \\kappa = \\frac{1}{2r} {\\cal G}(\\theta), \\quad {\\cal G}(\\theta) = {\\cal F}(\\theta) + {\\partial^2{\\cal F}\\over\\partial\\theta^2} . \\eeq Thus eq.~(\\ref{eq:phi}) describes a family of scale-free models that includes both the SIEP and SIED models but allows for general angular structure; in other words, these are generalized isothermal models.  We now show that this family is extremely useful in deriving the time delays between images.  The lens equation is given by \\beq \\label{eq:lens} \\xi  = x - \\phix, \\quad \\eta = y - \\phiy. \\eeq Substituting eqs.~(\\ref{eq:phi}) and (\\ref{eq:lens}) into the time delay expression in eq.~(\\ref{eq:time}), we obtain \\beq \\Delta t (x,y) = \\frac{D}{2c} (1+z_d) \\left[ (x-\\xi)^2 + (y-\\eta)^2 \\\\ - 2 x (x-\\xi) - 2 y (y-\\eta) \\right], \\eeq which can be further simplified into \\beq \\Delta t (x,y) = \\frac{D}{2c} (1+z_d)\\, (\\xi^2 + \\eta^2 - x^2 -y^2). \\eeq Since only the relative time delay is observable, we compute the time delay between two images $i$ and $j$, which is given by \\beq \\label{sedDelay} \\Delta t_{i,j} = \\frac{D}{2c} (1+z_d) (r_j^2 - r_i^2), \\eeq where $r_i = (x_i^2 + y_i^2)^{1/2}$ is simply the distance of image $i$ from the center of the galaxy.  This surprisingly simple expression is valid for all lens potentials satisfying eq.~(\\ref{eq:phi}), including both the SIEP and SIED models as well as more general angular structures.\\footnote{After the completion of our paper, we learned that eq.~(\\ref{sedDelay}) has been derived for particular cases by Koopmans, de Bruyn \\& Jackson (1998) and Zhao \\& Pronk (2000) using more complicated methods.}  Since $r_i^2$ and $r_j^2$ are rotationally invariant, eq.~(\\ref{sedDelay}) is valid even after an arbitrary rotation. In other words, provided we know the center of the lens galaxy, the predicted time delay can be computed using eq.~(\\ref{sedDelay}) irrespective of the lens orientation and without any need to search for the best fit model parameters.  Since the predicted time delay scales as $\\propto \\H0^{-1}$, it can be compared with a measured time delay to determine $\\H0$; we demonstrate this method in \\S 4.  For completeness we note that the following relations hold for a potential which satisfies eq.~(\\ref{eq:phi}): \\beq x \\phi_{xx} + y \\phi_{xy} = 0,~ x \\phi_{xy} + y \\phi_{yy} = 0,~ \\phi_{xx} \\phi_{yy} - \\phi_{xy}^2 = 0. \\eeq The last expression can be used to simplify the magnification of the images: \\beq \\mu^{-1} = (1-\\phi_{xx}) (1-\\phi_{yy}) - \\phi_{xy}^2 = 1- \\phi_{xx} -\\phi_{yy} = 1 - 2 \\kappa, \\eeq revealing a simple relation between the magnification and the surface density in this class of potentials. ",
        "conclusions": "We have shown that for generalized isothermal models (with a potential of the form $\\phi = r {\\cal F}(\\theta)$), the time delay between images can be expressed in a common and surprisingly simple form involving only the observed image positions (eq.~\\ref{sedDelay}).  An external shear changes the time delay by an amount proportional to the shear strength, and affects 4-image lenses more than 2-image lenses.  Changing the radial profile of the lens galaxy changes the time delay approximately by a multiplicative factor.  In this paper, we have only considered singular potential and density distributions.  This is justified because so far no convincing faint central image has been observed, and this sets stringent limits on the central core radius (Wallington \\& Narayan 1993; Kochanek 1996).  Numerically we find that such small core radius changes the time delays negligibly ($\\lesssim 1\\%$).  Our simple expression for the time delay can be combined with directly observable quantities to give an estimate for the Hubble constant $\\H0$ without the need for any modeling.  The analytic $\\H0$ estimates for four of five time delay lenses agree surprisingly well with each other and with distance ladder measurements of $\\H0$.  The outlier, Q~0957+561, is known to be a complicated system where the lensing potential is inconsistent with our assumptions of a pure isothermal model.  Thus it appears that simple isothermal lens models are often reasonable for obtaining lensing estimates of $\\H0$ (see also Koopmans \\& Fassnacht 1999).  Still, it is crucial to test this conclusion more rigorously, and there are three ways to do so.  First, the astrometry for two of the time delay lenses (B~0218+357 and PKS~1830$-$211) must be significantly improved to reduce the uncertainties and verify the $\\H0$ concordance.  Second, more time delays should be measured (and good astrometric precision obtained for those lenses) to test whether all of the $\\H0$ estimates derived from isothermal lens models are mutually consistent. Increasing the sample size will also reduce the random uncertainties due to large-scale structure.  Third, more individual lenses should undergo detailed modeling to check the $\\H0$ estimates obtained from isothermal models and understand the systematic uncertainties.  These detailed models are most instructive when they are constrained by observations of the host galaxies of the lenses quasars, because requiring models to reproduce the distortion of the host galaxy significantly improves the ability to find a robust and well-constrained model (see Keeton et al.\\ 2000ab; Blandford et al.\\ 2000).  Finally, we recall that the dependence of the time delay on the density profile of the lens galaxy affects the $\\H0$ estimates (also see Keeton \\& Kochanek 1997; Koopmans \\& Fassnacht 1999).  If lens galaxies tend to have profiles steeper than isothermal ($\\phi(r) \\propto r^\\beta$ with $\\beta<1$), then by insisting on using isothermal models we tend to underestimate $\\H0$ (or vice versa).  There are two possible situations: \\begin{itemize} \\item {\\it Galaxies have a common density profile (little or no scatter in $\\beta$).\\/}  If so, the $\\H0$ estimates obtained by applying isothermal models to numerous lenses would show little scatter.  If these estimates were to agree with distance ladder measurements of $\\H0$, it would suggest that galaxies generally have isothermal profiles. \\item {\\it Galaxies have a scatter in $\\beta$.\\/}  If so, isothermal lens $\\H0$ estimates would show a significant scatter.  It is possible that lenses are biased tracers of $\\beta$: on the one hand, lenses with larger values of $\\beta$ tend to produce higher magnifications and hence have a larger magnification bias, while on the other hand they have smaller cross sections for lensing (e.g., Kochanek 1991; Wambsganss \\& Paczy\\'nski 1994; Witt et al.\\ 1995; Witt \\& Mao 2000).  However, when lenses are normalized to reproduce the image separations in observed lenses, these two effects essentially cancel (see Kochanek 1996).  Thus it appears that lenses are not strongly biased with regard to $\\beta$, so the scatter in lensing $\\H0$ estimates would be similar to the scatter in $\\beta$. \\end{itemize} The present data suggest the first situation, namely that lenses are generally consistent with isothermal galaxies. By collecting more and better data, however, we will be able to test this conclusion more rigorously and obtain a robust and independent measurement of $\\H0$."
    },
    "0004/astro-ph0004296_arXiv.txt": {
        "abstract": "The running-mass inflation model, which has strong motivation from particle physics, predicts density perturbations whose spectral index is strongly scale-dependent. For a large part of parameter space the spectrum rises sharply to short scales. In this paper we compute the production of primordial black holes, using both analytic and numerical calculation of the density perturbation spectra. Observational constraints from black hole production are shown to exclude a large region of otherwise permissible parameter space. ",
        "introduction": "Particle physics models of inflation based on supergravity theories are plagued by the so-called $\\eta$-problem \\cite{LR}, which states that the mass-squared of any scalar field, including the putative inflaton field, is typically of order $H^2$ ($H$ being the Hubble parameter) which ruins slow-roll inflation.  An elegant proposal to circumvent this is the running-mass model of inflation, introduced by Stewart \\cite{S1,S2}, where the flatness of the potential arises because of the quantum corrections, which serve to flatten the potential over a significant region where inflation can then take place.  Because the flatness is brought about by a cancellation of the intrinsic curvature of the potential against the quantum corrections, only a limited portion of the potential can support slow-roll inflation, as is necessary to generate the approximately flat power spectrum seen by the COBE satellite. Well away from COBE scales, one typically expects to see dramatic deviations from near scale-invariance as the slow-roll regime breaks down. Copeland et al.~\\cite{CGL} examined the possibility that this breakdown of scale-invariance might be detectable through scale-dependence of the spectral index of primordial perturbations, and more recently the model has been extensively explored by Covi, Lyth and collaborators \\cite{CLR,CL,LC} in a series of papers investigating its  viability both from a theoretical standpoint and in confrontation against large-scale structure data.  In this paper we examine constraints arising from the more radical departures from scale-invariance which may take place towards the end of inflation. In much of parameter space, the spectrum rises sharply on short scales, which can give rise to production of primordial black holes (PBHs). These are strongly constrained by observation and, as we will see, a significant region of otherwise viable parameter space is excluded. ",
        "conclusions": "We have investigated primordial black hole constraints within the context of the failure of slow-roll inflation, focusing on the well-motivated running-mass model of inflation which features a strong scale-dependence of the spectral index \\cite{LC}. Although applying to the amplitude of perturbations on very short scales, to a good approximation the constraint can be represented as a constraint on $n_{{\\rm COBE}}$, the spectral index on the largest observable scales.  The constraint depends strongly on the number of $e$-foldings $N_{{\\rm COBE}}$ between the production of those perturbations and the end of inflation, and, as with models with a constant spectral index, the constraint becomes weaker as $N_{{\\rm COBE}}$ is reduced.  We have shown that a significant region of the parameter space of the model, viable under other constraints, is excluded by excess production of black holes. This demonstrates the importance of evaluating the density perturbation spectrum not just across astrophysical scales but also right to the end of inflation. In models where the slow-roll approximation holds accurately only over a limited range of scales, such as the running-mass model, there will be strong deviations from scale-invariance towards short scales. In models where the deviation takes the form of a strongly blue spectrum, excessive black hole production is always likely to be a danger."
    },
    "0004/astro-ph0004405_arXiv.txt": {
        "abstract": "We present deep optical $B$,$V\\!$,$R$ images of a sample of 10 interacting systems which were selected for their resemblance to disturbed galaxies at high redshift.  Photometry is performed on knots in the tidal features of the galaxies. We calculate a grid of evolutionary synthesis models with two metallicities and various burst strengths for systems consisting of some fraction of the stellar population of a progenitor spiral plus starburst.  By comparison with two-color diagrams we interpret the photometric data, select from a total of about 100 condensations 36 star-forming objects that are located in the tidal features and predict their further evolution. Being more luminous by 4 mag than normal \\ion{H}{ii} regions we argue that these objects could be tidal dwarf galaxies or their progenitors, although they differ in number and mean luminosity from the already known tidal dwarf galaxies typically located at the end of tidal tails in nearby giant interacting systems. From comparison with our models we note that all objects show young burst ages. The young stellar component formed in these tidal dwarf candidates contributes up to 18\\% to the total stellar mass at the end of the starburst and dominates the optical luminosity. This may result in fading by up to 2.5 mag in $B$ during the next 200 Myrs after the burst. ",
        "introduction": "Most investigations of interacting and merging galaxies have essentially focussed on the phenomena occuring in their inner regions. However, for a few spectacular cases like the Antennae \\citep{MLM91}, Arp 105 \\citep{DM94,DBW97} and NGC 7252 \\citep{HM95} the enormous tidal tails were investigated in some detail. Massive condensations of stars and \\ion{H}{i} were found in these tails, the so called Tidal Dwarf Galaxies (TDGs).  TDGs are characterized by a luminosity comparable with that of typical dwarf galaxies, but which span a range in oxygen abundance of $8.3 \\lesssim 12+\\log(O/H) \\lesssim 8.6$ equivalent to metallicities of $4\\cdot10^{-3} \\lesssim Z_\\mathrm{TDG} \\lesssim 8\\cdot10^{-3}$ with a mean of $Z \\approx 7\\cdot10^{-3} \\approx Z_{\\sun} / 2.6$.\\footnote{For metallicity we use the notation of $Z$ (as the proportion of heavy elements given in the Geneva stellar tracks which our models are based on) throughout this paper and give a comparison with the solar value $Z_{\\sun} = 0.018$ where appropriate.} Metallicities of TDGs are therefore higher than those of other dwarf galaxies with comparable luminosity \\citep{DM97Mess,DM98}.  TDGs have blue colors as a result of an active starburst. Most known TDGs have high \\ion{H}{i} masses from which total masses of the order of 10$^9$ M$_{\\sun}$ have been estimated.  They are gravitationally bound and in a few interacting systems evidence for their kinematical independence has been found \\citep[e.g.][]{DBW97,DM98}.  These systems together with the strong increase of the galaxy merger rate with redshift gave rise to speculations about a possible contribution of TDGs to the faint blue galaxy excess at high redshift \\citep{FvAM98}. \\citet{BH92} observed the formation of TDGs in their numerical simulations. Their 3D N-body/SPH-code created several massive bound condensations along the tidal tails. At first, their mass is dominated by the stellar component.  Later on, gas from the tail might fall into these condensations.  \\citet{EKT93} proposed a different scenario.  They have used a 2D N-body code including dissipation for the gas component in which gaseous condensations form first.  Observations in most cases show both components, neutral gas and stars, along the tails as well as in their condensations. In an optically indentified condensation, determination of the relative fraction of young stars -- born in situ from the collapse of tidally extracted \\ion{H}{i} -- to old stars -- pulled out from the progenitor disk(s) -- might constrain the above mentioned numerical models. Because the systems studied here are too distant to be resolved into stars, estimating these relative stellar fractions can only be done from multi-band integrated photometric measurements in comparison with an evolutionary synthesis model.  Photometric work on the tidal features of several classical interacting systems has been carried out by \\citet{SWS90}.  Studies by \\citet{AM97} and \\citet{BAM98} concentrated on ring galaxies.  \\citet{DMT98} have systematically catalogued all faint non-stellar objects around disturbed galaxies. Part of them might be TDG candidates.  These previous investigations have focused on disk-disk interacting systems with a well-understood morphology, e.g. well defined tails. However locally the majority of disturbed galaxy systems, for instance in the Catalogue of Southern Peculiar Galaxies \\citep{AM87}, do not have long tails.  In particular disk-spheroid collisions and 3-body systems leading to chaotic morphologies are quite common.  HST deep surveys with the WFPC2 show that the percentage of disturbed (i.e.~interacting, merging or peculiar) galaxies has been much higher in the early universe than today \\citep{vdB96}.  NICMOS observations in the near-infrared have shown that these disturbances are not only due to irregularities in the distribution of star forming regions, but reflect real morphological features of interactions \\citep{CBD98}. These galaxies at high redshift often do not show long tidal tails either because they are absent or because they have too low a surface brightness. Examining a local sample which resembles these high redshift interacting galaxies is therefore important to learn about their evolution.  In this paper we extend previous investigations of a few individual systems to deep optical photometry of a sample of peculiar galaxies. The objects from our sample mostly do not possess prominent tidal tails, but definitely show signs of merging, as e.g.~multiple nuclei or disturbed morphology and nearby companions.  The sample is extracted from the Catalogue of Southern Peculiar Galaxies \\citep{AM87}.  It includes disk-disk, disk-elliptical, dwarf-dwarf and 3-body systems observed at different stages of the interaction.  Spectrophotometric evolutionary synthesis models have been used for various applications of normal galaxy evolution. More recently starburst events were modeled as well. These models assume that the standard star-formation rate (SFR) of an underlying galaxy is increased to simulate a starburst \\citep{FvAG94a,LH95}. This technique has been applied to reproduce the star formation history of various classes of objects from merging galaxies \\citep{FvAG94b} to blue compact galaxies \\citep{KFL95}.  We take an approach similar to the latter and compute a grid of models specifically aimed to describe the properties found in previous investigations of TDGs.  We do not use solar-abundance tracks in our models but instead two more realistic metallicities in the range expected for TDGs ($Z = 1\\cdot10^{-3}, 8\\cdot10^{-3}$). As a primary goal, these models are used to select candidate TDGs or TDG progenitors.  Based on their colors only, background objects can be excluded not only with statistical methods but also one by one if they disagree with any kind of TDG model.  Furthermore, this technique allows to derive the present ratios of old to young stellar mass in these TDG candidates, to predict its evolution and hence to constrain the above-mentioned formation scenarios of TDGs.  Obviously spectrophotometric data is required to confirm the photometric results.  The structure of the paper is as follows. First we present our data reduction and analysis techniques (Sect.~\\ref{ObsData-Sec}). We describe in Sect.~\\ref{Model-Sec} the evolutionary synthesis models and the parameters used to interpret the photometric data.  In Sect.~\\ref{Results-Sec} we present our sample, compare with the model grid and comment on individual features of the interacting systems. In Sect.~\\ref{Disc-Sec} we discuss more general results and in particular compare the tidal objects identified in our sample with those produced in numerical models and with the few typical TDGs studied so far. ",
        "conclusions": "\\label{Conc-Sec} We have presented imaging data of 10 southern interacting systems from the Arp \\& Madore Catalog selected for their resemblance to disturbed distant galaxies. Our high quality images reveal a heterogeneous set of morphologies from loosely interacting dwarf galaxies to disk-disk mergers.  Through detailed photometry based on polygonal apertures in the three optical bands $B$,$V\\!$,$R$ we were able to measure the colors of about one hundred condensations in or near the tidal features of the interacting galaxies.  We have computed a grid of evolutionary synthesis models simulating starbursts of various strengths in an underlying stellar population of mixed ages typical of spiral galaxies. Such models reproduce well the situation of TDGs and their progenitors that are composed of an old population pulled out from the parent galaxy and young stars formed in situ.  Our models take into account the emission lines and the continuum from the photoionized gas with appropriate metallicities. Comparing our photometric data with these models, we selected 36 blue condensations as possible TDGs.  Follow-up spectroscopy is required in some cases to assess their physical association with the interacting system.  We discussed the overall properties of the 36 TDG candidates, noting that they have luminosities brighter than \\ion{H}{ii} regions in spiral galaxies and more typical for that of dwarf galaxies.  Their color is much bluer than those of their parent galaxies and mostly bluer than the surrounding material in the tidal features. The contribution of young stars to their mass is at least 1\\% and may reach 20\\% or more for the most active condensations. In any case their $B$-band luminosity is dominated by this young population. The models also indicate that after the starburst when the young stars cease to be the dominating source the condensations may fade by up to 2.5 mag in $B$ within 200 Myrs after the burst. Although the TDG candidates seem to have young burst ages, we did not find any case of condensations solely made of old stars from the parent galaxies although the formation of such objects is also predicted in numerical simulations.  Finally the TDG candidates studied here have properties different from those of the more massive TDGs typically found at the tip of giant tidal tails of nearby interacting systems.  We discuss whether our TDG candidates could be progenitors of TDGs or a class of less luminous objects."
    },
    "0004/astro-ph0004319_arXiv.txt": {
        "abstract": "The Hubble space telescope observations of the northern Hubble deep field, and more recently its counterpart in the south, provide detections and photometry for stars and field galaxies to the faintest levels currently achievable, reaching magnitudes $V \\sim 30$. Since 1995, the northern Hubble deep field has been the focus of deep surveys at nearly all wavelengths. These observations have revealed many properties of high redshift galaxies, and have contributed to important data on the stellar mass function in the Galactic halo. ",
        "introduction": "Deep surveys have a long history in optical astronomy. Indeed, nearly every telescope ever built has at some point in its lifetime been pushed to its practical limit for source detection. An important historical motivation for such surveys has been the desire to test cosmological models via the classical number--magnitude and angular-diameter--magnitude relations \\citep{Sandage61,Sandage88}. However interpretation has always been hampered by the difficulty of disentangling the effects of galaxy selection and evolution from the effects of cosmic geometry. The modern era of deep surveys in optical astronomy began with the advent of CCD detectors, which allowed 4-m telescopes to reach to nearly the confusion limit allowed by the point spread function \\citep{Tyson88}. These images revealed a high surface density of faint, blue galaxies, and the combination of their number--magnitude relation, number--redshift relation, and angular correlation function was difficult to reconcile with the critical-density universe $(\\Omega_M = \\Omega_{\\rm tot} = 1)$ that seemed at the time to be a robust expectation from inflationary models \\citep{Kron80,FYTY90,GRV90,EBTKG91}.   Prior to launch, simulations based on fairly conservative assumptions suggested that the Hubble space telescope (HST) would not provide an overwhelming advantage over ground-based telescopes for studies of distant galaxies \\citep{BGS90}: \\begin{quotation} {\\it Our results show that the most sensitive exposures achieved so far from the ground reveal more galaxies per unit area than will be seen by planned HST observations unless galaxy sizes decrease with the maximum rate consistent with ground-based observations. In this, the most favorable case for HST, the space exposures will show almost as many galaxy images as have been observed so far in the most sensitive ground-based data.} \\end{quotation} It was apparent from the images returned immediately after the HST refurbishment in 1993 that distant galaxies had substantially higher surface brightnesses than predicted by such simulations \\citep{DOSL94,Dickinson95Gal}. Despite the pessimistic predictions, extensive field galaxy surveys were always a top priority for HST, and large amounts of observing time were awarded to deep, pointed surveys and to the medium-deep survey \\citep{MDSIAU}. These observations provided valuable information on galaxy scale lengths and morphologies, on the evolution of clustering and the galaxy merger rate, and on the presence of unusual types of galaxies \\citep{MWSPGRCIEGGS94,SLCHLT95,CHS95Nat,MDSIAU}. In this context, the idea of using the HST to do an ambitious deep field survey began to look quite attractive, and was eventually adopted as one of the primary uses of the HST director's discretionary time.  \\subsection{This Review}  The Hubble deep fields north and south (HDF-N and HDF-S, respectively) are by now undoubtedly among the most ``data rich'' portions of the celestial sphere.  The task of reviewing progress is complicated by the fact that much of the supporting data are still being gathered, and by the wide range of uses that have been found for the HDF.  It is not practical to provide a comprehensive review of all HDF-related research in a single review article. Instead we have chosen to divide the review into two parts. The first part discusses the data, both from HST and from other facilities, and summarizes measurements and phenomenology of the sources found in the field. The second part focuses specifically on distant galaxies, and attempts to provide a critical view of what the HDF has, and has not, taught us about galaxy evolution. \\citet{Ferguson98} has reviewed the HDF with a different focus, and the series of papers in the 1996 Herstmonceux conference and the 1997 STScI May Symposium provide a broad summary of the overall field \\citep{Herst96,STScIMay97}. Throughout this review, unless explicitly stated otherwise, we adopt cosmological parameters $\\Omega_M, \\Omega_\\Lambda, \\Omega_{\\rm tot} = 0.3, 0.7, 1.0$ and $H_0 = 65 \\kms {\\rm Mpc^{-1}}$. Where catalog numbers of galaxies are mentioned, they refer to those in \\citet{WBDDFFGGHKLLMPPAH96}. ",
        "conclusions": "The HDFs represent an important portion of the frontier in studies of the distant universe. Although we can point to few problems that were solved by the HDF alone, the HDF has contributed in a wide variety of ways to our current understanding of distant galaxies and to shaping the debate over issues such as the origin of elliptical galaxies and the importance of obscured star formation. The scientific impact of the HDF can be attributed in part to the wide and nearly immediate access to the data. The fact that many groups and observatories followed this precedent (both for HDF data and subsequently for data from other surveys) illustrates that the deeper understanding of the universe will come not from any one set of observations but from sharing and comparing different sets of observations. Not all kinds of surveys are amenable to this kind of shared effort, but the precedent set by the HDF in the social aspects of carrying out astronomical research may ultimately rival its significance in other areas."
    },
    "0004/astro-ph0004155_arXiv.txt": {
        "abstract": "Moderate resolution data for 40 quasi-stellar objects (QSOs) at z $\\approx$ 2 were combined with spectra of comparable resolution of 59 QSOs with redshifts greater than $1.7$ found in the literature to form a large, homogeneous sample of moderate resolution ($\\sim$1 $\\AA$) QSO spectra.  These spectra were presented and the statistics of the Lyman $\\alpha$ forest were discussed in Paper I. In this analysis, we demonstrate that a proximity effect is present in the data, ie. there exists a significant (5.5$\\sigma$) deficit of lines at $z_{abs} \\approx z_{em}$.  Within 1.5 $h^{-1}$ Mpc of the QSO emission redshift, the significance does depend on QSO luminosity, in accordance with the theory that this effect is caused by enhanced ionization of hydrogen in the vicinity of the QSO from UV photons from the QSO itself. The photoionization model of Bajtlik, Duncan, \\& Ostriker permits an estimate of the mean intensity of the extragalactic background radiation at the Lyman limit.   We compare the results of this standard analysis with those obtained using a maximum likelihood technique. If the spectrum of the background is assumed to be identical to that of each individual QSO, and if this background is assumed to be constant over the redshift range $1.7 < z < 3.8$, then the best fit value for $J(\\nu_{0})$ is found to be 1.4$^{+1.1}_{-0.5}$ $\\times$  10$^{-21}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, using QSO redshifts based on the Ly$\\alpha$ emission line. Systemic QSO redshifts based on the  [OIII] $\\lambda$5007 emission line for 19 objects in our sample show an average redshift of $\\sim$400 km s$^{-1}$ with respect to Ly$\\alpha$ emission. Using redshifts based on [OIII]  or Mg II for the 35 objects for which they are measured and adding 400 km s$^{-1}$ to the remaining QSO Ly$\\alpha$ redshifts  gives a lower value of $J(\\nu_{0})$, 7.0$^{+3.4}_{-4.4}$ $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. This value is in reasonable agreement with the predictions of various models of the ionizing background based on the integrated QSO luminosity function. Allowing for the fact that individual QSOs have different spectral indicies which may also be different from that of the background, we use the standard methods to solve for the HI photoionization rate, $\\Gamma$, and the parameters describing its evolution with redshift. The best fit value for the HI ionization rate we derive is 1.9$^{+1.2}_{-1.0}$ $\\times$ 10$^{-12}$ s$^{-1}$, in good agreement with models of the background which incorporate QSOs only. Finally,  we use simulated Lyman $\\alpha$ forest spectra including the proximity effect to investigate curve-of-growth effects in the photoionization model used in the analysis.  We find that the presence of lines on the saturated part of the curve-of-growth could cause our estimates of the background intensity to be overestimated by a factor of two to three. This large absorption line sample and these techniques for measuring the background and understanding the systematics involved allow us to place what we believe are the firmest limits on the background at these redshifts. ",
        "introduction": "The spectra of  quasi-stellar objects (QSOs) blueward of Ly$\\alpha$ emission show a prodigious number of absorption lines primarily due to Ly$\\alpha$ absorption by intervening neutral hydrogen along the line of sight to the QSO (Lynds, 1971; Sargent et al. 1980; Weymann, Carswell, \\& Smith, 1981). In models of these Ly$\\alpha$ systems, they are in photoionization equilibrium with a background ultraviolet radiation field. This radiation field is modeled  as the integrated emission from QSOs and young galaxies (Bechtold et al. 1987; Miralda-Escud\\'{e} \\& Ostriker 1990; Madau 1992, Meiksin \\& Madau 1993, Madau \\& Shull 1996, Haardt \\& Madau 1996, Fardal et al. 1998).  The basic trend in the number density evolution of Ly$\\alpha$ absorbers (12 $\\lesssim$ log(N$_{HI}$) $\\lesssim$ 16 cm$^{-2}$) at high redshift in a given QSO spectrum is a steep power law increase with redshift (Sargent et al. 1980). However, a decrease in the number of lines near the emission redshift of an individual QSO relative to the number of lines expected from the power law distribution has been observed (Carswell et al. 1982, Murdoch et al. 1986; Tytler 1987; Bajtlik, Duncan, \\& Ostriker 1988, hereafter BDO; Lu, Wolfe, \\& Turnshek 1991, hereafter LWT; Bechtold 1994, hereafter B94).  The simplest explanation for this $``$inverse\" effect is enhanced ionization of HI in the vicinity of the QSO by ultraviolet photons from the QSO itself.  Thus, the name $``$proximity effect\" is also used.  This interpretation, along with the assumptions about the spectrum of the background and the photoionization of the nearby intergalactic medium (IGM) by the QSOs, allows for a measurement of the mean intensity of the ionizing background at the Lyman limit of hydrogen (BDO), which can be compared to estimates of the integrated emission from QSOs.  BDO find that $J(\\nu_{0})$ is approximately constant over a redshift range $1.7 < z < 3.8$. Expressing $J(\\nu_{0})$ as $J_{-21}$ $\\times$ 10$^{-21}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, these authors find $J_{-21} \\approx 1$. Their best fit form for the dependence of $J(\\nu_{0})$ on redshift rules out a luminosity dependent cutoff in the QSO luminosity function (BWLM; Schmidt, Schneider, \\& Gunn 1986). Several other authors have carried out this analysis on other data sets (LWT, Kulkarni \\& Fall 1993, hereafter KF93, B94, Williger et al. 1994, Cristiani et al, 1995, Giallongo et al. 1996.) LWT found $J_{-21} \\approx 1$ for $1.7 < z < 3.8$; and  B94 found a value consistent with this, $J_{-21} \\approx 3$ for $1.6 < z < 4.1$. Srianand \\& Khare (1996) compile a sample of 69 QSOs from the literature (54 from B94) and obtain $J_{-21} \\approx 6$ for $1.7 < z < 4.1$. Williger et al. (1994) find a lower value, $J_{-21} \\approx 0.1-0.3$ for z $\\sim$4.2 from a single z=4.5 QSO; and Giallongo et al. (1996) find $J_{-21} \\approx 0.5$ for $2.8 < z < 4.1$, based on high resolution spectra of six QSOs. KF93 re-express the BDO formalism using maximum likelihood  techniques and use this to derive $J(\\nu_{0}) \\approx$ 6 $\\times$ 10$^{-24}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ for z $\\approx$ 0.5 from the sample of Bahcall et al. (1993).  In this paper, a homogeneous sample of QSO spectra is used to measure the mean intensity of the ionizing background via the standard proximity effect analysis and the maximum likelihood analysis of KF93. We have presented new Lyman $\\alpha$ forest data for 39 objects with $1.9 < z_{em} < 2.5$ in Paper I; and supplemented this sample with 59 objects from the literature, and one QSO from the Hamburg/CfA Bright Quasar Survey (Dobrzycki, Engels, \\& Hagen 1999). The spectra comprising our dataset are of moderate resolution, $\\sim$1 $\\AA$ FWHM; but this is not a disadvantage, as this analysis requires good absorption line statistics and therefore many QSO sight lines.  This is difficult to achieve at high resolution, the major reason for using a large set of moderate resolution spectra such as this one.  A major uncertainty in the proximity effect analysis is in the systemic redshifts of the QSOs. Redshifts measured from low ionization permitted lines (e.g. Balmer lines or Mg II) or forbidden lines (e.g. [OIII] $\\lambda\\lambda$4959,5007) lines have been shown to be redshifted with respect to Ly$\\alpha$ and C IV emission by up to $\\approx$250 km s$^{-1}$ (Boroson \\& Green 1992, Laor et al. 1995).   B94 found that increasing the values of the QSO redshifts by 1000 km s$^{-1}$ caused the best fit value of $J_{-21}$ to be decreased by a factor of 3. We therefore obtained emission line spectra for several objects in our sample in order to examine redshift differences between Ly$\\alpha$ and [OIII] $\\lambda\\lambda$4959,5007, Mg II or Balmer emission. We  investigate the effect of these shifts on the value of $J(\\nu_{0})$ derived.  In Section~\\ref{sec-data} the spectra used to measure Lyman limit fluxes for some of our sample objects and spectra used to perform systemic redshift measurements of several QSOs in our sample are presented. In Section~\\ref{sec-pe} the techniques used to measure the mean intensity of the HI ionizing background at the Lyman limit are discussed, which includes discussions of the QSO systemic redshifts in Section~\\ref{sec-zs} and the HI ionization rate measurements in Section~\\ref{sec-gamma}. In Section~\\ref{sec-cog} we present Lyman $\\alpha$ forest simulations and investigate curve-of-growth effects in the proximity effect analysis. Section~\\ref{sec-disn} is a summary and discussion of our results and the possible systematic effects entering into our analysis.  Throughout this paper, we assume values of 75 km s$^{-1}$ Mpc$^{-1}$ for H$_{0}$ and 0.5 for $q_{0}$. ",
        "conclusions": "\\label{sec-disn}  Table~\\ref{table-jnu} lists the best fit values of $J(\\nu_{0})$ found for various subsamples of this dataset using both the canonical BDO and the maximum likelihood methods. For the BDO method, the 1$\\sigma$ confidence limits are found from a $\\Delta \\chi^{2}$ of 8.18 for 7 degrees of freedom. The maximum likelihood method 1$\\sigma$ confidence limits derive from the fact that ln$(L/L_{max})$ is distributed as $\\chi^{2}/2$. The total sample consisting of 74 QSOs with all QSO redshifts based on the Ly$\\alpha$ emission line gives a best fit value of log[$J(\\nu_{0})$] of -20.90$^{+0.61}_{-0.48}$ for the BDO analysis and -20.83$^{+0.23}_{-0.20}$ for the maximum likelihood analysis.  As the results in  Table~\\ref{table-jnu} demonstrate, using narrow line redshifts for 35 of the 74 QSOs for which they have been directly measured and Ly$\\alpha$ redshifts for the rest does not change the result. However, when 400 km s$^{-1}$ is added to the Ly$\\alpha$ redshifts of the objects with no measured narrow line redshift, a value for log[$J(\\nu_{0})$] of -21.15$^{+0.17}_{-0.43}$ is derived using the BDO method and log[$J(\\nu_{0})$]=-21.17$^{+0.19}_{-0.15}$ is found using the maximum likelihood method. Recall that the mean blueshift of Ly$\\alpha$ with respect to [OIII] for the 19 objects in this paper with [OIII] $\\lambda$5007 measurements was found to be $\\sim$400 km s$^{-1}$.  This decrease in the mean intensity of the background derived when larger QSO redshifts are used is to be expected. (cf. Section~\\ref{sec-zs}) Because this measurement of the background accounts for the systematic blueshift of the Ly$\\alpha$ emission line with respect to the systemic redshift of each QSO, we consider it to be our best estimate for the mean intensity of the background at the Lyman limit.  These measurements have been made, however, using a photoionization model with somewhat unrealistic assumptions, particularly that Ly$\\alpha$ absorbers are isothermal and are composed of pure hydrogen. For clouds  with a primordial He abundance and which are in thermal and ionization equilibrium, Using CLOUDY to model the ionization state of absorbers with a metal adundance of 10$^{-2}$ solar (Cowie et al. 1995, Tytler \\& Fan 1994) as a function of $\\omega$, we find that the neutral fraction, $\\chi$, is proportional to $(1+\\omega)^{-1.21}$. This implies that \\begin{equation} \\frac{d{\\cal N}}{dz}= {\\cal A}_{0}(1+z)^{\\gamma}[1+\\omega(z)]^{-1.21(\\beta-1)} \\end{equation} In this scenario, the optimal value found for log[$J(\\nu_{0})$] is -21.10$^{+0.53}_{-0.28}$.  This value is marginally larger than the value discussed above, found under the assumption of absorbers composed of pure hydrogen; but it is not significantly different, so we conclude that the absence of metals in the BDO model has not drastically affected our measurement of the background.  It is worth noting that 16 objects in our sample of objects with no associated absorption show evidence for damped Ly$\\alpha$ absorption: 0058+019, 0100+130, 0334-204, 0913+072, 0938+119, 0952+338, 0955+472, 1009+299, 1017+280, 1215+333, 1247+267, 1548+092, 1946+770, 2126-158, 2233+131, and 2320+079. The dust in these systems could cause the intrinsic QSO fluxes to be underestimated. This in turn can cause log[$J(\\nu_{0})$] to be underestimated by up to a factor of 3, in addition to the sources of error discussed above (Srianand \\& Khare 1996). Only six of these objects, 0334-204, 0938+119, 0955+472, 1215+333, 2126-158, and 2233+131, appear in our low luminosity subsample, suggesting that this subsample is not preferentially heavily dust-obscured.  Nevertheless,  the BDO analysis was performed on all 16 objects exhibiting damped  Ly$\\alpha$ systems; and  found the best fit value for log[$J(\\nu_{0})$] to be -21.45$^{+0.40}_{-0.53}$, a factor of 1.9$^{+8.1}_{-1.6}$ lower than the value obtained for the sample as a whole. This does not allow us to say anything significant about the presence or absence of dust, so we will neglect its influence.  Dividing our line sample into subsamples of high ($z > 2.5$) and low ($z < 2.5$) redshift lines, we find marginal evidence for evolution in the intensity of the background, namely that the maximum likelihood background intensity is lower by a factor of about 1.9$^{+3.9}_{-1.4}$ at lower redshift. The BDO results corroborate this, but with larger uncertainties. The factor by which $J(\\nu_{0})$ is found to be lower at lower redshifts is 2.5$^{+27.7}_{-2.2}$. Gravitational lensing  could mimic a trend with redshift with about the same order of magnitude, if the high redshift subsample contains a significant number of unknown lenses. However, Figure~\\ref{fig:zlum} suggests little if any trend for high luminosity objects to exist at high redshifts in our sample; and the results of Section~\\ref{sec-npr} indicate that the high luminosity objects do show a somewhat stronger proximity effect despite the fact that the measured background at high redshift appears to be {\\it higher}. No other studies have found this evidence of redshift evolution in the background, so we regard it as tentative; and note that it will be interesting to  see in future work if this trend can be shown to be real and if it extends smoothly to the low values of $J(\\nu_{0})$ found at redshifts less than 1.5.  Since we find high luminosity objects do not exist preferentially at high redshift in our sample, a simple test can be done to determine whether or not there is a significant number of lensed objects in our sample. If the high luminosity QSOs are indeed intrinsically more luminous, and the proximity effect is a purely photoionization-driven phenomenon, these objects should show a more prominent proximity effect. The results of Section~\\ref{sec-npr} suggest this is the case. However, in the analysis, this  larger line deficit is normalized  to the higher Lyman limit luminosities of this subsample. Therefore, one expects these objects, when analyzed as a separate subsample, to yield a value of $J(\\nu_{0})$ that is consistent with that found for low luminosity objects if the values of the QSO fluxes are not in error due to lensing. If the high luminosity QSOs, or a subset of them, are lensed objects, then they are not necessarily intrinsically more luminous than the low luminosity QSOs.  In this case, the influence of the lensed objects on the surrounding IGM will be overestimated and given the observed line deficit, the background will also be overestimated. Table~\\ref{table-jnu} lists the results obtained for the high and low luminosity subsamples of our data set.  The values obtained for these subsamples are equal within the uncertainties. This is consistent with there being no significant effects from gravitational lensing in our sample.  \\subsection{HI Ionization Rate} We tested a range of values for $\\Gamma$, the HI ionization rate, using our data. The constant value found to fit the data the best is  1.9$^{+1.2}_{-1.0}$ $\\times$ 10$^{-12}$ s$^{-1}$. This value is in reasonable agreement with that predicted by the QSO-dominated model of Haardt \\& Madau (1996) at this redshift, 1.0 $\\times$ 10$^{-12}$ s$^{-1}$ ($q_{0}$=0.5). Using Equation~\\ref{eq:gamint} and $J^{Q}(\\nu)=J^{Q}(\\nu_{0})(\\frac{\\nu}{\\nu_{0}})^{-\\alpha}$, and assuming global QSO spectral indicies of 0, 1.5, and 2, the ionization rate found from our data corresponds to log[$J(\\nu_{0})$]= -21.34, -21.17, and -21.12, respectively.  The parameter set ($A, B, z_{c}, S$) found to give the best fit to the data is that of Fardal et al. (1998) for the Q2 luminosity function (1.2 $\\times$ 10$^{-12}$ s$^{-1}$, 0.58, 2.38, 2.77) which, for a redshift of 2.9 yields an ionization rate of 2.7 $\\times$ 10$^{-12}$ s$^{-1}$, in good agreement with our solution, and within a factor of $\\sim$3 of the Haardt \\& Madau result. Thus, we conclude that a significant contribution to the ionizing background from stellar UV emission is not required at this redshift.  \\subsection{Curve-of-Growth and Other Systematics} On the basis of a curve-of-growth argument, one might expect that weak lines would show a more prominent proximity effect than strong lines.  We have compared the results obtained for a constant equivalent width threshold of 0.32 $\\AA$ with that obtained for lines with 0.16 $\\AA < W <$ 0.32 $\\AA$. Instead of finding a more pronounced proximity effect for the weak lines, we find a less significant deficit of lines within 1.5 $h^{-1}$ Mpc of the QSOs. This deficit is 4.0$\\sigma$, versus 5.5$\\sigma$ for lines with $W >$ 0.32 $\\AA$.  As Table~\\ref{table-jnu} lists, the value of log[$J(\\nu_{0})$] recovered from these weak lines is correspondingly higher than that found using strong lines,  -20.45$^{+0.37}_{-0.90}$ versus -21.15$^{+0.17}_{-0.43}$. Cooke et al. (1997) point out that this could be the result of a higher degree of blending of weaker lines compared to strong ones in crowded spectral regions. The background flux measurement will be an overestimate because blending will cause fewer individual lines to be resolved further from the QSO.  Because the reduction in line density near the QSO will  work to reduce line blending, the overall effect of line blending will be to suppress the true magnitude of the proximity effect causing $J(\\nu_{0})$ to be overestimated, by a factor of 4.5 in this case.  It is difficult to ascertain whether this effect is as strong for lines with $W >$ 0.32 $\\AA$ or whether the curve-of-growth effect discussed in Section~\\ref{sec-cog} which also causes $J(\\nu_{0})$ to be overestimated, is more important. We expect that for lines with $W >$ 0.32 $\\AA$, the effects of blending are reduced somewhat, while the curve-of-growth effects will remain a factor.  We have addressed many of the systematics which could possibly have affected our analysis.   A treatment of the QSO systemic redshifts was integrated directly into our analysis and was found to influence the $J(\\nu_{0})$ found by up to a factor of $\\sim$2. Other effects, such as the influences of metals and dust, which can cause $J(\\nu_{0})$ to be underestimated, and the influences of lensing, line blending, and curve-of-growth effects, which can cause $J(\\nu_{0})$ to be overestimated, were treated after the fact in an attempt to understand the magnitude of their effects on the value of $J(\\nu_{0})$ derived. The CLOUDY simulations discussed above indicate that allowing for an absorber metal abundance of 10$^{-2}$ solar has little effect on the value of $J(\\nu_{0})$ found from the data. Dust in intervening absorption systems may have affected our result. Though we were unable to quantify this effect with high confidence, it could be on the order of a factor of 2. We assert that QSO flux amplification due to lensing has not significantly biassed our result; and we attempt to minimize the effect of blending discussed above by using only lines with $W >$ 0.32 $\\AA$. Our result may be susceptible to the curve-of-growth effect we addressed through the simulations in Section~\\ref{sec-cog}.  In those simulations, we found that the discrepancy between in the input and recovered values of $J(\\nu_{0})$ depended upong the input value of $J(\\nu_{0})$ itself.  The magnitude of the discrepancy corresponding to the $J(\\nu_{0})$ we found from the data was a factor of $\\sim$2.  We therefore suspect that if our result, log[$J(\\nu_{0})$]=-21.15$^{+0.17}_{-0.43}$, is systematically biased in any way, it is an overestimate of the true background and could be in error by up to a factor of 2; though this could be balanced somewhat by systematic error due to dust, which works in the opposite direction.  \\subsection{Comparison with Previous Measurements}  Our value for $J(\\nu_{0})$ agrees well with other measurements at similar redshift, with the exception of those of B94 and Fern\\'{a}ndez-Soto et al. (1995) who both derive values four times larger than our best value for $J(\\nu_{0})$, $\\sim$3 $\\times$ 10$^{-21}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$.  The measurement of B94 does not take into account QSO systemic redshifts, but she notes that if they are blueshifted with respect to Ly$\\alpha$ by 1000 km s$^{-1}$, this would lower the derived value of $J(\\nu_{0})$ by a factor of 3, bringing it into reasonable agreement with our result.  The Fern\\'{a}ndez-Soto et al. (1995) value is derived from 3 QSO spectra showing a proximity effect due to foreground QSOs. These authors are not able to place an upper limit on their measurement, but our value of 7.0 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$  for $J(\\nu_{0})$ is consistent with their lower limit of 1.6 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$. In fact, when these authors examine the proximity effect in a single QSO spectrum due to the background z$\\sim$2 QSO itself, they derive a value for $J(\\nu_{0})$ of  7.9$^{+23.}_{-6.0}$ $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, which brings their estimate into better agreement with our values for our total sample and for our low redshift subsample within their large errors. Direct measurements of the background at redshifts $\\sim$3-3.5 have been made using long-slit spectroscopy of fields containing optically thick Ly$\\alpha$ absorbers in efforts to detect fluorescent emission the absorbers produce from the ionizing radiation field incident upon them (Lowenthal et al. 1990, Mart\\'{i}nez-Gonz\\'{a}lez et al. 1995). Recent Keck telescope observations by Bunker et al. (1998) at 2.5 $<$ z $<$ 4.1 have achieved a factor of 2-10 higher sensitivity and place a firmer direct limit on the background than previous work. Their null signal in a 90-minute integration with a 3$\\arcmin$ slit sets an upper limit on $J(\\nu_{0})$ of 2 $\\times$ 10$^{-21}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$.  Cooke et al. (1997) claim that the value for the background at z$\\sim$4 is between their value of 8.0$^{+8.0}_{-4.0}$ $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$ and that of Williger et al. (1994), 1.0-3.0 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$.  Our best value of $J(\\nu_{0})$ at z$\\sim$3, 7.0 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, is in agreement with this, although within the uncertainty there is an allowance for the background to decrease as z approaches 4.  Table~\\ref{table-lit} lists these various measurements of $J(\\nu_{0})$ in the literature as well as the Kulkarni \\& Fall (1993) measurement at z$\\sim$0.5.  Figure~\\ref{fig:sum} also summarizes the literature measurements of $J(\\nu_{0})$ from z$\\sim$0.5 to z=4.5.  The solid curves in Figure~\\ref{fig:sum} delineate the evolution of the mean background intensity as a function of redshift for global background source spectral indicies between 0 and 2, derived from the Haardt \\& Madau (1996) model for the HI photoionization rate as a function of redshift discussed in Section~\\ref{sec-gamma}. Over 90\\% of our sample QSO redshifts lie within the FWHM of the Gaussian in the  Haardt \\& Madau (1996) expression using their best fit parameters. At these redshifts, the Haardt \\& Madau (1996) curves in Figure~\\ref{fig:sum} are turning over.  Nonetheless, for comparison with previous work (B94 and references therein), we investigate a power law redshift dependence of the background intensity: \\begin{equation} J(\\nu_{0},z)=J(\\nu_{0},0)(1+z)^{j}. \\end{equation} Using the BDO method, we executed a crude grid search in an attempt to constrain the power law index and normalization of this power law. The lowest $\\chi^{2}$ (3.86) between the binned data and the BDO photoionization model for a power law background was achieved by ($j$, log[$J(\\nu_{0},0)$])= (5.12, -23.97), shown by a dashed line in Figure~\\ref{fig:pl}. Extending this solution to low redshift gives log[$J(\\nu_{0}),0.5)$]=-23.0, in good agreement with the measurement of Kulkarni \\& Fall (1993). The solution ($j$, log[$J(\\nu_{0},0)$])= (-4.16, -18.76) gives the next lowest $\\chi^{2}$ (4.91); and though it also implies mean background intensities over four orders of magnitude too high at low redshift, it traces the Haardt \\& Madau model at high redshift, giving log[$J(\\nu_{0},4.5)$]=-21.8, in agreement with the Willigher et al. (1994) measurement.  It is also shown by a dashed line in Figure~\\ref{fig:pl}. Fitting parabolas to the regions near the  $\\chi^{2}$ minima in both $j$ and log[$J(\\nu_{0},0)$] gives the error in each parameter for both of these solutions, (5.12 $\\pm$ 1.96, -23.97 $\\pm$ 1.07) and (-4.16 $\\pm$ 2.36, -18.76 $\\pm$ 1.31). B94 found a similarly large range of acceptable solutions: $-7 < j < 4$ and $-16.5 < $ log[$J(\\nu_{0},0)$] $ < -23.0$. The large error bars on these fits indicate that the power law fit to the data is not well-constrained, due possibly to the fact that the mean intensity of the background is turning over at the redshifts of our sample objects, as the Haardt \\& Madau (1996) model predicts.  \\subsection{Comparison with Models for the Background} Recent models of the ionizing background include not only the integrated emission from QSOs but also a variety of other physical processes such as star formation in young, high redshift  galaxies and attenuation  of UV photons by  Ly$\\alpha$ absorbers and Lyman  limit systems (Miralda-Escud\\'{e} \\& Ostriker 1990, Madau 1991, 1992, Meiksin \\& Madau 1993, Haardt \\& Madau 1996,Fardal et al. 1998). Madau \\& Shull (1996) find that the production of metals in Ly$\\alpha$ absorbers may also be a significant contributor to the UV background at z $\\approx$ 3.  Their contribution may be up to 5 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, assuming that the bulk of the metals in the Lyman $\\alpha$ forest did not form at z $>>$ 3, and assuming a Lyman continuum escape fraction, $f_{esc}$, from a galaxy of $\\gtrsim$0.25.  They note, however, that $f_{esc}$ is essentially unconstrained.  Past debate about how the space density of QSOs evolves at high redshift  (Koo \\& Kron 1988; Boyle et al. 1991; Irwin et al. 1991; Schmidt et al. 1991; Warren et al. 1994; Kennefick et al. 1995) has been clarified by recent radio surveys (Hook et al. 1995, 1998; Shaver et al. 1996).  This work has demonstrated that the space density of radio-loud QSOs decreases rapidly with redshift beyond z$\\sim$3. Since these surveys are unaffected by any presence of dust in the intervening IGM; and since they confirm the behavior seen in optically selected surveys, they indicate that the QSO population is truly  declining at high redshift. Nevertheless, the discovery of QSOs with redshifts greater than 4 has brought better agreement between the values of $J(\\nu_{0})$ found via the proximity effect and the values predicted by the models with QSOs primarily contributing to the background (Madau 1992, Meiksin \\& Madau 1993, Haardt \\& Madau 1996).  Madau (1992) and Meiksin \\& Madau (1993) estimate the QSO UV background by integrating the QSO luminosity function (Boyle 1991) and including the effects of attenuation by hydrogen in the IGM.   Their estimates however, 1-3 $\\times$ 10$^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, are still somewhat lower than the values derived in this paper. The analysis of Haardt \\& Madau (1996) takes into account the effects of various atomic processes leading to the production of hydrogen-ionizing photons within Ly$\\alpha$ absorbers and Lyman limits systems themselves.  They conclude that observed QSOs can account for number of ionizing photons required by the proximity effect at z$\\lesssim$4.  These authors find a value of log[$J(\\nu_{0})$] equal to $\\sim$-21.4 at z=3, in good agreement with the value found in this paper at similar redshifts.  The solid lines in Figure~\\ref{fig:sum} show the results from the Haardt \\& Madau (1996) model for two different values of the global background source spectral index. The lower and upper curves show the evolution of the background for indicies of 0 and 2 respectively. The literature measurements at redshifts between 1.7 and 3.6 agree well with the model predictions.  The  z$\\sim$0.5 measurement of Kulkarni \\& Fall (1993) falls below both model curves and the z=4.5 measurement of Williger et al. (1994) falls above them.  Madau, Haardt, \\& Rees (1999) revisit the issue of the contribution of high redshift, star-forming galaxies to the ionizing background in light of recent work identifying such objects at 2 $<$ z $<$ 4. (Steidel et al. 1996a,b; Madau et al. 1996; Lowenthal et al. 1997) They calculate the critical photoionization rate necessary to reionize a non-uniform intergalactic medium as a function of redshift. This is compared to the expected contributions from QSOs and young, star-forming galaxies.  There are uncertainties in estimating both of these. The QSO luminosity function at z $>$ 4 must be extrapolated from that at lower redshifts. There is also still some debate between theory and observations, eg. of the {\\it Hubble Deep Field}, on the subject of a population of low-luminosity QSOs (see Madau et al. 1999 and references therein) which could cause the QSO luminosity function to steepen with lookback time, making up for the dearth of observed objects at z $>$ 4.  The estimation of the galaxy contribution of ionizing photons is limited by poor knowledge of luminosity function of Lyman-break galaxies at z $>$ 4 as well as by the lack of constraints upon $f_{esc}$. Nevertheless, the results are intriguing.  Assuming that $f_{esc}$=0.5, Madau et al. (1999) find that the contribution of hydrogen-ionizing photons from star-forming galaxies  z $\\sim$ 3 could exceed that from QSOs by a factor of more than 3. However, the QSO contribution at this redshift is sufficient, according to these estimates, to ionize the IGM at this redshift. Deharveng et al. (1997) estimate a much lower $f_{esc}$ at z=0, less than 1\\%, based on  the local galaxy H$\\alpha$ luminosity density. Furthermore, Devriendt et al. (1998) make an independent estimation of the galaxy contribution to $J(\\nu_{0})$ assuming damped Ly$\\alpha$ systems to be the progenitors of present day galaxies. Their semi-analytic models include a treatment of not only HI absorption of Lyman limit photons in the intervening IGM, but also of HI and dust absorption  in the interstellar medium of the photon-producing galaxies.  Their results show that constraining $f_{esc}$ in this way yields a much lower contribution to the UV background from galaxies at z $>$ 2.  At z $\\sim$ 2.5, their estimated QSO contribution to $J(\\nu_{0})$ is 3 orders of magnitude greater than that expected from galaxies. Our measurement of $J(\\nu_{0})$ is consistent with the UV background being QSO dominated in the models of both these authors and Haardt \\& Madau (1996).  In the models of Madau et al. (1999), the scenario changes at z $\\gtrsim$ 3.5.  At this redshift, the QSO contribution of ionizing photons falls below the critical limit needed to photoionize the IGM; and by z=5, it will fall short of the critical value by a factor of $\\sim$4. This implies that at high redshift, the contribution from young stars may become the dominant contributor to the background, with the caveat that the space density of star-forming galaxies would have to be maintained at the level observed at z $\\approx$ 3, and that most of their UV photons would have to be free to escape into the IGM.  The Devriendt et al. (1998) models lead to the conclusion, however, that the galaxy contribution to the UV background is negligible at high redshifts.  In conclusion, the proximity effect data at present reflect that the UV background at 2 $<$ z $<$ 4 is QSO dominated.  The discrepancies between this model at low and high redshifts (Kulkarni \\& Fall 1993, Williger et al. 1994) indicate that the contribution to the background from galaxies may be of larger relative importance.  We plan to undertake an analysis of the proximity effect at low redshifts from a large sample of QSO spectra taken with the Faint Object Spectrograph on the Hubble Space Telescope to place better constraints on the background at 0.5 $<$ z $<$ 2.  Further observations of objects at z $>$ 4 are also of particular interest to this subject."
    },
    "0004/astro-ph0004363_arXiv.txt": {
        "abstract": "We present the results of ISO SWS and LWS grating scans towards the embedded Young Stellar Objects (YSOs) S106 IR and Cep~A East. Emission from the pure rotational lines of H$_2$ and the infrared fine structure lines of [C\\,{\\sc ii}], [O\\,{\\sc i}], [S\\,{\\sc i}], [Si\\,{\\sc ii}] and [Fe\\,{\\sc ii}], as well as absorption bands due to H$_2$O, CO and CO$_2$ ice were detected toward Cep~A. In S106 we detected emission lines of H$_2$, CO, H\\,{\\sc i}, and a large number of ionized species including Fe, O, N, C, Si, S, Ne and Ar. S106 also shows many of the infrared PAH bands in emission. Excitation temperatures and molecular hydrogen masses were derived from the low-lying pure rotational levels of H$_2$ and are 500 and 730~K and 8 and 3 $\\times$ 10$^{-3}$~M$_\\odot$ for S106 and Cep~A, respectively. Since both objects are expected to have several solar masses of H$_2$ in their environment, we conclude that in both cases the bulk of the H$_2$ is cooler than a few hundred Kelvins. Excitation temperatures and line ratios were compared with those predicted by theoretical models for PDRs and dissociative and non-dissociative shocks. The [S\\,{\\sc i}] 25.2~$\\mu$m/[Si\\,{\\sc ii}] 34.8~$\\mu$m ratio is a particularly useful shock versus PDR discriminant and we conclude that S106 IR is dominated by PDR emission while Cep~A East has a large shock component. From an analysis of the ionic lines in S106 we conclude that the central star must have a temperature around 37,000~K, corresponding to a spectral type of O8. From its luminosity it is concluded that the driving source of Cep~A must also be a massive early-type star. The absence of strong high-ionization ionic lines in its ISO spectrum shows that Cep~A has not yet created a significant H\\,{\\sc ii} region and must be younger than S106, illustrating the process of the clearing of the surroundings of a massive young star. ",
        "introduction": "The infrared emission-line spectrum of a Young Stellar Object (YSO) is dominated by the interaction of the central star with the remnants of the clouds from which it formed. The intense UV radiation generated by accretion as well as by the central star itself causes dissociation of molecular material close to the YSO and ionizes much of the atomic material, giving rise to typical nebular and recombination lines. Furthermore, the interaction of the UV field with the remnants of the star's natal cloud will produce a photodissociation region (PDR; see Hollenbach \\& Tielens 1999 for a comprehensive review). In a PDR, heating of the gas occurs by collisions with electrons, photoelectrically ejected from grain surfaces. The gas in the surface regions of PDRs reaches temperatures of typically 500~K (e.g. Tielens \\& Hollenbach 1985), producing a distinctive infrared spectrum due to collisionally excited low-lying levels of many molecular and neutral species.  Another phenomenon associated with YSOs are strong stellar winds, often collimated into a bipolar outflow. These will cause a shock as they drive a supersonic wave into the surrounding molecular cloud, heating it in the process. Shocks are usually divided into two distinct categories: J- or Jump-shocks, and C- or Continuous-shocks. In J-shocks viscous heating of the neutrals occurs in a thin shock front in which radiative cooling is insignificant, while the post-shock gas is heated to several times 10$^4$ degrees, dissociating all molecular material (e.g. Hollenbach \\& McKee 1989). Cooling of the post-shock gas occurs through atomic fine-structure lines as well as through re-formation of molecules. In contrast, C-shocks are magnetized, non-dissociative shocks in which the physical conditions change more gradually from their pre- to post-shock values and cooling is mainly through radiation from molecular material (e.g. Kaufman \\& Neufeld 1996). If the temperatures in a C-shock become sufficiently high to start to dissociate molecules, the cooling through the molecular lines diminishes, and the shock temperature increases until it turns into a J-shock. Shocks with a shock velocity larger than 40~km~s$^{-1}$ are usually J-shocks, whereas slower shocks are usually of C-type (Chernoff et al. 1982). In this paper we will illustrate the power of infrared spectroscopy to study the above mentioned phenomena using two young bipolar nebulae, S106 and Cep~A. \\begin{figure}[t] \\centerline{\\psfig{figure=s106_iso.ps,width=8.7cm,angle=0}} \\caption[]{SWS AOT S01 (solid rectangles), AOT S02 (dashed rectangles), AOT S06 (dotted rectangles) and LWS (solid circle; beam average FWHM) aperture positions for our measurements of S106 superimposed on a K'-band image of the region. The rectangles indicate the apertures (in increasing size) for SWS bands 1A--2C (2.4--12.0~$\\mu$m), 3A--3D (12.0--27.5~$\\mu$m), 3E (27.5--29.5~$\\mu$m) and 4 (29.5--40.5~$\\mu$m).} \\end{figure} \\begin{figure}[t] \\centerline{\\psfig{figure=cepa_iso.ps,width=8.7cm,angle=0}} \\caption[]{Same as Fig.~2 for Cep~A East.} \\end{figure}  The bipolar nebula S106 is one of the most studied H\\,{\\sc ii} regions in our galaxy. The exciting source of the region is a very young massive stellar object, known as either IRS4 in the terminology of Gehrz et al. (1982) or IRS3 following Pipher et al. (1976). Around this central source, Hodapp \\& Rayner (1991) found a cluster of about 160 stars, embedded in the molecular cloud surrounding S106. Observed radio emission and H\\,{\\sc i} recombination lines from the region have been suggested to arise in a strong ($\\dot{M}$ $\\approx$ 2 $\\times$ 10$^{-5}$~M$_\\odot$~yr$^{-1}$), fast ($\\approx$ 200~km~s$^{-1}$), stellar wind, driving a shock into the surrounding extended molecular cloud (Hippelein \\& M\\\"unch 1981). A dark lane, largely devoid of any optical or infrared emission, separates the two lobes of the H\\,{\\sc ii} region. This dark lane has been quoted many times as a prime example of a large (30\\arcsec) circumstellar disk, consisting of cool gas and dust, seen nearly edge-on (Bieging 1984; Mezger et al. 1987). However, Loushin et al. (1990) showed this structure to be an expanding ring of molecular material rather than a protoplanetary disk. Near-infrared ro-vibrational emission of H$_2$ is seen to arise just outside the H\\,{\\sc ii} region, suggesting an origin in a PDR (Hayashi et al. 1990). \\begin{figure*} \\centerline{\\psfig{figure=s106_all.ps,width=16cm,angle=270}} \\caption[]{Combined ISO SWS/LWS full grating spectrum for S106 IR with the most prominent features identified. The insets show enlargements of the SWS spectrum in the 2.4--4.0 and 4.0--12.0 micron range.} \\end{figure*} \\begin{figure*} \\centerline{\\psfig{figure=cepa_all.ps,width=16cm,angle=270}} \\caption[]{Same as Fig.~3 for Cep~A East.} \\end{figure*}  Cep~A is another well-known site of recent star formation. It contains a luminous (2.5 $\\times$ 10$^4$~L$_\\odot$; Evans et al. 1981) far-infrared source, as well as several fainter infrared sources. The core of the Cep~A region remains obscured at optical and infrared wavelengths due to massive amounts of extinction, and is the source of an energetic, complex molecular outflow (Bally \\& Lane 1982). Radio-observations show that it contains two strings of about 13 ultracompact H\\,{\\sc ii} regions, arranged in a Y-shape (Hughes 1988). The source IRS 6a (Lenzen et al. 1984), located in the infrared nebula to the east of the core, is the dominant source of the eastern lobe at 20~$\\mu$m (Ellis et al. 1990), but polarization measurements show that the nebula is illuminated by the compact radio source HW-2 (Hughes \\& Wouterloot 1984), located $\\sim$~5\\arcsec~ south of IRS 6a (Casement \\& McLean 1996). The fact that it is not visible at 20~$\\mu$m implies an extremely large extinction towards HW-2. H$_2$ emission occurs in both the eastern and western parts of the nebula (Doyon \\& Nadeau 1988). Hartigan et al. (1996) showed that the molecular emission to the east appears as a regular jet, whereas that to the west concentrates in shells, which they proposed to arise in wakes from bow shocks surrounding Herbig-Haro objects in the region. ISO observations of the western part of the region were discussed by Wright et al. (1996), who modelled the observed infrared fine structure lines as arising in a planar J-shock with shock velocity 70--80~km~s$^{-1}$. The same authors modelled the observed molecular hydrogen emission as arising in a combination of several C-shocks. This multitude of shocks in the region agrees with the explanation of Narayanan \\& Walker (1996), who reported the presence of multiple episodes of outflow activity from the region. Recently, Goetz et al. (1998) presented strong evidence for this scenario through near-infrared H$_2$ and [Fe\\,{\\sc ii}] images of Cep~A East showing several distinct regions containing shocks.  In this paper we will present new {\\it Infrared Space Observatory} (ISO; Kessler et al. 1996) data on infrared fine-structure and molecular emission lines from the central region of S106, centered on IRS4, and the eastern part of Cep~A, centered on IRS 6a. These data give for the first time access to a large spectral range that was unaccessible from ground-based observations and contain lines which are decisive probes of the excitation conditions. We will show that the observed emission-line spectrum can be well explained as arising in the combination of a PDR and a H\\,{\\sc ii} region in the case of S106, whereas this emission is shock-excited in the case of Cep~A. ",
        "conclusions": "In the previous sections we saw that the infrared emission-line spectrum of S106 could be well explained by the presence of an H\\,{\\sc ii} region and a clumpy PDR. The average density in the H\\,{\\sc ii} region should be around 2 $\\times$ 10$^3$~cm$^{-3}$. The density in the PDR, which was estimated to be 10$^5$--10$^6$~cm$^{-3}$, is several orders of magnitude higher, in agreement with one would expect in a scenario in which a young massive star has photo-ionized its natal cloud in all but the densest clumps.  In Section~7 we concluded that the incident FUV flux on the S106 PDR should be rather high ($\\ge$ 10$^5$~G$_0$). We also showed that the central source in S106 is of spectral type O8, with a total luminosity of 4.2 $\\times$ 10$^4$~L$_\\odot$. From a Kurucz (1991) model for a stellar photosphere with $T_{\\rm eff}$ = 37,000~K and $\\log g$ = 4.0, we compute that such a star emits a total FUV (6--13.6~eV) flux of 3.6 $\\times$ 10$^{36}$~erg~s$^{-1}$~sr$^{-1}$. To dilute this stellar FUV field to a value of 10$^5$--10$^6$~G$_0$, the PDR must be at a location 3--10 $\\times$ 10$^3$~AU away from the central star, corresponding to a projected distance of 5--17\\arcsec. This projected distance is independent of the assumed distance towards S106. It is within the range allowed by the SWS entrance apertures and is compatible with an origin in the $\\approx$ 30\\arcsec~ diameter region of H$_2$ emission in the H$_2$ 1--0 S(1) image of S106 by Hayashi et al. (1990). We conclude that a central O8 star is sufficient to produce the FUV radiation field reaching the S106 PDR.  These results agree well with other determinations. Schneider et al. (2000) have spatially resolved the 157.7~$\\mu$m [C\\,{\\sc ii}] and 63.2~$\\mu$m [O\\,{\\sc i}] emission in S106 and concluded that there is a PDR region with density $\\simeq$~10$^6$~cm$^{-3}$ and strong (up to 8 $\\times$ 10$^5$~G$_0$) UV field in the immediate environment of S106 and a more extended dense (10$^5$--10$^6$~G$_0$) PDR region exposed to a lower intensity (300--500~G$_0$) UV field. We are clearly dominated by their first component. This same component may be identified with the high-density (10$^5$~cm$^{-3}$) fluorescent H$_2$ regions identified by Hayashi et al. (1990).  Both the presence of [S\\,{\\sc i}] 25.2~$\\mu$m emission and the H$_2$ spectrum point unambiguously to shocked gas as the origin of the infrared emission lines in Cep~A East. A combination of a slow (20~km~s$^{-1}$) non-dissociative shock and a faster (30--60~km~s$^{-1}$) dissociative shock are required to explain our ISO spectra. The density of the pre-shock gas in Cep~A East is similar to that in the S106 PDR: 10$^5$--10$^6$~cm$^{-3}$. Our J-shock component may be identified with the [Fe\\,{\\sc ii}] 1.64~$\\mu$m clumps observed by Goetz et al. (1998). The C-shocks could be due to their H$_2$ knots which do not show ionic emission.  It is interesting to compare the results derived here for the emission from the eastern lobe of Cep~A with the results by Wright et al. (1996), obtained using the same instruments on board ISO, for the western part of the nebula. They derived a $T_{\\rm rot}$ of 700 $\\pm$ 30~K for the H$_2$ lines with upper level energies up to 7000~K, and a temperature range up to 11,000~K for the higher upper level energies. They explained this H$_2$ emission as arising from a combination of at least two C-shocks with different pre-shock density, shock velocity and covering factor. In addition to this, they also reported emission from [Ne\\,{\\sc ii}], [S\\,{\\sc i}] and [Si\\,{\\sc ii}], which they explained as arising in a planar J-shock with pre-shock density of 10$^3$--10$^4$~cm$^{-3}$ and shock velocity 70--80~km~s$^{-1}$. Qualitatively, the detected lines in the eastern and western part of the nebula are well in agreement. The main difference between the observed lines seems to be that they are much more intense in the western part of Cep~A, indicating that the densities there are a factor of 100 higher than those obtained for the eastern lobe, but the shock velocities are comparable in both parts of the nebula. The similarity between these two parts of Cep~A are in agreement with a scenario in which the driving source of the molecular outflow, HW-2, went through multiple episodes of outflow activity, as suggested by Narayanan \\& Walker (1996). In this scenario, the J-shock component could be due to the most recent period of enhanced mass loss, whereas one or more C-shocks could be due to older generations of outflows. Alternatively, the stellar wind material could produce a (fast) J-shock while the surrounding molecular cloud might be swept up by a slower C-shock.  As was already remarked by Staude \\& Els\\\"asser (1993), the differences between the environments of these two massive embedded YSOs, S106 and Cep~A, are remarkable. The mid- and far-infrared observations presented in this paper show that the difference between these two sources cannot be due to a different orientation of their nebulae; it can only be a reflection of their different evolutionary status, with Cep~A being the younger of the two. In S106 the stellar wind and UV radiation of the exciting source have cleared and excited a sufficiently large region to create strong PDR emission, whereas in the case of Cep~A, the central source is still heavily embedded and we only observe the interaction of the outflow with its surroundings. In due time, Cep~A will clear its surroundings, ionize the hydrogen, and evolve into a bipolar nebula quite similar to S106."
    },
    "0004/astro-ph0004225_arXiv.txt": {
        "abstract": "s{I discuss some theoretical ideas concerning the representation of quantum gravity as a Lorentz-symmetry-violating `medium' with non-trivial optical properties, which include a refractive index in `vacuo' and stochastic effects associated with a spread in the arrival times of photons, growing linearly with the photon energy. Some of these properties may be experimentally detectable in future satellite facilities (e.g. GLAST or AMS), using as probes light from distant astrophysical sources such as gamma ray bursters. I also argue that such linear violations of Lorentz symmetry may not always be constrained by ultra-high-energy cosmic-ray data, as seems to be the case with a specific (stringy) model of space-time foam.}  \\begin{minipage}{4.2truein} {\\small Presented at the {\\it Lake Louise Winter Institute 2000: From Particles to the Universe}, Lake Louise, Alberta, Canada, 20-26 February 2000.} \\end{minipage} ",
        "introduction": "The suggestion made by J.A. Wheeler~\\cite{wheeler} that space time may acquire a discrete foamy structure at sub-Planckian scales has received considerable attention. The relevant works span a wide range of research fields, from phenomenological approaches~\\cite{phenom} to theoretical modelling of quantum gravity~\\cite{qg} and/or string theory~\\cite{string}. The purpose of this talk is to focus on a recent scenario on the emergence of a foamy space-time structure in the context of string theory~\\cite{emn}, and to discuss briefly its possible phenomenological consequences, especially in an astrophysical context using gamma-ray-bursters (GRB) as the relevant probes~\\cite{aemns,efmmn}.  An important feature of most models of quantum space time foam is the breaking of Lorentz Invariance (LI) by quantum gravity effects. In such an approach LI is only an approximate symmetry of the low-energy world.  In the context of the specific model of (stringy) space time foam proposed in \\cite{emn}, the basic idea may be summarised as follows: consider a closed-string state propagating in a $(D+1)$-dimensional space-time, which impacts a very massive $D$(irichlet) particle embedded in this space-time. In the modern view of string theory, $D$ particles must be included in the consistent formulation of the ground-state vacuum configuration. We argue that the scattering of the closed-string state on the $D$ particle induces recoil of the latter, which distorts the surrounding space-time in a stochastic manner.  From a string-theory point of view, the essential feature is the deviation from conformal invariance of the relevant world-sheet $\\sigma$ model that describes the recoil. This is compensated by the introduction of a Liouville field~\\cite{distler89}, which in turn is identified with the target time in the approach~\\cite{emn} adopted here. ",
        "conclusions": "We have discussed here some possible low-energy astrophysical probes of quantum gravity, concentrating on the possibility that the velocity of light might depend on its frequency, i.e., the corresponding photon energy. This idea is very speculative, and the model calculations that we have reviewed require justification and refinement. However, we feel that the suggestion is well motivated by the basic fact that gravity abhors rigid bodies, and the related intuition that the vacuum should exhibit back-reaction effects and act as a non-trivial medium. These features have appeared in several approaches to quantum gravity, including the canonical approach and ideas based on extra dimensions.  As could be expected, we have found no significant effect in the data available on GRBs~\\cite{efmmn}, either in the possible delay times of photons of higher energies, or in the possible stochastic spreads of velocities of photons with the same energy. However, it has been established that such probes may be sensitive to scales approaching the Planck mass, if these effects are linear in the photon energy. We expect that the redshifts of many more GRBs will become known in the near future.  Future observations of higher-energy photons from GRBs~\\cite{GLAST} would be very valuable, since they would provide a longer lever arm in the search for energy-dependent effects on photon propagation. As emphasised in the text, such observations are {\\it essential} for the stochastic quantum gravity effect (\\ref{stochastic}), and hence should be considered {\\it complementary} to cosmic ray or other astrophysical data, which may not be relevant for constraining this effect."
    },
    "0004/astro-ph0004048_arXiv.txt": {
        "abstract": "The ACIS front-illuminated CCDs onboard the Chandra X-ray Observatory were damaged in the extreme environment of the Earth's radiation belts, resulting in enhanced charge transfer inefficiency (CTI).  This produces a row dependence in gain, event grade, and energy resolution. We model the CTI as a function of input photon energy, including the effects of de-trapping (charge trailing), shielding within an event (charge in the leading pixels of the 3$\\times$3 event island protect the rest of the island by filling traps), and non-uniform spatial distribution of traps.  This technique cannot fully recover the degraded energy resolution, but it reduces the position dependence of gain and grade distributions.  By correcting the grade distributions as well as the event amplitudes, we can improve the instrument's quantum efficiency.  We outline our model for CTI correction and discuss how the corrector can improve astrophysical results derived from ACIS data. ",
        "introduction": "}  The Advanced CCD Imaging Spectrometer (ACIS) instrument on the Chandra X-ray Observatory \\citep{odell98} employs bulk front-illuminated (FI) and back-illuminated (BI) CCDs to give good spectral resolution and good 0.2--10~keV quantum efficiency (QE) \\citep{bautz98}.  These devices, designed and manufactured at MIT's Lincoln Laboratories \\citep{burke97}, couple with the Chandra mirrors to provide spatially-resolved X-ray spectroscopy on spatial scales comparable to ground-based optical astronomy.  Due to the unanticipated forward scattering of charged particles (probably mainly $\\sim$100 keV protons, \\cite{prigozhin00}) by the Chandra mirrors onto the ACIS CCDs, the FI devices have suffered degraded performance on-orbit, most pronounced at the top of the devices, near the aimpoint of the imaging array.  This radiation first encounters the gate structure of FI devices, causing charge traps that increase the charge transfer inefficiency (CTI) of these devices, resulting in gain variations, event grade distortions, and degraded energy resolution as a function of row number.  Due to the manufacturing process, the ACIS BI devices have always shown modest CTI effects, but the silicon covering the gates protects them from severe radiation damage on-orbit so they show no additional CTI.  Our ongoing efforts to mitigate the BI CTI by modeling and post-processing the data allow us to address the new problem of FI CTI in a timely fashion.  Figure~\\ref{fig:allenergy-cti} (upper panel) illustrates the effects of CTI on the ACIS I3 chip by showing the row-dependent gain and energy resolution for spectral lines in the External Calibration Source (ECS) (from low to high energy:  Al~K$\\alpha$, Ti~K$\\alpha$,$\\beta$, Mn~K$\\alpha$,$\\beta$).  The data were filtered to keep ASCA-like grades 0, 2, 3, 4, and 6 (``g02346'', \\cite{yamashita}) only.  Note that the charge loss is more severe at higher energies (the loci steepen with line energy).  This dataset, obtained after the damage was discovered and satellite operations were modified to prevent further degradation, is representative of the state of the CTI from mid-September 1999 to the end of January 2000, when the ACIS focal plane temperature was lowered from -110C to -120C.  At this new colder temperature, the trap population that causes the row-dependent energy resolution is partly suppressed \\citep{gallagher98}.   \\begin{figure} \\plotone{cti_I3_g02346_allenergy.ps} \\plotone{nocti_I3_g02346_allenergy.ps} \\figcaption{Upper panel:  CTI in the I3 chip, illustrated with data from the ECS.  Note the row-dependent gain and energy resolution and the energy-dependent charge loss.  Lower panel: CTI-corrected version of these data. \\label{fig:allenergy-cti}} \\end{figure}   CTI redistributes the charge in each event among the pixels in the 3$\\times$3 island (grade ``morphing'') as some charge traps release their charge on short timescales, so that charge reappears in adjacent pixels.  This grade morphing, coupled with the standard g02346 grade filtering, causes the QE to vary with event energy and position on the device.  We have developed and optimized a model of CTI that accounts for charge loss and the spatial redistribution of charge (``trailing''), in both the parallel and serial registers, thus reproducing the spatially-dependent gain, QE, and grade distribution of these devices. By forward modeling of these effects separately for each event, we produce a ``CTI-corrected'' event list that is useful for data analysis and astrophysical interpretation.  For the BI devices, we must model parallel and serial charge loss and trailing independently to match the data.  There are charge traps in the imaging array, framestore array, and serial register.  Since the FI devices were damaged by radiation that penetrates only a few microns into the device \\citep{prigozhin00}, their (shielded) framestore arrays and serial registers are protected.  So, although the CTI is more pronounced on FI devices, it is easier to model than in the BI devices.  Our techniques are described in more detail, with more emphasis on the BI devices, in a separate article \\citep{townsley00b}. ",
        "conclusions": ""
    },
    "0004/astro-ph0004338_arXiv.txt": {
        "abstract": "I present the first results from a HST/NICMOS imaging study of the most metal-poor blue compact dwarf  galaxy IZw18. The near infrared color--magnitude diagram (CMD) is dominated by two populations, one 10-20 Myr population of red super giants and one 0.5-5 Gyr population of asymptotic giant branch  stars. Stars older than 1 Gyr are required to explain the observed CMD at the adopted distance of 12.6 Mpc, showing that IZw18 is {\\emph {not}} a young galaxy. The results hold also if the distance to IZw18 is significantly larger. This rules out the possibility that IZw18 is a truly young galaxy formed recently in the local universe. ",
        "introduction": "The blue compact galaxy (BCG) IZw18 (also known as Markarian 116) was first described by Zwicky in 1966. In a seminal paper Searle and Sargent (\\cite{ss}) showed that IZw18 had a very low oxygen abundance, and they concluded that the galaxy either must be young or form stars in short bursts intervened by long quiescent epochs (cf. Searle et al. \\cite{ssb}). This has been followed by many studies of the chemical properties of IZw18 and BCGs in general. Alloin et al. (\\cite{alloin}) derived an oxygen abundance of 12+log(O/H)=7.2 (2\\% solar) for IZw18 in agreement with more recent values, e.g. Izotov  and Thuan (\\cite{it}). In spite of many searches for more metal-poor galaxies, IZw18 remains the most metal-poor one (when nebular oxygen abundances are considered) and  is often regarded as the prototypical BCG (see Kunth and \\\"Ostlin \\cite{ko} for a review).  If IZw18 was a genuinely young galaxy, that would imply that galaxies can form still at the present cosmic epoch and considerable effort has therefore been devoted to examine the ages of BCGs. Although it is now clear that most BCGs are old (see Kunth and \\\"Ostlin \\cite{ko}), IZw18 has remained among a small number of galaxies which show no clear signs for the presence of an  old stellar population. Kunth and Sargent (\\cite{ks}) argued that IZw18 could be young and that the low, but non-zero, metallicity could be due to self-pollution by massive stars. Recently an extensive study of chemical abundances of BCGs was published by Izotov and Thuan (\\cite{it}). They found that for very low oxygen abundances (less than 5\\% solar) the N/O and C/O ratios stay constant and have very little scatter. Their interpretation is that N and C are produced entirely by massive ($M \\ge 10 M_{\\odot}$) stars and that this implies that these galaxies must be younger than 40 Myr. While other interpretations are possible (Kunth and \\\"Ostlin \\cite{ko}), this an interesting idea that deserves to be tested. With its record low oxygen abundance, IZw18 provides the critical test case.  The chemical abundances of IZw18 have been used to infer both a high and a low age (Kunth and \\\"Ostlin \\cite{ko}), and it is clear that chemistry alone has limited value in constraining the age and star formation history  given the many uncertain physical parameters involved (winds, yields, infall, initial mass function, etc.). Stronger constraints would be obtained by direct observations of the resolved stellar population. Hunter and Thronson (\\cite{hunter:thronson}) used deep HST/WFPC2 imaging to investigate the stellar population of IZw18, and found a population of young massive stars. No evidence for old stars was found, but neither any against. Similar results were obtained  by Dufour et al. (\\cite{dufour}) from an independent WFPC2 data set. Lately, these two  WFPC2  data sets have been carefully reanalyzed by Aloisi, Tosi and Greggio (\\cite{aloisi}), who discovered a population of red stars. Their results indicate that IZw18 is at least 0.5 Gyr old. This was however disputed by Izotov et al. (\\cite{izotov}) who claim that the distance to IZw18 has been underestimated, and that hence the age has been overestimated.   In this paper I report on the first results from a near infrared imaging study of IZw18 using NICMOS  on board the HST. The goal was to search for an old stellar population, by identifying luminous red giants and asymptotic giant branch (AGB) stars. The data is however also valuable for studying the young massive stars, e.g. red super giants (RSGs) in the near infrared (NIR). A follow up program designed at investigating the presence of carbon stars by targeting molecular features were performed during {\\sl  Cycle 7-NICMOS}. The results of this investigation will be presented in a forthcoming paper. ",
        "conclusions": "I have analyzed the near infrared color-magnitude diagram from HST/NICMOS images of IZw18, the most metal-poor galaxy known. Assuming a  distance of 12.6 Mpc, I find support for the presence of AGB stars at least 1 Gyr old.  The results presented in this letter hold qualitatively even if the distance to IZw18 is significantly different. The latest star formation event appears to have peaked 14 Myr ago. My results are in good  agreement with Aloisi et al. (\\cite{aloisi}) who from a careful reanalysis of HST/WFPC2 archival data  concluded that IZw18 must be at least 0.5  Gyr old at their adopted distance of 10 Mpc. Thus, there is no need for galaxies with oxygen abundance less than 5 \\% of the solar value to be young as was claimed by Izotov and Thuan (\\cite{it}). This closes the question whether IZw18 is a genuinely young galaxy, still in the process of formation: the answer is -- No!"
    },
    "0004/astro-ph0004174_arXiv.txt": {
        "abstract": "We present near-infrared photometry and optical spectroscopy of very low-mass stars and brown dwarf candidates in the Pleiades open cluster. The membership status of these objects is assessed using color-magnitude diagrams, lithium and spectral types. Eight objects out of 45 appear to be non-members. A search for companions among 34 very low-mass Pleiades members (M$\\le$0.09~M$_\\odot$) in high-spatial resolution images obtained with the Hubble Space Telescope and the adaptive optics system of the Canada-France-Hawaii telescope produced no resolved binaries with separations larger than 0.2 arcsec (a$\\sim$27~AU; P$\\sim$444 years). Nevertheless, we find evidence for a binary sequence in the color-magnitude diagrams, in agreement with the results of Steele \\& Jameson (1995) for higher mass stars. We apply the lithium test to two objects: CFHT-Pl-16, which lies in the cluster binary sequence but is unresolved in images obtained with the Hubble Space Telescope; and CFHT-Pl-18, which is binary with 0\".33 separation (Mart\\'\\i n et al. 1998). The first object passes the test, but the second object does not. We conclude that CFHT-Pl-16 is an Pleiades brown dwarf binary with separation $<$11~AU, and that CFHT-Pl-18 is a foreground system. We compare the multiplicity statistics of the Pleiades very low-mass stars and brown dwarfs with that of G and K-type main sequence stars in the solar neighborhood (Duquennoy \\& Mayor 1991). We find that there is some evidence for a deficiency of wide binary systems (separation $>$27 AU) among the Pleiades very low-mass members. We briefly discuss how this result can fit with current scenarios of brown dwarf formation. We correct the Pleiades substellar mass function for the contamination of cluster non-members found in this work. We find a contamination level of 33\\% among the brown dwarf candidates identified by Bouvier et al. (1998). Assuming a power law IMF across the substellar boundary, we find a slope dN/dM$\\sim$M$^{-0.53}$, implying that the number of objects per mass bin is still rising but the contribution to the total mass of the cluster is declining in the brown dwarf regime. ",
        "introduction": "Brown dwarfs (BD) provide the opportunity of extending classical studies of stellar properties into the hitherto unexplored realm of objects with masses lower than the H-burning limit (0.075~M$_\\odot$; e.g. Baraffe et al. 1998). Stars are frequently associated with other stars in multiple systems. 57\\% of nearby G-type dwarfs have H-burning companions (Abt \\& Levy 1976; Duquennoy \\& Mayor 1991, hereafter DM91). The multiplicity frequency seems to be lower among nearby M-type dwarfs (38\\%, Henry \\& McCarthy 1990; Fisher \\& Marcy 1992). This difference could be due to the smaller mass range available for the stellar secondaries of M dwarfs. On the other hand, the distribution of binary separations is similar for G- and M-type stars. It has a broad maximum from 3 to 30 Astronomical Units (AU). There is no indication that stellar binary properties depend on primary mass.  Stellar clusters are important for studying the multiplicity properties as a function of age, metallicity and mass. One of the best studied open clusters is the Pleiades. It is nearby (d=125~pc), young (120~Myr), the metallicity is close to solar, and there are more than 800 known members (see Hambly 1998 for a review). The Pleiades mass function has been studied over a broad mass range, from 4~M$_\\odot$ to 0.04~M$_\\odot$ (Hambly, Hawkins \\& Jameson 1993; Meusinger, Schilbach \\& Souchay 1996; Bouvier et al. 1998, hereafter B98). Cluster BDs have been confirmed via the lithium test (Basri, Marcy \\& Graham 1996; Rebolo et al. 1996). Several imaging surveys have recently identified a numerous population of Pleiades BD candidates (B98; Festin 1998; Zapatero Osorio et al. 1999; Hambly et al. 1999). Follow-up observations of these objects is necessary for assessing their membership.  The search for BD binaries is interesting because of the following reasons: 1) Any fainter and cooler secondary of a BD should have even lower mass than the primary. 2) The binary frequency among BDs and the distribution of orbital periods and eccentricities is an important clue for understanding the formation of these objects. 3) BD binaries provide the opportunity to measure dynamical masses, which are necessary for calibrating evolutionary models. 4) The Pleiades substellar mass function needs to be corrected for binarity.  The paper is organized as follows: in Section 2 we present the observations and describe the data analysis. In Section 3 we search for companions and we construct color magnitude diagrams. In Section 4 we discuss the cluster membership of our objects, we derive parameters for the likely members, we estimate the binary frequency, and we discuss the implications of our results for the substellar mass function. ",
        "conclusions": ""
    },
    "0004/gr-qc0004075_arXiv.txt": {
        "abstract": "Cusps of cosmic strings emit strong beams of high-frequency gravitational waves (GW). As a consequence of these beams, the stochastic ensemble of gravitational waves generated by a cosmological network of oscillating loops is strongly non Gaussian, and includes occasional sharp bursts that stand above the rms GW background. These bursts might be detectable by the planned GW detectors LIGO/VIRGO and LISA for string tensions as small as $G \\mu \\sim 10^{-13}$. The GW bursts discussed here might be accompanied by Gamma Ray Bursts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004032_arXiv.txt": {
        "abstract": "The problem of the line intensities -- oxygen abundance calibration has been considered. We confirm the idea of McGaugh (1991) that the strong oxygen lines ($[OII] \\lambda \\lambda 3727, 3729$ and $[OIII] \\lambda \\lambda 4959, 5007$) contain the necessary information for determination of accurate abundances in low-metallicity (and may be also in high-metallicity) HII regions. It has been found that the excitation parameters $p_{3}$ or $p_{2}$ (which are defined here as contributions of the radiation in $[OIII] \\lambda \\lambda 4959, 5007$ lines and in $[OII] \\lambda \\lambda 3727, 3729$ lines to the \"total\" oxygen radiation respectively) allow to take into account the variations in $R_{23}$ values among HII regions with a given oxygen abundance. Based on this fact a new way of the oxygen abundance determination in HII regions (p -- method) has been constructed and corresponding relations between $[OII] \\lambda  \\lambda 3727, 3729$, $[OIII] \\lambda  \\lambda 4959, 5007$ line intensities and the oxygen abundance have been derived empirically using the available oxygen abundances determined via measurement of temperature-sensitive line ratios ($T_{e}$ -- method).  In parallel a new $R_{23}$ calibration has been derived on the base of recent data and compared to previous calibrations. For oxygen-rich HII regions the present $R_{23}$ calibration is close to that of Edmunds and Pagel (1984): their calibration has the same slope but is shifted towards higher oxygen abundances by around 0.07 dex as compared to the present calibration. ",
        "introduction": "The oxygen abundance is one of the fundamental characteristics of a galaxy. The radial distribution of the oxygen abundance in a galaxy (together with radial distributions of gas and star surface mass densities) provides a strong constraint on models of (chemical) evolution. In the general case the intensities of oxygen emission lines in spectra of HII regions depend not only on the oxygen abundance but also on physical conditions. The oxygen abundance can be derived from accurate measurement of temperature-sensitive line ratios, such as $[OIII] \\lambda \\lambda 4959, 5007 / \\lambda 4363$. This method will be referred to as the $T_{e}$ -- method.  Unfortunately, in oxygen--rich HII regions the temperature-sensitive lines like $[OIII] \\lambda 4363$ are too weak to be detected. For such HII regions empirical abundance indicators based on more readily observable lines were suggested (Pagel et al 1979, Alloin et al 1979). The empirical oxygen abundance indicator $R_{23} = (I_{[OII] \\lambda 3727,\\lambda 3729}  + I_{[OIII] \\lambda 4959, \\lambda 5007}) /I_{H\\beta }$ suggested by Pagel et al (1979) has found widespread acceptance and use for the oxygen abundance determination in HII regions where the temperature-sensitive lines are undetectable. This method  will be referred to as the $R_{23}$ -- method. Several workers have suggested calibrations of $R_{23}$ in terms of oxygen abundance (Edmunds \\& Pagel 1984, McCall \\& Rybski and Shields 1985, Dopita \\& Evans 1986, Zaritsky \\& Kennicutt and Huchra 1994, among others). Zaritsky et al's calibration is an average of the three former calibrations.  The oxygen abundances of HII regions in many irregular galaxies are derived with the $T_{e}$ -- method. Those in HII regions of the Milky Way have also been determined with the $T_{e}$ -- method (Shaver et al 1983). This data is a base for many investigations of chemical evolution of our Galaxy. Oxygen abundances in HII regions of many spiral galaxies have been derived with the $R_{23}$ -- method. As indicated by Skillman et al (1996), the precise choice of the O/H -- $R_{23}$ calibration is not critical in differential comparisons of the abundance properties of different objects if the oxygen abundances in all the objects are derived with the same O/H -- $R_{23}$ calibration. However, the comparison of the abundance properties of galaxies where the oxygen abundances were determined with the $T_{e}$ -- method (irregular galaxies, Milky Way  and a few other spiral galaxies) and galaxies where the oxygen abundances were determined with the $R_{23}$ -- method (spiral galaxies) is justified only if the two methods agree. The existing O/H -- $R_{23}$ calibrations (Edmunds \\& Pagel 1984, McCall \\& Rybski and Shields 1985, Dopita \\& Evans 1986, McGaugh 1991) are based on then-available oxygen abundance determinations through the $T_{e}$ -- method and HII region models, whereas more oxygen abundance determinations through the $T_{e}$ -- method are available now. Then the existing O/H -- $R_{23}$ calibrations should be verified in the light of more recent data and a new calibration derived if necessary.  The search for a line intensities -- O/H calibration which results in the same oxygen abundances as the $T_{e}$ -- method is the goal of this study. The line intensities -- O/H calibration has been derived in Section 2. A discussion will be presented in Section 3. Section 4 is a brief summary. ",
        "conclusions": "The problem of line intensities -- oxygen abundance calibration has been considered. It has been obtained that the oxygen abundances derived with the $R_{23}$ -- method involve a systematic error caused by the failure to take into account the differences in physical conditions in different HII regions.  We confirm the idea of McGaugh (1991) that the strong (readily observable) oxygen lines ($[OII] \\lambda \\lambda 3727, 3729$ and $[OIII] \\lambda \\lambda 4959, 5007$) contain the necessary information for determination of accurate abundances in low-metallicity (and may be also in high-metallicity) HII regions. It has been found that the excitation parameters $p_{3}$ or $p_{2}$ (which are defined here as contributions of the radiation in $[OIII] \\lambda \\lambda 4959, 5007$ lines and in $[OII] \\lambda \\lambda 3727, 3729$ lines to the \"total\" oxygen radiation respectively) allow to take into account the variations in $R_{23}$ values among HII regions with a given oxygen abundance. Based on this fact a new way of the oxygen abundance determination in HII regions (p -- method) has been constructed and corresponding relations between $[OII] \\lambda  \\lambda 3727, 3729$, $[OIII] \\lambda  \\lambda 4959, 5007$ line intensities and the oxygen abundance have been derived empirically using the available oxygen abundances determined via measurement of temperature-sensitive line ratios ($T_{e}$ -- method).  By comparing of oxygen abundances in HII regions derived with the $T_{e}$ -- method and those derived with the p -- method it has been found that for the low-metallicity HII regions the exactness of oxygen abundance determination with the p -- method is comparable to that obtained with the $T_{e}$ -- method. For the low-metallicity HII regions the p -- method provides a more robust diagnostic of oxygen abundance than the $R_{23}$ -- method.  For oxygen-rich HII regions both p and $R_{23}$ calibrations are less reliable. For the majority of HII regions in the present compilation  the differences between oxygen abundances derived with the p -- method and with the $R_{23}$ -- method are appreciably less than between those derived with the $R_{23}$ -- method (or the p -- method) and those from the $T_{e}$ -- method. In the general case, however, two HII regions can have the same $(O/H)_{R_{23}}$ while their $(O/H)_{P_{2}}$ can differ by $\\sim$ 0.3 dex.  The calibration of $R_{23}$ presented here is compared to previous calibrations. The calibration of Edmunds and Pagel (1984) has the same slope but is shifted towards higher oxygen abundances by about 0.07 dex as compared to the present calibration. Other previous calibrations (McCall \\& Rybski and Shields 1985, Dopita \\& Evans 1986, Zaritsky et al 1994) are shifted towards still higher abundances.  For the high-metallicity HII regions the exactness of oxygen abundance determination with the suggested p -- method cannot be firmly estimated due to the lack of the high-quality determinations of oxygen abundances with the $T_{e}$ -- method. Indirect arguments suggest that the p -- method provides as accurate oxygen abundances in oxygen-rich HII regions as the $T_{e}$ -- method taking into account the present-day state-of-art with the line intensity measurements."
    },
    "0004/astro-ph0004026_arXiv.txt": {
        "abstract": "The Canadian Network for Observational Cosmology (CNOC2) Field Galaxy Redshift Survey is a spectroscopic/photometric survey of faint galaxies over 1.5 square degrees of sky with a nominal spectroscopic limit of $R_c\\sim21.5$ mag. The primary goals of the survey are to investigate the evolution of galaxy clustering and galaxy populations over the redshift range of $\\sim 0.1$ to $\\sim0.6$. The survey area contains four widely separated patches on the sky with a total sample of over $6000$ redshifts, representing a sampling rate of about 45\\%. In addition, 5-color photometry (in $I_c$, $R_c$, $V$, $B$, and $U$) for a complete sample of approximately 40,000 galaxies to $R_c\\sim23.0$ mag is also available. We describe the survey and observational strategies, multi-object spectroscopy mask design procedure, and data reduction techniques for creating the spectroscopic-photometric catalogs. We also discuss the derivations of various  statistical weights, including corrections for the effects of limited spectral bandwidth, for the redshift sample which allow it to be used as a complete sample. As the initial release of the survey data, we present the full data set and some statistics for the Patch CNOC$\\,$0223+00. ",
        "introduction": "Fundamental to our understanding of the Universe is the formation and evolution of structures, from galaxies to clusters of galaxies to large-scale structures such as sheets, filaments, and voids. Theoretical advances, often in the form of increasingly larger and  more sophisticated N-body simulations, have laid much of the groundwork in interpreting the development and evolution of dark matter clustering (e.g., Davis et al.~1985; Colin, Carlberg, \\& Couchman 1997; and  Jenkins et al.~1998). The modeling of the connection between the clustering of galaxies (which are the most easily observed component) and dark matter (which provides the gravitational field), however, is complex and much less well understood. In essence, deriving such a connection requires a full understanding of galaxy formation and evolution.  On the observational side, progressively larger redshift surveys of galaxies have provided relatively robust measurements of the present-epoch galaxy correlation function or power spectrum (e.g., the CfA survey [Davis \\& Peebles 1983, Vogeley et al. 1992]; and the LCRS [Lin et al. 1996]). Although there has been a number of investigations attempting to measure the evolution of galaxy clustering out to $z\\sim$ 0.5 to 1.0, they were all based on small surveys which were not specifically designed for this purpose (e.g., LeF\\`evre et al. 1996, Shepherd et al. 1997, Carlberg et al. 1997). These surveys cover very small areas with sample sizes of 100's of objects. Furthermore, the substantial galaxy population evolution over this redshift range must be taken into account to ensure that similar samples of galaxies at different redshifts are being used to measure the clustering evolution.  The second Canadian Network for Observational Cosmology (CNOC2) Redshift Survey is the first large redshift survey of faint field galaxies carried out with the explicit goal of investigating the evolution of clustering of galaxies. Such an investigation requires a sample with a large number of galaxies spanning a significant redshift range and covering a sufficiently large area on the sky. The redshift range over which the evolution is to be measured is 0.1 to 0.6, chosen to maximize the efficiency of a 4m class telescope. The survey is also designed specifically to provide a large database of both spectroscopy and multicolor photometry data for the study of the evolution of galaxy populations over the same epochs. This is particularly important as such information will allow us to attempt to disentangle any inter-dependence between the correlation function and galaxy types and evolution.  The current knowledge of field galaxy evolution at $z\\sim0.5$ is based mostly on the measurement of the luminosity function (LF) (e.g., Lilly et al.~1995; Ellis et al.~1996; Lin et al.~1997), with the primary conclusion being that the most rapid evolution is seen in blue, star-forming galaxies. A much larger redshift sample of several thousand galaxies with multi-color photometry will vastly improve on the current results. The multicolor data can provide sufficient information to allow a classification of the spectral energy distributions of the sample galaxies. This, along with a large sample size, will make possible a much more detailed analysis of the evolution of the LF as a function of galaxy population, perhaps enabling a definitive discrimination between luminosity and density evolution.  In this paper, we describe the general strategy and design of the CNOC2 survey, the data reduction methods, and the creation of weight functions. Also presented is the first data catalog from the survey. In Section 2 we present the survey strategy, the field selections and the observations. Sections 3 and 4 describe the data reduction methods for the photometric and spectroscopic data, respectively. The issues of completeness and weights are discussed in Section 5. Finally, Section 6 presents the catalog for the patch CNOC$\\,$0223+00. First results on galaxy population evolution and clustering evolution are presented in papers by Lin et al.~(1999) and Carlberg et al.~(2000), respectively. Reports on close-pair merger evolution and the serendipitous active galactic nuclei sample from the survey are given in papers by Patton et al.~(2000) and Hall et al.~(2000). Additional papers on more detailed studies of galaxy evolution, the dependence of clustering evolution on galaxy types, and other related subjects are in preparation. ",
        "conclusions": ""
    },
    "0004/astro-ph0004210_arXiv.txt": {
        "abstract": "We re-examine the age-velocity dispersion relation in the solar neighborhood using improved stellar age estimates that are based on Hipparcos parallaxes and recent stellar evolution calculations. The resulting relation shows that the Milky Way stellar disk was relatively quiescent, suffering little heating or dispersion increase between 3 and 9 Gyrs.  However, at an age of $9\\pm 1$ Gyrs there is an abrupt increase in the stellar velocity dispersions. To explain the abrupt increase (by almost a factor of 2) in dispersions we propose that the Milky Way suffered a minor merger 9 Gyrs ago that created the thick disk. The quiescent phase is consistent with inefficient heating caused by scattering from tightly wound transient spiral structure. It is possible that many galaxy disks saturated radial velocity dispersions at the point where heating from spiral structure becomes inefficient. Surface brightness profiles can be used to predict spiral arm winding numbers at this saturation point. A comparison between the observed and predicted winding numbers could test this hypothesis and also search for evidence of merger induced disk heating in nearby galaxies. ",
        "introduction": "The observed correlation between the ages and velocity dispersions of solar neighborhood stars has been a subject of study since the work of Spitzer and Schwarzschild (1951) and (1953). Spitzer and Schwarzschild established that scattering of stars from initially circular orbits into more eccentric and inclined orbits was a likely explanation for the increase in velocity dispersion, $\\sigma$, with age $t$. They suggested that massive gas clouds (then undetected) were responsible. Molecular clouds were thought to be the sole scattering agents (e.g., \\cite{mihalas}) until \\cite{lacey84_} showed that the observed ratio of the dispersion in the direction perpendicular to the Galactic plane and that towards the Galactic center, ${\\sigma_z / \\sigma_R}$, was too low to be consistent with the predictions from this scattering process.  There was also a discrepancy between the predicted and then measured relation between $\\sigma$ and $t$: if $\\sigma \\propto t^\\alpha$, models predicted $\\alpha \\sim 0.20$, while observations of the time yielded $\\alpha \\sim 0.5 $ (\\cite{wielen}). This resulted in the development of models that incorporated the heating of the stellar disk from transient spiral structure (\\cite{barbanis}; \\cite{sellwood}; \\cite{carlberg}) in addition to scattering from molecular clouds (\\cite{jenkins}; \\cite{jenkins92}).  The measured value of $\\alpha$ depended on stellar age estimates that were necessarily highly uncertain. However, parallaxes from the Hipparcos satellite, along with improvements in stellar evolution models for low-mass stars, have made it possible in the past few years to estimate the ages of stars in the solar neighborhood more precisely than previously possible (e.g., \\cite{ng}). Imprecise ages not only add scatter to a plot of stellar dispersion versus age but also smooth over any underlying structure in this relation that might actually exist in the solar neighborhood. The theoretical work by \\cite{lacey84_} and \\cite{jenkins_} was motivated in part by the smooth form of the age dispersion relation advocated by \\cite{wielen_}.   % However \\cite{stromgren_} and \\cite{freeman} suggested that the observations were more consistent with little or no heating between 3 and 10 Gyrs and an abrupt increase in the dispersion occurring at about 10 Gyrs. \\cite{freeman} suggested that disk heating mechanisms saturated or became inefficient at $\\sigma  \\sim 30 $ km s$^{-1}$ and that another process, such as a merger or large scale bar, would be required to explain the high dispersions seen in the `thick disk' or in stars with ages greater than 10 Gyrs. The thick disk was proposed to explain an excess in the number of stars compared to predictions from an exponential disk model (\\cite{thickdisk}; \\cite{kuijken}). However, the discreteness of this component compared to Population I stars in age, metallicity or dispersion has been difficult to establish (e.g., \\cite{milkywaybook}). Stellar evolution models for the solar neighborhood have supported both forms for the age dispersion relation. \\cite{larsen} found that a jump at 8 Gyrs in the age dispersion relation was required to fit the stellar dispersions and other observational constraints.  However, \\cite{binney_} fit the color distribution of stars observed with Hipparcos from an evolution model with $\\sigma \\propto t^{0.33}$. In this paper we use improved age estimates to re-examine the age-dispersion correlation.  We then compare the resulting relations with the existing theoretical framework.  \\subsection{Age and velocity component measurements for stars in the solar neighborhood}  We have compiled new space velocity components for solar neighborhood stars from the metallicity study of \\cite{edvardsson93}. \\cite{edvardsson93} obtained metallicity and kinematic information for 189 nearby F and G star dwarf stars within 80 parsecs of the Sun. They  selected stars for metallicity and effective temperature based on Str\\\"omgren indices. Because the stars were selected from the magnitude-limited photometric catalogue of \\cite{olsen}, the sample is expected to be unbiased kinematically (\\cite{freeman_}). For the stars in this sample we computed velocity components U, V, and W from parallaxes and proper motions measured by the Hipparcos satellite and radial velocities obtained from SIMBAD. In most cases our velocity components were within 1 km s$^{-1}$ of those listed by \\cite{edvardsson93} which were derived from ground-based proper motions and distances. Eight stars, (HD 17548,  18768,  68284, 78747,  98553,  144172,  155358, and 218504) listed by \\cite{edvardsson93} have no published radial velocities, so for these stars we use the velocity components listed by \\cite{edvardsson93}. For the purposes of this study, we excluded the eleven stars % labeled as new spectroscopic binaries, which could introduce spurious scatter into the velocity dispersions.  New ages for the \\cite{edvardsson93} stars were taken from Tables 5 and 6 of \\cite{ng_}. These ages are derived in the same manner as those of \\cite{edvardsson93}, that is, by comparing each star's position in the HR diagram with theoretical stellar evolution tracks, but are based on Hipparcos parallax measurements and the isochrones of \\cite{bert94}, as opposed to the isochrones from \\cite{vandenberg85} used by \\cite{edvardsson93}. The precision in derived ages depends on how far a given star lies from the main sequence. Stars close to the main sequence have larger uncertainties because of the smaller spacing between evolution tracks in the HR diagram. The uncertainty is generally between 0.05 and 0.3 dex, though there are cases where the uncertainty is particularly large; for example, for HR 5459 ($\\alpha$ Cen A) the \\cite{edvardsson93} and the \\cite{ng_} ages differ by a factor of two! Another example, HD 148816, was assigned an age of 13.5 Gyr by \\cite{edvardsson93}, but was assigned an age of only 5.3 Gyr by \\cite{ng_}. Its V velocity of $-$220 km s$^{-1}$ suggests that HD 148816 is actually a halo star which would be consistent with the larger age. Nevertheless, we exclude HD 148816 from our analysis because of the uncertainty; if included at the \\cite{ng_} age, it increases the velocity dispersion in the log age = 0.7--0.8 bin from 42 km s$^{-1}$ to about 70 km s$^{-1}$. For this reason we excluded two additional stars for which the \\cite{ng_} and \\cite{edvardsson93} ages differed by more than 0.3 dex.  Although the actual uncertainties (particularly systematic errors) in the stellar ages are difficult to quantify, these ages are undoubtedly a major improvement over the ages estimated during the 1970s (e.g., \\cite{wielen}) which typically used mean ages for groups of stars of similar spectral types.  Figure 1 shows the complete set of U, V, and W velocities for all stars in our sample (where U is defined as positive toward the Galactic center, V is positive in the direction of Galactic rotation, and W is positive toward the North Galactic Pole).  Figure 2 shows the total and U, V, W velocity dispersions derived from these data as a function of stellar age (values are given in Table 1 for each age bin). The error bars are estimated from the scatter of the stars in each bin and so represent the statistical error in the velocity dispersions only. A comparison of the bottom panel of Figure 1 with the corresponding Figure 16b of \\cite{edvardsson93} and the similar figure of \\cite{freeman} shows that the we find a slight, but noticeable reduction in the scatter in the W velocities. A similar effect was noted by \\cite{garnett} for the scatter in the age-metallicity relation compared to that found by \\cite{edvardsson93}.  \\subsection{An abrupt increase in velocity dispersion at an age of $9$ Gyrs}  We see from Figure 2 that there is a statistically significant and abrupt increase in all velocity components of the velocity dispersion at a stellar age of $9\\pm 1$ Gyrs. If we use the size of the errorbars on the points to represent our uncertainty then the dispersion at 12 Gyrs is $\\sim 2 \\sigma$ in the individual components and $\\sim 3 \\sigma$ in the total dispersion above the dispersions from 5-9 Gyrs.  Note that the data in Figure 2 show velocity dispersions computed directly from the velocity data, with no weightings. \\cite{wielen_} advocated weighting the velocity differences by the amplitude of the W component, appropriate for stars moving in a harmonic potential. A harmonic potential is appropriate for the low velocity dispersion stars.  For stars which spend most of their orbit above the bulk of the stellar population the weighting should be somewhat larger. We show the results of weighting by $|W|$ in Figure 3.  In general, the results are the same, with the exception that the jump in the U dispersion is reduced.  Nevertheless, the jump in $\\sigma$ at 9 Gyr remains significant. We also note that our derived dispersions for the oldest (thick disk) stars are in excellent agreement with the values determined by \\cite{chibabeers}, ($\\sigma_U$, $\\sigma_V$, $\\sigma_W$) = (46$\\pm$4, 50$\\pm$4, 35$\\pm$3) km s$^{-1}$.   Figure 4 shows the UV plane distribution of the stars in 4 different age bins and illustrates that there is little difference in the UV plane distribution of the stars with ages between 2.5 and 5 Gyrs and those between 5 and 9 Gyrs.  We have been unable to imagine a selection effect that could lead to kinematic biases in the sample. Although the stellar selection criteria used by \\cite{edvardsson93} causes the sample to contain a larger fraction of metal-poor stars than is actually the case in the solar neighborhood, contamination by halo stars is unlikely because the halo contributes less than 1\\% to the local stellar space density (\\cite{schmidt75}). Only one strong candidate for a halo star (HD~148816) is present in the sample. The \\cite{edvardsson93} sample may also be biased against selection of old, metal-rich stars, but it is not clear that this would result in a bias against an old, low velocity dispersion component if it exists.  We therefore suspect that there is a genuine increase in the dispersion at an age of about 9 Gyrs, which we associate with the onset of the thick disk.  We can see from the figures that this increase happened on a timescale less than a few Gyrs. An abrupt large increase is inconsistent with scattering models involving either molecular clouds or spiral arms. The only process discussed in the literature that could cause such an abrupt and large increase in the velocity dispersion would be a merger that did not completely destroy the disk of the Milky Way (e.g., \\cite{toth}). Numerical simulations (e.g.,  \\cite{walker}; \\cite{huang}; \\cite{velazquez}) suggest that the Milky Way could have survived merging with a smaller galaxy of up to 1/3 of the Milky Way's mass (at that time). Such a merger would result in significant thickening of the disk. Based on the statistics of lopsided galaxies, \\cite{zaritsky} estimate that minor mergers occur  in an $L_*$ galaxy about every 10 Gyrs.  This is consistent with the possibility that the Milky Way suffered such an event $\\sim 10$ Gyrs ago. Cosmological models also predict that for an $L_*$ galaxy the merger rate would have been higher in the past (\\cite{laceyandcole}).   \\subsection{A quiescent disk between 3 and 9 Gyrs}  Our figures show that the velocity dispersion hardly increases between an age of 3 and 9 Gyrs rather than smoothly increasing during this time (as previously proposed; \\cite{wielen}). \\cite{freeman} suggested that the heating mechanisms become inefficient at $\\sigma  \\sim 30 $ km s$^{-1}$ and that another process % is required to explain the high dispersions seen in the `thick disk'.  Heating solely from molecular clouds predicts $\\sigma \\propto t^{0.20}$ (\\cite{lacey84}) and would be consistent with the quiescent disk or minimal increase in dispersion between 3 and 9 Gyrs ago. However the observed dispersion ratio, $\\sigma_z / \\sigma_R$, is too low to be consistent with this model. \\cite{jenkins_} showed that this ratio can be set by the ratio of the heating rate (or radial dispersion increase) caused by the two processes: 1) scattering from spiral structure and 2) scattering from molecular clouds.  To match the observed dispersion ratio the radial heating rate from transient spiral structure must be significantly larger than that from molecular clouds. However, their model that best fit the observations of the time predicted $\\sigma \\propto t^{0.5}$ which is inconsistent with what we see between 3 and 9 Gyrs.  \\cite{jenkins_} assumed that heating from spiral structure was efficient for dispersions up to 70 km s$^{-1}$. As pointed out by \\cite{carlberg_} once the epicyclic amplitude reaches a sizescale similar to that of the spiral arm wavenumbers, heating is much less efficient and $\\sigma \\sim t^{0.2}$. A star with such a large epicyclic amplitude would cross more than one spiral arm as it oscillates radially making it more difficult for the spiral structure to perturb the star. This lower exponent, $\\alpha \\sim 0.2$, is consistent with the nearly constant dispersions between 3 and 9 Gyrs shown in Figure 2 and 3.  The simplest explanation for the observations is that both heating processes are important (setting the ratio of the dispersions as modeled by \\cite{jenkins}) but that heating from spiral structure becomes inefficient once the radial dispersion reaches about 30~km~s$^{-1}$.  The saturation point or maximum dispersion reached can be compared to the limit predicted by the observed wavenumber of spiral structure in the Milky Way. Heating from tightly wound spiral structure becomes inefficient when $k a \\sim 1$ where $k$ is the radial wavenumber of the spiral structure and $a$ is the epicyclic amplitude (\\cite{carlberg}). The epicyclic amplitude is $a = \\sqrt{2 E_R}/\\kappa $ where $E_R$ is the energy in the epicycle, and $\\kappa$ is the epicyclic frequency. For a star with maximum radial velocity $V_R$, $E_R =  1/2 V_R^2$. $\\kappa = \\sqrt{2} v_c/R_0$ for a flat rotation curve with circular velocity $v_c$ at a radius $R_0$ from the Galactic Center. Setting $E_R \\sim \\sigma_R^2$ (corresponding to a typical particle with maximum radial velocity $\\sim \\sqrt{2} \\sigma_R$), we find that $k a \\sim  1$ when $k R_0 \\sim {v_c/\\sigma_R}$. For $v_c \\sim 200$ km s$^{-1}$ and heating diminishing at $\\sigma_R \\sim 30$ km s$^{-1}$, consistent with the dispersions between 3 and 9 Gyrs, heating from transient spiral structure in the solar neighborhood  becomes inefficient for $k R_0 \\gtrsim 7$. This is fairly tightly wound; more tightly wound than that seen in the numerical simulations of \\cite{carlberg87}, $k R \\sim 4.1$, which was adopted as limits by \\cite{jenkins_}. This approximate limit can be directly compared to the winding number of spiral structure observed in the Milky Way.  Tracers of spiral structure in the solar neighborhood suggest that the spiral structure is quite tightly wound (e.g., \\cite{vallee}). A two armed spiral is seen at K band in COBE/DIRBE data with pitch angle in the range $p\\sim 15.5^\\circ$~--~$19^\\circ$ (\\cite{drimmel}), giving $k R_0 = m \\cot{p}$ in the range 5.8~--~7.2. This range is consistent with the approximate limit predicted from the maximum radial dispersion reached between 3 and 9 Gyrs. However, the tightly wound spiral structure observed in the Milky Way is unlikely to efficiently heat the disk to the large velocities observed in the thick disk. Spiral arms with a significantly higher pitch angle of $\\sim 26^\\circ$ (corresponding to $k R_0 \\sim 4$) would be required to do this. We find that the slow increase in stellar dispersions between 3 and 9 Gyrs observed in Figure 2 is consistent with inefficient heating from the tightly wound spiral structure observed in the Milky Way. More detailed calculations are needed to more accurately determine if the slow observed increase in dispersions can be predicted from a model consistent with the observed spiral structure.  \\cite{jenkins92_} found that the model cloud heating rate that fit the dispersions at ages greater than 8 Gyrs (data from \\cite{wielen}) required an order of magnitude larger cloud number density than presently observed.  However, to fit the dispersions for the stars younger than 8 Gyrs (data from \\cite{stromgren}), he found that the heating required was consistent with the present molecular cloud number density. If an additional process is responsible for the largest dispersions at ages greater than 8 Gyrs,  we expect that there would be no discrepancy between the model heating rates and observed dispersions.  The nearly constant dispersion between 3 and 9 Gyrs does not imply that stars that are 8 Gyrs old were heated to 30 km s$^{-1}$ within 3 Gyrs. It may be that the oldest stars were heated initially to a lower value (which might have been more consistent with a less massive disk) then subsequently were slowly heated as the disk became more massive.  This would be consistent with a picture where the Toomre stability parameter, $Q$, remained constant as the disk increased in mass. ",
        "conclusions": "\\cite{freeman} and \\cite{stromgren_} suspected that the age-dispersion relation of solar neighborhood stars was not consistent with the smoothly increasing relation proposed by \\cite{wielen_}. In this paper we re-examine the age-velocity dispersion relation in the solar neighborhood using improved stellar age estimates which are based on Hipparcos parallaxes and recent stellar evolution calculations. We find that there is an abrupt statistically significant jump in all velocity components at age of about 9 Gyrs, confirming the statements made by \\cite{freeman} and \\cite{stromgren_} and modeling done by \\cite{larsen}. While the slow diffusion model proposed by \\cite{wielen_} is consistent with the theoretical models for ages up to about 9 Gyrs, it is unlikely to explain the abrupt increase in dispersions seen at this time.  The most likely explanation is that the Milky Way suffered a minor merger that formed the thick disk.  The best theoretical modeling of the age-dispersion relation (e.g., \\cite{jenkins92}) was able to fit the age-dispersion relation proposed by \\cite{wielen_} but was still inconsistent with the observed number density of molecular clouds. \\cite{jenkins92_} noted that this problem was resolved if he fitted only stars younger than 8 Gyrs. As previously discussed in the literature (\\cite{lacey84}, \\cite{jenkins}) a model incorporating both scattering from molecular clouds and spiral structure is required to match the ratio of the vertical to planar velocity dispersions. However, to be consistent with the nearly constant observed dispersions between 3 and 9 Gyrs a model with more tightly wound spiral structure than assumed by \\cite{jenkins_} is required. In such a model, once the epicyclic amplitude reaches a size scale similar to that of the spiral arm wavenumbers, heating is much less efficient and the velocity dispersion increases extremely slowly. The limit given by the observed plateau in dispersion is consistent with the tracers of spiral structure in the Milky Way which find that the local spiral structure is quite tightly wound (\\cite{vallee}; \\cite{drimmel}).  Since we find little evidence for an increase in velocity dispersion for stars between 3 and 9 Gyrs old, disk heating must occur primarily within 3 Gyrs of birth. However, 3 Gyrs is only about 12 rotation periods at the solar circle. This might suggest that a resonant heating model or bar induced heating model might be more appropriate (e.g., as explored in \\cite{dehnen}) to explain the rapid increase in dispersions in the solar neighborhood rather than a random diffusive scattering model. In such a model, the ratio of vertical to planar dispersions would also be set by the relative strength of the planar heating and vertical heating rates. If the ratio of vertical to planar dispersions and the Toomre stability parameter are nearly constant with radius (as suggested by \\cite{bottema}), then we might expect that spiral structure is more important than the Galactic bar in heating the stellar disk in the solar neighborhood. Alternatively, the short timescale implies that local heating rates could be highly sensitive to the recent history of spiral structure or the Galactic bar.  In this case we might expect strong variations in the heating rates and stellar dispersions as a function of radius.  It is possible that many galaxies have radial velocity dispersions at the maximum value where spiral structure ceases efficient heating of the disk.  Using surface brightness profiles and assumed or measured values for the mass-to-light ratio and Toomre $Q$ parameter, we suggest that the winding number of spiral structure consistent with a disk at this particular dispersion can be predicted. If the winding number of the observed spiral structure is above the predicted value, then the spiral structure is incapable of heating the stars to their estimated radial velocity dispersion. By comparing predicted winding numbers with those observed in a sample of galaxies spanning a range of properties the hypothesis that disks tend to lie near their maximum dispersions can be tested. This type of study could also provide evidence for past mergers in other galaxies. Indeed the large scatter in vertical scale heights seen in edge-on galaxies and tendency of interacting galaxies to have larger vertical scale heights (\\cite{schwarzkopf}) suggests that galaxy interactions and mergers may be a dominant cause of vertical heating in disk galaxies.  Much of the controversy concerning the age-dispersion relation is due to the difficulties in measuring accurate ages for a large number of solar neighborhood stars.  We expect that in the next few years, through more comprehensive metallicity studies, and with improved stellar evolution tracks, a larger number of stars from the Hipparcos catalog will have improved ages estimates.  The trends discussed in this paper can then be more thoroughly tested with a larger sample."
    },
    "0004/astro-ph0004160_arXiv.txt": {
        "abstract": "We study the gas emission of galaxies with active star formation, consisting mostly of H{\\sevensize II} galaxies and Starbursts, as well as some Seyfert 2's, and determine chemical and physical parameters. The data consist of 19 high signal-to-noise ratio optical templates, a result of grouping 185 emission line galaxy spectra. Underlying stellar population models (Raimann et al., 1999) were subtracted from the templates in order to isolate the pure emission component.  We analyse the distribution of these improved signal-to-noise ratio emission spectra in diagnostic diagrams and find that the H{\\sevensize II} templates show a smaller spread in $log$([O{\\sevensize III}]/H$\\beta$) values than the individual galaxies, apparently due to the population subtraction and better signal-to-noise ratio. We thus suggest the template sequence as a fiducial observational locus for H{\\sevensize II} galaxies which can be used as reference for models. The sequence of line ratios presented by the H{\\sevensize II} galaxies in the diagram $log($[O{\\sevensize III}]$\\lambda5007$/H$\\beta) \\times log($[N{\\sevensize II}]$\\lambda$6584/H$\\alpha)$ is primarily due to the gas metallicity, of which the $log($[N{\\sevensize II}]/H$\\alpha)$ ratio is a direct estimator. We also study the properties of the Starburst galaxies and those intermediate between H{\\sevensize II} and Starburst galaxies, which are more metal rich and sit on more massive galaxies.  We discuss the present results in the frame of a recently proposed equivalent width diagnostic diagram for emission line galaxies (Rola et al., 1997) and conclude that the observed ranges in W([O{\\sevensize II}])/W(H$\\beta)$ and W(H$\\beta)$ are mostly due to the non-ionizing stellar population contribution. W(H$\\beta)$ can be used as an estimator of this contribution to the continuum, and briefly discuss implications to the cosmological use of H{\\sevensize II} galaxies. ",
        "introduction": "The investigation of nearby star-forming galaxies plays an important role in the interpretation of the ever increasing data on distant galaxies, as the so-called Lyman-break galaxies seem to be well described by local Starburst galaxies (Meurer et al., 1999). Melnick et al. (1999) propose that H{\\sevensize II} galaxies can be used as distance estimators over a wide range of redshifts.  H{\\sevensize II} galaxies are among the less luminous star-forming galaxies. Their emission-line spectra and relatively low gaseous metallicities ($\\approx$ 0.1 to 0.25 solar, e.g. Pe\\~{n}a et al., 1991) would be consistent with the idea that they are very young galaxies undergoing their first episodes of star formation. On the other hand, in a recent study, Schulte-Ladbeck and Crone (1998) have concluded that in the blue compact dwarf galaxy VII Zw 403 there are older stellar population components.  In a previous study (Raimann et al., 1999, hereafter Paper I) we have investigated the stellar population properties of a sample dominated by nearby H{\\sevensize II} galaxies, but including also Starburst and Seyfert 2 galaxies for comparison purposes. We have considered their continuum and emission/absorption line properties, grouped them into high signal-to-noise templates, and performed stellar population syntheses using a base of stellar cluster spectra (Bica, 1988; Bica and Alloin, 1986). We concluded that most H{\\sevensize II} galaxies present important flux contributions from populations as old as 500 Myr.  In the present paper, we investigate the gaseous properties of the templates obtained in Paper I, after subtraction of the underlying stellar population. We use the emission-line fluxes to classify the spectra according to their excitation characteristics in diagnostic diagrams, to calculate the gaseous abundance and to obtain the age of the last generation of (ionizing) stars. The role of the underlying stellar population, including its effect on the interpretation of data from distant star-forming galaxies and its relation to the emission line properties, is also explored.  We present the data set in section 2. The diagnostic diagrams are discussed in section 3. The gas metallicity, the age of the ionizing stellar population and relation between the metallicity and the non-ionizing stellar population are studied in section 4. Possible uses of emission-line galaxies for cosmological studies are discussed in section 5. The concluding remarks are given in section 6. ",
        "conclusions": "We have analysed the emission-line spectra of 19 galaxy templates, obtained from the grouping of 185 emission-line galaxy spectra, as described in Paper I. After correction for internal reddening and subtraction of the synthesized population spectrum, both from Paper I, the emission-line fluxes of each template were measured and corrected for additional gaseous reddening, when present. According to the corresponding locus of each template in the BPT diagrams, the present sample comprises 10 groups of H{\\sevensize II} galaxies, 3 groups of Seyfert 2, 2 groups of Nuclear Starbursts and 4 groups of intermediate cases between Nuclear Starbursts and H{\\sevensize II} galaxies.  The H{\\sevensize II} galaxy groups define a much tighter sequence in the BPT diagrams than the individual galaxies, suggesting that the spread obtained in previous studies is due to a combination of lower S/N ratio spectra and of not taking into account the contribution of the underlying stellar population. With population subtraction in spectra of improved S/N ratio, the H{\\sevensize II} galaxy groups get closer to the theoretical sequences of H{\\sevensize II} regions presented by Evans and Dopita (1985) and McCall et al. (1985), suggesting that they are more similar to H{\\sevensize II} regions than concluded in previous works. The resulting sequence of H{\\sevensize II} templates in the BPT diagram is suggested as a fiducial observational locus for future models of H{\\sevensize II} galaxies.  From the emission-line ratios, we calculate the gas metallicity and age of the ionizing stellar population, and investigate the effect of these two parameters in the BPT diagram above. We conclude that the sequence defined by the H{\\sevensize II} galaxy templates is primarily due to metallicity, which spans the range 7.87 $<$ 12+$log$(O/H) $<$ 8.32 ($\\approx$ 1/11 to 1/4 solar), while the age of the ionizing stellar population ranges only from 2.7 to 5.0 Myr.  A connection with age is suggested when we relate the gas metallicity with the percentage contribution of stellar population components older than 100 Myr, which may indicate that the metal enrichment is mostly due to previous stellar generations, whose signatures are present in the spectral distribution even for the bluest H{\\sevensize II} galaxies, as discussed in Paper I. The larger the contribution of older stellar components, the more metal rich is the gas. We find a good correlation between $log($[N{\\sevensize II}]/H$\\alpha)$ and $12+log$(O/H) and propose a calibration to obtain the latter from the former.  We also explore the effect of the stellar population contribution to the equivalent width diagnostic diagrams of Rola et al. (1997). We conclude that the observed ranges in W([O{\\sevensize II}])/W(H$\\beta)$ and W(H$\\beta)$ are essentially due to the non-ionizing stellar population contribution. By relating $log($W(H$\\beta))$ to the percentage contribution of this population, we conclude that there is a tight correlation between these two quantities. We thus propose a calibration which can be used to estimate the  non-ionizing stellar population contribution to the spectra from the measured W(H$\\beta)$, which is particularly useful for probing properties of distant galaxies."
    },
    "0004/hep-ph0004043_arXiv.txt": {
        "abstract": "In recent work, it has been argued that multi-TeV masses for scalar superpartners are not unnatural. Indeed, they appear to have significant phenomenological virtues.  Here we explore the implications of such `focus point' supersymmetry for the dark matter problem.  We find that constraints on relic densities do not place upper bounds on neutralino or scalar masses. We demonstrate that, in the specific context of minimal supergravity, a cosmologically stable mixed gaugino-Higgsino state emerges as an excellent, robust dark matter candidate.  We estimate that, over a wide range of the unknown parameters, the spin-independent proton-neutralino cross sections fall in the range accessible to planned search experiments. ",
        "introduction": "\\label{sec:introduction}  The idea that supersymmetry is an approximate symmetry of Nature, broken at a `low' scale, is best and most concretely motivated by two general arguments.  First, it addresses the question of why radiative corrections to the electroweak symmetry breaking scale, due to the exchange of virtual particles at some much larger unification or Planck scale, do not pull the electroweak scale to a much higher value (the gauge hierarchy problem~\\cite{SM}).  Second, it brings the unification of couplings calculation, which does not quite work if we use only the particles of the Standard Model, into adequate agreement with experiment~\\cite{DRW}.  However neither of these arguments, on the face of it, yields very specific constraints on the spectrum of masses of the supersymmetric ($R$-parity odd) partners of ordinary ($R$-parity even) particles: the first because it is inherently semi-quantitative, the second because it depends on these masses only logarithmically.  On the other hand, it is becoming notorious that the existence of superpartners with masses below 1 TeV is difficult to reconcile with established limits on various flavor violating effects, unless one assumes accurate degeneracy of squarks and of sleptons.  Even if one assumes such degeneracy, supersymmetric contributions to $CP$ violation and (in unified theories) proton decay become uncomfortably large, unless the squarks and sleptons are heavy.  Therefore one is motivated to examine supersymmetric models with large scalar superpartner masses.  The issue immediately arises: To what extent can such models, when critically scrutinized, embody the motivating features we mentioned in our opening paragraph?  Recently it has been argued, using a seemingly reasonable objective definition of naturalness, that all squark and slepton masses can be taken well above 1 TeV with no loss of naturalness~\\cite{FMMshort,FMMlong}.  The mathematical basis of this result is the existence of focus points in renormalization group trajectories, which render the weak scale ({\\em i.e.}, the Higgs potential) largely insensitive to variations in unknown supersymmetry parameters~\\cite{FMMshort,FMMlong,FMMproc}.  We should note that while in these models the squark and slepton masses are unusually large (compared to conventional wisdom), the electroweak gaugino and Higgsino particle masses are generically well below 1 TeV.  In the focus point models, moreover, the quantitative aspect of the unification of coupling constants is actually slightly improved relative to more traditional supersymmetry models, reducing the need for significant high scale threshold corrections~\\cite{FMinprep}.  Besides preserving and sharpening the original motivations for supersymmetry, the focus point models with heavy scalars ameliorate the problems noted above~\\cite{FMinprep}.  Specifically, constraints on $CP$ violation from electric dipole moments can be satisfied even with ${\\cal O} (1)$ phases in the fundamental parameters, and the predicted rate for nucleon decay --- being proportional to the square of gaugino and inversely proportional to the fourth power of squark/slepton masses --- is suppressed and less dangerous, though still perhaps accessible.  In this note we will consider the implications of focus point supersymmetry for neutralino dark matter.  We will concentrate on the focus point mechanism in its simplest `default' incarnation, assuming a universal scalar mass within minimal supergravity. The focus point mechanism can also be realized in alternative models featuring gravity-~\\cite{FMMshort,FMMlong}, gauge-~\\cite{A}, and anomaly-mediated~\\cite{FM} supersymmetry breaking.  In addition, the requirement of a universal scalar mass, though sufficient, is not necessary.  For the precise set of requirements, see Ref.~\\cite{FMMlong}.  Of course, the cosmology of minimal supergravity (with universal scalar mass $m_0 \\leq 1 \\tev$) has been studied extensively. Prominent among conclusions one finds in the literature are the following:  \\bigskip  \\noindent (i) In essentially all of parameter space, the lightest neutralino $\\chi$ is Bino-like, with gaugino fraction $R_{\\chi} > 0.9$~\\cite{Snowmass}. (See below for the definition of $R_{\\chi}$.)  \\noindent (ii) The neutralino relic abundance $\\Omegachi h^2$ increases as $m_0$ increases.  The requirement that neutralinos not overclose the universe, along with the requirement that the lightest supersymmetric particle (LSP) be neutral, then typically leads to stringent upper bounds on scalar and/or gaugino masses, independent of naturalness considerations~\\cite{OS,GKT,LYN,KKRW,BB,AN3,EFO,EFOS,AC,EGO}.  \\noindent (iii) Dark matter detection rates decrease as $m_0$ increases.  The rates predicted range widely over parameter space, but the prospects for discovery at ongoing experiments grow increasingly dim as $m_0$ increases~\\cite{ANdetection}.  \\bigskip  In our analysis, we confirm these conclusions --- for $m_0 \\leq 1 \\tev$.  However, we find that they fail for focus point models with $m_0 > 1 \\tev$.  Instead, we find that in such models:  \\bigskip  \\noindent (i) The lightest neutralino is a gaugino-Higgsino mixture over much of parameter space.  \\noindent (ii) Neutralino relic abundances do not overclose the universe.  Rather, they tend to lie in the cosmologically interesting range.  Therefore, in particular, cosmological considerations alone do not place stringent upper limits on superpartner masses, even apart from possible conspiracies involving poles~\\cite{DNOmega,DNRY} or co-annihilation~\\cite{EFO,Dreesstop}.  \\noindent (iii) Direct detection rates may be fairly large, with proton-neutralino cross sections plausibly falling within the range $10^{-6}$-$10^{-8}$ pb.  (See the more detailed discussion and plots below.)  \\bigskip  These results are especially interesting since ongoing and planned dark matter detection experiments promise to cover the indicated range. ",
        "conclusions": "\\label{sec:discussion}  Several ongoing and planned experiments, using a variety of techniques, are devoted to the search for WIMP dark matter~\\cite{Marina}.  Assuming that the matter-WIMP interactions are dominated by spin-independent couplings, these experiments may be compared by scaling the matter-WIMP cross sections for each detector material to the proton-WIMP cross section $\\sigma_P$ and displaying the sensitivity in the $(m_{\\chi}, \\sigma_P)$ plane.  In Fig.~\\ref{fig:mchi0_sigma} we plot predictions of focus point supersymmetry in the $(m_{\\chi}, \\sigma_P)$ plane.  We find that, for a given $m_{\\chi}$, cross sections are typically maximized for large $m_0$, and large cross sections near current sensitivities are possible.  At present, the DAMA Collaboration has reported data favoring the existence of a WIMP signal in their search for annual modulation~\\cite{DAMA}.  When WIMP velocity uncertainties are included~\\cite{BR}, the preferred range of parameters is $10^{-6}\\ \\text{pb} \\alt \\sigma_P \\alt 10^{-5}\\ \\text{pb}$ and $30 \\gev \\alt m_{\\chi} \\alt 200 \\gev$ ($3\\sigma$ CL).  The DAMA result has been criticized~\\cite{DAMAcomments}, and a recent analysis of the CDMS Collaboration excludes the $3\\sigma$ DAMA region at $> 84 \\%$ CL~\\cite{CDMS}.  While the situation is at present unclear, in the near future both experiments will improve their sensitivities significantly.  It is noteworthy, in any case, that current experiments are on the verge of probing the region of parameter space arising in focus point models.  \\begin{figure}[t] \\postscript{mchi0_sigma.ps}{0.74} \\caption{Points in minimal supergravity parameter space in the $(m_{\\chi}, \\sigma_P)$ plane.  The parameters scanned, symbols, and assumptions are as in Fig.~\\ref{fig:omega_sigma}.} \\label{fig:mchi0_sigma} \\end{figure}  It is also interesting to compare the discovery potentials of dark matter searches and future colliders in focus point models.  In the focus point scenario, all squarks and sleptons are beyond the reach of the Tevatron and may present significant challenges even for the LHC. However, the lightest Higgs boson has mass $m_h \\alt 120 \\gev$, and so may be discovered in Run II of the Tevatron. The correlation between $m_h$ and $\\sigma_P$ is presented in Fig.~\\ref{fig:mh_sigma}.  We note that focus point models provide a natural setting for Higgs masses above the current bound, as the large squark masses raise the Higgs mass through their radiative corrections. At the same time, both Higgs and dark matter searches are promising, and will confront the focus point scenario over the next few years.  \\begin{figure}[tb] \\postscript{mh_sigma.ps}{0.74} \\caption{Points in minimal supergravity parameter space in the $(m_h, \\sigma_P)$ plane.  The parameters scanned, symbols, and assumptions are as in Fig.~\\ref{fig:omega_sigma}.} \\label{fig:mh_sigma} \\end{figure}  It has been argued that the regions of parameter space which give large dark matter detection rates may also predict large rates for other processes, such as proton decay~\\cite{protondecay} and $b\\to s \\gamma$~\\cite{bsg}.  Conversely, present bounds from such processes may exclude the largest possible rates for dark matter detection.  The proton decay rate, as noted in the introduction, scales as $m_{\\tilde{q}}^{-4}$, so it is relatively suppressed in the focus point models.  The process $b \\to s\\gamma$ provides a stringent constraint on large $\\tb$ only for one sign of $\\mu$.  For the positive sign (the one adopted in our plots), destructive interference between the chargino and charged Higgs diagrams typically leads to predictions within the experimental bounds.  Focus point scenarios can also imply interesting rates for anti-matter detectors and experiments sensitive to dark matter annihilation in the cores of the earth or sun~\\cite{Marina}.  Finally, of course, neutralinos and charginos themselves are targets for discovery at colliders.  If they are found, detailed measurements of their masses and compositions could then be made, and might even be used to exclude neutralinos as dark matter candidates~\\cite{collider}. However, colliders will never be able to distinguish unstable particles with long decay lengths from those stable on cosmological time scales.  Since some well-motivated models with low energy supersymmetry breaking predict LSPs which are stable on accelerator detector, but unstable on cosmological, scales, in principle identification of neutralinos as the dark matter of our universe cannot be made at accelerators.  Ultimately this identification requires detection, either directly or indirectly, of neutralinos permeating space."
    },
    "0004/astro-ph0004095_arXiv.txt": {
        "abstract": "This paper analyzes single-stellar-population (SSP) equivalent parameters for 50 local elliptical galaxies as a function of their structural parameters.  The galaxy sample is drawn from the high-quality spectroscopic surveys of Gonz\\'alez (1993) and Kuntschner (1999).  The basic data are central values of SSP-equivalent ages, $t$, metallicities, \\z, and ``enhancement'' ratios, \\enh, derived in Paper I, together with global structural parameters including velocity dispersions, radii, surface brightnesses, masses, and luminosities.  The galaxies fill a two-dimensional plane in the four-dimensional space of \\z, $\\logt$, $\\log\\sigma$, and \\enh.  SSP age, $t$ and velocity dispersion, $\\sigma$, can be taken as the two independent parameters that specify a galaxy's location in this ``hyperplane.'' The hyperplane can be decomposed into two sub-relations: (1) a ``\\zp-plane,'' in which \\z\\ is a linear function of $\\log\\sigma$ and $\\logt$; and (2) a relation between \\enh\\ and $\\sigma$ in which \\enh\\ is larger in high-$\\sigma$ galaxies.  Velocity dispersion is the only structural parameter that is found to modulate the stellar populations; adding other structural variables such as $I_e$ or $r_e$ does not predict \\z\\ or \\enh\\ more accurately.  Cluster and field ellipticals follow the same hyperplane, but their $(\\sigma,t)$ distributions within it differ.  Most Fornax and Virgo cluster galaxies are old, with a only a small sprinkling of galaxies to younger ages.  The field ellipticals span a larger range in SSP age, with a tendency for lower-$\\sigma$ galaxies to be younger.  The present sample thus suggests that the distribution of local ellipticals in the $(\\sigma,t)$ plane may depend on environment. Since the $(\\sigma,t)$ distribution affects all two-dimensional projections involving SSP parameters, many of the familiar scaling laws attributed to ellipticals may also depend on environment.  Some evidence for this is seen in the current sample.  For example, only Fornax ellipticals show the classic mass-metallicity relation, whereas other sub-samples do not.  The tight Mg--$\\sigma$ relations of these ellipticals can be understood as two-dimensional projections of the metallicity hyperplane showing it edge on.  At fixed $\\sigma$, young age tends to be offset by high \\z, preserving Mg nearly constant.  The tightness of the Mg--$\\sigma$ relations does not necessarily imply a narrow range of ages at fixed $\\sigma$.  Although SSP parameters are heavily weighted by young stars, modeling them still places tight constraints on the total star formation history of elliptical galaxies.  The relation between \\enh\\ and $\\sigma$ is consistent with a higher effective yield of Type II SNe elements at higher $\\sigma$.  This might occur if the IMF is enhanced in massive stars at high $\\sigma$, or if more SNe II-enriched gas is retained by deeper galactic potential wells.  Either way, modulating Type II yields vs. $\\sigma$ seems to fit the data better than modulating Type Ia yields.  The \\zp-plane is harder to explain and may be a powerful clue to star formation in elliptical galaxies if it proves to be general.  Present data favor a ``frosting'' model in which low apparent SSP ages are produced by adding a small frosting of younger stars to an older ``base'' population (assuming no change in $\\sigma$).  If the frosting abundances are close to or slightly greater than the base population, simple two-component models run along lines of constant $\\sigma$ in the \\zp-plane, as required. This favors star formation from well-mixed pre-enriched gas rather than unmixed low-metallicity gas from an accreted object. ",
        "introduction": "The star formation histories of elliptical galaxies, once thought to be very simple---old and metal-rich (\\cite{Baade63})---have come under increasing scrutiny in the last three decades (e.g., \\cite{ST71}; \\cite{Faber73}, 1977; \\cite{O'Connell76}, 1980; \\cite{Pickles85}; \\cite{P89}; \\cite{SS92}; \\cite{G93}; \\cite{Worthey94}; \\cite{Lee94}; \\cite{Renzini95}, 1998; \\cite{TCB98}; \\cite{K98}; \\cite{J99}). Currently there are two basic models for elliptical galaxy formation: hierarchical clustering of small objects into larger galaxy-sized units with accompanying star formation over time (e.g., \\cite{BFPR84}; \\cite{KWG93}), versus monolithic collapse and star formation in a nearly coeval single early burst (e.g., \\cite{ELS62}; \\cite{Larson74}; \\cite{AY87}).  Measurements of the spectral energy distributions (SEDs) and spectral features of elliptical galaxies can provide a test of these scenarios.  For example, evidence for substantial intermediate-age stellar populations (between 1 and 10 Gyr) might favor hierarchical models, which more naturally have extended star formation over time.  A goal of the present series is to assess the evidence for such intermediate-age populations.  The first paper of this series (\\cite{Paper1}, hereafter Paper I) used Lick absorption-line strengths for a sample of local elliptical galaxies observed by Gonz{\\'a}lez (1993, hereafter G93) to derive single-stellar-population (SSP) equivalent parameters $t$ (age), \\z\\ (metallicity), and \\enh\\ (``enhancement ratio,'' see below). Single-burst model line-strengths by Worthey (1994, hereafter W94) were corrected for the effect of non-solar abundance ratios using theoretical spectral calculations by Tripicco \\& Bell (1995, hereafter TB95).  The resultant SSP ages cover a range of 1 to 18 Gyr (including observational errors), while the ranges in \\z\\ and \\enh\\ are fairly narrow.  These parameters, particularly the ages, are based on the assumption that \\hbeta\\ faithfully traces the mean temperature of the main-sequence turnoff and is not seriously affected by other hot stellar populations.  Evidence supporting this assumption was presented in Paper I.  In deriving single-burst SSP parameters for elliptical galaxies, we do not mean to imply that their star formation histories were actually single bursts.  In fact, our favored ``frosting'' model (Section~\\ref{sec:disc_plane}) involves adding a minority of young stars to an older base population.  Our use of SSP parameters is simply a convenient way of condensing all the presently measured line strength data into just three numbers: light-weighted age, \\z, and \\enh.  For the moment, that is all the observations allow.  It is our hope that SSP parameters will be adopted by those who model the full evolutionary history of elliptical galaxies (e.g., \\cite{AY87}; \\cite{Vazdekis96}; \\cite{TCBMP98}) and that they will serve as as a convenient meeting ground between models and data.  We show below that, even though SSP parameters are heavily influenced by the light of any young stars that may be present, modeling them still places important constraints on the total history of star formation in ellipticals.  This paper explores the central stellar populations of a sample of local elliptical galaxies and develops correlations among them and with parent-galaxy structural parameters.  Many previous works have studied such correlations, but most have focused on \\emph{raw} line strengths.  Only three other studies, to our knowledge, have measured ages (using Balmer lines) and developed correlations based on underlying stellar populations.  Tantalo, Chiosi \\& Bressan (1998) studied the G93 galaxies using models based on the ``Padua'' isochrones.  Their correction for non-solar abundance ratios was approximate, however, leading to systematic errors in derived age, \\z, and \\enh\\ (Paper I).  Kuntschner (1999) studied ellipticals in Fornax using high-quality data, which we add to our sample here.  He found that Fornax ellipticals were mainly old, and also discovered a strong relation between \\enh\\ and $\\sigma$, which we confirm.  J{\\o}rgensen (1999) studied Coma ellipticals using line-strength models by Vazdekis et al.~(1996).  Her conclusions foreshadow ours in many respects, but some seem in retrospect to be the product of observational errors. All three of these papers are discussed in Section~\\ref{sec:sum_compare}.  The outline of this paper is as follows.  A brief review of the G93 and Kuntschner (1999) samples, line-strength data, SSP-equivalent stellar population parameters, and structural parameters is given in Section~\\ref{sec:data}.  Section~\\ref{sec:results} presents the sample distribution in the four-dimensional space spanned by \\z, $\\logt$, $\\log\\sigma$, and \\enh; this proves to be a highly flattened, two-dimensional ``hyperplane'' that in turn consists of two sub-relations, a ``\\zp-plane,'' plus a linear \\enh--$\\sigma$ relation. Section~\\ref{sec:other_views} shows how projections of this hyperplane depend on the distribution of points within it, and thus how the appearance of two-dimensional scaling laws can vary depending on this distribution. Section~\\ref{sec:scaling_laws} illustrates these effects using two classic scaling laws---the mass-metallicity relation and the Mg--$\\sigma$ relation.  Possible evidence for environmental variation in the former is presented.  Sections~\\ref{sec:disc_sz} and \\ref{sec:disc_plane} investigate the origins of the hyperplane.  The \\zp-plane in particular appears difficult to explain and, if it proves general, will place very tight constraints on the history of star formation in local elliptical galaxies.  Section~\\ref{sec:conc} summarizes our findings and conclusions. ",
        "conclusions": "\\label{sec:conc}  The centers of local elliptical galaxies appear to contain quite complex stellar populations.  The present sample of local ellipticals spans a wide range of stellar population parameters, most notably a large range in SSP-equivalent age (especially in, but not limited to, field ellipticals).  Despite their diversity, the central stellar populations of these galaxies are described by a few simple scaling relations.  (1) Abundance parameters \\z\\ and \\enh\\ are specified to high accuracy by SSP-equivalent age, $t$, and central velocity dispersion $\\sigma$; ellipticals thus occupy a ``metallicity hyperplane'' in $(\\z,\\logt,\\log\\sigma,\\enh)$-space.  (2) SSP-equivalent metallicity, \\z, is a function of both $t$ and $\\sigma$ (the ``\\zp-plane'').  At fixed $t$, \\z\\ increases with $\\sigma$; at fixed $\\sigma$, \\z\\ is larger at younger age.  (3) SSP-equivalent enhancement ratio, \\enh, is found to be a monotonically increasing function of $\\sigma$ only, in the sense that adding other structural parameters such as $I_e$ or $r_e$ does not predict either \\enh\\ or \\z\\ more accurately.  Our use of SSP-equivalent parameters is not meant to imply a single-burst origin for elliptical galaxies; in fact, the existence of the \\zp-plane seems to imply that the populations are largely old with a ``frosting'' extending to younger ages.  However, despite the fact that SSP-parameters are not true means but rather likely to be influenced by the light of younger stars, they still place very important constraints on the history of star formation in elliptical galaxies (see below).  We take $\\sigma$ and $t$ as the independent parameters that specify the distribution of galaxies in the hyperplane.  Any variation in this distribution will influence all other two-dimensional projections of SSP parameters, and thus many of the common scaling laws for elliptical galaxies.  Our sample shows a possible difference in the ($\\sigma, t$) distribution with environment---our field ellipticals span a wide range in SSP age, while the Fornax and Virgo ellipticals are generally old.  This results in a significant mass-metallicity trend for the cluster galaxies but not for the field galaxies.  Other correlations between stellar population parameters and structural parameters may also turn out to vary with environment.  The Mg--$\\sigma$ relations are edge-on projections of the metallicity hyperplane.  At a given $\\sigma$, young age is offset by a correspondingly high metallicity, preserving line strength.  The narrowness of the observed Mg--$\\sigma$ relations therefore does not imply a narrow range of ages at fixed velocity dispersion.  A more detailed look at the Mg--$\\sigma$ relations is the subject of a future paper.  Physical models to account for the hyperplane have been considered. The rise in \\enh\\ with $\\sigma$ and the mass-metallicity relation (at fixed $t$) is consistent with a higher effective yield of Type II SNe products at high $\\sigma$.  This trend has several possible explanations, for example, greater retention of outflow-driven gas or a flatter IMF at high $\\sigma$.  The existence of the \\zp-plane is more challenging.  A ``frosting'' scenario is favored, in which young stars are added to an old base population, resulting in a range of SSP-equivalent ages.  With a suitable choice of burst populations, the composite populations can be engineered to lie on lines of constant $\\sigma$ in the \\zp-plane after a few Gyr.  However, to preserve both the \\zp-plane and the \\enh--$\\sigma$ relation requires that abundances in the frosting population must be closely coupled to that of the base population---the metallicity, \\z, of the frosting must be somewhat higher than that of the base population, while the enhancement ratio, \\enh, must be nearly equal.  The frosting scenario therefore seems to favor star formation from gas that was pre-enriched in the same parent galaxy rather than from gas that was accreted in an unrelated merger. However, several merger remnants in the sample do indeed lie nicely on the \\zp-plane, in defiance of this logic.  The present picture of the hyperplane is preliminary and needs to be checked against a better local sample and a wide array of other data. For example, SSP mass-to-light ratios should be compared to dynamical $M/L$ measurements, and \\emph{global} SSP parameters should be analyzed, as they are more indicative of the global star formation history than the central SSP parameters used here.  A further interesting question is whether the color-magnitude relation and other scaling laws might also be near-edge-on projections of the hyperplane, like Mg--$\\sigma$.  Most important, the implications of the frosting model must be developed for lookback observations of distant elliptical galaxies. Many observations of distant cluster ellipticals suggest that their stellar populations formed very early, and this may be consistent with the generally old ages for cluster galaxies found here.  Our field ellipticals do show a wide spread of SSP ages, but we have noted that the sample is not volume-limited, and thus predictions for the evolution of distant field ellipticals cannot yet be drawn.  In short, a great deal more data must be gathered and reconciled before we can claim a solid understanding of the star formation histories of elliptical galaxies."
    },
    "0004/astro-ph0004380_arXiv.txt": {
        "abstract": "We present far-ultraviolet spectra of the Seyfert 1 galaxy Mrk~509 obtained in 1999 November with the Far Ultraviolet Spectroscopic Explorer (FUSE). Our data span the observed wavelength range 915--1185 \\AA\\ at a resolution of $\\sim20~\\kms$.  The spectrum shows a blue continuum, broad \\ovi\\ $\\lambda\\lambda1032,1038$ emission, and a broad \\ciii\\ $\\lambda977$ emission line. Superposed on these emission components, we resolve associated absorption lines of \\ovi\\ $\\lambda\\lambda1032,1038$, \\ciii\\ $\\lambda977$, and Lyman lines through L$\\zeta$.  Seven distinct kinematic components are present, spanning a velocity range of $-440~\\rm to +170~\\kms$ relative to the systemic velocity. The absorption is clustered in two groups, one centered at $-370~\\rm \\kms$ and another at the systemic velocity.  The blue-shifted cluster may be associated with the extended line emission visible in deep images of Mrk~509 obtained by Phillips et al.  Although several components appear to be saturated, they are not black at their centers. Partial covering or scattering permits $\\sim 7$\\% of the broad-line or continuum flux to be unaffected by absorption. Of the multiple components, only one has the same ionization state and column density as highly ionized gas that produces the \\ovii\\ and \\oviii\\ ionization edges in X-ray spectra of Mrk~509. ",
        "introduction": "Mrk 509 straddles the boundary in luminosity between Seyfert 1 nuclei and QSOs with an absolute magnitude of $M_B = -22.0$.  This gives it particular importance for understanding how the properties typical of the well-studied nearby active galactic nuclei (AGN) might scale with luminosity, and, therefore, be applied to the more distant and luminous QSOs. Like many Seyfert 1 galaxies, Mrk~509 exhibits intrinsic UV absorption lines (\\cite{York84}; \\cite{Crenshaw95}; \\cite{Savage97}; \\cite{Crenshaw99}) and X-ray absorption edges of \\ovii\\ and \\oviii\\ (\\cite{Reynolds97}; \\cite{George98}). The UV absorption is blue-shifted relative to the systemic velocity of 10,365 \\kms\\ (\\cite{Phillips83}), and Phillips et al.\\markcite{Phillips83} suggest that it is related to the expanding shell of high ionization gas visible in images and spectra out to a radius of 15\\arcsec\\ from the nucleus.  \\begin{figure*}[t] \\plotfiddle{\"fig1.ps\"}{2.55 in}{-90}{77}{77}{-300}{435} \\caption{ FUSE spectrum of Mrk 509 (binned by 20 pixels) is shown as a thin black line. Hopkins Ultraviolet Telescope (HUT) data obtained during the Astro-2 mission in 1995 March (scaled by a factor of 0.69) are shown as the green curve. \\vskip -14pt \\label{fig1.ps}} \\end{figure*}  At high resolution, UV absorption in Seyferts appears to be kinematically complex (NGC~3516: \\cite{Crenshaw98}; NGC~5548: \\cite{Mathur99}), with at most one of the UV components having any possible relation to the X-ray absorbing gas (\\cite{Mathur99}).  High resolution observations of UV absorption lines covering a range of ionization states is the key to determining which, if any, of the UV absorbers are associated with an X-ray warm absorber. The FUSE observations of Mrk~509 we present in this paper extend the far-UV spectral coverage to wavelengths shortward of 1200 \\AA. Our spectra include the \\ovi\\ $\\lambda\\lambda1032,1038$ resonance doublet and the high-order Lyman lines, down to the redshifted Lyman limit. Since the L$\\alpha$ lines resolved in earlier UV observations are saturated, the high-order lines visible with FUSE provide a better constraint on the total neutral hydrogen column density. The \\ovi\\ doublet is a crucial link for establishing a connection between the higher ionization absorption edges seen in the X-ray and the lower ionization absorption lines seen in earlier UV observations. Based on these new observations, we discuss the implications for the location of the absorbing gas in AGN, and how intrinsic absorption in low-$z$ AGN may be related to the broad-absorption line phenomenon in more luminous QSOs. ",
        "conclusions": "The far-UV spectrum of the Seyfert 1 galaxy Mrk 509 shows bright broad \\ovi\\ $\\lambda\\lambda$1032,1038 emission as well as broad \\ciii\\ $\\lambda$977 emission. Kinematically complex intrinsic absorption in Mrk 509 shows at least 7 distinct components in the Lyman lines, \\ovi\\ $\\lambda\\lambda$1032,1038, and \\ciii\\ $\\lambda$977. Although many of the \\ovi\\ and Lyman line components appear to be saturated, they are not black, implying that partial covering or scattering affects the absorption. Only one of the intrinsic absorption components in Mrk 509 is likely to be associated with the warm X-ray absorbing gas. Component 5 (near the systemic velocity) has an ionization state and column density that is in reasonable agreement with the \\ovii\\ and \\oviii\\ absorption edges seen in the ASCA X-ray spectrum. The high resolution of FUSE and its sensitivity in the \\ovi\\ band make it an ideal tool for identifying high-ionization UV absorbers that may correspond to X-ray warm absorbers."
    },
    "0004/astro-ph0004349_arXiv.txt": {
        "abstract": "We investigate the evolution of the metallicity of the intergalactic medium (IGM) with particular emphasis on its spatial distribution. We propose that metal enrichment occurs as a two step process. First, supernova (SN) explosions eject metals into relatively small regions confined to the surroundings of star-forming galaxies. From a comprehensive treatment of blowout we show that SNae by themselves fail by more than one order of magnitude to distribute the products of stellar nucleosynthesis over volumes large enough to pollute the whole IGM to the metallicity levels observed. Thus, a additional (but as yet unknown) physical mechanism must be invoked to mix the metals on scales comparable to the mean distance between the galaxies which are most efficient pollutants. From this simple hypothesis we derive a number of testable predictions for the evolution of the IGM metallicity. Specifically, we find that: {\\it (i)} the fraction of metals ejected over the star formation history of the universe is about 50\\% at $z=0$; that is, approximately half of the metals today are found in the IGM; {\\it (ii)} if the ejected metals were homogeneously mixed with the baryons in the universe, the average IGM metallicity would be $\\langle Z \\rangle=\\Omega_Z^{ej}/\\Omega_b \\simeq 1/25 Z_\\odot$ at $z=3$. However, due to spatial inhomogeneities,  the mean of the distribution of metallicities in the diffusive zones has a wide (more than 2 orders of magnitude) spread around this value; {\\it (iii)} if metals become more uniformly distributed at $z\\simlt 1$, as assumed, at $z = 0$ the metallicity of the IGM is narrowly confined within the range $Z \\approx 0.1 \\pm 0.03 Z_\\odot$. Finally, we point out that our results can account for the observed metal content of the intracluster medium. ",
        "introduction": "Primordial nucleosynthesis enriched gas in the universe with the light elements He, D, and Li. It is only when the first galaxies and their stars appeared that heavier elements could be synthesized and, in some cases, ejected into the intergalactic medium. In currently popular models of galaxy formation based on hierarchical clustering the first galaxies to form were low mass systems with such shallow potential wells that a few supernovae could deposit sufficient kinetic energy to expel the entire interstellar medium of the galaxy (Ferrara 1998). In this way, these initial episodes of star formation in the universe (sometimes referred to as Population III) may have enriched the intergalactic medium (IGM) to an average metallicity $Z_{\\rm IGM} \\approx 10^{-4} Z_\\odot$ (Miralda-Escud\\'e \\& Rees 1997, Nath \\& Trentham 1997, Gnedin \\& Ostriker 1997, Ciardi \\etal 2000b), comparable to that of the most metal-poor stars in the halo of our Galaxy (Ryan, Norris, \\& Beers 1996).  Most of the metals in the universe were presumably produced in larger collapsed structures at redshifts $z \\simlt 10$. In general such galaxies were better able to retain the products of stellar nucleosynthesis and could therefore achieve the abundance levels observed today. It is also likely that some metal-enriched gas escaped into the IGM, but it is far from clear by which process and to what degree the metals became distributed over large volumes, far from their production sites. These are the questions we consider in the present paper. Recent numerical hydrodynamic simulations of large-scale structure formation (Hellsten \\etal 1997; Rauch \\etal 1997; Zhang \\etal 1998) have shown that baryons in the universe are distributed in a network structure (the ``cosmic web'') with galaxies located at the high overdensity peaks. This implies that most of the volume is occupied by voids filled with gas at, or below, the mean cosmic density; such gas can be viewed as the long-searched-for true intergalactic medium. Since the voids are very large, with typical dimensions of several Mpc, their pollution by heavy elements produced in SN explosions is far from a trivial problem. In any case there must have been a protracted epoch when the metal content of the IGM was highly patchy.  Observationally, our best view of the IGM is still provided by the \\lya\\ forest in the spectra of distant QSOs (Sargent et al. 1980) and the detection of metals in the forest ranks as one of the most significant discoveries made possible by the Keck telescopes (Cowie \\etal 1995; Tytler \\etal 1995).  From an analysis of these data Hellsten et al. (1997) and Rauch, Haehnelt, \\& Steinmetz (1997) concluded that the measured column density ratios $N$(C~IV)/$N$(H~I) imply that typically [C/H]~$ \\simeq -2.5$ at $z \\simeq 3$, with a one order of magnitude dispersion in the metallicity of different clouds about this mean value.\\footnote{In the usual notation, [C/H] = log (C/H) $-$ log (C/H)$_{\\odot}$.} At lower redshifts ($z=0.3-0.8$) Barlow \\& Tytler (1998) using HST/FOS spectra, with reasonable assumptions concerning the ionization correction and the line clustering properties, conclude that metallicities are as high as [C/H]$\\simgt -1.3$, roughly an order of magnitude larger than the value at $z=2.5$. These measurements, however, still refer to overdense regions of the universe, traced by \\lya\\ clouds with column densities in excess of log~$N$(H~I) = 14.5\\,. The situation is far less clear-cut when we turn to the voids---or log~$N$(H~I)$ < 14.0$---where the observations are very challenging even with a 10-m telescope. Two studies have addressed this problem, with conflicting results. Lu \\etal (1998) applied a stacking technique to nearly 300 C~IV regions in QSO spectra but still found no composite signal. They interpreted this non-detection as evidence for a highly non-uniform degree of metal enrichment at $z \\simeq 3$, with the voids having [C/H]~$\\simlt -3.5$\\,. On the other hand, Cowie \\& Songaila (1998) used a pixel-to-pixel optical depth technique to conclude that the average C~IV/H~I ratio remains essentially constant over the full range of neutral hydrogen column densities tested, down to log~$N$(H~I)~$\\simeq 13.5$\\,. If this is indeed the case, the ejection and transport of metals away from galaxies must have been much more efficient than envisaged. Most recently, Ellison \\etal (1999, 2000) have re-examined the problem and found that both approaches suffer from limitations which had not been properly taken into account in previous analyses. In their view, whether there are metals in the voids remains an open question.  On the theoretical side too, only a limited amount of work has been done on this topic. The common assumption that supernova driven winds may be able to distribute metals over large distances has been shown to be too simplistic and does not stand up to close quantitative scrutiny (MacLow \\& Ferrara 1999; Murakami \\& Babul 1999; D'Ercole \\& Brighenti 1999). These studies have come to the conclusion that efficient blowout is likely to be inhibited by the galaxy ISM and, at least at high redshift where densities are higher, by the pressure of the surrounding intergalactic gas (Babul \\& Rees 1992; Ciardi \\& Ferrara 1997). Gnedin \\& Ostriker (1998) numerically simulated the IGM enrichment in a CDM+$\\Lambda$ cosmological model. They found that the mean metallicity of the universe at $z=4$ is about 1/200 $Z_\\odot$ but the variations around this value are large and dependent on the overdensity.  At this epoch the enrichment was incomplete, with some regions of the universe still of pristine composition. Their simulations, however, end at $z=4$ so that a direct connection with the nearby universe is difficult. In addition, quantitative conclusions on the transport mechanism are affected by uncertainties arising from insufficient numerical resolution so that, for example, they are unable to resolve the snowplough phase of explosions. To alleviate the problem of inefficient blowout Gnedin (1998) developed the idea that metals are predominantly transported following merger events, a mechanism which would also predict a highly inhomogeneous distribution of heavy elements in the IGM at $z = 4$. However, the merger history has not yet been followed to more recent epochs.  In this paper we make explicit predictions for the IGM metallicity evolution by suggesting a two-step mechanism for the transport and mixing of heavy elements. In our scenario, SN explosions first eject metals over a relatively small region in the surroundings of the host galaxy; subsequently, a yet unknown diffusion process transports and mixes these elements over much larger scales. The plan of the paper is as follows. In \\S2 we consider the conditions for gas ejection from galaxies, paying particular attention to the treatment of blowout, which we show to be quite a rare occurrence in massive galaxies. In \\S3 we explicitly derive the metallicity evolution of the IGM for a particular (Cold Dark Matter) cosmological model. The results are discussed in \\S4 and a brief summary (\\S5) concludes the paper. ",
        "conclusions": "The basic conclusion of this paper is that metal ejection driven by SN events fails, by more than one order of magnitude, to distribute the products of stellar nucleosynthesis over volumes large enough to pollute the whole IGM to the typical metallicity of \\lya\\ clouds, [C/H]~$ \\simeq -2.5$. We are therefore forced to conclude that some additional physical process must be at play, the nature of which remains to be determined. In our scheme transport of metals occurs in two sequential steps. First, SNae provide the initial kick necessary to eject heavy elements outside the potential well of a galaxy, but the ejecta are then confined by the IGM pressure to a relatively small bubble, of radius $R_e$. We then postulate that a second process is responsible for the diffusion of metals on a typical scale $R_d$, comparable to the mean distance between the galaxies which are the most efficient pollutants of the IGM.  Our scenario is qualitatively different from that proposed by Gnedin (1998), who attributed the mixing to more violent and rarer galaxy mergers. In the model proposed here SNae are of fundamental importance as they initiate the mixing process. However, galaxy collisions might well play a role in the subsequent phase during which metals are spread over larger scales, of order $R_d$. Other diffusive processes, such as thermal conduction and turbulent mixing layers at the interfaces between cosmological flows, as well as the peculiar motions of galaxies, may also be important in determining the spatial structure of the distribution of metals. A detailed study of such processes is beyond the scope of this paper and should, in any case, be based on cosmological simulations which are in progress. We will report of this extension of the work in a forthcoming paper.  In the present study we have stressed that the IGM enrichment proceeds in a very inhomogeneous manner, with pockets of metal rich gas gradually increasing both in number and in size until they overlap at $z\\approx 1$. The average metallicity of the IGM increases with time, a trend confirmed by the results in Barlow \\& Tytler (1998), who found an order of magnitude increase in the metallicity of the Ly$\\alpha$ forest between $z=2.5$ and $z=0.5$. Although our results are strictly applicable only to the `true' IGM with overdensities close to unity, we nevertheless regard this as an encouraging performance of our model.  The metallicity spread is predicted to decrease with the progress of time. At $z \\simlt 1$, when the entire volume of the universe has been exposed to metal pollution, the spread in metallicity is less than one order of magnitude, and at the present epoch $\\langle Z \\rangle = 0.1 \\pm 0.03 Z_\\odot$. Thus we predict that at $z < 1$ essentially all absorbers should have associated C~IV absorption, irrespectively of their column density. The effect should be very pronounced. Not only are the voids polluted by the overlap of metal-enriched diffusive spheres, as discussed above, but the decreasing intensity and hardness of the ionizing background lead to a further increase in the fraction of C which is triply ionized, so that the ratio $N$(C~IV)/$N$(H~I) increases for a fixed [C/H] (Rauch et al. 1997). It should be possible to test these predictions with forthcoming observations. UV-efficient echelle spectrographs, now available on the VLT and soon on the Gemini South telescope, will allow sensitive searches for C~IV doublets at significantly lower redshifts than probed so far, while STIS on {\\it HST} and FUSE will map the \\lya\\ forest with the required spectral resolution at wavelengths below 3000~\\AA, which are inaccessible from the ground.  While our study has concentrated on the IGM, it also has important consequences for the metallicity of the intracluster medium. which is found to have a rather uniform  value $\\approx 1/3 Z_\\odot$ with little, if any, evolution up to $z \\simeq 0.3$ (Fukazawa \\etal 1998, Renzini 1999). This has been interpreted as evidence supporting the view that the enrichment process in clusters was already completed by that epoch. In general we expect two sources to contribute to the build up of the intracluster medium, infalling IGM and gas stripped from cluster galaxies. In our models these two components have significantly different composition. If the clusters formed at $z \\simeq 1$, we expect the IGM to have $\\approx 0.1 Z_\\odot$ (roughly constant from $z \\simeq 1$ to the present), while the galaxies' ISM has approximately solar composition. The simplest estimate of the resulting metallicity is the geometric mean of these two values (as appropriate for hydrodynamical mixing problems, see Begelman \\& Fabian 1990), that is $\\langle Z \\rangle_{icm} = (0.1 \\times 1)^{1/2} Z_\\odot= 0.32 Z_\\odot$, as observed. While this may well be a fortunate coincidence, it is also true that this conclusion lends qualitative support to the model and the assumptions made.  We make a final point concerning the role of intergalactic dust. The gas phase abundances we have derived do not take into account the possibility that some fraction of heavy elements may be locked up into dust grains. If this were the case, clearly the distributions of metallicities in Figure 6 and 7 would shift towards lower values of $Z$. In addition, at redshifts were the enrichment is still inhomogeneous, dust associated with regions of high metal concentration may reprocess UV/optical photons into IR radiation and give rise small scale anisotropies in the Cosmic Microwave Background which may be detectable (Ferrara \\etal 1999). Having said this, however, we consider it highly speculative whether dust can survive at all in the hostile environments associated with outflowing superbubbles, where the temperatures are high and the gas has been shocked by SN explosions."
    },
    "0004/astro-ph0004039_arXiv.txt": {
        "abstract": "\\b12 is the only known binary system with a radio pulsar from which the non-pulsed radio and X-ray emission was detected. The companion star in this system is a Be star SS 2883. A rapidly rotating radio pulsar is expected to produce a wind of relativistic particles. Be stars are known to produce highly asymmetric mass loss. Due to the interaction of the pulsar wind and the Be star wind the system of two shocks  between the pulsar and the Be star forms. In this paper we show that the observed non-pulsed radio emission from the system is a result of the synchrotron emission of the relativistic particles in the outflow beyond the shock wave and that the non-pulsed X-ray emission is due to the inverse Compton scattering of the Be star photons on this particles. ",
        "introduction": "\\b12 is the only known binary system with a radio pulsar from which the non-pulsed radio and X-ray emission was detected. The radio pulsar \\b12 of spin period $P$=47.76ms was discovered in a high-frequency survey of the southern Galactic plane.  Spindown luminosity is equal to $L_p=9\\times10^{35}$erg/s. Pulsar timing confirmed by later observations revealed that the pulsar is in a high-eccentricity ($e$=0.87), long-orbital period ($T$=3.4year) binary system with a massive companion. Its companion  is a 10th mag Be star SS2883 of luminosity $L_*$=5.8$\\times$10$^4L_\\odot$, estimated radius $R_*\\sim$(6-10)$R_\\odot$, mass $M\\sim$10$M_\\odot$. The distance to the Earth is about  2 kpc,  the binary separation at periastron is equal to 10$^{13}$cm. Assuming a pulsar mass of 1.4$M_\\odot$ the implied inclination angle of the binary orbit to the plane of sky is 36$^\\circ$. The longitude of periastron is equal to 138$^\\circ$ (Johnston \\etal 1996).  Timing measurements have shown that the disc of the Be star is likely to be highly inclined to the orbital plane (Wex \\etal, 1988).   The extensive radio observations of \\b12  were carried out during the 1994 and 1997 periastron passage at frequencies 1.5, 2.3, 4.8 and 8.4 GHz (Johnston \\etal 1996, 1999).  It turned out that when the pulsar is far from the periastron the pulsar emission is highly linear polarized and its intensity is practically independent on the pulsar orbital position.  But as the pulsar approaches to the periastron the properties of the pulsar emission start to change.  The depolarization of pulsed emission, the increase of the dispersion measure and the absolute value of the rotation measure near the periastron occurs. The pulsar was not detected in 1.5 GHz data on 1993 December 20 (20 days before the periastron) and reappeared only on 1994 February 4 (24 days after the periastron).  During the 44 days eclipse no pulse emission was detected in extensive observations at 1.5 and 8.4 GHz. In the works of Lipunov \\etal (1994) and Johnston \\etal (1996) it was shown that the observed eclipse is due to the free-free absorption of the pulsar emission in the Be star disk.  During the 1994 and 1997 periastron passage the observations of the unpulsed radio emission coming from the system were made using the Molongo Observatory Syn\\-thesis Telescope (MOST) and the Australia Telescope Compact Array (ATCA) at five frequencies between 0.84 and 8.4 GHz (Johnston \\etal 1999). The first significant data at all frequencies were received 21 day before the periastron.  Since that time  the intensity of the non-pulsed emission by and large gradually increases  until the 20th day after the periastron. That day it reaches its maximum (57 mJy at 1.4 GHz) and then it starts to decline.  The source remained significant in the ATCA data until about 60 days after the periastron and until about hundred days after the periastron in the MOST data. At times when the pulsar emission is eclipsed, there is no evidence of polarized emission in any of the observations in either 1994 or 1997.  \\begin{figure} % \\centering\\epsfig{file=proc1.eps,width=12cm, height=20cm} \\vspace{-10cm} \\caption {The geometry under consideration. Point $P$ marks the position of the pulsar, point $Be$ - the position  of the Be star. Curves $S_1$, $S_2$ represent the shock waves, curve $c$ - the contact surface. Curve $U$ is a curve of the constant density (see text). Points $A$ and $B$ mark the intersection of the curve $U$ with the shocks $S_1$ and $S_2$, point $C$ is the point of the shock $S_2$ intersection with the binary axis.  $ED$ is a small arc centred at the point $P$.} \\end{figure}    Since its discovery the PSR B1259-63 system was observed several times by  X-ray instruments (see review in Tavani \\etal 1997, Hirayama \\etal 1999). It was shown that the X-ray spectrum of the system is consistent with a power-law of photon index $\\sim-1.7$. No significant X-ray pulsations with the pulsar spin period were detected.  The aim of this paper is to explain the origin of the non-pulsed radio and X-ray emission. It is well-known that pulsars lose rotational energy in the form of relativistic MHD winds (Michel 1969, Arons 1992). The wind carries energy flux in the form of electromagnetic fields and the kinetic energy of the relativistic electrons and positrons.  The energy flux in the pulsar wind ${\\cal F}_w$ at distance $r$ well over the light cylinder is equal  to \\begin{equation} {\\cal F}_w=\\frac{L_p}{4\\pi r^2}=\\frac{cB_w^2}{4\\pi\\sigma}. \\label{b0} \\end{equation} $B_w$ is a magnetic field in the wind, $\\sigma\\ll 1$ is a pulsar wind magnetization, i.e. the upstream ratio of electromagnetic energy density and particle kinetic energy density. For the Crab Nebula $\\sigma\\sim0.005$ (Kennel\\& Coronity, 1984).  Be stars are well-known to be the source of the strong highly anisotropic matter outflow. Both a dilute polar wind and a denser equatorial disk have been invoked to reconcile models for IR, UV and optical observations of Be stars (Waters \\etal 1988). Due to the interaction of the pulsar and Be-star winds the system of two shocks arises and the relativistic particles flowing at first radially from the pulsar turn after passing the shock and start to outflow along the contact surface. In the paper of Tavani \\etal (1997) the non-pulsed X-ray emission was explained as a result of the synchrotron emission of the highly relativistic pulsar wind particles with Lorentz factor $\\gamma \\sim 10^6$. In this model the origin of the unpulsed radio emission is unclear.  Here we propose a model in which Lorentz factor is supposed to be much more moderate, $\\gamma \\sim 10 $, and show that in this case the unpulsed radio emission is a result of the synchrotron radiation of the relativistic particles in the outflow beyond the shock and that the non-pulsed X-ray emission is due to the inverse Compton scattering of the Be star photons on this particles. ",
        "conclusions": ""
    },
    "0004/astro-ph0004043_arXiv.txt": {
        "abstract": "The star formation process in molecular clouds usually leads to the formation of multiple stellar systems, mostly binaries. Remaining disks around those stars may be located around individual stars (circumstellar disks) or around the entire binary system (circumbinary disk). We shall briefly review the present observational evidence for both types of disks in binary stars, in particular the properties of circumbinary disks.  We then present recent results of the theoretical modeling of the collapse and fragmentation of gravitationally unstable molecular cloud cores and their implications for binary and disk formation, and discuss the dynamical influence of the binary companions  on disk truncation and gap formation. The presence of binaries may have profound influence on the process of planet formation as well. We present results on the stability and evolution of orbits of planets in disks around binaries. ",
        "introduction": "\\label{obs_properties} This contribution is dealing primarily with theoretical aspects, however for completeness, we summarize first some basic observational features of disks and planets in binary stars. An excellent recent review on this topic has been given at PPIV by Mathieu et al. (2000). \\subsection{Binary Stars} The majority of stars is found in binary or multiple systems. To quantify the occurrence of multiplicity one usually defines the {\\it Binary Frequency} (BF) as \\[  BF = \\frac{Number~~of~~multiples}{Total~~Number~~of~~Systems} \\]  The seminal study of Duquennoy \\& Mayor (1991) found for field stars (not members of a stellar cluster) having an age larger than 1Gyr a binary frequency of $BF=60\\%$. The distributions of the orbital elements of binary stars such as a mass ratio $q$, periods $P$, or eccentricities $e$ are typically very broad and they have mean values of \\[   <q> = 0.3 \\] \\[     <\\log(P[days])> = 5.0 \\] \\[   <e> = 0.5 \\] For pre main-sequence stars one finds a bimodal binary frequency. Dense clusters in giant molecular clouds, such as the Trapezium cluster in Orion (with a stellar density of over 2000 stars per $pc^3$), tend to have a binary frequency similar to the field value, ($60\\%$). In low stellar density clusters (such as the Taurus or Ophiuchus T Tauri associations) the binary frequency is much higher than the field value, about $80-100\\%$ (Ghez et al. 1993, Leinert et al. 1993). This difference in the $BF$ may be the result of several mechanisms the most important of which are (see Mathieu et al. 2000): a) observational effects due to incompleteness may lead to a wrong determination of the BF, b) evolutionary effects may change the BF with time, c) in higher density clusters more frequent stellar encounters may lead to the disruption of systems and a lowering of the BF, and d) the BF may just depend on the environment, such as the temperature or density of the cloud core.  From these results one may infer that most of the field stars originated in dense clusters. \\subsection{Disks in Binaries} The classical signatures for the presence of disks around stars are an excess in IR to mm radiation, Balmer and forbidden lines, polarization, and optical veiling, see contribution by L. Hartmann (this volume).  Disks in binary star systems may be divided into two main configurations 1) In the case of a {\\it Circumbinary Disk} the complete binary system is surrounded by a disk which may have in inner clearing. 2) If the two disks surround each component of the binary individually the situation is similar to single stars and one refers to this configuration as {\\it Circumstellar Disks}  The {\\it mm-flux} from systems gives an indication for the mass of the disk. For binaries with larger separations ($d> 100AU$) and for spectroscopic binaries ($d<1 AU$) the mm-flux is similar to the field value. This indicates that in the first case (large separations) the two stars are surrounded individually by circumstellar disks behaving just as disks around unperturbed single stars. In the second case (very close stars) the system is surrounded by one circumbinary disk which, when observed form far away, resembles a disk around one single star.  In binaries with intermediate orbital separations $d$ between 1 and 100 $AU$ the mm-flux is mostly undetected, which may indicate that in this case the binary system has neither extended circumstellar disks, as a consequence of the gravitational disturbance of the companion, nor luminous circumbinary disks. However, there are classical TTau signatures for disks with $d=10-100AU$.  Additionally, speckle and mid-IR observations give evidence for circumstellar disks around the stars HK~Tau, HR~4796A, or GG~Tau3.  The most prominent observational example of a circumbinary disk is presented by GG~Tau, where the disk has been directly imaged in the CO 1.3mm line (Guilloteau 1999), in the J-Band at $1.2 \\mu m$ (Roddier et al. 1996) and in the optical. The central binary system consists of two stars with a total mass of $1.3 M_\\odot$ at a separation of $d=45AU$, which are surrounded by a disk with an inner radius of $r_{disk}=180AU$. The inferred mass of the disk (ring) is about $0.12 M_\\odot$. Additional systems with imaged circumbinary disks are UY~Aur and BD+31$^o$643.  In several spectroscopic binaries (eg. UZ Tau, DQ Tau or GW Ori) mm-observations have indicated disk masses of about $0.03$ to $0.05 M_\\odot$ At the same time there is evidence for inner disk clearing (lack of near IR-emission), but there are signs for hot gas near the stars which may indicate the flow of material through the cleared inner disk (leaking gaps).  An interesting example of the simultaneous occurrence of circumstellar and circumbinary disks in a quadruple system is given by UZ Tau (Jensen et al. 1996). \\subsection{Planets in Binaries} Among the known 32 extrasolar planets around main sequence stars three (16 Cyg B, 55 Cnc, $\\tau$ Boo) are definitely known to be in binary systems with relatively wide orbital separations (see table 1). Additionally, in one system (Gl 86) there is some evidence for a stellar companion with a distance which may be as close as $10 AU$. \\begin{table} \\caption{The Properties of known planets in binary stars} \\begin{center} \\begin{tabular}{cccc|c} Parameter  &   16Cyg B & 55 Cnc & $\\tau$ Boo & Gl~86\\\\ \\hline Binary $d$ (AU)  &  700  &  1150  &  240 & $> 10 $\\\\ M sin$i$ ($M_J$)  &  1.5  &  0.84  & 3.87 & 4 \\\\ Planet a  (AU)   &  1.72 &  0.11  & 0.046 & 0.11\\\\ Planet e         &  0.63 &  0.05  & 0.02 & 0.04 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} It has been argued that the high eccentricity of 16 Cyg B may be caused be the presence of a companion star. However, as the ratio of distance $d$ of the secondary over the semi-major axis $a$ of the planets falls into the range \\[  \\frac{d}{a} =  400 - 10000 \\] the companion has basically no influence on the eccentricity of the planet. ",
        "conclusions": ""
    },
    "0004/gr-qc0004004_arXiv.txt": {
        "abstract": "We investigate the quantum evolution of large black holes that nucleate spontaneously in de~Sitter space. By numerical computation in the s-wave and one-loop approximations, we verify claims that such black holes can initially ``anti-evaporate'' instead of shrink. We show, however, that this is a transitory effect. It is followed by an evaporating phase, which we are able to trace until the black holes are small enough to be treated as Schwarzschild.  Under generic perturbations, the nucleated geometry is shown to decay into a ring of de~Sitter regions connected by evaporating black holes. This confirms that de~Sitter space is globally unstable and fragments into disconnected daughter universes. ",
        "introduction": "\\subsection{Anti-evaporation and proliferation}  Schwarzschild-de~Sitter black holes have unusual quantum properties and instabilities% ~\\cite{GibHaw77a,GinPer83,BouHaw96,BouHaw97b,NojOdi98b,% EliNoj99,Bou98,Bou99}. They are of cosmological interest because they can be produced during inflation~\\cite{BouHaw95}. They are also theoretically significant because they change the global structure of de~Sitter space~\\cite{Bou98,Bou99} fundamentally.  Unlike their asymptotically flat cousins, Schwarzschild-de~Sitter black holes are surrounded by a cosmological horizon. This limits their size and puts them in a thermal bath. Their temperature is always larger than that of the cosmological horizon~\\cite{GibHaw77a,BouHaw96}.  Nevertheless, Hawking and one of the authors argued in Ref.~\\cite{BouHaw97b} that some black holes accrete so much quantum radiation that they will grow, or `anti-evaporate,' instead of evaporating (see also~\\cite{NojOdi98b,EliNoj99}).  In the maximal solution, the black hole and cosmological horizon are of equal size, and the spatial geometry will be $S^1 \\times S^2$, with constant two-sphere radius. This geometry can nucleate semiclassically in de~Sitter space through a gravitational tunneling process.  Its evolution is unstable to the formation of $n$ de~Sitter regions, distributed around the $S^1$ and connected by $n$ black hole throats~\\cite{Bou98}. This configuration may be visualized as a `doughnut' with $n$ `wobbles,' whose minima and maxima represent the black hole and cosmological horizons, respectively. For $n=1$, this reduces to the ordinary, submaximal Schwarzschild-de~Sitter solution. For higher $n$, it resembles a ring-like sequence of $n$ Schwarzschild-de~Sitter solutions. As the de~Sitter regions form between the cosmological horizons, the black holes are expected to evaporate and shrink.  When they finally disappear, the doughnut is pinched in $n$ places, leaving $n$ disconnected pieces. This corresponds to the fragmentation of space into $n$ large daughter universes~\\cite{Bou98}.  Thus de~Sitter space is globally unstable. Locally, however, the daughter universes are indistinguishable from de~Sitter.  Therefore, they will harbour more maximal black holes, whose evaporation can lead to further fragmentation. Iteratively, an unbounded number of disconnected de~Sitter universes is produced: de~Sitter space proliferates.  This drastic conclusion depends crucially on the nonlinear evolution of the instabilities of the maximal Schwarzschild-de~Sitter solution (also referred to as the `Nariai' solution below).   \\subsection{Analytical shortcomings}  In Refs.~\\cite{BouHaw97b} and~\\cite{Bou98,Bou99}, the anti-evaporation and proliferation effects were derived analytically, using a number of approximations. The stability of the Nariai geometry was studied using linear perturbation theory. The evolution of the horizons was found only for very early times, as a power series in the time variable. This led to the discovery of anti-evaporation for a class of perturbations. But a power series can prove only {\\em initial} anti-evaporation. It cannot answer an important question: Will such black holes continue to grow, asymptotically approaching the maximal (Nariai) solution? Or will their growth eventually stall, and evaporation set in after all? This problem is highlighted by recent claims that anti-evaporation may permanently stabilize black holes nucleated during inflation\\footnote{We point out, however, that even stable black holes nucleated semiclassically during inflation would be diluted by the cosmological expansion. Typically, they will lie outside the current horizon and will not be observable~\\cite{BouHaw96,Bou96b}.}~\\cite{NojOdi98b}.  At least for Euclidean boundary conditions, it was shown that the black holes do evaporate at late times~\\cite{BouHaw97b}. In fact, arguments might be made that the evaporating mode is attractive at late times for all initial conditions~\\cite{Bou99}. In any case, however, the intermediate period containing the putative transition between anti-evaporation and evaporation is beyond the reach of the analytical methods employed. Can there be a smooth transition, and how does it proceed? To understand this crucial phase of the black hole evolution, it should be explored numerically.  If anti-evaporation is only temporary, then how can we be sure that the evaporating phase, once entered, continues until the black hole disappears?  This question is particularly important in the context of multiple black hole formation and proliferation~\\cite{Bou98,Bou99}. The fragmentation of space is possible only if all black holes become small and disappear like Schwarzschild black holes.  To show that black holes really evaporate completely, it is necessary to trace the backreaction at least until they are significantly smaller (perhaps by a factor of 10) than the cosmological horizon. Then they can be assumed to behave like Schwarzschild black holes, which presumably radiate all of their mass away. But linear perturbation theory requires that the small two-spheres inside the black hole and the large two-spheres between the cosmological horizons still be of nearly equal size. Thus, the analytic `late-time' solution, showing black hole evaporation, is not really valid for arbitrarily late times. A full, non-linear numerical calculation is needed to interpolate to the Schwarzschild regime.  \\subsection{Numerical analysis}  In Sec.~\\ref{sec-antievap} we consider $n=1$ perturbations of the Nariai geometry, which can be parametrized by their initial amplitude, $\\epsilon$, and phase, $\\vartheta$.  Depending on $\\vartheta$, the black holes undergo initial periods of anti-evaporation or evaporation.  We demonstrate the transitory nature of anti-evaporation for neutral black holes by showing that evaporation sets in at sufficiently late times, independently of $\\vartheta$. The evaporation rate is compared to an approximate solution of the linear equations in order to verify that this solution is an attractor. Beyond the linear regime, the black holes are found to shrink to a size much smaller than the cosmological horizon before our computation breaks down. Then they live in approximately flat space, and should simply continue to evaporate until they disappear.  In Sec.~\\ref{sec-charged} we show that maximal black holes become stable if their charge exceeds a critical value, in agreement with thermodynamic expectations. This confirms analytic arguments given in Ref.~\\cite{Bou99}. In particular, we show that supercritical black holes, independently of their initial behavior, always end up anti-evaporating.  This is the condition for a novel type of daughter-universe production proposed in Ref.~\\cite{Bou99}.  In Sec.~\\ref{sec-proliferation} we show that a higher mode ($n>1$) perturbation behaves as predicted by analytical arguments: First it oscillates, then it freezes out, forming black hole interiors and growing de~Sitter regions. Depending on initial conditions, a period of anti-evaporation may follow. Finally, the black holes evaporate and become much smaller than the cosmological horizons. ",
        "conclusions": ""
    },
    "0004/astro-ph0004275_arXiv.txt": {
        "abstract": "During the 1996/1997 outburst of the X-ray transient GRO\\,J1655$-$40 the {\\em Wide Field Cameras} onboard BeppoSAX observed the system with nearly continuous coverage during some stages of this outburst. On two occasions we find clear evidence for frequent dipping behaviour, characterized by several absorption dips occurring within one binary orbit. However, during preceding and following binary orbits there is no sign of such dips at the same range of orbital phases. We also collected light curves of absorption dips as seen with other X-ray instruments. Although the 65 dips that were collected occur preferably between orbital phases 0.68 and 0.92, there is evidence that the morphology of the material causing the dips is changing over just one binary orbit. The short duration ($\\sim$minutes) and phasing of the dips are discussed in the framework of models where the stream from the companion star encounters a hot outer rim on the accretion disk. ",
        "introduction": "The black-hole candidate and soft X-ray transient GRO\\,J1655$-$40 (X-ray Nova Sco 1994) was discovered on 1994 July 27 by the Burst and Transient Source Experiment (BATSE) onboard the {\\em Compton Gamma Ray Observatory} (Zhang et al.\\ 1994, Harmon et al.\\ 1995). Soon thereafter the optical counterpart was found (Bailyn et al.\\ 1995a). It has since shown irregular outburst activity (e.g.\\ Zhang et al.\\ 1997). After the 1995 July/August hard X-ray outburst the source settled back to quiescence. However, on 1996 April 25 it became active again (Remillard et al.\\ 1996, Levine et al.\\ 1996, Orosz et al.\\ 1997) for a period which lasted $\\sim$16~months (see e.g.\\ Remillard et al.\\ 1999). GRO\\,J1655$-$40 has an orbital period of 2.62~days, has dynamically been shown to contain a black hole with a mass of $\\sim$7\\,M$_{\\odot}$, and is viewed at an inclination of $\\sim$70$^{\\circ}$ (Bailyn et al.\\ 1995b, Orosz \\&\\ Bailyn 1997, van der Hooft et al.\\ 1998, Shahbaz et al.\\ 1999). The object has become notorious for being one of the systems showing relativistic radio jets, during its first outburst in 1994 (Tingay et al.\\ 1995; Hjellming \\&\\ Rupen 1995).  Deep short dips in the X-ray light curves of GRO\\,J1655$-$40 were discovered during its 1996/1997 outburst These dips appeared periodically between orbital phases $\\sim$0.7 and $\\sim$0.85 and were shown to be due to a medium absorbing the radiation coming from the inner parts of the accretion disk (\\mbox{Kuulkers} et al.\\ 1998a). Such ``orbital'' dips are known to occur also in various other low-mass and high-mass X-ray binaries (see Parmar \\&\\ White 1988; Marshall et al.\\ 1993; Saraswat et al.\\ 1996; and references therein). They probe the interaction between the stream from the companion and the accretion disk, as well as the X-ray emitting region itself (see e.g.\\ \\mbox{Kuulkers} et al.\\ 1998a; Tomsick et al.\\ 1998).  We report here on observations made with the {\\it Wide Field Cameras} (WFC) onboard the {\\em Beppo Satellite per Astronomia X} (BeppoSAX) observatory, obtained during various stages of the last outburst in 1996/1997 of GRO\\,J1655$-$40. Due to the relatively long orbital period of GRO\\,J1655$-$40, the WFC is ideally suited for monitoring such a system at all orbital phases almost continuously, which is not possible or undertaken with other X-ray instruments carrying out short programs of typically less than a day or with all-sky monitors viewing sources typically up to ten times a day for a short period of order minutes. We summarize all dip events seen during the 1996/1997 outburst with the WFC and other instruments and discuss their origin. ",
        "conclusions": "\\subsection{Orbital phase dependence}  Combining all times of X-ray dip occurrences during the 1996/1997 outburst of GRO\\,J1655$-$40 we find that the dips occur between orbital phases 0.68 and 0.92, i.e.\\ extending the dip range with respect to that quoted by \\mbox{Kuulkers} et al.\\ (1998a). The occurrence of the observed dips concentrate near orbital phases 0.75--0.80. Similar deep short dips have now been seen in other X-ray transient systems: 4U\\,1630$-$47 (Tomsick et al.\\ 1998; \\mbox{Kuulkers} et al.\\ 1998a) and Cir X-1 (Shirey et al.\\ 1999). As noted by \\mbox{Kuulkers} et al.\\ (1998a) the duration of the dips is also of the same order as those seen in Cyg\\,X-1 and Her\\,X-1 (e.g., Kitamoto et al.\\ 1984, Leahy 1997), but is shorter than those typically seen in LMXB dip sources (see e.g.\\ Parmar \\&\\ White 1988).  For the first time, we were able to obtain a nearly-continuous view of the system for several binary orbits during different parts of the outburst. On one occasion we found long irregular dipping behaviour for a period of $\\sim$9.2\\,hr, which spans about $\\sim$0.15 in orbital phase. This duty cycle is much larger than quoted by Tomsick et al.\\ (1998) for GRO\\,J1655$-$40; it is similar to that seen in low-mass X-ray binary dip sources (`dippers'; e.g.\\ Parmar \\&\\ White 1988; White et al.\\ 1995).  Do the dips recur every binary orbit? At certain times they do, as was shown by \\mbox{Kuulkers} et al.\\ (1998b) to be the case during the early part of the 1996/1997 outburst using the RXTE/ASM light curves. Our collection of dip times suggests the same during other parts of the outburst (see Table 2). If one compares in detail the light curves from binary orbit to binary orbit, it is clear that the dips do not persist at exactly the same phase. Moreover, if the strong dipping behaviour seen in binary cycle 272 were present every orbit, we would certainly have seen this in the WFC light curves. In contrast most of the WFC light curves do not show evidence for any strong dipping behaviour. Such a transient nature of the dip activity has been noted before as well in the other LMXB dip sources, see e.g.\\ Parmar et al.\\ (1986) for the case of EXO\\,0748$-$676. We note that it looks as if dipping only occurred during the first halfs of the two stages of the outburst, i.e.\\ they appear during the rise and plateau phases of the outburst.  \\subsection{Stream-disk interaction}  \\begin{figure} \\centerline{ \\hspace{2.5cm} \\psfig{figure=1655_wfc_fig7a.ps,bbllx=78pt,bblly=258pt,bburx=534pt,bbury=534pt,width=8cm,clip=yes} \\hspace{.0cm} \\psfig{figure=1655_wfc_fig7b.ps,bbllx=78pt,bblly=258pt,bburx=534pt,bbury=534pt,width=8cm,clip=yes} } \\caption{View of GRO\\,J1655$-$40 at the extreme ends of the dip phase range, i.e.\\ $\\phi_{\\rm orb}\\sim 0.68$ (left panel) and $\\phi_{\\rm orb}\\sim 0.92$ (right panel). We set the following binary parameters: $q$=0.354 (Shahbaz et al.\\ 1999) and $i$=67.2$^{\\circ}$ (van der Hooft et al.\\ 1998). We arbitrarily choose a disk opening angle of 2$^{\\circ}$ (see Phillips et al.\\ 1999) and a disk radius of 0.85\\,R$_{\\rm L1}$ (see Orosz \\&\\ Bailyn 1997). The dashed line shows the ballistic path of a stream from the secondary to the compact object in the absence of an accretion disk. } \\end{figure}  The orbital phase dependence of the dip occurrences suggests the dip mechanism to be fixed in the binary orbit. For illustrative purposes (Fig.~7) we show views of the system at the extremes of the dip phase range. The most straightforward site is the interaction region between the mass transfer stream from the companion star and the accretion disk (e.g.\\ Frank et al.\\ 1987). However, it is not the bulge on the accretion disk itself which obscures the X-ray emitting region. This would result in a rather long X-ray dip as seen in e.g.\\ accretion disk corona sources (see e.g.\\ White et al.\\ 1995); moreover, the inclination at which we view GRO\\,J1655$-$40 is too low for such an effect to be seen (see Frank et al.\\ 1987). Rather the short duration of the dips points to an absorbing medium which is filamentary in nature and probably situated above and/or below the impact region (\\mbox{Kuulkers} et al.\\ 1998b). A similar filamentary nature has been proposed for the medium in Cyg\\,X-1 that causes the second peak in the distribution of (short) dips with orbital phase at $\\phi_{\\rm orb}$$\\sim$0.6 (where phase zero corresponds to superior conjunction of the black hole). This peak has been attributed also to a mass transfer stream from the O-supergiant interacting with the accretion disk (Ba\\l uci\\'nska-Church et al.\\ 2000). Note that the main peak occurs at $\\phi_{\\rm orb}$$\\sim$0.95, which is due to absorption in the wind of the supergiant.  The dynamics of mass transfer streams from the companion up to the place of impact onto the accretion disk was first investigated by Lubow \\&\\ Shu (1975, 1976). More recently it has become possible to study in more detail the impact region itself and what happens with the matter after it reaches this region (e.g.\\ Armitage \\&\\ Livio 1998, and references therein). The hydrodynamic simulations by Armitage \\&\\ Livio (1998) show that the fate of the stream after it impacts the disk depends on the efficiency of cooling in the shock-heated gas created by the impact. If the cooling is efficient, the upper and/or lower parts of the stream are able to freely overflow the disk, unless the disk rim is too thick. This ballistic stream will reimpact the disk near to its closest approach to the compact object (see also Frank et al.\\ 1987, Lubow 1989). Observationally this may lead to rather broad dips occurring near orbital phase 0.6 (Lubow 1989, see also Armitage \\&\\ Livio 1996). If cooling is inefficient than the stream `splashes' onto the disk rim, i.e.\\ the material overflowing the disk is a coherent stream but rather erratic. Moreover, more material is able to reach greater heights from the disk, which leads to substantially higher absorption columns with respect to the former case, for lines-of-sight well away from the disk plane. At such lines-of-sight one expects a plethora of dips to be observed from orbital phases $\\sim$0.7--1.0, with a peak near orbital phase 0.8 (Armitage \\&\\ Livio 1996, see also Frank et al.\\ 1987). Moreover, in the latter case one expects a bulge extending along the disk rim as well (Armitage \\&\\ Livio 1998).  The dips observed during the 1996/1997 outburst of GRO\\,J1655$-$40 all share the characteristics of the case of non-efficient cooling: they are short, show absorption columns up to $\\sim$10$^{24}$~atoms\\,cm$^{-2}$ (\\mbox{Kuulkers} et al.\\ 1998a) and the distribution of dip occurrence times extends from binary orbital phases $\\sim$0.7--0.9, with a peak near binary orbital 0.8.  Apart from the `splashing' material causing the deep short dips, the presence of an extended bulge along the disk rim is supported by the observed decrease in optical polarisation near orbital phase $\\sim$0.7--0.8 (Gliozzi et al.\\ 1998). The observations were obtained just after the peak of the second part of the outburst during 1997 July 2--9. The polarisation region must be rather extended and is located primarily around the inner disk regions. A thickening of the disk rim or the remnant stream passing under and over the disk then obscures smoothly this polarisation region (Gliozzi et al.\\ 1998).  ASCA observed a $\\sim$4\\,hr drop in intensity by $\\sim$75\\%\\ in the light curve of GRO\\,J1655$-$40 on 1994 August 23 (Ueda et al.\\ 1998). These observations were obtained just after the peak of the first hard X-ray and radio outburst (Harmon et al.\\ 1995; Tavani et al.\\ 1996). Using the ephemeris by van der Hooft et al.\\ (1998), the ASCA dip occurred between orbital phases $\\sim$0.53--0.61, i.e.\\ outside the phase range we see for the short and deep dips during the 1996/1997 outburst\\footnote{We note that an eclipse-like feature in the optical light curves was found on 1994, August 17 by Bailyn et al.\\ (1995a), i.e.\\ close in time to the ASCA dip; a possibly similar optical feature occurred one binary orbit earlier (Bianchini et al.\\ 1997). The eclipse-like light curve of Bailyn et al.\\ (1995a) lasted from orbital phase $\\sim$0.0--0.1. Bianchini et al.\\ (1997) suggested this to be due to a bright spot region, or an extended optically thick disk rim shielding part of the disk.}. No deep short dips have been reported for the outbursts which occurred in 1994/1995, however, the soft X-ray ($<$10\\,keV) coverage during these outbursts was very sparse. The long ASCA dip is in agreement with what is expected when the stream encounters a cold accretion disk rim, so that part of the stream can continue to flow above and under the accretion disk (see above). According to Ueda et al.\\ (1998), the material causing the dip by absorption must consist of slightly ionized (warm) and non-ionized (cold) material, which is also in line with the fact that the stream will re-impact the disk sufficiently close to the region where ionisation becomes important (Frank et al.\\ 1987).  We suggest that the difference between the observed X-ray dips during the 1994 and 1996/1997 outbursts may be due to the efficiency in irradiating the outer disk regions. In 1994 the disk rim may have had a shape such that it does not see the X-ray emitting regions (i.e.\\ self-screening by the disc of its outer regions; see Dubus et al.\\ 1999), possibly due to a large enhancement of mass transfer from the secondary, and the gas at the stream impact region could cool efficiently. During the 1996/1997 outburst the outer disk regions were less affected by the mass transfer stream from the secondary, which were therefore more easily prone to X-ray irradiation by the X-ray emitting region, keeping the outer disk hot."
    },
    "0004/astro-ph0004105_arXiv.txt": {
        "abstract": "We reconsider the general properties of gravitational lensing effects induced by cosmic string systems, taking into account their equation of state and motion equations.  We explicitly show that the deflection patterns induced by a string is equivalent to the one of a lineic and static distribution of matter.  We then rigorously show that the convergence part of the deformation field is always zero, except on the intersection of the string worldsheet and the observer past light cone, extending previous results obtained in peculiar cases.  Phenomenological consequences of this result on multiple image systems are investigated. ",
        "introduction": "It  is well  known  that topological  defects  may appear  whenever, in  the thermal history of the universe,  a symmetry breaking phase transition occurs \\cite{kibble76,mermin},   as  for   instance   in  grand--unified   theories \\cite{kibble76}  or in  some extensions  of the  standard  electro-weak model \\cite{peter92}.   Such  defects  represent  spacetime  positions  where  the underlying order parameter cannot relax, because of topological constraints, to its low energy  vacuum state \\cite{kibble76,vilenkin}.  They are expected to interact mainly gravitationally with  the ordinary matter so that they can induce (i) deflection and  redshifting of massless particles, (ii) accretion of massive  non--relativistic particles and (iii)  emission of gravitational waves (see e.g.  \\cite{vilenkin,hindmarsh,durrer99rev} for a review of these effects).  One very interesting example, from a high energy physics point of view as well as from a cosmological point of view, corresponds to the case of cosmic strings.  In this case, their phenomenological properties are determined by the energy density per unit length of a string, $U$.  For instance the deflection angle is of order $4\\pi G U$, $G$ being the Newton constant.  For defects formed at the grand--unification scale ($T_{\\rm GUT}\\sim 10^{16}\\,~{\\rm Gev}$), we expect effects of a magnitude of $GU\\sim 10^{-6}$.  This corresponds for instance to the magnitude of the cosmic microwave background (CMB) anisotropies induced through the Kaiser-Stebbins effects \\cite{kaiser}. Although topological defects may have clear observational signature on the CMB sky \\cite{kaiser,btest}, observations seem to disfavor such an origin \\cite{turok,battye,durrer99,uzan99}.  Nonetheless, it does not mean that topological defects do not exist. Their detection would be of dramatic importance both for astronomy and particle physics \\cite{kibble76,vilenkin,durrer99rev,up} since for instance estimation or bounds on their density will help constraining the high energy physics theories predicting their existence.  Definitive predictions for string properties are however difficult to obtain, because in particular of the complex evolution equations that may depend on their microscopic structure through their equation of state.  For the so-called Goto--Nambu strings (where the energy per unit length $U$ and the tension $T$ of the string are equal), it was shown that the spacetime around such straight cosmic string was conical \\cite{vilenkin81,vilenkin84,vilenkin86,gott85,hogan84}.  Such a cosmic string formed at GUT scale would therefore induce image pairs of distant objects with angular separation of order $\\theta\\sim5.2''\\times(GU/10^{-6})$. From a pure phenomenological point of view such a string is expected to produce lines of double images \\cite{hindmarsh2}.  Recognizing the peculiarity of such a system, it was later extended to straight cosmic string with different equation of state \\cite{hindmarsh90,peter94} and to a string with a lightlike current pulse\\cite{peter94b}. Moreover numerical simulations for Goto--Nambu string were also performed in the case of long strings \\cite{delaix1} and cosmic loops \\cite{delaix2,delaix3}.  More quantitatively the prospects for a direct detection of relic string via gravitational lensing, and in particular the expected number of events, was first discussed by Hindmarsh \\cite{hindmarsh2} who estimated the angular length of string per unit area on the sky out to a redshift $z$ to be of order $$ \\theta_{\\rm loops}\\sim 0.1\\nu z^2\\,{\\rm deg}^{-1}\\qquad \\theta_{\\rm long\\,string}\\sim 0.1A z^2\\,{\\rm deg}^{-1} $$ where  $\\nu$  and  $A$  are  two  coefficients  of  order  unity  (see  also \\cite{delaix1}).  In conclusion, it  is widely believed that the observation of  a  cosmic  string  can  be achieved  through  double  image  detections, although, in  practice, it might  be difficult to  be positive about  such a detection since  pairs with  the same angular  separation appearing  by pure coincidence can  be very high (as pointed  by Cowie and Hu  who reported for such a cosmic string lens candidate \\cite{cowie,emhu}).  The aim of this article and its companion \\cite{bu2} is to make a systematic investigation  of the gravitational  lensing effects  by cosmic  strings. We focus  our analysis  on the  deformation equation  of a  geodesic  bundle in presence  of cosmic  strings  (\\S~\\ref{I}). Standard  approximations of  the gravitational  lens theory  are also  discussed  in this  section.  After  a description  of  the cosmic  string  dynamics  in  \\S~\\ref{II}, we  show  in \\S~\\ref{III} that the deflecting potential  of a cosmic string is equivalent to  the one by  a static  distribution of  matter on  the projection  of the string worldsheet  onto the observer past  light cone. In the  course in this calculation we show that the deformation field induced by cosmic strings has a zero convergence (without  any approximation). Examples illustrating these results are discussed as well as the validity of the thin lens approximation and  the influence  of the  equation of  state. In  \\S~\\ref{IV} we investigate  the  phenomenological  consequences  of  the  zero  convergence property on multiple  image systems.  In a companion  paper, \\cite{bu2}, we propose  a  phenomenological  model   of  string  energy  distribution  that gives a more quantitative account of these results. ",
        "conclusions": "We have considered  the general lensing properties of  objects such as cosmic strings where relativistic motions  and non-trivial equation of state induce metric  perturbations of all  sorts.  We  demonstrated that  the deformation field of  a string system  on the image  plane is the same  as the one  of a static linear distribution  of matter projected on the  photon trajectory. A consequence of  this result  is that the  deformation field has  a vanishing convergence   ($\\kappa=0$)  everywhere   but  on   the  projection   of  the intersection of the  observer past light cone and  the string worldsheet.  We explicitly illustrate this result with  the case of a circular cosmic string loop in a plane perpendicular to the line-of-sight for which we generalize the results found  in \\cite{delaix2}  to string with  any equation of  state. We also showed that (i) the fact that the deformation field outside the loop is equivalent to the one obtained by a  massive point and (ii) that a light ray passing inside the loop is not deflected are due to the energy conservation of the string.  We also paid  attention to the validity of thin  lens approximation for this unusual  lens   system.   This  approximation is discussed in detail through the case of a  static tilted straight cosmic string. It lead us  to point  out that  for string  with a  general equation  of  state, the deflection may be  more important than for a  Goto--Nambu string, this being understood by  the fact that the  stress--energy tensor of  a general cosmic string can always be decomposed as the superposition of a Goto--Nambu string and  a lineic distribution  of non--relativistic  matter. The  deflection is then due  to the  combined effect  of the deficit  angle of  the Goto--Nambu string and of the curvature induced by the lineic distribution of matter.  We also studied  general phenomenological  consequences  of deformation fields with zero convergence on  multiple image systems, the main goal being to be able to  assess if two images are likely to form  an image pair of the same source.  For that purpose, we derived  all the image shapes  of a given source as well as all the source  shapes of a given image. These results may serve as a groundwork for the elaboration of string detection strategies.  We are aware that this latter part  is limited by the fact that we only took advantage  of the  zero-convergence property.   In practical  more intricate distortion properties of the images are  likely to be useful. This is one of the  aim   of  the  companion  paper  \\cite{bu2}   where  more  quantitative phenomenological  properties  are presented  that  take  into account  local string energy fluctuations."
    },
    "0004/astro-ph0004333_arXiv.txt": {
        "abstract": "We present estimates of the energy input from supernovae (SNe) into the intergalactic medium using (i) recent measurements of Si and Fe abundances in the intracluster medium (ICM) and (ii) self-consistent gasdynamical simulations that include processes of cooling, star formation, SNe feedback, and a multi-phase model of the interstellar medium. We estimate the energy input from observed abundances using two different assumptions: (i) spatial uniformity of metal abundances in the ICM and (ii) radial abundance gradients. We show that these two cases lead to energy input estimates which are different by an order of magnitude, highlighting a need for observational data on large-scale abundance gradients in clusters. Our analysis indicates that the SNe energy input can be important for heating of the {\\em entire} ICM (providing energy of $\\sim 1$ keV per particle) only if the ICM abundances are uniform {\\em and}\\/ the efficiency of gas heating by SN explosions is close to $100\\%$ ($\\epsilon_{SN}\\approx 1$, implying that all of the initial kinetic energy of the explosion goes into heating of the ICM).  The SNe energy input estimate made using simulations of galaxy formation is consistent with the above results derived from observed abundances, provided large-scale radial abundance gradients exist in clusters.  For the cluster AWM7, in which such a gradient has been observed, the energy input estimated using observed metal abundances is $\\sim 0.01$ and $\\sim 0.1$ keV per particle for $\\epsilon_{SN}=0.1$ and $\\epsilon_{SN}=1$, respectively.  These estimates fall far short of the required energy injection of $\\sim 0.5-3$ keV per particle that appears to be needed to bring models of cluster formation into agreement with observations.  Therefore, our results indicate that, unless the most favorable conditions are met, SNe alone are unlikely to provide sufficient energy input and need to be supplemented or even substituted by some other heating process(es). ",
        "introduction": "\\label{sec:intro}  \\footnotetext{Hubble Fellow} Hierarchical models of structure formation have been very successful in explaining many observed properties of galaxies and galaxy clusters.  Nevertheless, some puzzling problems remain open.  Several theoretical studies have demonstrated that some heating of gas, in addition to the heating during the gravitational collapse, is required to explain the observed properties of the intracluster medium (ICM). \\scite{kaiser91:preheat} first showed that an early injection of energy results in correlations and evolution of bulk cluster properties (X-ray luminosity, gas temperature, etc.) that match observations.  This conclusion was backed by numerical simulations: \\scite{evrard90:formation} was able to get a better fit to the X-ray luminosity of the Coma cluster in his cosmological simulation by preheating the gas to $10^7$ K ($\\approx 0.9$ keV), while \\scite{nfw95} showed that pre-heated clusters matched the observed slope of the correlation between X-ray luminosity and temperature. Recent semi-analytical studies of cluster evolution have reached similar conclusions (\\pcite{cavaliere97:lt,balogh99:preheat,valageassilk99:entropy,wu99:heating}).  The exact amount of required energy injection depends on the epoch, and have been argued to be in the range of $0.5-3$ keV per gas particle (\\pcite{nfw95,cavaliere97:lt,balogh99:preheat,wu99:heating}). Although the problem has been identified, it is not yet clear what processes can provide the required heating. It is clear that identification of these processes is crucial for a complete picture of cluster formation.  The candidate process which has been discussed most is supernovae-driven galactic winds. The gas of the galactic interstellar medium (ISM) can be heated by supernovae explosions and acquire energy comparable to or larger than its gravitational binding energy. This heated gas can then flow away and result in additional heating if the winds result in shocks when they encounter intergalactic gas. Although there is observational evidence for such winds in present-day galaxies (e.g., \\pcite{heckman90:superwinds}), theoretical models of winds are rather ill-constrained due to uncertainties in the efficiency of conversion of supernovae (SNe) explosion energy into thermal energy of the gas and other details.  Early estimates of possible energy input from SNe based on observed metal abundances showed that SNe are plausible candidates (e.g., \\pcite{white91:esn,david91:enrichment,loewenstein96}). However, in these estimates it was assumed that the distribution of metals in the ICM is uniform and that the efficiency with which energy of SNe explosions can be converted into thermal energy of the gas is close to $100\\%$. Therefore, there is a need for detailed estimates using new measurements of metal abundances in clusters and current galaxy formation models which have become much more advanced and sophisticated in the last several years.  Evaluations of SNe as a heating source have recently been performed by \\scite{valageassilk99:entropy} and \\scite{wu99:heating} using a semi-analytical approach to galaxy modelling. These authors conclude that it is unlikely that supernovae are the only source of heating. \\scite{valageassilk99:entropy} argue then that radiation from quasar population can provide the required heating much more easily.  In this paper we repeat previous estimates of the possible SNe energy input using updated values of observed ICM abundances and relaxing the assumption of abundance uniformity, motivated by recent observations (\\pcite{ezawa97:awm7}; \\pcite{finoguenov99}). We also make a separate estimate for the cluster AWM7, for which the radial abundance gradient was measured.  The details of this estimate are presented in \\S~\\ref{sec:esnmetals}. We complement this analysis with direct estimates of the energy input by counting the total number of supernovae exploded in all cluster galaxies throughout their evolution in self-consistent three-dimensional gasdynamical simulations of galaxy formation. The simulations include cooling, star formation, SNe feedback and a multi-phase model of the interstellar medium in galaxies and have been shown to match many fundamental observed correlations of galactic properties such as the galaxy luminosity function, the Tully-Fisher relation and its scatter, the color-magnitude sequence, and, perhaps most importantly, the evolution of the global star formation rate in the Universe.  The details of the simulations are described in \\S~\\ref{sec:esnsimul}.  The energy input estimates are presented and compared in \\S~\\ref{sec:results} and discussed in \\S~\\ref{sec:discussion}. We summarize our main results and conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented estimates of the possible energy input by supernovae into the intracluster medium. Although these estimates are prone to a number of uncertainties, we have defined conditions which determine whether SNe can be a significant source of ICM heating. The following main conclusions can be drawn from our analysis.  The SNe energy input, $\\esn^m$, estimated from observed ICM abundances of Si and Fe is only significant ($\\sim 1$ keV per particle) when we assumed that the distribution of metals in the ICM is uniform (no significant radial gradients) {\\em and}\\/ that $\\sim 100$\\% of individual SN explosion energy goes into heating the ambient gas followed by negligible cooling ($\\epsilon \\approx 1$) (see \\S~\\ref{sec:esnmetals}).  If large-scale metallicity gradients are assumed in clusters, the estimated energy input is $\\sim 0.1-0.5$ keV per particle for $\\epsilon=1$ and, correspondingly, $0.01-0.08$ keV per particle for a more realistic value of $\\epsilon=0.1$.  As an example, we present estimates of the energy input for the cluster AWM7 for which the abundance gradient has been measured. We find that the observed abundance of iron in this cluster implies a SNe energy input of $\\lesssim 0.01$ and $\\lesssim 0.1$ keV per particle for $\\epsilon=0.1$ and $\\epsilon=1$, respectively.  The energy input, $\\esng$, estimated using self-consistent three-dimensional numerical simulations of galaxy formation which include effects of shock heating, cooling, SN feedback, and multi-phase model of ISM, are $\\approx 0.01$ and $\\approx 0.1$ keV per gas particle for values of efficiency parameter $\\epsilon=0.1$ and $\\epsilon=1$, respectively.  These values are somewhat lower than the values of $\\esn^m$ (but are in good agreement with estimates for the AWM7). Nevertheless, the two estimates agree reasonably well if the existence of large-scale abundance gradients is assumed in clusters. We therefore emphasize the importance of new measurements of large-scale metallicity gradients for testing the theoretical models.  Our estimates of the SN energy input in all cases, except the case of uniform ICM abundances and $\\epsilon=1$, fall short of the energy injection of $\\sim 0.5-3$ keV per particle required to bring theoretical models of cluster formation in accord with observations. This suggests that supernovae are unlikely to be the only source of the IGM heating and should possibly be supplemented (or substituted) by some other heating mechanism. Similar conclusions have been reached in recent studies of \\scite{balogh99:preheat}, \\scite{valageassilk99:entropy}, and \\scite{wu99:heating}. \\scite{valageassilk99:entropy} propose radiation from quasars as an alternative heating mechanism. This opens discussion of new possible processes for what appears to be a required high-redshift preheating of the intergalactic medium."
    },
    "0004/astro-ph0004111_arXiv.txt": {
        "abstract": "We analyze the star formation history (SFH) of the Galactic disk by using an infall model. Based on the observed SFH of the Galactic disk, we first determine the timescales of the gas infall into the Galactic disk ($t_{\\rm in}$) and that of the gas consumption to form stars ($t_{\\rm sf}$). Since each of the two timescales does not prove to be determined independently from the SFH, we first fix $t_{\\rm sf}$. Then, $t_{\\rm in}$ is determined so that we minimize $\\chi^2$. Consequently, we choose three parameter sets: $(t_{\\rm sf}\\,{\\rm [Gyr]},\\, t_{\\rm in}\\,{\\rm [Gyr]})= (6.0,\\, 23)$, (11, 12), and  (15, 9.0), where we set the Galactic age as 15 Gyr. All of the three cases predict almost identical star formation history. Next, we test the intermittency (or variability) of the star formation rate (SFR) along with the smooth SFH suggested from the infall model. The large value of the $\\chi^2$ statistic supports the violent time-variation of the SFH. If we interpret the observed SFH with smooth and variable components, the amplitude of the variable component is comparable to the smooth component. Thus, intermittent SFH of the Galactic disk is strongly suggested. We also examined the metallicity distribution of G-dwarfs. We found that the true parameter set lies between $(t_{\\rm sf}\\,{\\rm [Gyr]}, \\, t_{\\rm in}\\,{\\rm [Gyr]})= (6, 23) \\mbox{ and }(11, 12) $, though we should need a more sophisticated model including the process of metal enrichment within the Galactic halo. ",
        "introduction": "\\label{sec:intro}  In the field of the Galactic chemical evolution, \\citet{Eggen62} have inspired the modeling. From the correlation between the ultraviolet excess and orbital eccentricity of stars, they concluded that the Galaxy formed by collapse on a free-fall timescale from a single protogalactic cloud. An alternative picture of halo formation is proposed by \\citet{Searle78}. They argued that Galactic system is formed from the capture of fragments such as dwarf galaxies over a longer timescale than that proposed by Eggen et al. In any case, determining the timescale of the infall of matter and the chemical enrichment is an important problem to resolve the formation mechanism of the Galactic system. Indeed, there have been a number of papers that investigated the formation and chemical evolution of the Galaxy (e.g., \\citealt{Matteucci89}) and spiral galaxies (e.g., \\citealt{LB75}; \\citealt{Sommer96}).  Many models of the star formation history (SFH) of the Galaxy includes the formation of the Galactic disk by gas infall. This scenario (so-called the infall model) is consistent with the age-metallicity relation of the disk stars (e.g., \\citealt{Twarog80}), if a reasonable SFH is applied. Moreover, the infall model provides a physically reasonable way of solving the G-dwarf problem (e.g., \\citealt{Pagel97}, p.236), contrary to the closed-box model which tends to overpredict the number of the low-metallicity stars.  One of the factors that determine the star formation rate (SFR) is the gas content of galaxies. Indeed, the SFR and the gas density is closely related to each other \\citep{Kennicutt98}, and this relation is generally called the Schmidt law \\citep{Schmidt59}. The SFR and the gas density, $\\rho$, are related as $\\SFR\\propto\\rho^n$, where $n=1 - 2$. In the one-zone model, the relation is simply assumed to be described with $\\SFR\\propto M_{\\rm g}^n$, where $M_{\\rm g}$ is the gas mass in the galaxy.  Though the ``classical'' infall model is widely accepted, there are observational data that suggest the intermittent star formation activities in spiral galaxies. Kennicutt, Tamblyn, \\& Congdon (1994) showed the ratio of present-to-past SFR in spiral sample has a significant scatter. More recently, Tomita, Tomita \\& Sait\\={o} (1996) analyzed the far-infrared to $B$-band flux ratio $f_{\\rm FIR}/f_B$ of 1681 spiral galaxies (see also \\citealt{Devereux97}). The indicator $f_{\\rm FIR}/f_B$ also represents the ratio between the present SFR and the averaged SFR during the recent Gyr. They showed the order-of-magnitude spread of $f_{\\rm FIR}/f_B$ and suggested the violent temporal variation of the SFR.  The intermittency of the SFH in the Galactic disk is recently suggested by Rocha-Pinto et al.\\ (2000; hereafter R00). They derived the SFH of the Galactic disk from observed age distribution of the late-type stars and suggested that the star formation activity of the disk is intermittent or violently variable. They tested to show that the variation is of statistical significance. Their test are based on the null hypothesis that the SFR of the Galaxy is temporally ``constant'', but their derived SFH indicates that the SFR around 10--15 Gyr ago is significantly smaller than that around 0--10 Gyr ago. This trend of the Galactic SFR can be interpreted along with the infall scenario of the Galactic disk formation: Gas-infall to the Galactic disk occurred in the first $\\sim$ 5 Gyr. Thus it is an important work to interpret their data through the infall model and to examine whether the variation of SFR is still significant even in the infall scenario.  In this paper, we re-analyze the data by R00 in the context of the infall model. In the next section, we briefly review the derivation of the SFH in the Galactic disk by R00. Next, in \\S\\ref{sec:infall}, we present the formulation of the infall model, based on which we test the data by R00 statistically. In \\S\\ref{sec:discuss}, we interpret the observed SFH by two components; underlying infall-model component and violently variable component. In the same section, possible physical mechanisms for the variation of star formation are also presented. Finally, we summarize the content in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Metallicity and G-dwarf Problem}  One of the prime motivation for infall models is that they provide a physically reasonable way of solving the G-dwarf problem (\\citealt{Pagel97} and references therein). Thus, in this section, we examine the chemical evolution of the Galaxy.  Under the one-zone treatment and the instantaneous recycling approximation, the time evolution of the abundance of the heavy element, whose species is labeled by $i$ ($i={\\oxy }$, C, Si, Mg, Fe, ...), is expressed as \\begin{eqnarray} \\frac{\\sd M_i}{\\sd t}=-X_i\\psi +E_i+\\xif F(t),\\label{basic_metal} \\end{eqnarray} where $M_i$, $X_i$, $E_i$ and $\\xif$ are the total mass of heavy element $i$, the abundance of $i$ (i.e., $X_i\\equiv M_i/M_{\\rm g}$), the total injection rate of element $i$ from stars, and the abundance of the infalling gas, respectively \\citep{Tinsley80}. Here, $E_i$ is expressed as \\begin{eqnarray} E_i=(RX_i+Y_i)\\psi ,\\label{ira_metal} \\end{eqnarray} where $R$ is defined in equation (\\ref{recycle}) and $Y_i$ is the mass fraction of the element $i$ newly produced and ejected by stars; in other words, \\begin{eqnarray} Y_i = \\int_{m_t}^{m_u} mp_i(m)\\phi (m)\\,\\sd m, \\end{eqnarray} where $p_i(m)$ is the fraction of mass converted into the element $i$ in a star of mass $m$. Adopting the Salpeter's IMF and the stellar yield by \\citet{maeder92}, $Y_{\\oxy }=1.8\\times 10^{-2}$, where we adopt the oxygen as a tracer of heavy elements (i.e., $i={\\oxy }$; \\citealt{lisenfeld98,Hirashita99}). Combining equations (\\ref{basic_gas}), (\\ref{infall}), (\\ref{def_tilde}), (\\ref{basic_metal}), and (\\ref{ira_metal}), we obtain \\begin{eqnarray} f_\\sg\\frac{\\sd X_i}{\\sd\\tau}=\\frac{f_\\sg Y_i}{\\tau_{\\rm sf}}- \\frac{X_i-\\xif}{\\tau_{\\rm in}}e^{-\\tau /\\tau_{\\rm in}}. \\label{basic_metal2} \\end{eqnarray} Since the time-evolution of $f_\\sg$ is solved in equation (\\ref{gas_fraction}), equation (\\ref{basic_metal2}) can be integrated to obtain $X_i$ as a function of $\\tau (=t/T_{\\rm G})$. We hereafter choose oxygen as a tracer of the metallicity; i.e., $i={\\oxy }$ and assume $X_\\oxy^{\\rm f}=0.01X_{\\oxy\\, ,\\odot}$. For a more detailed modeling, we should include the metal enrichment within the Galactic halo (e.g., \\citealt{Ikuta99}).  The result is shown in Figure \\ref{fig3} for the three parameter sets in Table \\ref{tab1}, where $X_{{\\oxy },\\,\\odot}=0.013$ represents the solar oxygen abundance (\\citealt{Whittet92}, p.42). The solid, dashed, and dash-dotted lines present the cases of $(\\tau_{\\rm sf},\\,\\tau_{\\rm in})=(0.4, 1.5), (0.7, 0.8), \\mbox{ and } (1.0, 0.6)$, respectively. We also present the observational data of age-metallicity relation by \\citet{Rocha00b} (see their Table 3), where we assume that ${\\rm 0.7\\,[Fe/H]=[O/H]} \\equiv \\log (X_{\\oxy }/X_{{\\oxy },\\,\\odot})$ by using \\citet{Edvardsson93}. The true parameter set which reproduce the observational data points seems to lie between the solid line $(\\tau_{\\rm sf},\\, \\tau_{\\rm in}) = (0.4,1.5)$ and $(\\tau_{\\rm sf},\\,\\tau_{\\rm in}) = (0.7, 0.8)$, though the discrepancy at the lower-metallicity side is prominent. But if the infalling gas is more enriched than have been assumed, this discrepancy might be resolved. We should also note the large scatter of the relation (Fig.~13 of \\citealt{Rocha00b}) and the uncertainty in the above relation between [Fe/H] and [O/H].  The infall timescale is larger than that in \\citet{Sommer93}, who gave 3.4 Gyr. The difference comes not only from the different way of modeling but also from the no prominent decline of the recent SFR as presented in Figure~\\ref{fig2}.  The G-dwarf problem is also tested along with our model. The probability distribution function $P(\\log X_{\\oxy })$ of the metallicity is calculated from our model as \\begin{eqnarray} P(\\log X_{\\oxy })\\,\\sd\\log X_{\\oxy }=\\Psi\\,\\sd\\tau .\\label{prob_X} \\end{eqnarray} Using equations (\\ref{basic_metal2}) and (\\ref{prob_X}), we obtain the following analytical expression for $P$: \\begin{eqnarray} P(\\log X_{\\oxy}) = (\\ln 10)\\; \\Psi (\\tau )X_{\\oxy }(\\tau )\\left[ \\frac{Y_{\\oxy }}{\\tau_{\\rm sf}}- \\frac{X_{\\oxy}(\\tau )-X_\\oxy^{\\rm f}} {\\tau_{\\rm in}f_\\sg (\\tau )}e^{-\\tau /\\tau_{\\rm in}} \\right]^{-1} \\;, \\end{eqnarray} where $\\tau$ is a function of $X_{\\oxy }$ and the functional form is determined by solving equation (\\ref{basic_metal2}). In comparing the distribution function with the observational data, we should take into account the scatter of the data. Here, we simply convolve $P$ with a Gaussian kernel as \\begin{eqnarray} P_{\\rm conv}(\\log X_\\oxy )\\equiv\\int_{-\\infty}^\\infty P(u)\\, \\frac{1}{\\sqrt{2\\pi}\\,\\sigma}\\, \\exp\\left[ -\\frac{(\\log X_\\oxy -u)^2}{2\\sigma^2}\\right]\\,\\sd u, \\end{eqnarray} where we adopt $\\sigma =0.1$ to compare with \\citet{Rocha96}. We adopt these data because we would like to use a sample of G-dwarfs, whose lifetime is as long as the age of the universe. The result is shown in Figure~\\ref{fig4}. The solid, dashed, and dash-dotted lines present the cases of $(\\tau_{\\rm sf},\\,\\tau_{\\rm in}) = (0.4,1.5), (0.7, 0.8), \\mbox{ and } (1.0, 0.6)$, respectively. The histogram shows the data by \\citet{Rocha96}, where ${\\rm 0.7\\,[Fe/H]=[O/H]}$ is assumed and the integrated number is normalized to unity. We see the same trend as Figure \\ref{fig3}: Again we see that the true parameters lie between the solid line and the dashed line. As mentioned above, we should include the detailed enrichment process within the Galactic halo as e.g., in \\citet{Ikuta99}. However, since the aim of this paper is to examine the intermittent SFH proposed by R00, we do not examine the chemical evolution further.  \\subsection{Comment on the Instantaneous Recycling Treatment} \\label{subsec:non_ira}  The instantaneous recycling approximation is adopted only because it provides a very convenient way to obtain an analytic SFH, which is easily applied to the statistical analysis in \\S\\ref{subsec:statistics}. However, if we fix $t_{\\rm sf}$, the instantaneous recycling treatment tends to overestimate the rate of gas consumption to form stars due to the instantaneously recycled gas, which should be recycled with a delay in the realistic situation. Thus, we here examine the effect of the delayed recycling. We examine the opposite extreme case where the effect of the delayed recycling is significant: The gas is never recycled. In other words, we examine the case of $R=0$.  In Figure~\\ref{fig5}, we present the result for $R=0$ for the three parameters $(\\tau_{\\rm sf},\\,\\tau_{\\rm in})=(0.4,\\, 1.5),$ (0.7, 0.8), and (1.0, 0.6) with the thick lines. The thin lines represent the case where $R=0.32$; i.e., the same as Figure~\\ref{fig3}. We adopt the same normalization for both sets of lines (equation~\\ref{normalization}). Comparing the thick and thin lines, we see the more decline of the SFR in the recent epoch in the non-recycling case than the thin lines. This is because the gas consumption timescale is shorter in the non-recycling case \\citep{Kennicutt94}. However, we stress that the variation of the observed SFH (R00) is so large that the difference between the instantaneous and non-instantaneous treatments does not affect the conclusion that the SFR is violently variable.  \\subsection{Two-Component Model of the Star Formation Rate}  Since the observed SFH cannot be reproduced by a simple infall model, we introduce a stochastically varying residual and modify the model as follows: \\begin{eqnarray} \\Psi_{\\rm obs} (\\tau) = \\Psi_{\\rm infall}(\\tau) + \\varepsilon \\; , \\end{eqnarray} where $\\Psi_{\\rm infall} (t)$ is the estimated best-fit infall model in the above discussion, and the residual, $\\varepsilon$, is produced by a probability distribution with zero mean and dispersion $\\sigma^2$. We simply assume that $\\sigma^2$ is time-independent. Since $\\varepsilon$ is estimated for each of the 37 bins, we can show a distribution of $\\varepsilon$ (Figure~\\ref{fig6}). The estimated sample dispersion $\\sigma^2 = 0.22$ ($\\sigma = 0.47$). Considering the uncertainty in the age estimation in R00, the time-variation of the SFR is blurred with the uncertainty, as commented in R00. Thus, the above value of $\\sigma$ can still be underestimated. Therefore, we conclude that the $\\sigma$ is at least comparable to (perhaps larger than) $\\Psi_{\\rm infall}$.  If we assume that the large variation of the SFR is typical of spiral galaxies, the star formation activities of them should show a variety. The large value of the variance ($\\sigma$) is consistent with previous works that suggested the variety of star formation activities of spiral galaxies \\citep{Kennicutt94,Tomita96,Devereux97}. Furthermore, the kurtosis of the residual, $K = -0.81$, which means that $\\varepsilon$ is distributed flatly and is not strongly concentrated around the mean. Indeed, in the Figure~8 of \\citealt{Tomita96}, there seems to be little concentration of star formation activity around the mean, which is consistent with the flat distribution in Figure~\\ref{fig6}.  \\subsection{Possible Mechanisms for the Variation of SFH} \\label{subsec:mechanism}  Here, we consider what mechanisms are possible for the violent time-variation of SFR in the Galactic disk. We mention the following two mechanisms.  The first possibility is that the infall into the Galactic disk is a stochastic process. If an infalling gas is in the form of a cloud or a small-sized galaxy, such gas may induce a burst of star formation and increase SFR instantaneously. Indeed, infall of small galaxy seems to occur frequently seeing that the Sagittarius dwarf galaxy is now infalling to the Galaxy \\citep{Ibata94}. The high-velocity clouds \\citep{Wakker97} may fall into the Galactic disk stochastically and induce stochastic bursts of star formation.  The second possibility is that the interstellar medium is a non-linear open system. A non-linear open system often shows a limit-cycle behavior of physical quantities \\citep{Nicolis77}. The application of the non-linear-open-system model to the interstellar medium is described in \\citet{Ikeuchi83}. According to their model, the fractional mass of the cold component ($X_{\\rm c}$) can oscillatory change. Recently, \\citet{Kamaya97} suggested that the SFR varies oscillatory if the Schmidt law of $\\SFR\\propto X_{\\rm c}^2$ is assumed, where $X_{\\rm c}$ is the mass ratio of the cold gas to the whole gas (see also \\citealt{Ikeuchi88} for a review and \\citealt{Hirashita00} for a recent development on this theme).    In the context of the infall model, we re-analyzed the SFH of the Galaxy derived observationally by R00. We test to examine whether variation of the star formation rate proposed by R00 is significant.  We first statistically infer the timescales of the gas infall into the Galactic disk ($t_{\\rm in}$) and that of the gas consumption to form stars ($t_{\\rm sf}$). Since each of two timescales are not determined independently, we first fix $t_{\\rm sf}$. Then, $t_{\\rm in}$ is determined so that we minimize $\\chi^2$. Consequently, we choose three parameter sets: $(t_{\\rm sf}\\,{\\rm [Gyr]},\\, t_{\\rm in}\\,{\\rm [Gyr]})= (4.5, 33), (10.5, 9.0), \\mbox{ and } (15, 7.5)$, where we set the Galactic age as 15 Gyr. All of the three cases predict an almost identical SFH. The parameter set that seems to fit best to the age-metallicity relation and the metallicity distribution of the G-dwarfs is $(t_{\\rm sf}\\,{\\rm [Gyr]},\\, t_{\\rm in}\\,{\\rm [Gyr]})= (10.5,\\, 9.0)$. The infall timescale is larger than that in \\citet{Sommer93}, who gave 3.4 Gyr. The difference comes not only from the different way of modeling but also from the no prominent decline of the recent SFR as presented in Figure~\\ref{fig2}.  Next, we test the intermittency (or violent variability) along with the smooth SFH suggested from the infall model. The large value of $\\chi^2$ statistic supports the violent variation of the SFH. Then, we interpret the observed SFH with the two components; underlying smooth component described by the infall model, and violently variable component. We find that the variation of the latter component is comparable to the former.  As a test of the models, we also examine the age-metallicity relation and the metallicity distribution of the Galactic stars. Consequently, we observe that the degeneracy of the three parameters are resolved and find that the true parameter set seems to lie between $(t_{\\rm sf}\\,{\\rm [Gyr]},\\, t_{\\rm in}\\,{\\rm [Gyr]})= (10.5,\\, 9.0) \\mbox{ and } (4.5, 33)$.  Finally, two physical mechanisms for the variable SFH are suggested. One is the stochastic infall of the clouds or small galaxies, and the other is the non-linear oscillation of the SFR due to the limit-cycle behavior of the fractional mass of the cold-gas component."
    },
    "0004/astro-ph0004261_arXiv.txt": {
        "abstract": "Astrometry, on the International Celestial Reference Frame (epoch J2000.0), is presented for the Walker (1994, \\pasp, 106, 828) stars in the $\\omega$ Centauri (=NGC\\,5139=C\\,1323-1472) {\\sl{Hubble Space Telescope}} Wide Field/Planetary Camera (WF/PC) calibration field of Harris \\et (1993, \\aj, 105, 1196). Harris \\et stars were first identified on a WFPC2 observation of the $\\omega$ Cen {\\sl HST} calibration field. Relative astrometry of the Walker stars in this field was then obtained using Walker's CCD positions and astrometry derived using the {\\sc{stsdas}} {\\sc{metric}} task on the positions of the Harris \\et stars on the WFPC2 observation. Finally, the relative astrometry, which was based on the {\\sl{HST}} Guide Star Catalog, is placed on the International Celestial Reference Frame with astrometry from the USNO-A2.0 catalog. An ASCII text version of the astrometric data of the Walker stars in the $\\omega$ Cen {\\sl HST} calibration field is available electronically in the online version of the article. ",
        "introduction": "Walker (\\cite{W94}; hereafter W94) presented ground-based $BV$ (Johnson) $RI$ (Kron-Cousins) photometry of the Harris \\et (\\cite{H93}; hereafter H93) {\\sl{Hubble Space Telescope (HST)}} Wide-Field/Planetary Camera (WF/PC) calibration field in the Galactic globular cluster $\\omega$ Centauri (=NGC\\,5139=C\\,1323-1472). This small (100\\arcsec$\\!\\times$100\\arcsec) field is located 13.2 arcmin southwest of the center of $\\omega$ Cen at 13$^{\\rm{h}}$\\,25$^{\\rm{m}}$\\,37$^{\\rm{s}}$ -47\\arcdeg\\,35\\arcmin\\,38\\arcsec\\ (epoch J2000.0) [see Fig.\\ 1 (Plate 59) of H93, Figs.\\ 1 and 2 of W94]. Although this field has been observed over 1100 times with the {\\sl{HST}}, there is no astrometry available for the W94 stars in the published literature. In an effort to encourage various ongoing WFPC2 calibration efforts, I now present astrometry (J2000.0) of the W94 stars in the $\\omega$ Cen {\\sl HST} calibration field on the International Celestial Reference Frame. ",
        "conclusions": ""
    },
    "0004/astro-ph0004057_arXiv.txt": {
        "abstract": "We present near-infrared observations of the counterpart of GRB 000301c.  The K' filter ($2.1 \\micron$) light curve shows a well-sampled break in the decay slope at $t \\approx 3.5$ days post-burst.  The early time slope is very shallow ($\\sim -0.1$), while the late time slope is steep ($-2.2$). Comparison with the optical (R band) light curve shows marginally significant differences, especially in the early time decay slope (which is steeper in the optical) and the break time (which occurs {\\it later} in the optical). This is contrary to the general expectation that light curve breaks should either be achromatic (e.g., for breaks due to collimation effects) or should occur later at longer wavelengths (for most other breaks). The observed color variations might be intrinsic to the afterglow, or might indicate systematic errors of $\\ga 0.08$ magnitude in all fluxes. Even if the break is achromatic, we argue that its sharpness poses difficulties for explanations that depend on collimated ejecta. The R light curve shows further signs of fairly rapid variability (a bump, steep drop, and plateau) that are not apparent in the K' light curve. In addition, by combining the IR-optical-UV data with millimeter and radio fluxes, we are able to constrain the locations of the self-absorption break and cooling break and to infer the location of the spectral peak at $t = 3$ days, $f_\\nu \\approx 3.4\\mJy$ at $\\nu \\approx 10^{12} \\Hz$. Using the multiwavelength spectral energy distribution, we are able to constrain the blast wave energy, which was $E \\ga 3 \\times 10^{53} \\erg$ if the explosion was isotropic.  This implies a maximum gamma ray production efficiency of $\\sim 0.15$ for GRB 000301C. ",
        "introduction": "Infrared observations can be used to improve our understanding of gamma ray burst afterglows in several ways.  First, they can be combined with optical measurements to obtain spectral slope measurements with a much wider wavelength baseline, and hence yield a more accurate spectral slope than optical data alone.  Second, they can be used to test observed light curve breaks for wavelength dependence, which is an important discriminant between breaks due to ejecta collimation and other possible causes.  Finally, it has been suggested that bursts are preferentially located in dusty regions, e.g., under ``hypernova'' scenarios where bursters are a final evolutionary stage of some class of massive stars. If so, and if the dusty region extends beyond the expected dust destruction distance (Waxman \\& Draine 1999), then near-infrared observations will detect some afterglows that are obscured at optical wavelengths.  At less extreme dust optical depths, near-infrared data help to characterize the host galaxy extinction and so infer both extinction corrected fluxes and properties of dust and gas in high redshift GRB host galaxies.  We present here near-infrared (NIR) photometric observations of the afterglow of the gamma ray burst GRB 000301c.  These constitute the best-sampled near infrared light curve for any afterglow to date.  GRB 000301c was detected independently by the All-Sky Monitor on the Rossi X-Ray Timing Explorer and by two spacecraft (Ulysses and NEAR) of the current Interplanetary Network on 2000 March 1.4108 UT.  The event was a single peaked GRB lasting approximately 10 seconds (Smith, Hurley, \\& Cline 2000) at low energies ($< 10 \\keV$) and 2 seconds at higher energies ($> 25 \\keV$) (Jensen et al 2000).  Coordinates were available approximately 36 hours after the burst, and an optical counterpart was reported by Fynbo et al (2000) based on observations at March 3.21 UT.  The redshift of the burst was first reported as $z = 1.95 \\pm 0.1$ (Smette et al 2000a) and subsequently refined to $z = 2.028 \\pm 0.025$ (Smette et al 2000b) using the observed Lyman break in a near-UV spectrum of the GRB afterglow.  Weak metal lines in absorption yielded further improvements to $z = 2.0335 \\pm 0.0003$ (Castro et al 2000) or $z=2.0404 \\pm 0.0008$ (Jensen et al 2000). ",
        "conclusions": "GRB 000301c is the third burst for which a strong break in the light curve is clearly observed.  Several classes of breaks are predicted by fireball models.  The most basic of these are due to features in the synchrotron spectrum moving through the observed bandpass (e.g., Paczy\\'{n}ski \\& Rhoads 1993; Sari, Piran, \\& Narayan 1998).  However, this class of features predicts relatively modest changes in light curve slope, with the break occurring first at short wavelengths and evolving to longer ones.  Jetlike burst ejecta, on the other hand, are expected to give strong breaks that are essentially independent of wavelength (Rhoads 1997, 1999; Sari, Piran, \\& Halpern 1999), and the observed breaks have generally been interpreted as evidence for collimation of the GRB ejecta (e.g., in GRB 990123, by Castro-Tirado et al 1999, Kulkarni et al 1999, Fruchter et al 1999, and Galama et al 1999; and in GRB 990510, by Stanek et al 1999 and Harrison et al 1999).  A difficulty with this model is that the predicted break is quite gradual (Rhoads 1999; Panaitescu \\& Meszaros 1999; Moderski, Sikora, \\& Bulik 2000; Kumar \\& Panaitescu 2000a), while observed breaks are rather sharp.  The current burst is no exception.  The prediction of Rhoads (1999) for the light curve around the break time for a collimated jet is \\begin{equation} f_\\nu \\propto {  [t/t_b]^{-3(p-1)/4} \\over \\left(1 + 3.72 \\, [t/t_b]^{2/5}\\right)^{5/2} * \\left(1 + 2.07 \\, [t/t_b]^{5/12}\\right)^{3(p-1)/5} } \\end{equation} where $t_b$ is the fiducial break time defined in Rhoads (1999). The break in this predicted light curve is extremely broad, and would give $\\chi^2$ little better than a single power law in fitting the observed break in either $K'$ or $R$ band.  The model curve is based on numerical integration of the remnant's dynamical equations, and ignores differences in light travel time between the center and edge of the remnant, which will only smooth the break further (e.g., Moderski et al 2000; Panaitescu \\& Meszaros 1999).  If we ignore the issue of break sharpness and fit a collimated jet model to the observed R band light curve, we can infer the opening angle from the measured break time.  To do so, we need a reasonable measurement of $t_b$ and crude estimates of $\\Omega/E$ and $n$, since the inferred opening angle scales as $(t_b^3 n \\Omega/E )^{1/8}$ (Rhoads 1999). We use $t_b = 5.1$ days and $E = 3 \\times 10^{53} \\erg \\times \\Omega / (4 \\pi)$.  To estimate $n$ more precisely, we use the column density $N(HI) \\approx 10^{21.2 \\pm 0.5}$ inferred from Lyman $\\alpha$ absorption (Jensen et al 2000) and estimate the linear size of the source as $\\la 0.2''$ based on its nondetection in late HST images. This implies a number density $n \\ga 0.4$, consistent with our earlier estimate. Using $n \\approx 1$, the inferred opening angle becomes $2.5^\\circ$, or $10^{-3}$ of the sky if the jet is bipolar.  The transition to the nonrelativistic regime has been proposed as another mechanism for light curve breaks both in this burst (Dai \\& Lu 2000) and others (Dai \\& Lu 1999).  However, we do not know of a detailed calculation of the sharpness of this break, making a fair evaluation of this possibility difficult.  Light travel time effects seem likely to broaden this feature to $\\delta t / t \\sim 1$, as with most other features.  A final possible cause for sharp breaks in GRB afterglow light curves is discontinuities in the ambient density distribution.  Assuming that the density is a function of radius alone, a minimum timescale for breaks due to such discontinuities is $\\Delta t \\ga t$, where $t$ is the time elapsed in the observer's frame since the burst and $\\Delta t$ the characteristic duration of a light curve feature.  This duration is set by differential light travel time effects between material moving along the line of sight and off-axis material moving in direction $1/\\Gamma$, and is a rough minimum for any afterglow light curve feature provided the ambient medium density is approximately independent of angle from the line of sight.  The time required for material already in the expanding blast wave to cool by adiabatic losses is also relevant for determining the sharpness of a density discontinuity break, as is the emission from angles $> 1/\\Gamma$ off the line of sight (Kumar \\& Panaitescu 2000b).  If we believe that the observed $R-K'$ variations are real, then the greatest difficulty posed by the observations of GRB 000301c is in finding a model whose light curve steepens in $K'$ {\\it before} it steepens in R.  For most mechanisms, breaks will occur either first at short wavelengths (e.g. the cooling break), or simultaneously at all wavelengths (e.g. ``beaming'' breaks).  One speculative way out is to suppose that a discontinuity in the ambient density is encountered while the cooling break is between the R and K' filters.  The predicted appearance of an afterglow at frequencies above and below this break is expected to differ qualitatively: A high frequency image would show an annular structure and a low frequency image a more nearly filled disk.  This is caused by the difference in the apparent dynamical age of the remnant along the line of sight (where we see things changing quickly) and near the edge of the observed afterglow (where light travel time is larger, and we see material at an earlier and hotter stage of its evolution).  (E.g., Granot, Piran, \\& Sari 1999.)  Now, the same variation of ``lookback time'' from the center to edge of the afterglow implies that we see the effect of an ambient density drop first in the middle of the afterglow, and that the fractional effect of such a discontinuity on the afterglow flux will initially be larger at long wavelengths than short ones.  This mechanism can reproduce the sign of the observed effect.  Detailed calculations would be necessary to see if it can approach the observed magnitude, given that the two filters are separated by only a factor of $3$ in wavelength."
    },
    "0004/astro-ph0004130_arXiv.txt": {
        "abstract": "We present a sample of 148 \\RRc\\ selected from Sloan Digital Sky Survey (SDSS) commissioning data for about 100 deg$^2$ of sky surveyed twice with $\\Delta t$ =1.9946 days. Although the faint magnitude limit of the SDSS allows us to detect RR Lyrae stars to large galactocentric distances ($\\sim$ 100 kpc, or $r^* \\sim$ 21), we find no candidates fainter than $r^* \\sim$ 20, i.e. further than  $\\sim$ 65 kpc from the Galactic center. On the assumption that all 148 candidates are indeed RR Lyrae stars (contamination by other species of variable star is probably less than 10\\%), we find that their volume density has roughly a power-law dependence on galactocentric radius, $R^{-2.7\\pm0.2}$, between 10 and 50 kpc, and drops abruptly at $R \\sim$ 50--60 kpc, possibly indicating a sharp edge to the stellar halo as traced by RR Lyrae stars.  The Galactic distribution of stars in this sample is very inhomogeneous and shows a clump of over 70 stars at about 45 kpc from the Galactic center. This clump is also detected in the distribution of nonvariable objects with RR Lyrae star colors. When sources in the clump are excluded, the best power-law fit becomes consistent with the $R^{-3}$ distribution found from surveys of bright RR Lyrae stars. These results imply that the halo contains clumpy overdensities inhomogeneously distributed within a smooth $R^{-3}$ background, with a possible cutoff at $\\sim$ 50 kpc. ",
        "introduction": "The SDSS is a digital photometric and spectroscopic survey which will cover one quarter of the Celestial sphere in the north Galactic cap and produce a smaller area ($\\sim$ 225 deg$^2$) but much deeper survey in the southern Galactic hemisphere (Gunn \\& Weinberg 1995, York {\\em et al.} 2000\\footnote{see also http://www.astro.princeton.edu/PBOOK/welcome.htm}). The flux densities of detected objects are measured almost simultaneously in five bands ($u'$, $g'$, $r'$, $i'$, and $z'$, \\cite{F96}) with effective wavelengths of 3540~\\AA, 4760~\\AA, 6280~\\AA, 7690~\\AA, and 9250~\\AA, complete to limiting (5:1 signal-to-noise) point source magnitudes of 22.3, 23.3, 23.1, 22.3, and 20.8 in the North Galactic cap\\footnote{We refer to the measured magnitudes as $u^*, g^*, r^*, i^*,$ and $z^*$ because the absolute calibration of the SDSS photometric system is still uncertain at the $\\sim 0.05^m$ level. The SDSS filters themselves are referred to as $u', g', r', i',$ and $z'$. All magnitudes are given on the AB$_\\nu$ system (Oke \\& Gunn, 1983; for additional discussion regarding the SDSS photometric system see \\cite{F96} and Fan 1999).}. The survey sky coverage of about $\\pi$ steradians will result in photometric measurements to the above detection limits for about 10$^8$ stars. The pixel size (0.4 arcsec) and optical quality of the telescope are such that the resolution is limited by atmospheric seeing. Astrometric positions are accurate to about 0.1 arcsec for sources brighter than 20.5$^m$, and the morphological and color information from the images allows robust star-galaxy separation to $\\sim$ 21.5$^m$.  The survey is being done with a dedicated, special-purpose 2.5 meter telescope (\\cite{Siegmund00}). It has a wide well-corrected field (3 deg) and is equipped with a large mosaic CCD camera and a pair of fiber-fed spectrographs. The camera utilizes thirty large-area (2048 x 2048) CCDs (\\cite{Gunnetal}, Doi {\\em et al.} 2000) which take the data in drift-scanning (time-delay-and-integrate, or TDI) mode with a total integration time of 54.1 s. The imaging data are obtained using the data acquisition system at the Apache Point Observatory (Petravick {\\em et al.} 1994, Annis {\\em et al.} 2000) and recorded on DLT tapes. These tapes are shipped to Fermilab by express courier and the data are automatically reduced through a set of software pipelines operating in a common computing environment (Kent {\\em et al.} 2000). The photometric pipeline (\\cite{Lupton00}) reduces the imaging data, measuring positions, magnitudes, and shape parameters for all detected objects. The photometric pipeline uses position calibration information from the astrometric pipeline (\\cite{Pier00}) and photometric calibration data from the photometric telescope (Smith {\\em et al.} 2000, Uomoto {\\em et al.} 2000), reduced through the photometric telescope pipeline (Tucker {\\em et al.} 2000). Final calibrations are applied by the final calibration pipeline which allows refinements in the positional and photometric calibration to be applied as the survey progresses. The outputs, together with all the observing and processing information, are loaded into the operational data base (Yanny {\\em et al.} 2000b), which is the central repository of scientific and bookkeeping data used to run the survey.  About 40\\% of the sky in the northern survey will be surveyed twice (because of the scan overlaps), and all of the southern survey dozens of times (to search for variable objects and, by stacking the frames, to go deeper). Although two observations are normally insufficient to characterize a variable  object, the multi-color nature of the photometric data helps enormously. Close to 1000 square degrees of sky along the Celestial Equator have been observed during the SDSS commissioning phase (e.g. \\cite{F99}). About 100 deg$^2$ of sky have been observed more than once; we use these data to search for variable objects. In this paper we describe the detection and analysis of $\\sim$ 150 variable objects, probably RR Lyrae stars. Detecting RR Lyrae stars is important for Galactic structure studies because:  \\begin{itemize} \\item They are believed to be an unbiased tracer of the low-metallicity halo population for kinematic studies (\\cite{H84}), \\item They are nearly standard candles ($\\langle M_V \\rangle$ = 0.7 $\\pm$ 0.1, \\cite{L96}) and thus it is straightforward to determine their distance, and \\item They are sufficiently bright to be detected at large distances ($\\sim$ 100 kpc for $r'$ $\\sim$ 21$^m$) and are thus especially well-suited for studies of the outer halo (\\cite{S84}). \\end{itemize}  For a comprehensive review of RR Lyrae stars we refer the reader to Smith (1996).   Wetterer \\& McGraw (1996) recently compiled data from several available deep searches for RR Lyrae stars in the outer Galactic halo. They pointed out that there are only nine RR Lyrae stars discovered at galactocentric distances larger than 30 kpc, and only a few more have been found since then (e.g. Margon \\& Deutsch 1999, and references therein). This paper presents a sample of candidate RR Lyrae stars with 81 stars estimated to be at such distances. The following section describes the selection procedure used in the search for \\RRc\\ and the resulting sample. We analyze the Galactic distribution of selected stars in \\S3, and compare the sample to two other surveys of variable stars in \\S4.    \\section         {        Selection Procedure        }  We utilize imaging data from four runs (77, 745, 752, and 756) obtained during the SDSS commissioning phase. The data were obtained in six parallel scanlines\\footnote{See also http://www.astro.princeton.edu/PBOOK/strategy/strategy.htm}, each 13.5 arcmin wide, along the Celestial Equator ($-1.2687^\\circ <$ Dec $< 1.2676^\\circ$). The seeing in all runs was variable between 1 and 2 arcsec (FWHM) with the median value typically 1.5 arcsec. Runs 745 and 756, which are used as the main data set, were taken 1.9946 days apart on March 20 and 22, 1999. Their overlap extends from RA = 10$^h$ 42$^m$ ($l=248.6^\\circ$, $b=48.7^\\circ$) to RA = 15$^h$ 46$^m$ ($l$=7.3$^\\circ$, $b$=40.2$^\\circ$) and covers 97.5 deg$^2$ of sky (all coordinates are given as J2000).  As auxiliary control data sets we use overlaps between runs 77 and 745 (33.7 deg$^2$), 77 and 756 (20.5 deg$^2$), and 752 and 756 (11.1 deg$^2$). The RA ranges for these overlaps are summarized in Table 1. Run 77 was obtained on June 27, 1998 and provides a $\\sim 9$ month (266 and 268 days) baseline when compared to runs 745 and 756. Run 752 was obtained on March 21, 1999 and provides a baseline of 0.9976 days when combined with run 756. The overlap between runs 752 and 756 perpendicular to the scan direction (i.e. declination overlap) is only 12.4\\% of the scanline width (these are two six-column strips which are later interleaved to make a filled stripe of imaging data), while the overlaps between runs 77-745, 77-756 and 745-756 are 67.7\\%, 63.7\\% and 95.9\\% of the scanline width, respectively (scanline width = 6 $\\times$ 0.2253 deg = 1.3517 deg).   RR Lyrae stars have the colors of A and F stars (Preston 1959; for SDSS colors see \\cite{F96}, \\cite{KMS98}; note also Figure 2 below) and thus could be selected by appropriately constraining all four SDSS colors, and then searching for variability. However, this would produce complicated selection effects because not all five bands have the same sensitivity. In order to avoid such effects we follow a two-step procedure; the first step uses only data from the most sensitive $g'$ and $r'$ bands, and only in the second step do we introduce color cuts based on data in all five bands. Throughout this work we use the ``point spread function'' (PSF) magnitudes (measured by fitting a PSF model of a double Gaussian) as computed by the photometric pipeline (``photo'', version v5\\_1, for details see \\cite{Lupton00}).   Since RR Lyrae stars are bluer in $g^* - r^*$ than most other stars, we started our search by selecting the 90569 unresolved and unsaturated sources with $-0.1 < g^* - r^* < 0.4$ from $\\sim$ 930,000 stars in the overlap of runs 745 and 756, and then required that candidate variable sources satisfied the following conditions:  \\begin{enumerate} \\item The difference between the magnitudes in the two runs in both $g'$ and $r'$ bands is at least 0.15$^m$. \\item The difference between the magnitudes in the two runs in both $g'$ and $r'$ bands is at least 5$\\sigma$. Here the errors are taken as estimated by the photometric pipeline and do not include systematic calibration errors. However, because of the above requirement (variability of at least 0.15$^m$), this condition becomes relevant only at the faint end where the errors are dominated by photon statistics (see Figure 1 below). \\item Candidates are brighter in $r^*$ when they are bluer in $g^* - r^*$ (since RR Lyrae stars are pulsating variable stars). This is equivalent to the condition that the difference between the magnitudes in $g'$ band is larger than the difference between the magnitudes in $r'$ band. This condition is implemented without accounting for photometric errors. \\end{enumerate}  These selection criteria yield 186 candidates. In Figure 1a, the large dots show the observed change in the $r^*$ magnitude for the selected candidates plotted as a function of the mean $r^*$ magnitude. The small dots mark the remaining 90569 sources with $-0.1 < g^* - r^* < 0.4$. The two dashed lines show the boundary of the observational cutoff (the combination of items 1 and 2 above). The mean errors (for the magnitude difference) are about 0.03$^m$ up to 20$^m$, increasing to about 0.07$^m$ at 21$^m$, and to 0.2$^m$ at 22$^m$; the observed error distribution is in good agreement with the errors quoted by the photometric pipeline).  There are 21 stars ($\\about 10 \\%$) rejected by the third condition $|g^*_2-g^*_1| > |r^*_2-r^*_1|$, where indices 1 and 2 mark data from each epoch. This condition is shown by the diagonal dashed lines in Figure 1b, where for clarity only stars with mean $r^* < 20$ are plotted (using the same symbols as in Figure 1a). Note that Figures 1a and 1b show two projections of the three-dimensional selection volume spanned by $g^*_2-g^*_1$, $r^*_2-r^*_1$, and the mean $r^*$ magnitude. The rejected candidates are probably W UMa stars, but may also be RR Lyrae stars scattered across the selection boundary by photometric errors. The former probably dominate in the rejected subsample because more than $80\\%$ are brighter than $r^*=20$, and thus have very small photometric errors. The diagonal solid line in Figure 1b shows a best-fit relation $|g^*_2-g^*_1| = 1.4 \\, |r^*_2-r^*_1|$. Similar analysis based on data from other bands yields $|u^*_2-u^*_1| = 1.0 \\, |g^*_2-g^*_1|$, $|r^*_2-r^*_1| = 1.2 \\, |i^*_2-i^*_1|$, and $|i^*_2-r^*_1| = 1.1 \\, |z^*_2-i^*_1|$, where the uncertainty of best-fit coefficients is $\\la 0.05$.   A striking feature in Figure 1a is the lack of faint objects with large variability amplitudes in the regions outlined by the two ellipses. The magnitude distribution of the RR Lyrae candidates turns off rather sharply at r$^* \\sim$ 20, even though nonvariable objects are detected 2.5 magnitudes fainter, and the errors are sufficiently small that variable sources with similar amplitudes could be detected to at least r$^* \\sim$ 21 (see below). This indicates that we are detecting the faint end of the RR Lyrae magnitude distribution, and hence the limit of their distance distribution. The absence of faint variable sources appears not to be due to our selection criteria, since sources in those two regions are already practically absent in the starting sample (small dots in Figure 1a). To repeat, the starting sample was selected from the full data set by simply extracting unresolved\\footnote{Ignoring the requirement that candidates must be unresolved has practically no effect on the resulting sample. Thus possible star-galaxy misclassification cannot be invoked as an explanation for the observed cutoff.} sources in the appropriate $g^* - r^*$ color range.   Additional support for the reality of this cutoff comes from the analysis of the overlap between runs 77 and 745. These runs were obtained 9 months apart, which is a sufficiently long baseline to detect quasar (QSO) variability. Although low-redshift QSOs ($z <$ 2) have $g^* - r^*$ colors similar to RR Lyrae stars, they are easily distinguished by their bluer $u^*-g^*$ colors (\\cite{Fan99}). Since variable QSOs should not have a faint magnitude cutoff, they can be used to test whether the data allow the detection of variable sources fainter than r$^* \\sim$ 20.  We searched for variable objects in the overlap of runs 77 and 745 ($\\sim$ 35 deg$^2$) using the same criteria outlined above. This new search is summarized in the right column panels in Figure 2, and compared to the results from the first search shown in the panels on the left. The top two panels are analogous to Figure 1a (indeed, the top left panel is the same). Vertical lines at $r^* = 20.2$ are added to guide the eye, and mark the apparent faint end of the RR Lyrae magnitude distribution detected in runs 745-756. QSOs are selected by requiring $u^*-g^* < 0.8$ and are marked by crosses\\footnote{Detailed analysis of the variable QSO sample is outside the scope of this work and will be presented in a separate publication.}. The color difference between QSOs and RR Lyrae stars can be easily seen in the $g^* - r^*$ vs. $u^* - g^*$ color-color diagrams shown in the two middle panels\\footnote{In all color-color diagrams, blue is towards the lower left corner and red is towards the upper right corner.}. In these diagrams variable sources are marked by lines which connect photometric measurements at the two epochs, and dots represent a subsample of 5000 nonvariable unresolved objects which outline the stellar locus and the position of low-z QSOs ($u^*-g^* < 0.8$). Note that there are no variable sources with $u^*-g^* < 0.8$ in the left panel since QSOs do not vary much on a two day timescale. The lower two panels display $r^*$ vs. $g^* - r^*$ color-magnitude diagrams and show that we detect no RR Lyrae stars fainter than $r^* \\sim 20$ even though variable QSOs, which are selected by identical criteria, and without using the $u'$ band data, are detected to $r^* > 21$.   The described search procedure deliberately used only $g'$ and $r'$ band data in the first step in order to simplify selection effects, and thus to show that the faint magnitude limit of selected \\RRc\\ is real. For subsequent analysis, however, we further constrain the sample to sources with colors appropriate for RR Lyrae stars ($1.0 <u^*-g^* < 1.5$, $-0.1 <g^*-r^* < 0.4$, $-0.2 <r^*-i^* < 0.2$, $-0.2 <i^*-z^* < 0.2$). These limits are shown as boxes in the color-color diagrams displayed in Figure 3 and result in the final sample of 148 stars. Thus, 80\\% of the sources in the initial sample of 186 stars pass these tight additional criteria, showing that \\RRc\\ can be quite efficiently selected with only two-epoch two-band  ($g'$ and $r'$) data with a sufficiently short baseline to avoid contamination by variable QSOs. About two-thirds of the 38 rejected sources narrowly fail one or two of the imposed color limits, and the remaining third usually fails by more than 0.5 magnitudes in a single color, most often in $u^*-g^*$ and $i^*-z^*$. While the former may be RR Lyrae stars, the latter are more likely to belong to various types of variable binary stars. The inclusion of the rejected sources in the subsequent analysis does not significantly change the resulting volume density of the selected candidates. However, it affects the estimated statistical significance of the observed cut-off in the candidates' magnitude distribution at $r^* \\about 20$ (\\S 4.1). ",
        "conclusions": "\\subsection { Statistical Significance of the Observed Cut-off at R $\\sim$ 50 kpc }  The statistical significance of the observed lack of candidate RR Lyrae stars with $r^* > 20$ can be determined from the expected number of candidates with such magnitudes. However, it is difficult to estimate this number with certainty because of the observed clumpy spatial distribution of selected candidates, and because of only weakly constrained selection effects at the faint end. We estimate the expected number of candidates with $r^* > 20$ by extrapolating the $R^{-3}$ power-law density to infinity (this power-law implies a flat magnitude distribution for sources with uniform luminosity).  A Monte Carlo study of the selection effects described in the Appendix finds that the selection efficiency starts to fall off slowly for $r^* > 19$, reduces to about 50\\% at $r^* \\about 20$, and drops to 0 for $r^* \\about 21.5$. Confirmation of this fall-off comes from analyzing the variable QSO sample. Assuming that the fraction of variable QSOs does not depend on their magnitude (we find this fraction to be $\\about$ 15\\%), the efficiency determined from simulations is in agreement with that implied by the observed numbers of variable QSOs (which, however, may have very different light curves). We adopt 25\\% as the mean efficiency in the $20 < r^* < 21.5$ magnitude range. In order to avoid the effects of the clumpy spatial distribution, \\ we determine the expected density of candidates by considering only those satisfying $160^\\circ < {\\rm RA} < 200^\\circ$ and $15 < r^* < 19$ (c.f. Figure 4). We estimate a density of 18 mag$^{-1}$, normalized to the entire RA  range. With the adopted efficiency, we find that the expected number of candidate RR Lyrae stars with $r^* > 20$ is 7, while we have selected none. The Poisson probability for this outcome is $\\about 10^{-3}$.  A more conservative approach may be taken by considering all 186 candidates from the first selection step (i.e. before the color cuts described in \\S 2 were imposed to yield the final sample of 148 sources). Due to a larger number of candidates, the expected number of sources with $r^* > 20$ is increased to 9. As can be seen in Figures 1a and 2, there are 4 sources which barely missed the cuts and could perhaps be RR Lyrae stars. The Poisson probability that 4 or fewer sources are observed, given the expectation value of 9, is $\\about 0.02$. As evident, this approach significantly reduces the implied statistical significance of the observed cut-off at $r^* \\about 20$.  We conclude that the significance of the observed cut-off is at the level of 2-3 $\\sigma$ (in terms of equivalent Gaussian probability). The best way to improve this estimate is to obtain follow-up observations of the rejected sources to establish whether they are RR Lyrae stars, and of course to analyze data for a significantly larger sky area.    \\subsection {       Comparison with FASTT and LONEOS data      }   The final sample of \\RRc\\ presented here is based on commissioning data taken at only two epochs. In order to estimate the level of spurious variability detections we have cross-referenced our list with the list of variable objects found from the FASTT (Flagstaff Astrometric Scanning Transit Telescope) data (\\cite{HS98}) and with the LONEOS (Lowell Observatory Near Earth Object Search\\footnote{LONEOS homepage is http://asteroid.lowell.edu/asteroid/loneos/loneos\\_disc.html}) database. The LONEOS data fully cover the overlap of runs 745 and 756, while the overlap with the FASTT fields is only partial ($\\sim$ 40 deg$^2$). Both data sets can detect variable stars brighter than $r^* \\sim 17$.   Henden \\& Stone (1998) require that a candidate variable star show night-to-night scatter (typically 8-12 epochs) larger than 3 times the expected error in the magnitude. Based on a random sampling of a set of model light curves with realistic amplitude and period distributions (see Appendix A), we determine that the mean root-mean-square scatter for RR Lyrae stars is $\\sim$ 0.3$^m$. The photometric accuracy of FASTT data (see Figure 1a in \\cite{HS98}) implies that the detection efficiency for RR Lyrae stars in FASTT data should drop sharply for $r^* > 17^m$. The list of selected \\RRc\\ includes 16 sources which are sufficiently bright and in the FASTT fields. We find that 14 of them are indeed in the Henden \\& Stone list as published, and one was detected later from additional observations (A. Henden, priv. comm.). This implies a reliability of our selection procedure of 94\\%. The remaining source, with a mean $r^*=14.9$ (SDSSp J110838.26$-$000514.3, see Table 2 below), is definitely variable in SDSS data, with amplitudes exceeding 0.15$^m$ in all 5 bands (0.34$^m$ in $g^*$). To estimate the efficiency of our two-epoch selection procedure, we matched all sources with appropriate colors (c.f. \\S 2) detected in the SDSS data (8786 sources) to the Henden \\& Stone list and found 33 matches (note that here no magnitude limit is imposed). As we recovered 17 of them as \\RRc, this implies an efficiency of (52$\\pm$15)\\%, assuming that all Henden \\& Stone variable sources with appropriate SDSS colors are RR Lyrae stars.   The sensitivity of the LONEOS data for variability detection drops sharply for $r^* > 17^m$, and 33 stars in our sample are brighter than this limit. For two of these stars no LONEOS data are available (both stars are too close to much brighter objects). Of the remaining 31 stars the LONEOS data clearly indicate variability for 29 (94\\%) and are inconclusive for the other two. Unfortunately, because of their coarse sampling in time, neither the LONEOS data nor the FASTT data can be used to produce light curves which could definitively identify the variable objects as RR Lyrae stars.    \\subsection {Contamination by Other Variable Stars }  While the SDSS multi-color photometry indicates that all candidates discussed here have colors appropriate for RR Lyrae stars, there exist other types of variable star with similar colors. For example, W UMa stars can have F type spectra but they should be efficiently screened out by the selection requirement that candidates must be brighter in $r^*$ when they are bluer in $g^* - r^*$.  The most significant contaminants are probably $\\delta$ Scu stars, another pulsating variable type with approximately the same color range as the RR Lyrae stars, periods of 1-3 hours, and magnitudes about 0.5$^m$ fainter than RR Lyrae stars (Hoffmeister {\\em et al.} 1985). For example, the Tycho photometric survey finds that for $m_B < 11^m$, the numbers of RR Lyrae stars and $\\delta$ Scu stars are comparable. However, the Population I $\\delta$ Scu stars should start to run out at a distance of $\\sim$ 2 kpc from the Galactic plane, or equivalently at $r^* < 13^m$ for $b > 30^\\circ$.  On the other hand, the Population II $\\delta$ Scu stars, also known as SX Phe stars, contribute only 10\\% to the $\\delta$ Scu population (G. Burks, priv. comm.). Thus, the likely contamination of our sample should not be more than 10\\%, which would not qualitatively change our results. We note that this fraction may be somewhat larger if the number of SX Phe stars decreases more slowly with Galactocentric radius than does the number of RR Lyrae stars. In the following section we discuss multi-epoch photometric observations for seven SDSS-selected candidate RR Lyrae stars, and show that all observed light curves have shapes and periods typical for RR Lyrae stars. This result is consistent with the above estimate for the contamination fraction of $\\la$ 10\\%.   Another independent piece of evidence that our sample is dominated by the low-gravity RR Lyrae stars is that the $u^*-g^*$ colors for the selected candidates are redder than those of the parent population. It has been long known that Balmer jumps are bigger for Horizontal Branch stars than for Main Sequence stars (e.g. Oke, Giver \\& Searle 1962, Pier 1983). The resulting $u^*-g^*$ colors for Horizontal Branch stars are redder than for Main Sequence Stars, given the same $g^*-r^*$ color (Lenz et al. 1998). Figure 10 shows the $u^*-g^*$ color distribution for stars selected by $0.0 < g^*-r^* < 0.1$ (c.f. Figure 2, this narrow range of $g^*-r^*$ color selects stars with similar effective temperature, $\\about$ 7000-8000 K, Lenz {\\em et al.} 1998) from the 90569 stars shown in Figure 1a (dashed line) and for stars selected by the same criterion from the resulting sample of candidate RR Lyrae stars (solid line). The $u^*-g^*$ color of candidate RR Lyrae stars is 0.2 mag redder, on average, than that of nonvariable stars within the same narrow range of $g^*-r^*$ color. For Main Sequence stars  the expected $u^*-g^*$ range is 0.85--1.0 (indicated as a horizontal line marked as log(g)=4.5 in the figure) and for Horizontal Branch stars the expected $u^*-g^*$ range is 1.1--1.3 (horizontal line marked as log(g)=2.5), as found by Lenz {\\em et al.} 1998. The intrinsic spread of the $u^*-g^*$ color for a given gravity is due to varying metallicity; the plotted values correspond to the range $-2.0 < [{\\rm M}/{\\rm H}] < 0$. The observed distribution of the $u^*-g^*$ colors of nonvariable stars is consistent with them being a mixture of low-gravity and high-gravity stars (for a detailed study of A stars detected in SDSS commissioning data see Yanny {\\em et al.} 2000a).  The number of candidates found here agrees within 1$\\sigma$ with the normalization given by Wetterer \\& McGraw (1996). While this supports the low level of contamination by other types of variable star, the large uncertainty of their normalization and our efficiency prevents an accurate determination of such contamination in our sample. As already discussed in \\S3, our normalization assumes an efficiency for detecting RR Lyrae stars from two-epoch data of 56\\% determined by the procedure described in Appendix A. This estimate agrees well with the efficiency determined from the comparison with the FASTT data (52$\\pm$15 \\%). Another way to estimate efficiency is to compare the subsamples detected in the 20.9 deg$^2$ overlap of runs 77, 745 and 756. There are 41 candidates selected from runs 745 and 756, and 37 candidates from runs 77 and 745. These two subsamples have 16 sources in common, implying an efficiency of (41$\\pm$12)\\%, in good agreement with the above estimates.     \\subsection { Light Curves for a Subsample of Candidate RR Lyrae Stars }  The accepted identifying characteristics of an RR Lyrae star are its light curve and period. Here we present preliminary results of follow-up observations for several candidate RR Lyrae stars selected from SDSS commissioning data.  One of the faintest candidates in the sample, SDSSp J113049.26$-$005918.2 ($r^* \\about 19.4$), was monitored for five nights during March 2000 with the SDSS 20-inch photometric telescope at Apache Point Observatory, New Mexico. 25 individual measurements in $g'$ band were taken over the five nights; the resulting light curve is shown in Figure 11a. The preliminary estimate for the period is 0.46379 d, and the the first maximum is at HJD 2451606.2323. Both the light curve shape and the period confirm that this candidate is a bona fide RR Lyrae star. The observed apparent magnitude of this star places it at $\\about$ 55 kpc from the Galactic center, and at $\\about$ 44 kpc from the Galactic plane.  Another sample of six candidate RR Lyrae stars was selected from SDSS commissioning data obtained in the southern Galactic hemisphere 59 days apart during the fall of 1998. While these data are not used in this work, the candidates were selected by the identical selection procedures. The candidates were observed in the Johnson V band by the 0.76-m reflector of the University of Washington's Manastash Ridge Observatory during the fall of 1999. These observations were supplemented with data obtained from the LONEOS database and the resulting light curves are shown in Figure 11b, together with the candidates' names and periods. All six candidates have light curves with shapes and periods characteristic of RR Lyrae stars. Their average Galactocentric distance is 30 kpc.  These prefatory follow-up observations show that the majority of stars in our sample are probably RR Lyrae stars, and support the estimate that the contamination by other types of variable star is $\\la$ 10\\%.    \\subsection{               Future Work                  }   These preliminary results obtained with a small sample of SDSS data indicate its potential for various Galactic structure studies. For example, the SDSS-FASTT and SDSS-LONEOS comparisons, and the light curves presented in Figure 11, demonstrate that even two-epoch SDSS photometric data are sufficient for efficient detection of variable stars. The sample of \\RRc\\ presented here shows that such a deep and wide area survey may significantly contribute to studies of the outer Galactic halo. Nevertheless, it is probable that this sample does not contain only RR Lyrae stars, and that some of the assumptions in our analysis may not be valid. The most straightforward approach to determine the contamination level by variable stars other than RR Lyrae stars is to obtain light curves for all candidates (so far only a few hundred, but over the next few years SDSS will produce several thousand \\RRc). To facilitate such observations, Table 2 lists coordinates and two-epoch SDSS photometry in $g'$ and $r'$ bands for 148 candidates discussed in this work\\footnote{The finding charts are available from the authors upon request.}. For the confirmed RR Lyrae stars the measurements of their radial velocity offer the exciting possibility of measuring the distribution of dark matter throughout the halo (\\cite{H84}).    \\vskip 0.4in \\leftline"
    },
    "0004/astro-ph0004306_arXiv.txt": {
        "abstract": "In this paper we present   high spatial resolution radio images of the puzzling binary system \\bly{} obtained with MERLIN at 5\\,GHz. We find a nebula surrounding the binary with a brightness temperature of $(11\\,000\\pm700)$\\,K approximately 40~AU across. This definitively confirms the thermal origin of the radio emission, which is consistent with emission from the wind of the B6-8\\,II component (mass loss of order of $ 10^{-7}~M_{\\sun} \\mathrm{yr}^{-1}$), ionized by the radiation field of the hotter companion. This nebula, surrounding the binary, is the proof that \\bly evolved in a non-conservative way, i. e. not all the mass lost by the primary is accretted by the secondary, and present measurements indicate that  almost 0.015 $M_{\\sun}$ had been lost from the system since the onset of the Roche lobe overflow phase.  Moreover, the nebula is aligned with the jet-like structures inferred from recent optical measurements, indicating a possible connection among them. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004240_arXiv.txt": {
        "abstract": "We show that main-sequence stars in dense stellar cusps around massive black holes are likely to rotate at a significant fraction of the centrifugal breakup velocity due to spin-up by hyperbolic tidal encounters. We use realistic stellar structure models to calculate analytically the tidal spin-up in soft encounters, and extrapolate these results to close and penetrating collisions using smoothed particle hydrodynamics simulations. We find that the spin-up falls off only slowly with distance from the black hole because the increased tidal coupling in slower collisions at larger distances compensates for the decrease in the stellar density. We apply our results to the stars near the massive black hole in the Galactic Center. Over their lifetime, \\( \\sim \\! 1\\, M_{\\odot } \\) main sequence stars in the inner 0.3 pc of the Galactic Center are spun-up on average to \\( \\sim \\! 10\\% \\)--\\( 30\\% \\) of the centrifugal breakup limit. Such rotation is \\( \\sim \\! 20 \\)--\\( 60 \\) times higher than is usual for such stars and may affect their subsequent evolution and their observed properties. ",
        "introduction": "\\label{sec: intro}  It is now widely accepted that super-massive black holes (BHs) exist in many, if not all galactic centers (Magorrian et al. \\cite{Mag98})\\@. Dynamical models of the evolution of such systems generically predict the formation of a dense stellar cusp near the BH (e.g. Bahcall \\& Wolf \\cite{Bah77}; Young \\cite{You80}). Stars moving rapidly in the dense stellar cusp near the BH will suffer numerous high-velocity close tidal encounters over their lifetimes. Although such encounters transfer some energy and angular momentum from the hyperbolic orbit to the colliding stars, they rarely remove enough energy for tidal capture. This is in marked contrast to the situation in the high density cores of globular clusters, where the colliding stars are on nearly zero-energy orbits and close collisions lead to the formation of tight binaries. The effects of hyperbolic encounters on the stars are mostly transient. The dynamical and thermal relaxation timescales are very short compared to the stellar lifetime, and thus apart from some mass-loss in very close collisions, the star is largely unaffected. It is however more difficult for the star to shed the excess angular momentum since magnetic breaking operates on timescales of the order of the stellar lifetime (Gray \\cite{Gra92}). High rotation is therefore the longest lasting dynamical after-effect of a close encounter.  The possibility that stars in BH cusps are rapid rotators may have interesting implications for their evolution and the interpretation of their observed properties. The effects of rotation and rotationally induced mixing in a main-sequence (MS) star on its subsequent evolution has been studied by Sweigart (\\cite{Swe97}) and VandenBerg, Larson \\& De Propris (\\cite{Van98}) in the context of globular clusters, without specifying the origin of the rotation, and by Sills, Pinsonneault \\& Terndrup (\\cite{Sil00}) in the context of young open clusters. Rotationally induced mixing may reveal itself in the spectral line ratios, and rotation may be directly observed in the spectral line profiles. Detection of such signatures in the spectra of the observed giants can provide additional evidence for the existence of an underlying cusp of MS stars, which at present cannot be directly observed.  The goal of this study is to estimate the magnitude of the tidal spin-up, with particular emphasis on the Galactic Center (GC). Present-day observations can already resolve individual giant stars very close to the BH in the GC (Genzel et al. \\cite{Gen97}; Ghez \\cite{Ghe98}) and high resolution infrared spectroscopy is possible for the brighter giants (Carr, Sellgren \\& Balachandran \\cite{Car99}; Ramirez et al. \\cite{Ram00}). The depletion of luminous giants in the inner \\( 2'' \\) (\\( \\sim \\! 0.1 \\) pc) around the supermassive BH in the Galactic Center (GC) was interpreted by Alexander (\\cite{Ale99a}) as evidence for collisional destruction in an extreme density of a sharp stellar cusp. It is inevitable that where the stellar density is high enough to destroy giants, smaller stars that escape destruction will suffer very close collisions. The inner GC is therefore a promising environment for studying the spin-up effect.  The paper is organized as follows. In Sec.~\\ref{sec: soft} we present the formalism for spin-up in the linear regime of soft hyperbolic encounters. In Sec.~\\ref{sec: strong} we discuss results from SPH simulations of non-linear close encounters and incorporate them in our spin-up calculations. In Sec.~\\ref{sec: results} we calculate the spin-up of stars in the inner parsec of the GC. We discuss and summarize our results in Sec.~\\ref{sec: discuss}. An appendix describes the analytic calculation of the tidal coupling constants for hyperbolic encounters using realistic stellar structure models of a MS dwarf and of a giant. ",
        "conclusions": "\\label{sec: discuss}  Dense stellar cusps around massive BHs are environments where stellar collisions are frequent and energetic. Hyperbolic head-on collisions may destroy the colliding stars, but are rare. For every head-on collision there are many more close tidal encounters and grazing collisions. The cumulative effects of such encounters are more subtle. Energy and angular momentum are transfered from the orbit to the stars, during the fly-by the stellar structure may be significantly disturbed by the tidal forces and some mass may be lost. However, the stars survive the collision, and because of the high initial orbital energy, they rarely form a bound binary system. This is in contrast to the situation in dense globular clusters. The stellar dynamical and thermal relaxation timescales are very short relative to the stellar lifetime and the star can radiate the excess energy quickly. It is harder to shed the excess angular momentum, since magnetic breaking operates on timescales of order of the stellar lifetime. Thus the direct long term effects of the collision are increased rotation, and possibly some mass-loss and some mixing of the stellar envelope.  In this study we calculated the magnitude of the rotation that is built up in a random walk fashion as the star undergoes multiple hyperbolic tidal encounters. This is of interest because of the effects high rotation may have on the star's evolution and on its observed properties and because, unlike destructive collisions, the spin-up affects the entire stellar population and extends over a much larger volume of the galactic nucleus. Many, if not most galaxies have a super-massive BH in their nucleus, and so high stellar rotation in galactic nuclei may be common.  Our approach to the problem was to use detailed stellar structure models to calculate, for the first time, the tidal coupling constants for arbitrary hyperbolic orbits in the linear regime of soft encounters. Because the tidal energy and angular momentum fall off as a high power of the peri-separation, it was necessary to extend the calculation to the strongly non-linear regime of grazing and penetrating collisions. We carried out a suite of SPH simulations to study the qualitative behavior of such collisions and then conservatively extrapolated the exact linear calculations to the non-linear regime by several simple prescriptions. The simulations indicated the following. (1) At peri-separations closer than about twice the target star's radius, the linear results are smaller than the actual spin-up by a factor of at least two. (2) Once mass-loss becomes significant, the ejecta carries away a large fraction of the angular momentum that is extracted from the orbit, and therefore the spin-up saturates at its level just before the onset of mass-loss. (3) Stars can survive deep collisions down to a peri-separation of half the target star's radius. (4) A simple tidal disruption criterion for hyperbolic encounters can roughly indicate the point where the fractional mass-loss increases to order unity.  We calculated the spin-up of stars in the inner parsec of the Galactic Center, which is of special interest because of the high quality and wealth of details of the observed stellar data. We find that over 10 Gyr, \\( \\sim \\! 1\\, M_{\\odot } \\) MS stars in the inner 0.3 pc are stochastically spun-up to 10--30\\% of the centrifugal break-up velocity (20--60 times higher than is usual for such stars in the field). This effect decreases only weakly with distance from the BH because the increased tidal coupling at lower collision velocities largely compensates for the decrease in the collision rate at lower stellar densities. We estimated also the stellar survival probability against head-on collisions over 10 Gyr, and found that it was significantly large even very close to the center. The fact that the spin-up is roughly constant over the volume of the inner GC implies that no large error was introduced by neglecting the fact that stars on non-circular orbits sample a varying stellar density over their lifetime. The spin-up of giant stars over the giant phase is much smaller, or order 1--2\\% of the centrifugal break-up velocity (doubling the rotation that is usual for such stars in the field).The effect is smaller because of their short lifetime and because of the large ratio between the collision velocity and the stellar escape velocity, which decreases the tidal coupling.  These results suggest that stochastic spin-up is an important stellar effect in BH cusps. However, several caveats apply. This results depend on the validity of our extrapolation to the non-linear regime. A real star is not well represented by an ideal gas \\( n=1.5 \\) polytrope, and more realistic and extensive hydrodynamical simulations, including the effects of radiation and nuclear burning, will be needed to verify our results. We note however that the \\( n=1.5 \\) polytrope model is significantly less bound than the realistic stellar models (Table~\\ref{tab: stype}). This implies that our estimates for the minimal peri-separation and the tidal radius may be overly conservative. In addition, because a smaller binding energy, larger tidal coupling coefficients (Fig.~\\ref{fig:Tl}) and a larger moment of inertia go hand in hand, the spin-up (\\( \\dW \\sim T_{l}/\\MOI  \\)) is less sensitive to the details of the stellar structure than any of these quantities separately. Another uncertainty in applying our results to any specific system, such as the GC, lies in modeling the PMF of the impactors. Generally, mass segregation will work towards increasing the spin-up effect by pushing the ineffective low mass projectiles out of the central region. Dynamical models of the evolution of the galactic nucleus are required to put this on a quantitative footing. We considered the survival probability of the stars against head-on collisions, but there are also other effects that may compete against the spin-up, which were not included in our estimate. Magnetic breaking, although slow, may be effective over 10 Gyr. Unfortunately, at present the details of this process can not be modeled with any certainty. The cumulative mass-loss in the course of many non-linear collisions was also not considered.  Detailed predictions for the observational consequences of high rotation are outside the scope of this work. We limit our comments on this matter to noting that rotation lowers the effective temperature and luminosity of a star (Kippenhahn, Meyer-Hofmeister \\& Thomas \\cite{Kip70}), but does not significantly affect the spectral classification of a star until it is close to break-up (Gray \\cite{Gra92}). The long term effects of rotation on stellar evolution may be more significant. We calculated the effective angular velocity, assuming solid body rotation. The actual distribution of angular momentum could become stratified over time and lead to rotational support of the core (VandenBerg et al. \\cite{Van98}) or to the replenishment the hydrogen in the core by large scale deep mixing (Sweigart \\cite{Swe97}). These effects will manifest themselves in the giant phase of the stars. Finally, We note that even in the GC only the giant stars can be presently observed. Rotational broadening in the giant spectra may be marginally detectable with high resolution spectroscopy, and could bolster the case for the existence of an underlying very dense population of faint stars.  To summarize, we have shown that MS stars in a substantial volume of the dense cusps around massive black holes are likely to rotate at a significant fraction of the centrifugal break-up velocity due to stochastic spin-up by hyperbolic tidal encounters.  \\appendix"
    },
    "0004/astro-ph0004289_arXiv.txt": {
        "abstract": "We show that the gas in growing density perturbations is vulnerable to the influence of winds outflowing from nearby collapsed galaxies that  have already formed stars.  This suggests that the formation of nearby galaxies with masses $\\lesssim 10^9 M_\\odot$ is likely to be suppressed, irrespective of the details of galaxy formation. An impinging wind may shock heat the gas of a nearby perturbation to above the virial temperature, thereby mechanically evaporating the gas, or the baryons may be stripped from the perturbation entirely if they are accelerated to above the escape velocity. We show that baryonic stripping is the most effective of these two processes, because shock-heated clouds that are too large to be stripped are able to radiatively cool within a sound-crossing time, limiting evaporation.  The intergalactic medium temperatures and star-formation rates required for outflows to have a significant influence on the formation of low-mass galaxies are consistent with current observations, but may soon be examined directly via associated distortions in the cosmic microwave background, and with near-infrared observations from the Next Generation Space Telescope, which may detect the supernovae from early-forming stars. ",
        "introduction": "Outflows from dwarf starbusting galaxies have long been the subject of theoretical and observational investigation.  Since the 1970s, theoretical work has shown that supernovae (SNe) and OB winds in these low-mass objects produce energetic outflows that shock and enrich the intergalactic medium (IGM) in which they are embedded (Larson 1974; Dekel \\& Silk 1986; Vader 1986).  This behavior has been clearly identified in studies of both local starbursting galaxies (Axon \\& Taylor 1978; Heckman 1997; Martin 1998) and their environments (della Ceca et al.\\ 1996; Bomans, Chu, \\& Hopp 1997) as well as with spectroscopy of high-$z$ galaxies (Franx et al.\\ 1997; Pettini et al.\\ 1998; Frye, Broadhurst \\& Sprinrad 1999).  Whether these outflows lead to a catastrophic loss of the interstellar gas in these objects is much more uncertain, with recent simulations and observations suggesting that dwarfs are very inefficient in removing gas from their cores (Mac Low \\& Ferrara 1999; Murakami \\& Babul 1999).  However, the issue of whether dwarfs retain a sizeable fraction of their gas is to some degree decoupled from the formation of outflows. In galaxies in the mass range $10^7 M_\\odot \\lesssim M \\lesssim 10^9 M_\\odot$, Mac Low \\& Ferrara (1999) have shown that a ``blowout'' occurs in which the superbubbles produced by multiple SNII explosions punch out of the galaxy, shocking the surrounding IGM while failing to excavate the interstellar medium of the galaxy as a whole.  In the generally investigated hierarchical models of structure formation, low-mass galaxies form in large numbers at early times (e.g. White \\& Frenk 1991) are highly clustered (Kaiser 1984).  The existence of a large number of small and clustered galaxies at early times is also favored observationally by the steep number counts and low luminosities of faint galaxies (Broadhurst, Ellis \\& Glazebrook 1992), the sizes of faint galaxies in the Hubble Deep Field (HDF) (Bouwens, Broadhurst, \\& Silk 1999a,b), and the clustering properties of Lyman-break galaxies (Adelberger et al.\\ 1998). Locally however, the number density of low-mass objects is far less than predicted theoretically (Ferguson \\& Binggeli 1994; Klypin et al.\\ 1999; Moore et al.\\ 1999), leading to theoretical studies into the disruption of dwarves by tidal forces from neighboring objects (Moore, Lake, \\& Katz 1998) and external UV radiation ( Norman \\& Spaans 1997; Corbelli, Galli, \\& Palla 1997; Ferrara \\& Tolstoy 2000).   Little attention has been directed however, toward the influence of outflows on the formation of neighboring galaxies, despite the fact that typical outflow velocities exceed the virial velocities of objects in the mass range in which the overabundance is most severe. In a companion article (Scannapieco \\& Broadhurst 2000) we describe the effects of heating and enrichment by outflows on galaxy formation using a Monte Carlo treatment of hierarchical structure formation. Here we show using simple scaling arguments that suppression of galaxy formation by outflows from dwarf galaxies is important over a large range of halo masses, irrespective of the details of cosmological simulations.  The structure of this work is as follows.  In \\S2 we consider two scenarios for the suppression of dwarf galaxy formation by outflows. In \\S3 we show that the outflow models considered fall well within the bounds of current observations.  Conclusions are given in \\S4. ",
        "conclusions": "Using simple scaling arguments, we have shown that outflows from low-mass galaxies play an important role in the suppression of the formation of nearby dwarf galaxies.  While the details of this process depend on the spatial distribution of forming galaxies, and can be only be studied in the context of cosmological simulations (Scannapieco \\& Broadhurst 2000), the overall implications of this mechanism can be understood from the general perspective of hierarchical structure formation.  In principle, outflows can suppress the formation of nearby galaxies both by shock heating and evaporation and by stripping of the baryonic matter from collapsing dark mater halos; in practice, the short cooling times for most dwarf-scale collapsing objects suggest that the baryonic stripping scenario is almost always dominant.  This mechanism has the largest impact in forming dwarves in the $\\lesssim 10^9 M_\\odot$ range which is sufficiently large to resist photoevaporation by UV radiation, but too small to avoid being swept up by nearby dwarf outflows.  It is interesting to note that numerical studies working from a different perspective have similarly shown that momentum transfer is the most important feedback mechanism in determining  the internal structure of larger galaxies. While SN heating is relatively ineffective at regulating star formation in numerical models, accounting for the kinetic component of SN feedback has resulted in quasi-stationary models of disc galaxies with reasonable star formation rates (see eg.\\ Katz 1992; Navarro \\& White 1993; Mihos \\& Hernquist 1994; Springel 2000).  Various analytical and N-body studies (Kauffmann, White, \\& Guiderdoni 1993; Klypin et al.\\ 1999; Moore et al.\\ 1999) have shown that $\\sim 50$ satellites with circular velocities $\\sim 20$ km/s and masses $\\sim 10^9 M_\\odot$ should be found within 600 kpc of the Galaxy, while only 11 are observed.  While the fraction of objects suppressed is dependent on the assumed cosmology and spatial distribution of galaxies, the partial suppression of the formation of objects in this mass range is a general consequence of baryonic stripping by outflows. Note that dark matter halos which are subject to sweeping are likely to accrete some gas at later times and thus this scenario provides a natural mechanism for the formation of ``dark'' Milky Way satellites, which may be associated the abundant High-Velocity Clouds as discussed by Blitz et al.\\ (1999).  Also, this population of halos could be identified with the low mass tail distribution of the dark galaxies that reveal their presence through gravitational lensing of quasars (Hawkins 1997).   The existence of an era of widespread IGM heating through outflows is consistent with current CMB constraints, but causes a mean spectral distortion in the range that will be probed by the next generation of experiments.  Similarly, the luminosities of the majority of the starbursting dwarfs that contribute to this process are beyond the limiting magnitudes of the Hubble Deep Field, but may be observable with of the Next Generation Space Telescope.  Thus it may be in the near future that we can directly observe the era in which dwarf galaxies were first formed and the impact of galactic outflows on this process."
    },
    "0004/astro-ph0004183_arXiv.txt": {
        "abstract": "The bulk viscosity of the neutron star matter due to the direct Urca processes involving nucleons, electrons and muons is studied taking into account possible superfluidity of nucleons in the neutron star cores. The cases of singlet-state pairing or triplet-state pairing (without and with nodes of the superfluid gap at the Fermi surface) of nucleons are considered. It is shown that the superfluidity may strongly reduce the bulk viscosity. The practical expressions for the superfluid reduction factors are obtained. For illustration, the bulk viscosity is calculated for two models of dense matter composed of neutrons, protons, electrons and muons. The presence of muons affects the bulk viscosity due to the direct Urca reactions involving electrons and produces additional comparable contribution due to the direct Urca reactions involving muons. The results can be useful for studying damping of vibrations of neutron stars with superfluid cores. ",
        "introduction": "The dissipative processes in neutron stars play an important role for some dynamical properties of these unique objects. Shear viscosity  damps differential rotation of neutron stars, leading to their uniform rigid-body rotation. Quite generally, the viscosity of neutron star matter implies damping of pulsations of neutron stars. Such pulsations could be excited during the process of neutron star formation. They could also be triggered by instabilities appearing during neutron star evolution, or by some external perturbations. The corresponding damping timescales involve  the shear and bulk viscosities of neutron star interior. Both viscosities depend on density, temperature and composition of dense matter. Calculations of damping timescales of pulsations for various models of neutron star interiors have been done by Cutler et al.\\ (\\cite{cls90}). Viscous  damping of pulsations of newly born hot neutron stars turns out to be due to the bulk viscosity.  Another role of the viscosity of neutron star matter is that it can damp gravitational radiation driven instabilities in rotating neutron stars and, therefore, could be important for determination of the maximum rotation frequency of neutron stars. In the absence of  viscosity all  rotating neutron stars would be driven unstable by the emission of gravitational waves. Viscous damping timescales enter explicitly the criteria for the appearance  of these  instabilities. Similarly, as in pulsating non-rotating neutron stars, viscous damping of gravitational radiation driven instabilities in rapidly rotating newly born neutron stars is dominated by bulk viscosity of neutron star interiors (e.g., Lindblom \\cite{l95}, Zdunik 1996, Lindblom et al.\\ \\cite{lom98}).  In this paper, we focus on the viscosity of matter in the neutron star cores (which extend from the layers of density $\\rho \\simeq 1.5 \\times 10^{14}$ g cm$^{-3}$ to the stellar centers). It is well known that the cores consist of baryons (neutrons $n$, protons $p$ and possibly hyperons) and leptons (electrons $e$ and muons $\\mu$). All constituents of matter are strongly degenerate fermions. The electrons and muons form almost ideal Fermi gases. The electrons are ultrarelativistic while the muons may be non-relativistic or relativistic depending on density. The nucleons are, to a good approximation, non-relativistic and constitute strongly interacting Fermi liquid. At the densities close to the normal nuclear density (baryon number density $n_0=0.16$ fm$^{-3}$ which corresponds to the mass density $\\rho_0=2.8 \\times 10^{14}$ g cm$^{-3}$), neutron star matter is composed of $n$, $p$, $e$, and $\\mu$. At higher densities [$\\rho \\!\\ga(3$--$4)\\,\\rho_0$] some models of dense matter predict appearance of hyperons. At still higher densities, some calculations indicate possible presence of exotic phases (pion condensate, kaon condensate, deconfined quark matter). We will not consider the hyperonic or exotic phases but restrict ourselves to the study of the $npe\\mu$ matter.  Our analysis is additionally complicated by possible superfluidity of nucleons in the neutron star cores. The superfluidity is thought to be produced by Cooper pairing of nucleons due to attractive parts of nucleon-nucleon interaction. The superfluidity of nucleons in the neutron star cores has been the subject of numerous papers (as reviewed, for instance, by  Yakovlev et al.\\  \\cite{yls99}). Various microscopic theories predict very different superfluid gaps (critical temperatures $T_{cn}$ and $T_{cp}$) of neutrons and protons depending on specific model of strong interaction employed and specific many-body theory used to account for medium effects. However, all these results have important common features. In particular, the proton pairing occurs mainly in the $^1$S$_0$-state since the $pp$ interaction is attractive in this state everywhere in the neutron star core due to not too high number density of protons. As for the neutrons, their interaction in the $^1$S$_0$ state turns from attraction to repulsion at densities $\\rho \\ga \\rho_0$ but the interaction in the $^3$P$_2$ state may be attractive and may lead to Cooper pairing. The critical temperatures $T_{cn}$ and $T_{cp}$ in the neutron star cores predicted by different microscopic theories depend on density and scatter in a wide range from about $10^8$ to $10^{10}$ K. Under these conditions we will not rely on any specific microscopic theory of nucleon superfluidity, but will treat $T_{cn}$ and $T_{cp}$ as free parameters.  The viscosity, we are interested in, is well known to consist of the shear viscosity and bulk viscosity. The standard source of the shear viscosity of the neutron star matter is scattering between its constituents. Classical calculations of shear viscosity for the $npe$ model of non-superfluid matter were done by Flowers \\& Itoh (\\cite{fi79}). Their results were used in the studies of damping of neutron star pulsations by Cutler et al.\\ (\\cite{cls90}). In the superfluid core of a rotating neutron star, there is an additional viscous mechanism, called mutual friction,  resulting from the scattering of electrons off the magnetic field trapped in the cores of superfluid neutron vortices (Lindblom \\& Mendell \\cite{lm95}).  The bulk viscosity may partly be determined by particle scattering. However, this component of bulk viscosity is usually much smaller than the shear viscosity (e.g., Baym \\& Pethick \\cite{bp91}). The main contribution into the bulk viscosity of sufficiently hot $npe\\mu$ matter comes from the neutrino processes of Urca type associated with electron and muon emission and capture by nucleons. We will focus on such bulk viscosity. Generally, the neutrino processes in question are divided into the {\\it direct Urca} and the {\\it modified Urca} processes. A direct Urca process is a sequence of two reactions (direct and inverse one) and can be written as \\begin{equation} n \\to p + l + \\bar{\\nu}_l, \\quad p + l \\to n + \\nu_l, \\label{baryon-Durca} \\end{equation} where $l$ is either electron or muon, and $\\nu_l$ is an associated neutrino. The most important is  the process (Lattimer et al.\\ \\cite{lpph91}) involving electrons ($l=e$); it consist of the beta decay of neutron and subsequent beta capture. It should be emphasized that the both direct Urca processes are forbidden by momentum conservation of reacting particles for the simplest model of dense matter as a gas of noninteracting Fermi particles (e.g., Shapiro \\& Teukolsky \\cite{st83}) at any density $\\rho$ in the neutron star cores. Nevertheless they are allowed (Lattimer et al.\\ \\cite{lpph91}) for some realistic equations of state at densities higher than some threshold densities (of several $\\rho_0$). Thus, the direct Urca processes may be open in the inner cores of rather massive neutron stars. The threshold density for the muon process is always higher than for the electron one.  If allowed, the direct Urca processes produce the most powerful neutrino emission from the neutron star cores (Lattimer et al. \\cite{lpph91}). Corresponding neutrino emissivities were calculated by Lattimer et al.\\ (\\cite{lpph91}) and used in numerous simulations of the neutron star cooling as reviewed, for instance, by Yakovlev et al.\\ (\\cite{yls99}). In the absence of nucleon superfluidity, the direct Urca processes lead to the {\\it fast} cooling of neutron stars. If allowed, the direct Urca processes produce the main contribution into the bulk viscosity we are interested in.  However, for many equations of state the direct Urca processes are forbidden by momentum conservation at any density in the neutron star cores. Moreover, they are prohibited at $\\rho\\la 3\\, \\rho_0$ for the majority of equations of state. In such cases, they do not operate in the low and medium-mass neutron stars and in the outer cores of all neutron star models constructed using these equations of state. If so, the bulk viscosity is determined by the the reactions of the modified Urca processes \\begin{equation} n + N \\to p + N + l + \\bar{\\nu}_l, \\quad p + N + l \\to n + N + \\nu_l, \\label{baryon-Murca} \\end{equation} where $N$ is an additional nucleon required to conserve momentum of the reacting particles. For instance, in $npe$ matter one has two modified Urca processes corresponding to $N=n$ and $N=p$, respectively, which can be referred to as the neutron and proton branches of the modified Urca process (e.g., Friman \\& Maxwell \\cite{fm79}, Yakovlev \\& Levenfish \\cite{yl95}). The rates of the modified Urca processes are typically 3--5 orders of magnitude lower than the rates of the direct Urca processes. The modified Urca processes either have no density threshold (as the neutron branch in $npe$ matter) or have much lower density thresholds than the direct Urca processes. Thus they operate in the entire neutron star core. If the direct Urca processes are forbidden and matter is non-superfluid, the modified Urca processes produce the main neutrino emission from the neutron star cores leading to {\\it slow} ({\\it standard}) cooling of neutron stars. Their role in the neutron star cooling theory has been studied in many papers (see, e.g., Yakovlev et al.\\ \\cite{yls99}, for review).  Thus, the problem of calculating the bulk viscosity due to neutrino processes is quite complicated: there are several neutrino processes involved influenced by possible nucleon superfluidity. So far the bulk viscosity has been studied only for non-superfluid $npe$ matter. The viscosity due to the neutron branch of the modified Urca process was analyzed by Sawyer (\\cite{s89}) while the viscosity produced by the nucleon direct Urca process was considered by Haensel \\& Schaeffer (\\cite{hs92}). The effects of superfluidity have not been analyzed for the problem of bulk viscosity but studied thoroughly for the neutrino emissivity produced in different reactions (e.g., Yakovlev et al.\\ \\cite{yls99}, and references therein).  The relative importance of the bulk viscosity produced by neutrino reactions with respect to the shear viscosity produced by collisions can be estimated by comparing the results by Sawyer (\\cite{s89}) and Haensel \\& Schaeffer (\\cite{hs92}) with the values of the shear viscosity calculated by Flowers \\& Itoh (\\cite{fi79}). The comparison shows that the neutrino bulk viscosity dominates in the neutron star cores for temperatures $T \\ga 10^8$ K if the direct Urca processes are switched on and for $T \\ga 10^9$ K if the direct Urca processes are forbidden. In superfluid matter, the bulk viscosity can be even more important.  In this paper, we consider the bulk viscosity produced by the direct Urca processes in $npe\\mu$ matter of the neutron star cores. In analogy with the effects of superfluidity on the neutrino emissivity, we will analyze the effects of superfluidity of nucleons on the bulk viscosity. ",
        "conclusions": "We have derived practical expressions for the bulk viscosity of matter in the cores of neutron stars under conditions in which the bulk viscosity is determined by the direct Urca processes. We have paid special attention to the case in which dense matter consists of neutrons, protons, electrons and muons (Sects.\\ 2.1--2.4). In addition, we have studied the reduction of the bulk viscosity by superfluidity of neutrons and protons (Sect.\\ 3). We have analyzed the cases of singlet-state superfluidity of protons, and triplet-state superfluidity of neutrons (without and with the nodes of superfluid gaps on the nucleon Fermi surface). These cases are most interesting for applications (Sect.\\ 1). We have obtained the practical expressions for the bulk viscosity of superfluid $npe\\mu$ matter. The results can be used for studying the damping of neutron star pulsations (Sect.\\ 1) and in the studies of the gravitational radiation driven instabilities in rotating neutron stars. We will analyze the bulk viscosity induced by the weaker modified Urca processes (\\ref{baryon-Murca}) in a separate paper."
    },
    "0004/astro-ph0004368_arXiv.txt": {
        "abstract": "We report the detection of a new H$_2$O megamaser in the Seyfert 2 galaxy Mrk~348 with the MPIfR 100 m telescope in Effelsberg. With an apparent isotropic luminosity of $L_{\\rm H_2O}\\simeq420L_\\odot$ the maser is the third most luminous maser discovered so far. The detected line is unusually broad ($\\Delta v\\simeq130$ km s$^{-1}$), is redshifted by $\\sim130$ km s$^{-1}$ from the systemic velocity, and is asymmetric with a pronounced blue wing. While evidence for obscuring material towards the nucleus of this galaxy was found earlier, this detection is the first direct observation of molecular material in the vicinity of the AGN. We also searched for absorption from ammonia (NH$_3$) and cyclopropenylidene (C$_3$H$_2$) against the bright radio nucleus. The H$_2$O line was only marginally detected in archival data indicating that the maser flared recently in conjunction with a huge radio continuum flare. Continuum and line flux density increased by a factor of three, suggesting an unsaturated maser. The radio continuum flare has made Mrk~348 the most radio luminous megamaser galaxy known. It is pointed out that megamaser galaxies contain a rather large fraction of galaxies with prominent radio cores and it is speculated that the flare in the maser emission in Mrk~348 is related to the flare in the nuclear jet. ",
        "introduction": "Emission from H$_2$O masers has been found in a few galaxies, exhibiting apparent isotropic luminosities a million times higher than in typical stellar masers \\cite{DosSantosLepine1979,GardnerWhiteoak1982,ClaussenHeiligmanLo1984,HenkelGuestenWilson1984,HaschickBaan1985,BraatzWilsonHenkel1994}. The detection rate of these megamasers is very low, i.e.~about 5\\% among Seyfert galaxies \\cite{BraatzWilsonHenkel1997} and almost zero among radio galaxies (e.g.,~\\citeNP{HenkelWangFalcke1998}). The maser is associated with dense and warm material, possibly a molecular torus or disk, around an active galactic nucleus (AGN). The AGN apparently produces the seed radio photons and the X-ray photons needed to pump the masing material \\cite{NeufeldMaloneyConger1994}.  With the help of Very Long Baseline Interferometry (VLBI) megamasers can be used to investigate the small-scale structure of an AGN in great detail. In the case of NGC~4258 this has helped to establish the presence of a thin, warped disk around the nucleus, to determine the black hole mass, and even to measure the precise distance to this galaxy \\cite{MiyoshiMoranHerrnstein1995,HerrnsteinMoranGreenhill1999}.  Finding new megamaser galaxies is therefore of prime interest. The only clear trend that has emerged in recent years is that megamasers are exclusively found in type 2 AGN, i.e.~those Seyferts and LINERs which are expected to be obscured by a molecular torus according to the unified scheme \\cite{Antonucci1993}. Many have high absorbing column depths inferred from X-ray spectroscopy. There is also an indication of an excess of megamasers in highly inclined galaxies \\cite{BraatzWilsonHenkel1997,FalckeWilsonHenkel2000}. Here we report the discovery of a hitherto undetected and very luminous megamaser in the Seyfert galaxy Mrk~348 during a radio flare of the AGN.  Mrk~348 (NGC~262, $z=0.01503$ \\citeNP{HuchraVogeleyGeller1999}, luminosity distance $D=62.5$ Mpc for $z$ converted into the Galactic Standard of Rest and $H_{0}=75$ km s$^{-1}$ Mpc$^{-1}$), is a well-studied Seyfert 2 galaxy with broad emission-lines in polarized light \\cite{MillerGoodrich1990}. The galaxy is classified as an S0 with a rather low inclination ($i=16^\\circ$, see \\citeNP{BraatzWilsonHenkel1997}).  Ground-based \\cite{SimpsonMulchaeyWilson1996} and Hubble Space Telescope imaging \\cite{FalckeWilsonSimpson1998} have revealed a dust lane crossing the nucleus and an excitation cone in emission-lines. Ginga observations found hard X-ray emission and a high absorbing column depth of $N_{\\rm H}=10^{23.1}\\,{\\rm cm}^{-2}$ towards the nucleus \\cite{WarwickKoyamaInoue1989}. All this suggests the presence of an obscuring torus in Mrk~348. Attempts to detect the obscuring material through radio spectroscopy have failed so far (e.g.~H I, \\citeNP{GallimoreBaumO'Dea1999}).  What makes this galaxy stand out among Seyfert galaxies is its bright and variable radio nucleus. \\citeN{NeffdeBruyn1983} found a compact radio core on VLBI scales with a flux density of several hundred milli-Jansky and a flat to inverted spectrum. \\citeN{UlvestadWrobelRoy1999} presented more recent VLBI observations for two epochs, showing a two-component structure expanding with sub-relativistic speeds. They also noted a flare of the radio continuum emission at 15 GHz with the flux rising from 120 mJy to 570 mJy between 1997.10 and 1998.75. In the following we will present and discuss results of K-band radio spectroscopy of this galaxy. ",
        "conclusions": "We have clearly discovered a new megamaser in Mrk~348.  Its luminosity is comparable to the emission from NGC~3079 which contains the second most luminous H$_2$O maser after TXS~2226-184 \\cite{KoekemoerHenkelGreenhill1995}. The line width is among the broadest found for a megamaser, similar to the masers in TXS~2226-184 and NGC~1052. Assuming the emission is associated with material close to the center this is the first spectroscopic evidence for molecular gas possibly obscuring the nucleus.  Despite the bright radio continuum and the molecular maser line we have found no absorption lines from either NH$_3$ or C$_3$H$_2$. This is in line with the earlier non-detection of H {\\sc I} absorption \\cite{GallimoreBaumO'Dea1999}. C$_3$H$_2$ is an organic ring molecule which is widespread in the Galaxy and is associated with diffuse gas in the ISM \\cite{MatthewsIrvine1985}. It was also found in absorption against the nucleus of the radio galaxy Centaurus A \\cite{SeaquistBell1986}. From our non-detection we find a 1$\\sigma$ upper limit for the optical depth $\\tau$ times the covering factor $f$ of $\\tau\\cdot f<0.0075$. This is seven times smaller than the value found for Centaurus A and might be related to the fact that we see Mrk~348 almost face on.  On the other hand, with its type 2 AGN, polarized broad emission-lines, and a nuclear dust lane seen, Mrk~348 falls right into the roster of typical megamaser galaxies, where it is suggested that the masing material is part of the obscuring `torus' in the unified scheme of AGN. The face-on orientation of Mrk~348 then would suggest that the axis of this torus and the galaxy disk axis are severely misaligned.  In October 1997 -- February 1998 the flux density of the line was three times lower than it was in March 2000. In the earlier survey by \\citeN{BraatzWilsonHenkel1996} the maser line was not detected and given its line width and low flux density it would have resulted only in a broad feature at the $3\\sigma$ level. Linear interpolation with time of the continuum flux density given in \\citeN{UlvestadWrobelRoy1999} suggests a continuum flux density for Mrk~348 around 310 mJy in October 1997 at 15 GHz. The level we measure is roughly a factor of three higher -- assuming a flat spectrum -- and this increase in flux density is similar to the increase in flux density of the line. This could indicate a correlation between continuum and maser flux density, implying an unsaturated maser. With its current radio luminosity the galaxy is now the most radio luminous megamaser galaxy ever discovered. The response of the line to the continuum flare within about 2 years sets an upper limit to the distance of the masers from the nucleus of $\\la$0.6 pc, which is of similar order as the size scale of the molecular disk found in NGC~4258.  Indeed, \\citeN{BraatzWilsonHenkel1996} noted a certain excess of detected megamaser galaxies with large radio powers. Mrk~348 certainly adds to this trend. The distribution of radio powers at 6 cm of the parent samples of AGN selected by \\citeN{BraatzWilsonHenkel1996} has a peak around $10^{21.75}$ Watt/Hz. If we add Mrk~348 and more recently discovered megamasers \\cite{GreenhillHerrnsteinMoran1997,HagiwaraKohnoKawabe1997} and also complement the radio data in \\citeN{BraatzWilsonHenkel1996} with more recent data from the NASA Extragalactic Database, we find that 12 out of 19 detected megamaser galaxies are at or above the peak in the distribution of radio power for galaxies without megamaser detections.  Because of the possibly biased selection of the detected megamasers this is not highly significant. However, it highlights an apparently necessary prerequisite for megamaser emission, namely an AGN with a compact radio core to provide seed photons. So far all detected megamasers have compact ($<1\\arcsec$), mostly flat-spectrum, radio emission at a level of a few milli-Jansky. In some cases, like Mrk~348 \\cite{NeffdeBruyn1983}, NGC~1052 \\cite{ShafferMarscher1979}, Mrk~1210 \\cite{HeislerNorrisJauncey1998}, NGC~2639 \\cite{WilsonRoyUlvestad1998}, NGC~3079 \\cite{TrotterGreenhillMoran1998}, NGC~5793 \\cite{WhiteoakGardner1987,GardnerWhiteoakNorris1992}, NGC~5506 \\cite{SadlerSleeReynolds1995}, and possibly NGC~4945 \\cite{ElmouttieHaynesJones1997} the radio cores can even reach several tens to hundreds of milli-Jansky.  While compact radio cores in Seyfert and LINER galaxies are not uncommon, only very few are so prominent as those in some of the radio-bright megamaser sources. We find that all megamasers mentioned above, i.e.~more than a third of known megamasers, have compact radio cores above a fiducial limit of 25 mJy at 5 GHz. On the other hand, in a survey of spiral galaxies \\citeN{SadlerSleeReynolds1995} find only 3 out of 54 galaxies (22 of which are Seyfert galaxies) with compact cores above 25 mJy at 5 GHz. Similarly, in a survey of nearby AGN \\citeN{NagarFalckeWilson2000} find roughly 40\\% of Seyfert and LINER galaxies to contain compact flat-spectrum radio cores. However, only three out of 48 galaxies have flux densities $>$25 mJy.  Known megamasers therefore seem to prefer galaxies with relatively bright compact radio emission.  Mrk~348 currently has the brightest and most prominent radio core among megamaser galaxies.  Morphology, spectrum, and variability of the core are very similar to the radio core in III~Zw~2 which was the first Seyfert galaxy discovered to contain a superluminal jet. This galaxy has a millimeter-peaked spectrum and a jet which shows a stop-and-go behavior indicative of a strong interaction with dense material on the sub-parsec scale \\cite{FalckeBowerLobanov1999,BrunthalerFalckeBower2000}. \\citeN{UlvestadWrobelRoy1999} therefore speculate whether the bright inverted radio core in Mrk~348 could be interpreted similarly to those in Compact Symmetric Objects (CSOs) with a Gigahertz-Peaked-Spectrum (GPS, see \\citeNP{O'Dea1998}). In these galaxies bright hotspots are formed in a jet that terminates already on the parsec scale. In III~Zw~2 and Mrk~348 this seems to happen on even smaller scales, leading to higher peak frequencies and could be due to frustration of the jet by a molecular cloud or even a warped or misaligned torus.  Since the masers in NGC~1052, which have similar broad line widths as in Mrk~348, are found along the radio jet \\cite{ClaussenDiamondBraatz1998} it should be checked whether in Mrk~348 one has an analogous situation. One can speculate that in such a case the evolution of the radio flare and the evolution of the maser flare and its blue wing could be related, possibly providing a unique diagnostic tool to study jet-ISM interactions.  In any case, with its bright radio core Mrk~348 provides an ideal opportunity to observe the maser lines in this galaxy at high resolution with VLBI during this flare even though the lines still have a rather low flux. Since radio and maser emission seem to be highly variable both should be monitored frequently. Given that Mrk~348 was not discovered in an earlier survey this finding also suggests that existing samples should be revisited to search for more flaring megamasers."
    },
    "0004/astro-ph0004197_arXiv.txt": {
        "abstract": "Recent VLBA observation indicates the existence of an elongated (jet) structure in the compact radio source Sgr A$^*$. This is hard to explain in the context of advection-dominated accretion flow (ADAF) model for this source. On the other hand, the mass accretion rate favored by ADAF is 10-20 times smaller than that favored by the hydrodynamical simulation based on Bondi capture. If the latter were adopted, the predicted radio flux would significantly exceed the observation. A similar situation exists in the case of nearby giant ellipticals, where the canonical ADAF model -- the widely assumed standard model for these sources -- also significantly overpredict the radio flux. Based on these facts, in this paper we propose a truncated ADAF model for Sgr A$^*$ and three ellipticals M87, NGC 4649 and NGC 4636. We assume that the accretion disk is truncated at a certain radius $R_{\\rm tr}$ within which the jet forms by extracting the energy of the disk. The radio flux is greatly suppressed due to the radiative truncation of the disk and the fits to the observational data are excellent. For example, for Sgr A$^*$, the model fits the observational spectrum very well from radio including the ``excess'' below the break frequency to hard X-ray under a high accretion rate near the simulation value, and the predicted size-frequency relationship is also in excellent agreement with the observation; for M87, the predicted upper limit of the jet location is $24 R_{\\rm g}$, in excellent agreement with the observational result that the jet is formed on scales smaller than $30 R_{\\rm g}$, and the $\\approx 20 \\%$ variability at $\\sim$ 1 Kev -- which is hard to explain in another model succeeded in explaining the low radio flux of M87 -- is also marginally interpreted. The success of the model supplies a possible evidence for the disk rather than the hole origin for the powering of jets. ",
        "introduction": "The energetic radio source Sgr A$^*$ located at the center of our Galaxy is widely believed to be a black hole with mass $M = 2.5 \\times 10^6 \\msun$ (Mezger, Duschl \\& Zylka 1996). Due to its proximity, the observational data are abundant, which makes it an ideal laboratory to test our accretion theories. Any successful model needs to explain its relative low-luminosity, the emergent spectrum from radio to X-ray wavelength, and the frequency-size relationship.  Several models have been put forward for Sgr A$^*$. They include Bondi accretion (Melia 1992; 1994), emission from mono-energetic electrons (Duschl \\& Lesch 1994), jet model (Falcke, Mannheim \\& Biermann 1993; Falcke \\& Biermann 1999), and the most recent two temperature advection-dominated accretion flows (ADAFs; Narayan, Yi \\& Mahadevan 1995; Manmoto, Mineshige, \\& Kusunose 1997; Narayan et al. 1998). Among them, ADAF model is the most dynamically self-consistent. However, two puzzles still remain in this model. First, as pointed out by Quataert \\& Narayan (1999), the mass accretion rate required in this model is more than 10 times smaller than that favored by three-dimensional hydrodynamical simulation based on the Bondi capture, with the former being $ \\approx 6.8 \\times 10^{-5} \\dot{M}_{\\rm Edd}$ (Quataert \\& Narayan 1999) and the latter $ \\approx 9 \\times 10^{-4} \\dot{M}_{\\rm Edd}$ (Coker \\& Melia 1997). Here $\\dot{M}_{\\rm Edd} =2.2 \\times 10^{-8} {\\rm M}{\\rm yr}^{-1}$ denotes the Eddington accretion rate and $M$ is the mass of the centre hole. Very recent Chandra observations of Sgr A$^*$ show that its $0.1 - 10$ kev luminosity is $L_X \\approx 4 \\times 10^{33} \\ergsec$ (Baganoff et al. 2000), considerably smaller than the upper limit used by Quataert \\& Narayan (1999). Taking this into account, the accretion rate required in ADAF model would be smaller hence the discrepancy between the two accretion rates will be larger. If a higher accretion rate were adopted in an ADAF model, the bremsstrahlung and synchrotron radiations would yield an X-ray and a radio fluxes well above the observational limits.  Yuan (1999) find that the outer boundary condition of the accretion flow plays an important role in the dynamics and radiation of an optically thin accretion flow. This point was neglected in previous ADAF models in the literature. After considering this effect, it was found that since the specific angular momentum of the flow at the outer boundary is low, $\\Omega_{\\rm out} \\sim 0.15 \\Omega_{\\rm Kepler}$ as shown by the numerical simulation (Coker \\& Melia 1997), the accretion in Sgr A$^*$ should belong to Bondi-like (not Bondi type), characterized by a much larger sonic radius and a low surface density compared with the conventional disk-type ADAF. The bremsstrahlung radiation therefore is greatly suppressed (Yuan et al. 2000). However, the synchrotron radiation would still produce too much radio flux were the mass accretion rate favored by the simulation adopted.  The second puzzle is that ADAF model strongly underpredicts the radio flux in the radio spectrum below a ``break'' frequency (Narayan et al. 1998). Mahadevan (1998) argued that the charged pion produced by the energetic proton collisions in ADAF will subsequently decay into $e^-$$e^+$, which will further produce synchrotron emission and could serve as the origin of the radio excess below the break. However, the result depends on the power-law index of the proton energy spectrum, which we are not clear. In fact, due to the existence of the threshold energy needed for pion production, these extra photons should be principally produced in the innermost region of the disk. This is in conflict with the observed intrinsic size of $\\sim 72 R_{\\rm g}$ at 43 GHz (Lo et al. 1998; see also Figure 2), here $R_{\\rm g}= 2GM/c^2$ is the Schwarzschild radius of the black hole.  More importantly, recent excellent near-simultaneous VLBA measurements show that the intrinsic source structure at 43 GHz is elongated along an essentially north-south direction, with an axial ratio of less than 0.3 (Lo et al. 1998). This is hard to explain in an ADAF model, while it provides a convincing evidence for the existence of a jet in the jet-nozzle model of Falcke et al. (1993). In this model the radio spectrum below the break is produced by a jet while those above the break is assumed to be produced by a ``nozzle'' located at the lower base of the jet. The question is, however, this model can not account for the entire spectral energy distribution from radio to X-ray, and we lack any detailed understanding to the ``nozzle'' which is crucial to the fit of the spectrum above the break. Considering this situation, recently Falcke (1999) and Donea, Falcke, \\& Biermann (1999) suggested that maybe we should combine the jet model with an accretion model, e.g., ADAF, Bondi-type accretion or others.  A similar situation exists in the case of nearby elliptical galaxies. The existence of a massive black hole and the hot gas pervading the galactic center is in sharp contrast to their low luminosity. Recently it was suggested that ADAF could serve as a feasible explanation for the quiescence of the galaxies (Fabian \\& Rees 1995). Recent observation, however, shows that the canonical ADAF model overpredicts the radio flux by more than two orders of magnitude (Di Matteo et al. 1999). The current explanation is to drive a wind from the accretion disk, assuming the accretion rate $\\dot{M}=(R/R_{\\rm out})^p\\dot{M}_{\\rm out}$ (Di Matteo et al. 2000), where $R_{\\rm out}$ is the radius from which the wind starts to play an important role and $\\dot{M}_{\\rm out}$ is the mass accretion rate there. Since the electron is basically adiabatic, its temperature is determined by the compression work. When a strong wind is introduced, the density profile would become flatter which results in a significant decrease of the temperature and further the synchrotron radiation. The presence of a strong wind is based on the assumption that the Bernoulli parameter of an ADAF is positive therefore the flow is in some sense unstable (Blandford \\& Begelman 1999). As shown by Nakamura (1998) and Abramowicz, Lasota, \\& Igumenshchev (2000), however, the Bernoulli parameter is negative in most of the reasonable cases. Numerous numerical simulations also confirm the absence of winds from the disks (but jets do present in the innermost region of the disk in some cases) (Eggum, Coroniti, \\& Katz 1988; Molteni, Lanzafame, \\& Chakrabarti 1994; Igumenshchev \\& Abramowicz 1998; Stone, Pringle, \\& Begelman 1999; Koide et al. 2000). Thus the scenario of strong winds from ADAFs might be questionable. Even so, the fit is not very satisfactory in some aspects, as illustrated below. What is interesting is, all the sources in Di Matteo et al. (2000)'s sample have the extended and often highly collimated radio structures which are again the evidence for the existence of jets (e.g. Stanger \\& Warwick 1986).  Based on the above facts, in this paper, we propose an ADAF-jet model for Sgr A$^*$ and nearby elliptical galaxies. ",
        "conclusions": "The predicted radio flux for nearby ellipticals by the canonical ADAF model is well above the observations. If the mass accretion rate favored by the numerical simulation were adopted, the ADAF model would also produce too much radio flux in the case of Sgr A$^*$. On the observational side, strong evidence for the existence of jets in Sgr A$^*$ and these ellipticals are observed. Based on these facts, in this paper we propose an ADAF-jet model for Sgr A$^*$ and the nearby giant elliptical galaxies. We assume that  the advection-dominated accretion disk is truncated by a jet at a certain radius $R_{tr}$ within which most of the energy of the disk (if it still exists) is extracted for the powering of the jet therefore its radiation is completely suppressed. We calculate the spectrum of the ADAF-jet system for Sgr A$^*$ and three ellipticals and find that the radio overprediction are solved and the fits to the observation are excellent. For example, for Sgr A$^*$, we can reproduce the break around the 50 GHz point in the radio wavelength, and the predicted size-frequency relation is in good agreement with the observation. For M87, the truncation radius favored by our model, $R_{\\rm tr} \\approx 24 R_{\\rm g}$, agrees with the most recent joint observation by VLBI, VLA and other instruments which shows the jet should originate from within $\\sim 30 R_{\\rm g}$ to the center black hole.  In our model we assume that the disk is radiatively truncated at radius $R_{tr}$ by the jet. It is interesting to note that this idea has been invented in the study to the black hole X-ray binary GRS 1915 +105 where we have much better observational constrains. Radio-infrared oscillation with periods of 10-60 min have been found to follow X-ray dips (Mirabel et al. 1998). The oscillation has been found to be due to synchrotron emission from repeated small ejections, i.e., jet, possibly resulted from suffering strong adiabatic expansion losses (Fender et al. 1997; Eikenberry et al. 1998), while the X-ray dips have been interpreted as the repeated {\\em disappearance} of the inner accretion disk (Belloni et al. 1997).  The suppression of the radio flux also can be accessed by ADAF with strong winds (Blandford \\& Begelman 1999; Quataert \\& Narayan 1999; Di Matteo et al. 2000). Our truncated ADAF model is different with their wind model. The wind model is based on the assumption that the Bernoulli parameter of ADAF is positive therefore strong winds are inevitable. The validity of this assumption is in debate at present. On the other hand, our model is based on the observation that jets really exist in the inner region of these sources therefore may has more reliable physical basis. The predictions of these two models are also different. Comparisons between the prediction and observation on the location of the jet origin in M87, the $\\approx 20\\%$ variability at $\\sim 1$ keV of M87, and the frequency-size relationship of Sgr A$^*$, incline to support our model.  Noting that in nearby ellipticals and Sgr A$^*$, two kinds of sources where ADAF is assumed to exist, jets are observed. In NGC 4258, another source where ADAF is assumed to exist (Gammie, Narayan, \\& Blandford 1999), a jet also exists (e.g., Herrnstein et al. 1998). Moreover,  in black hole X-ray binary a steady jet is usually present at the hard/low state, while this state is typically interpreted as an inner advection dominated accretion disk plus an outer standard thin disk (Fender 1999; Esin, McClintock, \\& Narayan 1997). Thus, it seems that a jet always be present in an ADAF. If it is the case, we should check some result obtained under the frame of an canonical ADAF. An example is the radio/X-ray luminosity relation obtained by Yi \\& Boughn (1998) in the frame of the standard ADAF. Since the radio spectrum below a possibly existed break frequency $\\nu_{\\rm break}$ (as in the sources presented in this paper) may be dominated by the jet while that above a frequency the spectrum may disappear due to the truncation of the disk at $r_{\\rm tr}$, their conclusion holds only if the radio frequency $\\nu$ satisfies, \\be \\nu_{\\rm break} \\la \\nu \\la 20 m_7^{-1/2}\\dot{m}_{-3}^{1/2}T_{e9}^2 \\left( \\frac{r_{\\rm tr}}{20}\\right)^{-5/4} {\\rm GHz}. \\ee Here $r_{\\rm tr}=R_{\\rm tr}/R_g$.  The crucial factor in our model is to assume that the powering of the jet is the energy output from the disk rather than from the hole. The excellent agreement between the predictions of our model and the observations for both Sgr A$^*$ and elliptical galaxies in turn confirms the correctness of our assumption, and supports the ``disk'' origin of jets. In this context we note Livio, Ogilvie \\& Pringle (1999) recently argue that since there is no reason to suppose that the magnetic field threading the central spinning black hole differs significantly in strength from that threading the central regions of the disk, and the hole is a very poor conductor compared to the surrounding disk material, the electromagnetic output from the inner disk region is expected to dominate over B-Z mechanism. We want to note that this conclusion doesn't contradict the possibility that the rotation of the hole may play a role in the powering of the jets. This is because  a strong magnetic field may be created due to the frame dragging in the ergo-sphere of the Kerr black hole, therefore the energy extraction from the disk will be more efficient compared with the Schwarzschild black hole case (Koide et al. 2000)."
    },
    "0004/astro-ph0004018_arXiv.txt": {
        "abstract": "We present the long term evolution of the fre\\-quency-depen\\-dent X-ray time lags of the black hole candidate \\mbox{\\object{Cygnus X-1}} as measured in 1996 and 1998 with the Rossi X-ray Timing Explorer (RXTE). Lag spectra measured during the 1996 June soft state are very similar to those seen during 1996 December and most of 1998 while \\mbox{Cyg~X-1} was in its hard state. During state transitions, however, the shape and magnitude of the X-ray lag is highly variable and tends to be much larger than outside of the state transitions. This behavior is most obvious in the 1--10\\,Hz band. The increase of the X-ray lag during the state transitions might be related to the formation and destruction of the synchrotron radiation emitting outflows present during the hard state. ",
        "introduction": "\\label{sec:intro}  Galactic black hole candidates (BHC) are predominantly found in two generic states: the hard state, in which the X-ray spectrum is a Comptonization spectrum emerging from a hot electron cloud with a typical electron temperature of $\\sim$150\\,keV \\citep{dove:97b,poutanen:98a}, and the soft state, in which the X-ray spectrum is thermal with a characteristic temperature of $kT_{\\rm BB}\\lesssim 1$\\,keV to which a steep power-law is added \\citep[and references therein]{cui:96a,gierlinski:99a}. Transitions between these states have been seen in all persistent galactic BHC, with the exception of \\mbox{\\object{LMC~X-1}}. Optically thick radio emission is observed during the hard state, during transitions between the hard and the soft states the radio emission tends to be optically thin and more highly variable. Finally, during the soft state, galactic black hole candidates tend to be radio quiet \\citep[see, e.g.,][for a review]{fender:00b}. Although the phenomenology of the states is rather well understood, the accretion geometry in these sources is still a matter of debate. Most current models for the hard state posit an accretion disk corona with a large covering factor that Comptonizes most of the accretion disk radiation \\citep[and references therein]{poutanen:98a}. Its physical size is also assumed to be rather large. On the other hand, the corona is assumed to have almost vanished during the soft state, where the X-ray luminosity is dominated by thermal radiation \\citep{gierlinski:99a}.  In recent years, several attempts have been made to use X-ray timing methods to constrain these models. In addition to the power spectrum analysis \\citep{belloni:90a,gilfanov:99b}, higher order statistics like the frequency-dependent coherence function and time lags have proven to be useful in evaluating physical accretion models \\citep{hua:98a,nowak:98a} by providing combined spectral and temporal information. For example, the maximum expected time delay between hard and soft photons is roughly the size of the Comptonizing medium divided by the slowest propagation speed of a disturbance, while the minimum time lag is roughly the photon diffusion time through the corona \\citep{nowak:98c}.  The canonical BHC, \\mbox{Cygnus~X-1}, stays predominantly in the hard state, but occasionally transits into the soft state for a few months \\citep[][see also Fig.~\\ref{fig:softlag}]{gierlinski:99a,cui:96a,cui:97c}. The apparent decrease of the X-ray lags during the 1996 state transitions was considered evidence that the size of the accretion region during the soft state is smaller than during the hard state \\citep{cui:97c}. First comparisons of transition and soft state lags to hard state lags, however, indicated that the physical interpretation has to be more complex \\citep{cui:99a}.  In addition to the state transitions, quasi-regular variations on a $\\sim$150\\,d timescale are observed in the X-ray spectral shape and soft X-ray flux as well as in other wavelength bands \\citep{pooley:98a,brocksopp:99b}. In 1998 we initiated a monitoring campaign with the Rossi X-ray Timing Explorer (RXTE) to systematically study the multi-wavelength long term variation of \\mbox{Cyg~X-1} over a period of several years. During 1998, weekly pointings of $\\sim$3\\,ks duration were performed. In later years, longer exposure times (but larger sampling intervals) were used \\citep{pottschmidt:00a}. In this Letter we present results of the first year of the campaign, focusing on the X-ray time lags. We describe our data analysis methods (Section~\\ref{sec:data}) and compare the temporal behavior of \\mbox{Cyg~X-1} during 1998 with that seen during 1996 (Section~\\ref{sec:tempevol}). Specifically, we show that large X-ray lags appear to be associated with \\emph{transitions} between the soft and hard state, but \\emph{not} with the state itself. In Section~\\ref{sec:disc} we discuss implications for the accretion models in galactic black hole candidates.  \\begin{figure}  \\resizebox{\\hsize}{!}{\\includegraphics{lag03_vs_freq_10236_30157.eps}}  \\caption{Comparison of the X-ray lag spectrum for one of the 3\\,ks monitoring observations of 1998 with that obtained for $\\sim$30\\,ks of the RXTE observation of 1996 December 16. The lags have been measured between $\\lesssim$4\\,keV and $\\sim$8--13\\,keV. Both examples are typical hard state observations.} \\label{fig:cmpoldnew} \\end{figure} ",
        "conclusions": "\\label{sec:disc} In Fig.~\\ref{fig:lagshape}a we display examples for the whole range of lag spectra present in this analysis.  Taking the shape of the typical hard state lag spectrum as a baseline, the lag is significantly longer during the 1998~July failed state transition. The frequency range from 1 to 10\\,Hz exhibits these changes most prominently. During the hard to soft transition in 1996~May, the lag is longer by almost a factor 10 at 6\\,Hz. As already shown in section~\\ref{sec:tempevol}, these results indicate that the magnitude and shape of the X-ray lag spectrum in \\mbox{Cyg~X-1} is related to the state transitions.  Previous models for the generation of the X-ray lags assumed more or less static media to produce the X-ray lag by scattering of seed photons in a Comptonizing medium. The size of the region required to produce the observed lags in such a model is large \\citep[$\\gtrsim$300 gravitational radii;][]{nowak:98c}. Such a large size simplifies models in which the soft state Compton corona is much smaller than the hard state corona, as had been initially inferred from the 1996 soft state lags and spectral shape. Indirect evidence for the change of the size of the X-ray emitting corona has also been presented by \\citet{zdz:99a}, \\citet{gilfanov:99b}, and commented on by \\citet{dimatteo:99a}, with the latter authors suggesting an upper limit for the hard state coronal radius of $\\sim$30 gravitational radii. \\citet{zdz:98a} suggested that in the hard state of GX~339$-$4 there is a correlation between the X-ray power law photon index and the fraction of this power law that is reflected by cold material. Specifically, they suggested that softer power laws implied greater reflection, which implies smaller coronae in certain models. Within the hard state, softer power laws are associated with higher luminosities, implying again that the corona is shrinking as the source goes from the hard to the soft state. Our data clearly show that the soft and hard state lag spectra are very similar. This makes the geometrical interpretation of the lags in terms of Comptonization models alone difficult. We therefore need to look for other models to explain the observed lags. In the following we present a qualitative picture, based on recent observational results. Note, however, that a detailed theoretical model is beyond the scope of this Letter.  In recent years, evidence has emerged that the state transitions are not solely an X-ray phenomenon. Studies of \\mbox{Cyg~X-1} and \\mbox{GX~339$-$4} have revealed that in the hard state the X-ray and radio behavior is correlated \\citep{brocksopp:99b,corbel:00a}. In \\mbox{GX~339$-$4}, there is strong evidence that the source is radio quiet during the soft state \\citep{fender:99b}. At least in one case, optically thin radio flares accompanied the state transition. In \\mbox{Cyg~X-1}, no radio data are available for the soft state. The end of the 1996 soft state as well as the 1998 July ``failed state transition'', however, coincided with radio flares \\citep{zhang:97e,brocksopp:99b}. Such flaring events are typically associated with the ejection of a synchrotron emitting cloud from the central, X-ray emitting region \\citep{corbel:00a,fender:99b}.  We suggest that the scattering of primary X-rays in ejected material may be responsible for the enhanced transition X-ray time lags in \\mbox{Cyg~X-1}: In the hard state, a stable, presumably partially collimated, radio emitting outflow (``jet'') exists, while the soft state shows no outflow \\citep{brocksopp:99b}. During the formation and (failed) destruction of the jet, when radio flaring is observed, outflows that are uncollimated and much larger than those of the hard state might be present. Assuming that the hard state lag spectrum is produced in the accretion disk \\citep[which still poses a problem for most models;][]{nowak:98c} and that the rather weak hard state jet does not significantly affect these intrinsic lags, the prolonged lags during the transitions could be produced in these large ejected outflows. Additional scattering might also be responsible for the reduced X-ray coherence that was reported by \\citet{cui:97c} for the 1996 state transitions. Such a model wherein a fraction of the observed lags is created near the base of a radio emitting jet or wind has been previously suggested in analogy to blazar jet emission models \\citep[van Paradijs 1999, priv.\\ comm., see also ][]{fender:99b}. The state change following the flaring/ejection could then lead to a new accretion disk configuration with different X-ray spectral emission characteristics, different power spectra, but similar inherent lag spectrum."
    },
    "0004/astro-ph0004254_arXiv.txt": {
        "abstract": "The Andromeda galaxy (M31) provides a new line of sight for Galactic MACHO studies and also a signature, near-far asymmetry, which may establish the existence of MACHOs in the M31 halo. We outline the principles behind the so-called ``pixel-lensing'' experiments monitoring unresolved stars in M31. We present detailed simulations of the POINT-AGAPE survey now underway, which is using the INT wide-field camera to map the microlensing distribution over a substantial fraction of the M31 disc. We address the extent to which pixel-lensing observables, which differ from those of classical microlensing, can be used to determine the contribution and mass of M31 MACHOs. We also present some preliminary light-curves obtained from the first season (1999/2000) of INT data. ",
        "introduction": "This meeting is in part a celebration of the success of the current pioneering generation of microlensing surveys, the discoveries of which are reported throughout these proceedings. However, as with all pioneering work, a number of puzzles remain to be solved including, crucially, whether MACHOs exist or whether current results towards the Large and Small Magellanic Clouds (LMC and SMC) can be explained by ordinary stellar populations.  Whilst the MACHO Collaboration interprets the excess of events it finds as evidence of a $\\sim 20\\%$ halo contribution in objects of $\\sim 0.5~\\sm$ (Alcock et al. 2000) the EROS Collaboration has chosen to place only upper limits on the halo fraction from its sample of LMC/SMC events, even though the EROS dataset is formally consistent with that of MACHO (Lasserre et al. 2000). This dichotomy in approach reflects a community-wide uncertainty as to the contribution to the microlensing rate from the Magellanic Clouds themselves. As the current LMC/SMC surveys come to an end it is possible that the Magellanic satellite galaxies will continue to represent clouds of uncertainty for MACHOs.  Crotts (1992) and Baillon et al. (1993) proposed searching for MACHOs towards the Andromeda galaxy (M31). There are several aspects which make M31 very appealing for microlensing studies. Firstly, it is a large external galaxy, with ${\\cal O}(10^{10})$ sources, and has a halo of its own, providing an additional site for MACHOs. We should therefore expect many more events than towards the Magellanic Clouds. Small number statistics ought not be an issue for M31 if the MACHO abundance is significant. Secondly, the additional line of sight through our own halo that M31 provides may help to decide the viability of some of the stellar self-lensing scenarios, which have been proposed as non-MACHO alternatives to explain the LMC/SMC events. Alcock et al. (1995) have also noted that the ratio of microlensing rates between M31 and the LMC could be used in principle to probe the outer rotation curve of our Galaxy, though in practice the background from events in M31 makes this difficult. Lastly, our external viewpoint is advantageous in that we can map the MACHO distribution across the face of the M31 disc. The importance of this has been emphasized by Crotts (1992), who noted that the $77\\deg$ inclination of the M31 disc should give rise to a noticeable gradient in the microlensing rate if MACHOs occupy a spheroidal distribution, as depicted in Figure~1. Such an asymmetry does not arise naturally for disc lensing or variable stars, and so would provide strong evidence for the existence of MACHOs if detected. \\begin{figure} \\begin{minipage}{6.5cm} \\plotfiddle{kerins1.ps}{0cm}{270}{30}{30}{80}{-20} \\caption{The concept of near-far asymmetry. Due to the tilt of the M31 disc the microlensing rate is larger towards the far disc than towards the near side if MACHOs are present in a spheroidal dark halo.  The effect is less pronounced if the halo is flattened. The distributions of disc events and variable stars are symmetric.} \\end{minipage} \\end{figure} ",
        "conclusions": ""
    },
    "0004/astro-ph0004062_arXiv.txt": {
        "abstract": "s{ Spontaneous magnetization of quark liquid is examined on the analogy with that in electron gas. It is pointed out that quark liquid has potential to be ferromagnetic at rather low densities, around nuclear saturation density. Some comments are given as for implications on magnetars. } ",
        "introduction": "Recently a new type of neutron stars with extraordinary magnetic field, usually called magnetars, has attracted much attention in connection with pulsars associated with soft-gamma-ray repeaters (SGR) and anomalous X-ray pulsars (AXP). There have been known several magnetar candidates such as SGR 1806-20 and SGR 1900+14 \\cite{ko}. Various analysis including the $P-\\dot P$ curve have indicated an intense magnetic field of $O(10^{14 - 15})$ G, while ordinary radio pulsars have a magnetic field of $O(10^{12 - 13})$G.  The origin of the strong magnetic field in compact stars, especially neutron stars, has been an open problem. Recent discovery of magnetars seems to renew this problem. Conservation of the magnetic flux during the collapse of a main sequence star has been a naive idea to understand the magnetic field in neutron stars \\cite{cha}. Then the strength $B$ should be proportional to $R^{-2}$, where $R$ is a radius of a star; for example, the sun, a typical main sequence star, has a magnetic field of $O(10^3)$G with the radius $R\\sim 10^{10-11}$cm. By decreasing the radius to $10^6$cm for neutron stars $B=O(10^{12})$G, which is consistent with observations for radio pulsars. However, if this argument is extrapolated to explain the magnetic field for magnetars, we are lead to a contradiction: their radius should be $O(10^4)$cm to get an increase in $B$ by a factor of $\\sim 10^{12}$, which is much less than the Schwartzschild radius of neutron stars with the canonical mass $M=1.4M_\\odot$, $R_{Sch}=2GM/c^2=4\\times 10^5$cm.   When we compare the energy scales for systems such as atomic system ($e^-$), nucleon system ($p$) and quark system ($q$), we can get a hint about the origin of the magnetic field.  In Table 1. we list the interaction energy, $E_{int}=\\mu_iB$, of the magnetic field $B=10^{15}$G and each constituent with the Dirac magnetic moment, $\\mu_i=e_i\\hbar/2m_ic$. We also list a typical energy scale $E_{typ}$ for each system. \\begin{table} \\begin{center} \\begin{tabular}{cccc} &$e^-$ & p & q\\\\ \\hline\\hline $m_i$[{\\rm MeV}]& 0.5 & $10^3$ & 1- 100\\\\ $E_{int}$[{\\rm MeV}]~~~ & 5 - 6  &~~~$2.5\\times 10^{-3}$~~~& $2.5\\times 10^{-2}$ - 2.5\\\\ \\hline $E_{typ}$ & $O({\\rm KeV})$ & $\\geq O({\\rm MeV})$ & $\\geq O({\\rm MeV})$ \\end{tabular} \\end{center} \\caption{} \\end{table} Then we can see that $E_{typ}\\ll E_{int}$ for the electron system, while $E_{typ}>E_{int}$ for the nucleon and the quark systems; that is, the strength of $O(10^{15})$G is very large for the former system with the electromagnetic interaction, while it is not large for the latter systems with the strong interaction. Hence it may be conceivable that the strong interaction should easily produce the magnetic field of the above magnitude. Since there is a bulk hadronic matter beyond nuclear saturation density ($n_0\\sim 0.16$fm$^{-3}$) inside neutron stars, it should be interesting to consider the hadronic origin of the magnetic field; ferromagnetism or spin-polarization of hadronic matter may give such magnetic field.  In 70's, just after the first discovery of pulsars, there have been done many works about the possibility of the ferromagnetic transition in dense neutron matter, using $G-$matrix calculations or variational calculations with the realistic nuclear forces. Through these works there seems to be a consensus that ferromagnetic phase, if it exists, should be at very high densities, and there is no transition at rather low densities relevant to neutron stars \\cite{pan}.   We consider here the possibility of ferromagnetism of quark liquid interacting with the one-gluon-exchange (OGE) interaction \\cite{ta}. One believes that there are deconfinement transition and chiral symmetry restoration at finite baryon density, while their critical densities have not been fixed yet. One interesting suggestion is that three-flavor symmetric quark matter (strange matter) may be the true ground state of QCD at finite baryon density \\cite{chi,far}. If this is the case, quark stars (strange quark stars), can exist in a different branch from the neutron-star branch in the mass-radius plane \\cite{mad}. Usually one implicitly assumes that the ground state of quark matter is unpolarized. We examine here the possibility of polarization of quark matter. We shall see our results should give an origin of the strong magnetic field for magnetars in the context of strange quark-star scenario. ",
        "conclusions": "We have seen that the ferromagnetic phase is realized at low densities and the metastable state is possible up to rather high densities for a reasonable range of the QCD parameters. \\begin{figure}[t] \\centerline{ \\epsfysize=4.5cm \\epsfbox{qstar.eps} } \\caption{A model of strange quark star with $M\\sim 1.4M_\\odot$ and $R\\sim 10$Km. Almost all the portion is occupied by strange matter and a small vacuum gap may separate the quark core from the outer crust, which is composed of usual solid below the neutron-drip density.} \\end{figure}  We have found that ferromagnetic instability is feasible not only in the massive quark system but also in the light quark system: the spin-nonflip contribution is dominant in the nonrelativistic case as in electron gas, while a novel mechanism appears as a result of the large spin-flip contribution in the relativistic case.  If a ferromagnetic quark liquid exists stably or metastably around or above nuclear saturation density, strange stars may have a strong magnetic field, which strength is estimated to be strong enough for magnetars. Thereby it might be interesting to model SGR or AXP using our idea.  Our calculation is basically a perturbative one and the Fermi sea remains in a spherical shape. However, if we want to get more insight about the ferromagnetic phase, we must solve the Hartree-Fock equation and thereby derive a self-consistent mean-field for quark liquid. Moreover, we need to examine the long range correlation among quarks by looking into the ring diagrams, which has been known to be important in the calculation of the susceptibility of electron gas.  Recently, there have been done many works about the color superconductivity of quark matter. The order of the energy gap amounts to $O(100)$MeV, while the energy gain per particle is rather small and several MeV around $n_0$ , which should be the same order of magnitude as that for the ferromagnetism \\cite{bai}. Hence it may be interesting to explore the phase diagram for ferromagnetic phase and superconducting phase."
    },
    "0004/astro-ph0004312_arXiv.txt": {
        "abstract": "We have obtained new near-infrared images of both GG\\,Tau and UY\\,Aur circumbinary disks with the polarimetric modes of NICMOS aboard the Hubble Space Telescope. The 1\\micron\\ intensity map of GG\\,Tau reveals a more complete elliptical shape than in previous ground-based images. Its eastern and western sides are definitely asymmetric. Our image strongly supports the ring geometry proposed by Guilloteau, Dutrey \\& Simon (1999) on the basis of their millimetre interferometry images: a geometrically thick and sharply edged ring surrounding an empty gap around the binary. Around UY\\,Aur, we identify structures that are in excellent agreement with the optical images of M\\'enard et al. (1999), which confirms that the inclination of the system to the line-of-sight is about 60 degrees. We also find tentative new structures closer to the stars. ",
        "introduction": "The presence of circumstellar disks around T\\,Tauri stars has been suspected for a long time, but it is only very recently that these disks were directly detected, using high-angular millimetre imaging. These images revealed extended gas structures which appeared to be in Keplerian rotation around the central object. Among the rare detections so far, two disks were found around binary T\\,Tauri stars: GG\\,Tau (Dutrey, Guilloteau \\& Simon 1994) and UY\\,Aur (Duvert et al. 1998). The separations of the binaries are 0\\farcs25 and 0\\farcs89 respectively, which correspond to projected physical separations of 35 and 125\\,AU at the distance of the Taurus star-forming region (140\\,pc).  In both cases, light scattered off the surface of the disks has been detected afterwards with adaptive optics imaging at near-infrared wavelengthes. Roddier et al. (1996) found that the GG\\,Tau ring has a clumpy appearance and that several radial spokes of material extend from the ring onto the central stars. The ring is brighter in its northern part, but is detected in all directions. They interpret this brightness difference as being due to the scattering geometry. The UY\\,Aur case is very different, as Close et al. (1998) only detected the disk on one side of the binary. Furthermore, they found evidences that a ``spiral arm'' splits from the main disk and gets closer to the star. Deconvolution processes were applied in both studies to retrieve the highest spatial resolution allowed by adaptive optics devices, and this may lead to some artifacts in the final images.  More recently, the first visible wavelength images of UY\\,Aur were obtained by M\\'enard et al. (1999) at 600 and 800\\,nm with HST/WFPC2. The PSF-subtracted images revealed a more complicated structure that was found by Close et al. (1998): a large ``clump'' appears to be independent from the disk itself. If true, this implies that the inclination of the system to the line-of-sight is larger than was first thought (about 60\\deg\\ instead of about 40\\deg).  To improve our knowledge of these two circumbinary disks, we have performed new observations at 1\\micron\\ and 2\\micron\\ of these systems with HST/NICMOS. We used the polarimetric modes, and we obtained both intensity and polarization maps, which do not need to be deconvolved. The GG\\,Tau polarization maps are the first ever obtained of this system, while Potter et al. (1998) already presented a deconvolved J-band polarization map of UY\\,Aur which revealed a nice centrosymetric pattern. Polarization maps are powerfull tools to investigate the dust grain properties and the geometry and structure of the disks.  In section\\,2, we summarize our observations and data processing steps, and the maps of both systems are presented and commented in section\\,3. Section\\,4 describes some implications of our results on the properties of these disks. ",
        "conclusions": ""
    },
    "0004/astro-ph0004124_arXiv.txt": {
        "abstract": "We report new determinations of the radius and mass of the pulsating helium-rich, hydrogen-deficient star \\object{LSS~3184} (\\object{BX Cir}) using measurements of radial velocity and angular radius throughout its pulsation cycle. Measurements of radial velocity, and thus changes in stellar radius ($\\Delta R_\\star$), were made using Anglo-Australian Telescope echelle spectra.  {\\it Hubble Space Telescope} ultraviolet spectra and ground-based {\\it BV} photometry were used to find temperatures and fluxes throughout the pulsation cycle.  The temperatures and fluxes were used to find the angular radius of the star ($\\alpha$).  The $\\alpha$, $\\Delta \\alpha$, and $\\Delta R$ values thus found were used to calculate the mean stellar radius $\\langle R_\\star \\rangle = 2.31 \\pm 0.10 R_\\odot $. If we use the previously determined $\\log g = 3.35 \\pm 0.1$ for \\object{LSS~3184} and our radius estimate, we find its mass to be $M_\\star = 0.42 \\pm 0.12 M_\\odot$. ",
        "introduction": "Pulsations in stars provide tools for researchers in several fields of astronomy.  They provide standard candles for measuring Galactic and extragalactic distances and they provide methods for measuring stellar parameters, even below the directly observable photosphere.  Pulsations in hydrogen-deficient, helium-rich stars such as extreme helium, R Coronae Borealis, and hydrogen-deficient carbon stars have not been studied in as much detail as those in stars with more `normal' chemical abundances.  While this is partly understandable since most pulsating stars have normal compositions, studying stars with little or no hydrogen is important to allow tests of pulsation theory.  Extreme helium stars, as their name would imply, have weak or non-existent hydrogen absorption lines and very large helium abundances ($\\ga 99$ per cent). Two extreme helium stars, \\object{V652~Her} and \\object{LSS~3184}, are known to pulsate. Saio (\\cite{s93}) showed that in extreme helium stars with temperatures around $2 \\times 10^5$~K, like \\object{V652~Her}, the $\\kappa$-mechanism caused by iron-group (Z-bump) opacity can excite the observed pulsations. Saio (\\cite{s94}, \\cite{s95}) predicted that \\object{LSS~3184} should pulsate because of its location in the Z-bump instability finger.  Kilkenny \\&\\ Koen (\\cite{kk95}) discovered that \\object{LSS~3184} shows photometric variations with a period of about 0.107~d.  The similarity of \\object{V652~Her} and \\object{LSS~3184} in temperature, surface gravity, and pulsation period implies that they are very similar in other physical parameters.  Kilkenny et~al. (\\cite{kk99}) have recently reported an observational analysis of \\object{LSS~3184}, including a determination of its photometric period (0.1066~d).  In addition, they used medium resolution ($\\lambda / \\Delta \\lambda \\approx 4000$) spectra to measure radial velocity variations and to show that the photometric  variability is caused by pulsations.  Drilling, Jeffery, \\& Heber (\\cite{d98}) reported an analysis of \\object{LSS~3184} in which they found $T_{\\rm eff} = 23\\,300 \\pm 700$~K, $\\log g = 3.35 \\pm 0.1$ and $n_{\\rm H}/n_{\\rm He} \\le 0.00015$. Kilkenny et~al. (\\cite{kk99}) report that using the radius they determined and $\\log g$ from Drilling et~al. (\\cite{d98}) gives a mass of $0.15 M_\\odot$, which is much smaller than $0.7 M_\\odot$, the mass accepted for \\object{V652~Her} (Lynas-Gray et~al. \\cite{l84}), and small enough to imply that some input parameter or procedure is in error.  In this paper we report an analysis of \\object{LSS~3184} using high resolution optical spectra for radial velocity measurement and ultraviolet {\\it Hubble Space Telescope} ({\\it HST}) spectra and ground-based {\\it BV} photometry for temperature, luminosity, and angular radius measurement.  The new data provide better temperature and angular radius estimates and a much cleaner radial velocity curve for the star, allowing a more reliable estimate of its radius and mass than was possible previously. ",
        "conclusions": "In this paper we report new determinations of the radius ($\\langle R_\\star \\rangle = 2.31 \\pm 0.10 R_\\odot$) and mass ($M_\\star = 0.42 \\pm 0.12 M_\\odot$) of the pulsating hydrogen-deficient star \\object{LSS~3184}.  The radial velocity data and temperature measurements are more reliable than those used for previous radius and mass determinations for this star, so our estimates are likely to be closer to the star's actual parameters.  Further improvements can be accomplished by improving the $\\log g$ and chemical composition estimates for \\object{LSS~3184}. In addition, there needs to be further study of the difference between the shapes of the angular radius and physical radius curves based on temperature and flux measurements, and radial velocity measurements, respectively.  It is important to determine if the difference in shapes is a sign that some of our data or a step in our analysis is flawed, or if there is something happening in the star itself that causes the two measures of radius to behave differently with pulsational phase."
    },
    "0004/astro-ph0004202_arXiv.txt": {
        "abstract": "We investigate the Gaussianity of the 4-year COBE-DMR data (in HEALPix pixelisation) using an analysis based on spherical Haar wavelets. We use all the pixels lying outside the Galactic cut and compute the skewness, kurtosis and scale-scale correlation spectra for the wavelet coefficients at each scale. We also take into account the sensitivity of the method to the orientation of the input signal. We find a detection of non-Gaussianity at $> 99$ per cent level in just one of our statistics. Taking into account the total number of statistics computed, we estimate that the probability of obtaining such a detection by chance for an underlying Gaussian field is 0.69. Therefore, we conclude that the spherical wavelet technique shows no strong evidence of non-Gaussianity in the COBE-DMR data. ",
        "introduction": "The cosmic microwave background (CMB) provides a unique tool for investigating the formation of structure in the Universe. In particular, studying the Gaussianity of the CMB temperature fluctuations allows us to distinguish between two competing theories of structure formation: the standard inflationary model that predicts Gaussian fluctuations and topological defects that give rise to non-Gaussian signatures in the CMB. In order to test the Gaussianity of the CMB, a large number of methods have already been proposed in the literature (see e.g. Barreiro 2000 and references therein). In particular, many of them have been applied to the 4-year COBE-DMR data, although most have not yielded any detection of non-Gaussianity (see e.g. Mukherjee, Hobson \\& Lasenby 2000).  Nevertheless, three apparently robust detections of non-Gaussianity in the 4-year COBE-DMR data have recently been reported. Ferreira, Magueijo \\& G\\'orski (1998) studied the distribution of an estimator for the normalised bispectrum, finding that Gaussianity is ruled out at the $98$ per cent confidence level with the non-Gaussian signal mainly concentrated on the multipole $l=16$. Although this non-Gaussian signal is certainly present in the data, Banday, Zaroubi \\& G\\'orski (1999) have shown in a recent work that this non-Gaussian signal is most probably due to a systematic artefact. However, Magueijo (1999) has found a new non-Gaussian signal above the $97$ per cent confidence level using an extended bispectrum analysis, which is present even when the artefacts found by Banday et al. are removed.  Complementary to the above bispectrum analysis, Pando, Valls-Gabaud \\& Fang (1998) applied a technique based on the discrete wavelet transform to Face 0 and Face 5 of the QuadCube pixelisation of COBE-DMR data. On computing the scale-scale correlations of the wavelet coefficients in certain domains of the wavelet transform, they found a significant non-Gaussian signal at the 99 per cent confidence level in Face 0 corresponding to scales of $11-22$ degrees. However, no significant deviation from Gaussianity was found using the skewness and kurtosis of the wavelet coefficients at each of the considered domains. Mukherjee, Hobson \\& Lasenby (2000) (hereinafter MHL) have recently revised the previous work, taking into account that a large number of the computed statistics show no evidence of non-Gaussianity, and pointing out that the results depend critically on the orientation of the data. They find that Gaussianity can only be ruled out at the 41 per cent confidence level in the DSMB data and at the 72 per cent level in the 53+90 GHz coadded data and therefore that this analysis does not provide strong evidence for non-Gaussianity in the COBE-DMR data.  The above wavelet analyses were performed by applying planar (Daubechies) wavelets to Face 0 and Face 5 of the COBE-DMR QuadCube pixelisation in Galactic coordinates (i.e. the faces centred on the North and South Galactic poles respectively). This procedure thus uses only one-third of the COBE-DMR data and, furthermore, can lead to distortions of the CMB fluctuations when moving from the sphere to the planar faces of the QuadCube. Therefore, in this paper, we use orthogonal spherical Haar wavelets (SHW) (Girardi \\& Sweldens 1995, Sweldens 1995), which are better suited to analysing data over large regions of the sky (such as the COBE-DMR maps). These wavelets were introduced as a generalisation of planar Haar wavelets to more general spaces than $R^n$. Indeed, Tenorio et al. (1999) have used SHW as a tool for studying the spatial structure, denoising and compression of CMB maps. We use SHW to perform a similar analysis to those of Pando et al and MHL, based on the skewness, kurtosis and scale-scale correlation of the wavelet coefficients, in order to search for evidence of non-Gaussianity in the 4-year COBE-DMR data. A hierarchical pixelisation scheme which is particularly well-suited to the application of such a wavelet decomposition is HEALPix (G\\'orski, Hivon \\& Wandelt 1999). Therefore, we apply the former analysis to the COBE-DMR data in HEALPix pixelisation (Banday et al. 2000). We also include a Galactic cut derived by a simple propagation of the customised Galactic cut of Banday et al. (1997) and Bennett et al. (1996) to the HEALPix pixelisation. Owing to the characteristics of the SHW transform, it is straightforward to use {\\em all} the pixels lying outside the Galactic cut, which constitutes approximately two-thirds of the total number of COBE-DMR pixels. ",
        "conclusions": "We presented an analysis of the 4-year COBE-DMR data (in HEALPix pixelisation) based on spherical Haar wavelets (SHW) in order to look for large scale non-Gaussianity of the CMB. This analysis is performed using all the available data lying outside the Galactic cut. This constitutes about two-thirds of the pixels, as compared to just one-third in former works that use only Face 0 and Face 5 of the QuadCube pixelisation. We take into account the sensitivity of the method to the orientation of the original signal and present the results for three different orientations of the data. We also find that the choice of a different set of orientations, at least in our case, does not alter the general conclusions regarding non-Gaussianity.  We find that the value of the kurtosis for $j=2,m=3$ for one of the chosen orientations lies outside the $99$ per cent confidence level derived from 10000 CDM/inflationary realisations, whereas the rest of the statistics show no evidence of non-Gaussianity. However, since we have a total of 126 different statistics, the probability of finding at least one of them falling outside the 99 per cent confidence level even in the case of an underlying Gaussian field is as high as 0.69. Therefore, we conclude that an analysis based on SHW of the 4-year COBE-DMR data show no evidence for non-Gaussianity."
    },
    "0004/astro-ph0004034_arXiv.txt": {
        "abstract": "s{ An analysis of cosmological models in spatially flat Friedmann Universe with cosmic gravitational wave background and zero $\\Lambda$-term is presented. The number of free parameters is equal to 5, they are $\\sigma_8$, $n$, $\\Omega_\\nu$, $\\Omega_b$, and $h$. The normalization of the spectrum of density perturbations on galaxy cluster abundance ($\\sigma_8 = 0.52\\pm 0.04$) has been used to calculate numerically the value of the large scale CMB anisotropy ($\\ell\\simeq 10$) and the relative contribution of cosmological gravitational waves T/S. Increasing $\\Omega_\\nu$ weaken the requirements to the value of T/S, however even for $\\Omega_\\nu\\le 0.4$ the models with $h+n\\ge 1.5$ suggest considerable abundance of gravitational waves: T/S${}^>_\\sim 0.3$. In models with $\\Omega_\\nu\\le 0.4$ and scale-invariant spectrum of density perturbations ($n=1$): T/S${}^>_\\sim 10(h-0.47)$. Minimization of the value T/S is possible only in the range of the red spectra ($n<1$) and small $h$ ($<0.6$). It is shown that the models with T/S$\\in [0, 3]$ admit both moderate red and blue spectra of density perturbations, $n\\in[0.9,1.2]$, with rather high abundance hot dark matter, $\\Omega_\\nu\\in [0.2,0.4]$. Any condition, $n<0.9$ or $\\Omega_\\nu<0.2$, decreases the relative amplitude of the first acoustic peak for more than 30$\\%$ in comparison with its hight in the standard CDM normalized by COBE data.} ",
        "introduction": " ",
        "conclusions": "\\begin{itemize} \\item The data on the galaxy cluster abundance determine the value $\\sigma_8$ with a high accuracy, the other parameters ($n, \\Omega_\\nu, \\Omega_b, h$) remain free within their ranges. \\item None of the MDM models with $n=1$ and T/S$=0$ satisfies both normalizations, on the galaxy cluster abundance and large-scale $\\Delta T/T$ anisotropy, which leads either to rejection from the flat spectrum or to the introduction of a non-zero T/S (or both). \\item Small values of T/S are realised for the red spectra ($n<1$) and moderate $h(<0.6)$, the violation of these conditions leads to a high T/S(${}^>_\\sim 1$). \\item Increasing $\\Omega_\\nu$ weaken the requirement to the value of T/S, however even for $\\Omega_\\nu\\le 0.4$ the models with $h+n\\ge 1.5$ suggest considerable abundance of gravitational waves: T/S${}^>_\\sim 0.3$. \\item In models with $\\Omega_\\nu\\le 0.4$ and scale-invariant spectrum of density perturbations ($n=1$): T/S${}^>_\\sim 10(h-0.47)$. \\item In models with $\\Omega_b=0.05$ and $h=0.5$ we have the following approximation for the primordial gravitational waves (the accuracy is better than 11 \\% for $0.1\\le$T/S$\\le 3$): $$ \\frac{\\rm T}{\\rm S}=\\frac{30(n-0.7)^2}{10\\Omega_\\nu+1}+10\\Omega_\\nu\\left( n^{3/2}-1.06\\right). $$ \\item In double-normalized models with T/S$>$0 the hight of the acoustic peak is less than its `standard' value ($\\Re=5.1$). The deacrease of the parameter $\\Re$ does not exceed 30\\% in models with large $\\Omega_\\nu \\in [0.2,0.4]$ and any spectrum slope, $n\\in [0.9,1.2]$. Any condition, $n<0.9$ or $\\Omega_\\nu<0.2$, decreases the relative amplitude of the first acoustic peak for more than 30$\\%$ (i.e. $\\Re<3.5$ in models with $\\Omega_b=0.05$, $h=0.5$). The acoustic peak practically disappears in CDM models. \\item When increasing the baryonic abundance the difference between $\\Re$ in our models and that in sCDM decreases. The amplitude of the acoustic peak coincides with its 'standard' value ($\\Re\\ge 4.5$) in models with $\\Omega_b=0.1$ and either 'blue' spectrum $n\\in[1,1.2]$ and $\\Omega_\\nu\\ge 0.3$, or moderate 'red' spectrum $n\\in[0.9,1]$ and $\\Omega_\\nu\\ge 0.2$. \\item Thus, rising the parameter $\\Omega_\\nu$ up to the values in the interval $[0.2,0.4]$ solves the problem of the first acoustic peak in $\\Delta T/T$, leaving the baryonic density within the primordial nucleosynthesis constraints. \\end{itemize} \\begin{figure} \\centerline{\\psfig{figure=ffig2.eps,height=2.5in}} \\caption{$\\Re$ as a function $n$ for different values of $\\Omega_\\nu$ ($\\Omega_b=0.05$, $h=0.5$, $\\sigma_8=0.52$). Models in the non-shaded region have T/S$\\in[0,3]$.} \\end{figure}"
    },
    "0004/astro-ph0004344_arXiv.txt": {
        "abstract": "We obtained R-- and I--band CCD photometry of the soft X--ray transient/neutron-star binary Aql~X-1 in 1998 June while it was at quiescence. We find that its light curve is dominated by ellipsoidal variations, although the ellipsoidal variations are severely distorted and have unequal maxima. After we correct for the contaminating flux from a field star located only 0\\farcs46 away, the peak-to-peak amplitude of the modulation is $\\approx 0.25$ mag in the R band, which requires the orbital inclination to be greater than  $36^\\circ$. The orbital period we measure is consistent with the 18.95~h period measured by Chevalier \\& Ilovaisky (1998). During its outbursts the light curve of Aql~X-1 becomes single humped. The outburst light curve observed by Garcia et al.\\ (1999) agrees in phase with our quiescent light curve. We show that the single humped variation is caused by a ``reflection effect,'' that is, by heating of the side of the secondary star facing towards the neutron star. ",
        "introduction": "Soft X--ray Transients (SXTs, also called X--ray nova) are distinguished from other X--ray binaries by their transient outbursts and soft X--ray spectra (Tanaka \\& Lewin 1995; Tanaka \\& Shibazaki 1996; Chen, Schrader \\& Livio 1997). During the outbursts, which recur irregularly at intervals of 1 to 50 years, their X--ray luminosity increases by a factor of $10^4$--$10^6$ in a few days and then decays back to minimum on an e-folding time scale of $\\sim 30$ days. There is a concomitant increase in their optical luminosity by a factor of $\\sim 10^2$. The leading model for the outbursts is the Disk Instability Model (e.g., Cannizzo et al.\\ 1998). In this model, gas transferred from the normal star in the binary builds up in a disk around the the compact star. As the temperature and optical depth in the disk increase, the viscosity increases, eventually setting off a global instability that increases the mass flow through the disk, releasing accretion energy and causing the X--ray outburst. Between the X--ray outbursts, when the X--ray flux is low, the masses and dimensions of SXTs can be measured with some confidence. A startling result of these measurements has been the discovery that the compact star is a black hole in $\\sim$75\\% of the SXTs (e.g. see Charles 1998).  Aql~X-1 is a particularly interesting SXT because of its frequent outbursts --- the recurrence time is $\\sim$ 1 year (Priedhorsky \\& Terrell 1984; Kitamoto et al. 1993) -- and because its X--ray light curve displays richly complex behavior: (1) It shows ``type I'' X--ray bursts, demonstrating that the compact star is a neutron star (Koyama et al.\\ 1981, Zhang et al.\\ 1998b, Yu et al.\\ 1999); (2) Nearly coherent oscillations at $\\sim$548 Hz were detected during one type I burst and are likely related to the spin period of the neutron star (Zhang et al.\\ 1998b); (3) kHz QPOs have been discovered with frequencies in the range 670--870 Hz  (Zhang et al.\\ 1998b; Yu et al.\\ 1999); (4) A hard power--law tail extending to at least $>$100 keV has been detected during outbursts (Harmon et al.\\ 1996); and (5) A sudden spectral hardening of the X--rays has been observed during the decline from an outburst and has been interpreted as the onset of a ``magnetic propeller'' state (Zhang et al.\\ 1998a; Campana et al.\\ 1998). The inferred strength of the magnetic field, $\\sim1$--$6 \\times 10^{8}$~Gauss, is remarkably low for a neutron star. The spectral hardening may also be enhanced by the transition into an ADAF state (Shahbaz et al.\\ 1998b; Zhang et al.\\ 1998a).  The optical counterpart to Aql~X-1 is V1333~Aql (Thorstensen, Charles \\& Bowyer 1978). Two estimates for its orbital period have appeared in the literature, 18.9479 h (Chevalier \\& Ilovaisky 1991, 1998) and 19.30 h (Shahbaz et al.\\ 1998a). The former was derived from the outburst light curve, which has a single maximum per orbit; the later period was derived from the quiescent light curve, which has two humps per orbit. The single hump seen during outburst has been interpreted  as being due to reprocessed X--rays in an asymmetric disk and/or the secondary star. The double hump modulation seen in quiescence is interpreted as ellipsoidal variations from the Roche--lobe filling secondary star. If correct, the difference in periods is puzzling; photometric modulations caused by a non--axisymmetric disk during an outburst tend to have a longer period than the orbital period.  Garcia et al.\\ (1999) obtained V--band photometry of Aql~X-1 in outburst and K$^{\\prime}$ photometry in quiescence. Their outburst data confirm the strong ($\\sim$0.6 mag) single--humped light curve seen by Chevalier \\& Ilovaisky (1991) and Robinson \\& Young (1998), though their data were not sufficient to refine the orbital period. The quiescent K$^{\\prime}$ photometry did not show any clear evidence for a modulation, suggesting a low orbital inclination. They also obtained spectroscopy in quiescence that revealed a weak, low--amplitude modulation of the absorption line radial velocities, also implying a low orbital inclination. While this agrees with the analysis of the quiescent photometric modulations by Shahbaz et al.\\ (1998a), it is again puzzling: the large amplitude modulations seen in the outburst state suggest a high inclination. Furthermore, Garcia et al.\\ (1999) note that the relative phasing of the outburst maxima they observed  and the quiescent minima from the Shahbaz et al.\\ study disagree. Clearly, the interpretation of the modulations in the light curves is in a confused state.  Some of this confusion was resolved when Callanan et al.\\ (1999) obtained high-resolution images of Aql~X-1 showing that the object previously thought to be Aql~X-1 is actually two stars separated by only 0\\farcs46. Chevalier et al.\\ (1999) subsequently showed that the {\\em fainter} of the two stars is the counterpart of Aql~X-1, and that 88\\% of the V--band light is from the line--of--sight contaminating star. The presence of a large amount of contaminating light has two immediate consequences: (1) The fractional amplitude of any photometric variations have been substantially underestimated; this leads to an underestimate of the inclination based on the amplitude of the ellipsoidal variations; and (2) previous determinations of the spectral type of the secondary star (K0 by Thorstensen et al 1978; K1 by Shahbaz et al. 1997) are incorrect. Based on its photometric colors, the spectral type of the secondary star is K6--M0 (Chevalier et al.\\ 1999).  In this paper we present new optical photometry of Aql~X-1 obtained during quiescence, and use these data to revise the orbital period and constrain the inclination of the binary system. In \\S2 we describe the data acquisition and calibration. In \\S3 we derive a revised orbital period from the light curve modulations, and compare it with previous published values. Using the estimates of Chevalier et al.\\ (1999), we then correct the light curves for the contaminating light from the nearby field star. In \\S4, the contamination--corrected quiescent light curves are modeled as ellipsoidal variations and an estimate for the binary system inclination is derived. The outburst light curve of Garcia et al.\\ (1999) is also investigated; we model the outburst photometry with a reflection/heating effect of the companion star. Section 5 summarizes our results. ",
        "conclusions": "We have confirmed the presence of a double--humped orbital modulation in the quiescent light curve of Aql~X-1. Our orbital period estimate of 18.71~h is slightly shorter than, but consistent with, the 18.95~h period found by Chevalier \\& Ilovaisky (1998). A single orbital period, 0.789498~d, fits all the available light curves during both quiescence and outburst.  After subtracting the contaminating flux from a field star 0\\farcs46 away from Aql~X-1, the peak-to-peak amplitude of the double--humped modulation is $\\sim 0.25$ mag. The modulation is caused by ellipsoidal variations of the K star in the system. The ellipsoidal variations are distorted but do, nevertheless, yield a lower limit to the orbital inclination of Aql~X-1. We find that the inclination must be greater than $36^\\circ$.  In outburst, the orbital light curve contains a single, large--amplitude hump. A synthetic light curve produced by an irradiated, late--K star provides a good match to the outburst light curve observed by Garcia et al. (1999). Further support for the heated secondary star model comes from the excellent phase agreement between the outburst and quiescence photometric variations. During outbursts, the optical light is most likely dominated by the heating of the K--type secondary star."
    },
    "0004/astro-ph0004108_arXiv.txt": {
        "abstract": "We present sub-arcsecond resolution ground-based near-infrared images of the central regions of a sample of twelve barred galaxies with circumnuclear star formation activity, which is organized in ring-like regions typically one kiloparsec in diameter. We also present {\\it Hubble Space Telescope} near-infrared images of ten of our sample galaxies, and compare them with our ground-based data. Although our sample galaxies were selected for the presence of circumnuclear star formation activity, our broad-band near-infrared images are heterogeneous, showing a substantial amount of small-scale structure in some galaxies, and practically none in others.  We argue that, where it exists, this structure is caused by young stars, which also cause the characteristic bumps or changes in slope in the radial profiles of ellipticity, major axis position angle, surface brightness and colour at the radius of the circumnuclear ring in most of our sample galaxies. In 7 out of 10 {\\it HST} images, star formation in the nuclear ring is clearly visible as a large number of small emitting regions, organised into spiral arm fragments, which are accompanied by dust lanes.  NIR colour index maps show much more clearly the location of dust lanes and, in certain cases, regions of star formation than single broad-band images.  Circumnuclear spiral structure thus outlined appears to be common in barred spiral galaxies with circumnuclear star formation. ",
        "introduction": "Observations and modelling of the circumnuclear regions (CNRs) of barred galaxies give important clues about the nature of active galactic nuclei (AGN), (circum)nuclear starbursts and the gas dynamics in the central kpc (e.g. Knapen et al. 1995a,b; Buta \\& Combes 1996; Elmegreen et al. 1997; Quillen et al. 1999; Regan \\& Mulchaey 1999; see also reviews by Knapen 1999 and Shlosman 1999).  The presence of non-axisymmetries, such as oval distortions or bars, in the gravitational potential of a disc galaxy can lead to the infall of gaseous material, from galactocentric radii of a few kpc, into the central kpc. The nonaxisymmetry facilitates the removal of angular momentum from the gas, and directs it into the CNR or possibly the nucleus (Shlosman et al.  1990; Athanassoula 1992; Phinney 1994). This gas fuelling process is expected to induce morphological signatures in the bars and CNRs, which are useful tools in the understanding of the dynamics and kinematics of galaxies. Dust lanes and nuclear rings appear to be the dominant circumnuclear features.  \\begin{table*} \\hfill \\begin{tabular}{ccrrrcccc} \\hline Object & Type & Velocity & Distance & Scale & CNR & Nuclear & Date of & Passbands\\\\ &      & km s$^{-1}$  &  Mpc   & pc/arcsec & Features & activity & observation &\\\\ \\hline NGC 1300 & (R')SB(s)bc & 1568 & 20.9 & 101 & Ring&---&05/11/95& $JHK$\\\\ NGC 1530 & SB(rs)b     & 2461 & 32.8 & 159 & Ring+spiral&---&05/11/95& $JHK$\\\\ NGC 2903 & SAB(rs)bc   &  556 &  7.4 &  36 & Ring&Starburst&05/11/95& $JHK$\\\\ NGC 3351 & SB(r)b      &  778 & 10.1 &  49 & Ring&Starburst&06/11/95& $K$\\\\ NGC 3504 & (R)SAB(s)ab & 1539 & 20.5 &  99 & Ring+spiral&Starburst&04/02/96 & $JHK$\\\\ NGC 3516 & (R)SB(s)    & 2649 & 35.3 & 171 & Ring&Seyfert 1.5&06/11/95& $JHK$\\\\ NGC 3982 & SAB(r)b     & 1109 & 14.8 &  72 & Spiral&Seyfert 2&04/02/96 & $JHK$\\\\ NGC 4303 & SAB(rs)a    & 1566 & 20.9 & 101 & Ring&Seyfert&05/02/96& $JHK$\\\\ NGC 4314 & SB(rs)a     &  963 & 12.8 &  62 & Ring+spiral &--- &03/02/96& $JHK$\\\\ NGC 4321 & SAB(s)bc    & 1571 & 16.1 &  70 & Ring+spiral&--- &10/06/94& $K$\\\\ &             &      &      &     &      &    &04/02/96& $JH$\\\\ NGC 5248 & SAB(rs)bc   & 1153 & 15.4 &  74 & Ring+spiral&---&04/02/96& $JHK$\\\\ NGC 6951 & SAB(rs)bc   & 1424 & 19.0 &  92 & Ring+spiral&LINER/Sy&26/09/94& $JHK$\\\\  \\hline  \\end{tabular}  \\caption{Properties of our sample galaxies: morphological classification (col. 2; data from de Vaucouleurs et al. 1991), velocity (col.3; data from NED, the NASA/IPAC Extragalactic Database), distance and scale (col. 4 and 5, from Ferrarese et al. 1996 for NGC~4321 and Graham et al. 1997 for NGC~3351; from the velocity and H$_0$=75 km~s$^{-1}$ Mpc$^{-1}$ for all other objects), the dominant circumnuclear feature as determined from our data (col. 6), class of nuclear activity (col. 7; data from NED), date of run and passbands (col. 8 and 9). All images were obtained with MONICA on the {\\it CFHT} except the $K$-image of M100, which was obtained at UKIRT (Knapen et al. 1995a).}  \\end{table*}  Dust lanes have been interpreted as the location of shocks in the gas flow. Thus, the morphology of the dust distribution can reveal characteristics of the dynamics of these galaxies (Athanassoula 1992) and of the kinematics of the gas. According to Athanassoula's models, the degree of curvature of the dust lanes is a direct indicator of the strength of the bar for bars which are not very strong. Nuclear rings, which are usually sites of active star formation (SF), occur at the location of strong density enhancements in the gas, where the bar-driven inflow of gas slows down in the vicinity of inner Lindblad resonances (ILRs; Athanassoula 1992; Heller \\& Shlosman 1994; Buta \\& Combes 1996; Shlosman 1999). SF may result from the gas becoming gravitationally unstable (Elmegreen 1994, 1997), or from triggering in miniature spiral arms (Knapen et al. 1995b, 2000).  Whereas imaging in the blue, ultraviolet, or in spectral lines like H$\\alpha$ traces the massive SF (e.g. Sersic \\& Pastoriza 1967; Pogge 1989a,b; Knapen et al. 1995a; Maoz et al.  1996; Colina et al. 1997), near-infrared (NIR) imaging offers very important advantages. Firstly, the presence of substantial and mostly unquantified amounts of extinguishing (absorbing and scattering) dust in the CNRs hampers the interpretation of imaging at especially the UV and optical wavelengths. NIR, particularly $K$-band, emission is much less susceptible to extinction by dust (a factor of 10 when comparing $K$ to $V$), whereas colour index images such as $I-K$ (Knapen et al. 1995a) or $J-H$ are clear morphological dust indicators due to the relatively modest changes in those colours caused by different stellar populations. Secondly, NIR observations are of particular importance to the study of the CNRs because they allow a better determination of the mass distribution (excluding dark matter). NIR imaging is more sensitive to light primarily from cool giants and dwarfs which dominate the mass or are at least directly proportional to it. However, as shown in the literature and by us in this paper, there can be significant, or even dominant, contributions from young stars to the NIR emission even in the $K$-band, so for objects with strong SF any estimate of the mass distribution made from NIR maps must take the varying $M/L$ ratio into account (Knapen et al. 1995a,b; Elmegreen et al. 1997; Wada, Sakamoto \\& Minezaki 1998; Ryder \\& Knapen 1999).  The atlas of images presented in this paper (Paper I) represents the first results of a programme which studies the circumnuclear structures that appear at small scales (a few hundred parsecs to a kpc) and their connection with the global disc structure in a sample of 12 barred spiral galaxies. We present $J, H$ and $K$-images at subarcsec resolution for all 12 sample galaxies, and {\\it Hubble Space Telescope} ({\\it HST}) archive $H$-band images for 10 of them. In Paper II (P\\'erez--Ram\\'\\i rez \\& Knapen 2000a), we present an accompanying optical data set of broad-band and H$\\alpha$ images of the complete discs of all our sample galaxies, while in Paper III (P\\'erez--Ram\\'\\i rez \\& Knapen 2000b) the morphological information is interpreted in terms of the structure and evolution of CNRs of barred galaxies.  In Section 2, we describe the observations and data reduction techniques. The imaging data are presented as a series of multi-panel figures, which show broad-band images, colour index images highlighting dust and SF features against the dominant stellar luminosity, and radial profiles of ellipticity, position angle, surface brightness and colour. The results are summarized in Section 3 for individual galaxies, and briefly discussed in Section 4, while concluding remarks are given in Section 5.  \\setcounter{figure}{0}  \\begin{figure*} \\vspace{1cm} \\caption{{\\it HST} NICMOS images of the central areas of our sample galaxies obtained using the F160W filter (comparable to $H$-band). The regions shown are 19\\sec\\ squared in all cases. The grey scales are stretched in such a way as to emphasize the morphology of the CNR. Only NGC~3516 is shown with a logarithmic grey scale stretch, all others linear. Black and white border regions are due to rotation of the images and the orientation of the object on the array. The orientation of the images is the same as of those in Fig.~2, i.e., N is up, E to the left.} \\end{figure*} ",
        "conclusions": "This is the first paper in a series exploring the relationships between ring-like star-forming regions in the central few kpc regions of barred galaxies and the morphological properties of their hosts. In this paper, we present new sub-arcsec resolution NIR images obtained with the {\\it CFHT} of the central regions of a sample of 12 barred galaxies with circumnuclear SF activity, and {\\it HST} archival NIR images of most of the sample objects. In Paper~III, we will combine these data with optical imaging of the complete hosts, as presented in Paper II.  Our NIR images reveal a wealth of structure, caused by young stars and dust.  In colour index maps, structures such as SF sites and dust lanes are generally apparent, even in those galaxies where the broad-band images appear featureless. The colour maps of most galaxies show evidence of an often tightly wound spiral structure. In some cases, this spiral structure is visible also in the broad-band images, and as expected due to the higher spatial resolution, more often so in the {\\it HST} than in the ground-based images. Where circumnuclear structure is visible in the {\\it HST} images (in about three quarters of all cases), it shows large numbers of small emitting regions, presumably SF regions, often outlined in spiral arm fragments, and accompanied by dust lanes. Circumnuclear spiral structure thus appears to be common in barred spiral galaxies with circumnuclear SF.  In most of our sample galaxies, and in all those where the star-forming nuclear ring is well defined, radial profiles of surface brightness, colour, ellipticity and position angle show a characteristic bump or change in slope at the radius where the circumnuclear ring is located.  {\\it Acknowledgements} 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."
    },
    "0004/astro-ph0004278_arXiv.txt": {
        "abstract": "We compute the maximum mass of moderately rotating strange stars as a function of the strange quark mass, of the QCD coupling constant, $\\alpha_c$, and of the bag constant (vacuum energy density), $B$, in the MIT bag model of quark matter with lowest order quark-gluon interactions. For a fixed value of $B$, the maximum stellar mass depends only weakly on  $\\alpha_c$, and is independent of this coupling in the limit of massless  quarks. However, if it is the value of the chemical potential of quark matter at zero pressure which is held constant, for example at the value corresponding to the stability limit of nucleons against conversion to quark matter, the  maximum mass of the strange star is higher by up to 25\\% for $\\alpha_c=0.6$, than for non-interacting quarks, and this may be relevant in the discussion of kHz QPO sources. The maximum mass of a non-rotating strange star could be sufficiently high to allow an orbital frequency as low as  1.0 kHz in the marginally stable orbit. However, for all $\\alpha_c<0.6$, the stellar mass cannot exceed $2.6M_\\odot$ at any rotational period $\\ge1.6\\,$ms. ",
        "introduction": "The discovery of kHz QPOs \\citep{1996IAUC.6320....1S,1996IAUC.6319....1V} in several low-mass X-ray binaries (LMXBs) has renewed interest in the maximum mass of neutron stars, as its value limits the maximum observable orbital frequency \\citep{1990ApJ...358..538K}. The maximum masses for neutron stars modeled with various equations of state (e.o.s.) are well established \\citep{1977ApJS...33..415A,1986ApJ...304..115F,1994ApJ...424..823C,Salgado}, and have been examined in the context of kHz QPOs \\citep{1998ApJ...509L..37K}. In principle, LMXBs could contain strange stars instead (Cheng and Dai 1996\\nocite{1996ApJ...468..819C}), and it has been asked whether the observed values of kHz QPOs are compatible with the masses of such stars (Bulik et al., 1999a,b). \\nocite{1999A&A...344L..71B}  The existence of strange stars would shed light on the issue of stability of quark matter. It has been suggested that a quark fluid, composed of roughly equal number of up, down and strange quarks, may be the ground state of bulk matter \\citep{1971Bodmer,1984Witten} and a detailed descripion of such strange matter has been given by \\cite{FarJaf}. This idea has not been universally accepted, and it was argued that for realistic values of the QCD coupling constant, the phase transition to quark matter would occur at densities too high to be of interest (e.g., Bethe, Brown and Cooperstein, 1987). \\nocite{Bethe}  Strange stars may be very hard to distinguish from neutron stars, particularly in LMXBs, as they may have a crust of normal matter \\citep{1986ApJ...310..261A,1986A&A...160..121H}. The crust will contribute little to the mass ($\\le 10^{-5}M_\\odot$), but is expected to have a thickness sufficient to support nuclear burning, including helium flashes responsible for X-ray bursts, thus mimicking neutron stars. The heat released in conversion (at the bottom of the crust) of nuclei into strange matter is directed into the strange matter core \\citep{1990ApJ...362..572M,1990A&A...229..117H}. At any rate, steady release of energy from nuclear conversion is very difficult to distinguish from the gravitational binding energy released in steady accretion.  However, in young radio pulsars the crust thickness does make a difference---the crust in strange stars would be far too thin to allow the redistribution of angular momentum necessary to explain the glitch phenomenon (sudden spin-up of a pulsar), therefore glitching pulsars are though to be neutron stars and not quark stars \\citep{Alpar}. It is likely that the coalescence of two strange stars would lead to the contamination of our Galaxy with chunks of strange matter, and would thus preclude the formation of young neutron stars \\citep{Madsen,Caldwell}. These arguments make unlikely the presence of strange stars in the population of ordinary pulsars (including binary Hulse-Taylor type pulsars) or their accreting counterparts, the high-mass X-ray binaries. However, the presence of strange stars among the millisecond pulsars or LMXBs seems to be allowed \\citep{Kluz94,1996ApJ...468..819C}. In these old systems, the stellar mass itself may yield clues as to the nature of the compact object.  We are led to consider the maximum mass of strange stars in the expectation that LMXB masses will become available in the near future, either through a better understanding of accretion phenomena, including the observed quasi-periodic oscillations (QPOs) in the X-ray light curve, or through classical optical determinations of binary parameters for newly discovered transient sources.  The maximum frequency of stable orbital motion is attained in the innermost (marginally) stable circular orbit (ISCO) allowed by general relativity, and if its value for some X-ray source were that of the observed maximum QPO frequency---ranging from 1.0 kHz to 1.2 kHz in twelve LMXBs---conclusions could be reached about the nature of the compact object. [For a review of QPOs see e.g. \\cite{QPOrev98}.]  The question, whether ISCO frequencies as low as the maximum observed QPO frequencies can be attained outside quark stars, has been answered in the affirmative for rapidly rotating strange stars, with an equation of state (e.o.s.) based on the MIT bag model of quark matter---strange stars rotating close to the equatorial mass-shedding limit can have ISCO frequencies below 1~kHz for masses as low as $1.4 M_\\odot$ \\citep{1999A&A...352L.116S}. However, the same question has not yet been answered for moderately rotating strange stars, of periods $P\\sim 3\\,$ms or more, for which such ISCO frequencies would imply larger masses: $M=2.2M_\\odot(1+0.75j)(1.0\\,{\\rm kHz}/f_{max})$, where $f_{max}$ is the ISCO frequency, and $j\\le0.3$ is the dimensionless angular momentum \\citep{1990ApJ...358..538K,1998ApJ...509L..37K}.  There is a compelling reason to consider stars rotating at relatively low rates. The  period of the compact star in the the transient LMXB SAX J~1808.4-3658 \\citep{1998Natur.394..344W}, the first of possibly many such transients to be discovered, has been measured to be 2.5 ms. Whether or not kHz QPOs will be discovered in that source, its mass may eventually be determined by optical studies of the binary companion.  It has also been argued that the oscillations seen in some X-ray bursters imply a stellar rotational frequency of $\\sim 300\\,$Hz \\citep{1997ApJ...486..355S}. Finally, it is not yet clear whether strange stars can attain the equatorial mass-shedding limit because of the unusually high value of $T/W>0.2$ calculated for their models---note that for Newtonian stars the secular instability to a bar-mode deformation  sets in already at $0.1275\\le T/W\\le 0.1375$ \\citep{1996ApJ...460..379B}. For a discussion of other modes, also in general relativity, see, e.g., Gourgoulhon et al. (1999), and references therein.  For all these reasons, we considered exact numerical models of strange stars rotating with frequencies up to 700~Hz, and found that, to within a few per cent, the maximum mass of such strange stars is the same as that of the static configurations. Therefore, we investigate the maximum mass of static (non-rotating) strange stars, which till now has been discussed in the context of QPOs only in the simplest case of ideal quark gas in the bag model  \\citep{1999A&A...344L..71B}. Here, we use the e.o.s. of interacting, massive quarks within the MIT bag model of self-bound quark matter.    \\begin{figure} \\includegraphics[angle=-90,width=0.95\\columnwidth]{figa.ps} \\caption{The maximum mass of a strange star as a function of the strange quark mass, for various values of the QCD coupling constant. The stellar mass scales as the inverse square root of the bag constant, $B$, provided that the quark mass is scaled with $B^{1/4}$. Here,  $B_1=58.9\\,$MeV$\\,{\\rm fm}^{-3}$. } \\end{figure} ",
        "conclusions": "We have investigated the question whether moderately rotating strange stars with  masses somewhat higher than $2M_\\odot$ are allowed by relativistic equations of stellar equilibrium, and found that the answer could, in principle, be positive within the MIT bag model of beta-stable quark matter. An extension of this discussion to other models of strange matter will be given elsewhere.  The physical constraints on the bag constant following from the basic hypothesis of stability of self-bound quark matter (eq. [\\ref{bmin}]) allow $B$ to be so small, that the corresponding mass of the star could be as high as $2.5M_\\odot$ (eq. [3], Fig.~3). However, the strange star mass cannot be higher than $2.6M_\\odot$, even for stars of rotational period as short as 1.6 ms.  The perturbative approach used here is sensible only if the value of the QCD coupling constant $\\alpha_c$ is small---we used $0\\le\\alpha_c\\le0.6$. In this range, the direct dependence of $M_{max}$ on the coupling is practically negligible (Fig.~1), but the window of allowed values of $B$ does depend on $\\alpha_c$ (Fig.~3). The actual value of $B$ is subject to a very large uncertainty. Fits to the hadronic mass spectrum (DeGrand et al., 1975) gave $B=59~{\\rm MeV~fm^{-3}}$.  The stellar mass decreases with increasing mass, $m_s$, of the strange quark, and is lower by $10\\%$ to $20\\%$ than the one for massless-quark matter for the typical range considered in the literature, $150\\le m_s c^2/{\\rm MeV} \\le 300$ (Madsen 1999). Finally, the maximum stellar mass for strange stars rotating with a frequency up to  $700\\,$Hz, is larger than the one for non-rotating stars by less than 4\\% (Fig.~2).   The physical reason for which large values of mass are obtained for a low value of the bag constant, is that the latter is proportional to the density of quark matter at zero pressure, $\\rho_0$, while the maximum mass of the star is proportional to  $\\rho_0^{-1/2}$. Thus, within the MIT bag model of quark matter, strange stars with mass $M>2.0 M_\\odot$ must have surface densities $\\rho_0< \\rho_1\\equiv 4.2\\times10^{14}\\,$g/cm$^3$ (unless rotating with periods $\\le1.5\\,$ms).  \\bigskip This research was supported in part by KBN grants 2P03D00418, 2P03D02117, 2P03D04013. The numerical calculations have been performed in part on DARC computers purchased thanks to a special grant from the SPM and SDU departments of CNRS.  \\newcommand{\\apjl}{ApJ Let.} \\newcommand{\\apj}{ApJ} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\nat}{Nature} \\newcommand{\\apjs}{ApJ Supp.} \\newcommand{\\aap}{A\\&A} \\newcommand{\\iaucirc}{IAU Circ.}"
    },
    "0004/astro-ph0004087_arXiv.txt": {
        "abstract": "We compute the bispectrum for the galaxy distribution in the \\IRAS QDOT, 2~Jy, and 1.2~Jy redshift catalogs for wavenumbers $ 0.05 \\leq k \\leq 0.2 \\kMpc $ and compare the results with predictions from gravitational instability in perturbation theory.  Taking into account redshift space distortions, nonlinear evolution, the survey selection function, and discreteness and finite volume effects, all three catalogs show evidence for the dependence of the bispectrum on configuration shape predicted by gravitational instability.  Assuming Gaussian initial conditions and local biasing parametrized by linear and non-linear bias parameters $b_1$ and $b_2$, a likelihood analysis yields $1/b_1 = 1.32^{+0.36}_{-0.58}$, $1.15^{+0.39}_{-0.39}$ and $b_2/b_1^2=-0.57^{+0.45}_{-0.30}$, $-0.50^{+0.31}_{-0.51}$, for the for the 2~Jy and 1.2~Jy samples, respectively. This implies that \\IRAS galaxies trace dark matter increasingly weakly as the density contrast increases, consistent with their being under-represented in clusters. In a model with $\\chi^2$ non-Gaussian initial conditions, the bispectrum displays an amplitude and scale dependence different than that found in the Gaussian case; if \\IRAS galaxies do not have bias $ b_{1}> 1$ at large scales, $\\chi^2$ non-Gaussian initial conditions are ruled out at the 95\\% confidence level.  The \\IRAS data do not distinguish between Lagrangian or Eulerian local bias. ",
        "introduction": "Observations suggest that the distribution of matter on cosmological scales is nearly homogeneous, with structure that in our present understanding has grown by gravitational instability from small, nearly Gaussian, initial fluctuations.  On large scales, departures from homogeneity are small, and their evolution can be studied in perturbation theory.  To linear order in perturbation theory, an initially Gaussian fluctuation distribution remains Gaussian, but at higher orders this is no longer true.  Details of the departures depend on the nonlinear couplings; thus, non-Gaussianity can probe the dynamics of structure formation.  One useful non-Gaussian measure is the three-point correlation function, or the bispectrum in the Fourier domain, a quantity that vanishes identically for a Gaussian distribution but is nonzero and has a characteristic behavior as a function of the shape of the triangle defined by its three wavevector arguments if large-scale structure is formed as a result of gravitational instability \\citep{F84}.  When applied to the galaxies, the bispectrum also depends on galaxy biasing, i.e. on the relation between the galaxy and mass distributions \\citep{FG93,F94} and on any non-Gaussianity of the primordial density fluctuations \\citep{FS94}.  Early measurements of the bispectrum in the Lick galaxy catalog did not find the expected gravitational instability signature \\citep{FS82}.  In angular catalogs such as Lick, the density field is integrated over the radial selection function, which in principle can change the dependence on shape, although it appears that in practice the effects of projection on the shape dependence of the bispectrum are not large \\citep{FT99,BKJ00}.  Results on the moderate scales probed by this catalog are also affected by nonlinear evolution beyond leading order in perturbation theory, which weakens the configuration dependence of the bispectrum in a way that mimics the effects of bias \\citep{SCFFHM98}.  A more recent determination of the angular three-point correlation function in the APM survey \\citep{FG99}, which probes significantly larger scales than Lick, shows an amplitude and configuration dependence consistent with gravitational instability from Gaussian initial conditions and no bias, though with large errors.  Recent galaxy redshift surveys have mapped the distribution of galaxies out to large distances.  The largest distances are obtained for deep pencil beam and slice surveys; however, surveys with highly anisotropic one- and two-dimensional geometries have window functions which extend to large wavenumber in one or two directions, corresponding to the smallest survey dimension(s). As a consequence, nonlinearity on this scale tends to wash out the configuration dependence of three-point statistics.  Thus, although the three-point correlation function has been measured in the Las Campanas Redshift Survey \\citep{JB98}, the survey geometry is effectively two-dimensional, with redshift slices of thickness about $ 10 \\Mpc $; this scale is not large enough to allow a direct comparison with perturbation theory.  Thus, it is very difficult to extract quantitative information from this survey about biasing and the nature of the initial conditions.  To compare with predictions of the configuration dependence of the bispectrum from perturbation theory, we wish to observe the bispectrum directly in an analysis that to the extent possible is not affected by the complications of projection and nonlinear evolution.  In this paper, we report measurement of the bispectrum in several galaxy redshift catalogs derived from the \\IRAS Point Source Catalog \\citep{IRASPSC}.  Working with three-dimensional (redshift) data avoids the effects of radial projection.  Moreover, the large volume and nearly full sky coverage of these surveys allow us to make measurements at small wavenumber (large scales), where clustering is well-described by leading order perturbation theory.  Although working in redshift space introduces distortions arising from the modulation of position induced by peculiar velocities, these can be computed \\citep{HBCJ95,VHMM98,SCF99}.  In Section~2 we review the leading order nonlinear perturbation theory predictions for the bispectrum, including its dependence on configuration shape, the distortions introduced by peculiar velocities in redshift space, and the effects of a bias in the galaxy distribution relative to matter.  In Section~3 we summarize properties of the {\\it IRAS}-based redshift catalogs, and we present our estimators in Fourier space, using the optimal weighting of \\citet[FKP]{FKP94}.  In Section~4 we summarize the data analysis method based on comparison with mock catalogs developed in the companion paper \\citep[Paper I]{S00b}, and in Section~5 we present the results of our analysis.  Section~6 contains a final discussion. ",
        "conclusions": "We have presented results for the reduced bispectrum $Q$ of three \\IRAS redshift catalogs, computed using the optimal weighting scheme of FKP, after subtracting contributions to the power from discreteness and from the shape of the selection function. That $Q$ is nonzero, positive, and generally of the shape predicted by gravitational instability is unmistakable. The ability to extract values for bias parameters is marginal; however, after taking into account non-linear redshift distortions and finite volume effects, likelihood analysis yields reasonable values for the parameters of a local Eulerian bias model, $1/b_1 =1.32^{+0.36}_{-0.58}$, $1.15^{+0.39}_{-0.39}$ and $b_2/b_1^2=-0.57^{+0.45}_{-0.30}$, $-0.50^{+0.31}_{-0.51}$, for 2~Jy and 1.2~Jy respectively.  This implies that \\IRAS galaxies trace dark matter with decreasing effective bias as a function of the mass density fluctuation amplitude, consistent with their being under-represented in the cores of clusters. We have also considered constraints on non-Gaussian initial conditions, using as an example the $\\chi^2$ model, in which the bispectrum develops a different amplitude and scale dependence than in the Gaussian case.  If \\IRAS galaxies are {\\it not} positively biased at large scales (i.e., if $b_1^{IRAS} \\leq 1$), then $\\chi^2$ initial conditions are ruled out at the 95\\% confidence level.  The likelihood contours show that one combination of the bias parameters is known to better accuracy: the long axis of the likelihood contours in Figures \\ref{likeG}--\\ref{like_LAG} follows tracks of roughly constant skewness, $ S_3 = 3.25/b_1 + 3 b_2/b^2 $ for $ n = -1.4 $. Our values inferred for $S_3$ are consistent with previous work on the skewness of \\IRAS galaxies. For the 1.2~Jy sample, \\citet{BSDFYH93} obtained $ S_3 = 1.5 \\pm 0.5 $; for the same catalog, \\citet{KS98}, who include a finite volume correction but use a different estimator, obtained $ S_3 = 2.83 \\pm 0.09 $. Our results translate via \\eq{S3} (assuming $n=-1.4$) into a finite-volume corrected value $S_3 \\approx 2.25 \\pm 0.9$, comparable to that of \\citet{KS98}. To estimate the error bar in the skewness we assume that $S_3$ corresponds to sum of $Q$ over all triangles. The error is then roughly given by three times the width in the best determined direction in the ellipse in Figure~\\ref{likeG}, $\\pm 0.3$. An approximate estimate of $S_3$ that does not include the finite volume corrections can be obtained directly from Figure~\\ref{likeG}, by using the fact that at the $68\\%$ limit, $Q/b_1+b_2/b_1^2$ is in the range 0.2--0.8; this implies $S_3 \\approx 3 Q_g \\approx 1.5 \\pm 0.9$, consistent with \\citet{BSDFYH93}. These are only rough estimates, as they ignore the skewness of the bispectrum (and $S_3$) PDF.  Our results have interesting implications for the relative clustering of \\IRAS and optically selected galaxies. From the ratio of power spectrum amplitudes between the Stromlo-APM survey and the combined QDOT-1.2Jy survey, \\citet{TE96} conclude that $ b_1^{opt}/b_1^{IRAS} \\approx 1.2 $, while from the analysis of the APM projected three-point function, \\citet{FG99} conclude that $b_1^{opt} \\approx 1$.  These results are consistent with our findings for $b_1^{IRAS}$.  Future galaxy surveys will be able to quantify this issue with a greater level of accuracy."
    },
    "0004/astro-ph0004166_arXiv.txt": {
        "abstract": "We review the evidence for cosmic ray acceleration in the superbubble/hot phase of the interstellar medium, and discuss the implications for the composition of cosmic rays and the structure and evolution of the interstellar medium (ISM). We show that the bulk of the galactic supernovae, their expanding remnants, together with their metal-rich grain and gas ejecta, and their cosmic ray accelerating shocks, are all confined within the interiors of hot, low-density superbubbles, generated by the multiple supernova explosions of massive stars formed in giant OB associations. This superbubble/hot phase of the ISM provides throughout the age of the Galaxy a cosmic ray source of essentially constant metallicity for acceleration by the shocks of many supernovae over time scales of a few Myr, consistent with both Be/Fe evolution and ACE observations of $^{59}$Ni/$^{59}$Co. We show that the expected metallicity ($>$ 2 times Solar) and filling factor ($>$ 50\\%) of the superbubble/hot phase is high enough that the composition of cosmic rays accelerated from fast, supernova grains in these superbubbles is quite consistent with the both Be/Fe and cosmic ray data, while their acceleration from grains in the well-mixed cooler phases of the ISM is {\\em not} consistent with observations. We also show that if the refractory cosmic ray metals come from the sputtering of fast refractory grains then the accompanying scattering of ambient gas by these fast grains can also account for the relative abundance of cosmic ray volatiles. ",
        "introduction": "ACE measurements by Wiedenbeck et al. (53) of the cosmic ray $^{59}$Ni/$^{59}$Co abundance ratio and optical measurements of the Be/Fe abundance ratios in old stars, e.g. Molaro et al. (34) and Boesgaard et al. (4), have shown us that the making of cosmic rays is like the making of fine wine. Both have to be aged and blended. Cosmic rays can not be too young and fresh, because their low ratio of K-capture $^{59}$Ni to its daughter $^{59}$Co requires an age of $> 10^5$ years before acceleration, and they can not be too old and diluted, because the roughly constant ratio of cosmic ray produced Be to supernova produced Fe requires the metallicity in the matter from which the cosmic rays are accelerated to be roughly constant within a factor of $\\sim$2. This can be achieved if the metallicity is dominated by supernova ejecta.  The cellar where all of this aging and blending happens quite naturally is the vast metal-enriched superbubble/hot phase of the interstellar medium (SB/HISM), where most supernovae occur. This can be seen from the extensive studies of Galactic and extragalactic supernovae, their progenitors and their cosmic ray accelerating remnants, and of the chemical enrichment and evolution of the interstellar medium, that we will review here with special emphasis on the metallicity and filling factor of the SB/HISM. We also discuss how ACE and other measurements of the cosmic ray composition can define these two properties in the region of cosmic ray acceleration.  The source of energy for cosmic ray acceleration is thought to be shock waves driven by the expansion energy of supernova ejecta, e.g. Blandford \\& Ostriker (3) and Axford (1). The power required to maintain the Galactic cosmic rays is about 10$^{41}$ ergs$^{-1}$, e.g. Lingenfelter (23). The Galactic supernova rate is about 3 supernovae per century, e.g. van den Bergh \\& McClure (49), most of which (80 to 90\\%) are core-collapse (Type II and Ib/c) supernovae of relatively young ($<$ few 10$^7$ yrs) massive O and B stars; the remainder are Type Ia thermonuclear explosions of much older accreting white dwarfs. Thus, the average cosmic ray energy needed per supernova is about 10$^{50}$ ergs, which requires an acceleration efficiency of about 10\\% from the blast wave shocks of the supernovae, since supernovae all seem to have similar ejecta kinetic energies of about 10$^{51}$ ergs, e.g. Woosley \\& Weaver (56) and Nomoto et al. (36).  The source of the particles that are accelerated as cosmic rays and the site of their acceleration, however, are still debated. But if the energy comes from supernova shocks, the site and source of the particles clearly must be in the material through which the shocks pass. New clues to origin of the particles come from recent measurements, e.g. Molaro et al. (34) and Boesgaard et al. (4), of Be/Fe abundances in old halo stars that show that the ratio of cosmic ray spallation produced Be relative to core-collapse supernova-produced Fe has remained roughly constant throughout the evolution of the Galaxy. This constancy requires (Ramaty, Kozlovsky \\& Lingenfelter (38); Ramaty, Lingenfelter \\& Kozlovsky (40)) that the cosmic rays be accelerated out of matter that is only partially diluted by mixing with the interstellar medium (ISM) so that it is still sufficiently enriched in supernova-synthesized metals that its metallicity did not changed by more than $\\sim$2 over galactic evolution. Detailed calculations by Ramaty et al. (42) and Ramaty, Lingenfelter \\& Kozlovsky (41) of the production and evolution of the Galactic Be/Fe ratio clearly show that the bulk of the cosmic rays can not be accelerated from the well-mixed ISM, as has been recently assumed, e.g. Meyer et al. (33) and Ellison et al. (11).  Just such a metal-enriched environment is where most supernovae do in fact occur and through which most of the cosmic ray accelerating, supernova shocks propagate. As we have discussed  in Higdon, Lingenfelter \\& Ramaty (17,18), extensive observations show that the bulk of the core-collapse supernova progenitors are formed in OB associations in giant molecular clouds and that the combined winds and supernova ejecta of these stars form hot, low density superbubbles, that reach dimensions of several hundred pc and last for tens of Myr. During this time the bulk of the supernova ejecta, the supernova shocks, and their cosmic ray acceleration are all confined within the superbubble/hot phase of the ISM. ",
        "conclusions": ""
    },
    "0004/astro-ph0004350_arXiv.txt": {
        "abstract": "We report high-resolution, high-signal-to-noise, observations of the extremely metal-poor double-lined spectroscopic binary CS~22876--032.  The system has a long period : P = 424.7 $\\pm$ 0.6 days.  It comprises two main sequence stars having effective temperatures 6300~K and 5600~K, with a ratio of secondary to primary mass of 0.89 $\\pm$ 0.04.  The metallicity of the system is [Fe/H] = --3.71 $\\pm$ 0.11 $\\pm$ 0.12 (random and systematic errors) -- somewhat higher than previous estimates.  We find [Mg/Fe] = 0.50, typical of values of less extreme halo material. [Si/Fe], [Ca/Fe], and [Ti/Fe], however, all have significantly lower values, $\\sim$~0.0--0.1, suggesting that the heavier elements might have been underproduced relative to Mg in the material from which this object formed. In the context of the hypothesis that the abundance patterns of extremely metal-poor stars are driven by individual enrichment events and the models of Woosley and Weaver (1995), the data for CS~22876--032 are consistent with its having been enriched by a zero-metallicity supernova of mass 30 M$_{\\odot}$.  As the most metal-poor near-main-sequence-turnoff star currently known, the primary of the system has the potential to strongly constrain the primordial lithium abundance.  We find A(Li) (= log (N(Li)/N(H))~+~12.00)~=~2.03~$\\pm$~0.07, which is consistent with the finding of Ryan et al. (1999) that for stars of extremely low metallicity A(Li) is a function of [Fe/H]. ",
        "introduction": "In the past two decades a handful of stars with heavy element abundance less than 1/1000 that of the sun ([Fe/H]\\footnote{[Fe/H] = log~(N$_{\\rm Fe}$/N$_{\\rm H}$)$_{\\rm star}$--log~(N$_{\\rm Fe}$/N$_{\\rm H}$)$_{\\odot}$} $<$ --3.0) have been analyzed at high spectral resolution, with a view to understanding the production of the elements at the earliest times.  Such stars have the potential to constrain Big Bang Nucleosynthesis (Ryan, Norris \\& Beers 1999), the nature of the first supernovae (McWilliam et al. 1995; Ryan, Norris \\& Beers 1996; Nakamura et al. 1999), and the manner in which the ejecta from the first generations were incorporated into subsequent ones (Ryan et al. 1996; Shigeyama \\& Tsujimoto 1998; Tsujimoto \\& Shigeyama 1998; Ikuta \\& Arimoto 1999; Tsujimoto, Shigeyama \\& Yoshii 1999).  In special cases, they can even be used to determine the age of the Galaxy (Cowan et al. 1999).  As relatively simple objects, formed at redshifts $\\gtrsim$~4--5, they nicely complement and constrain abundance results from the more complicated and less well-understood Lyman--$\\alpha$ clouds and Damped Lyman--$\\alpha$ systems currently studied at redshifts z $<$ 3.5 (see eg Ryan 2000).  The most metal-poor objects are generally faint, and studies to date have been undertaken with somewhat limited signal-to-noise (S/N).  For the six stars having [Fe/H] $<$ --3.5 and so far analyzed at high resolution (see Table 6 of Ryan et al. 1996), perusal of the sources suggests that the data were obtained with a representative S/N per equivalent 0.04 {\\AA} pixel of roughly 30, corresponding to a 3$\\sigma$ detection limit for weak lines of $\\sim$~10 m{\\AA}.  In some cases this has led to uncertainty as to the existence of important elements in these objects\\footnote{For example, from Norris et al. (1996) : ``Ba 4554 {\\AA} measurements of CD--38$^{\\rm o}$245 are rather discordant, ranging from detection at 18 m{\\AA} and 19 m{\\AA} (Bessell \\& Norris 1984; Primas et al. 1994) to upper limits $<$~10 m{\\AA} and $<$~7 m{\\AA} by Molaro \\& Castelli (1990) and Peterson \\& Carney (1989).''}.  As one moves to the lowest abundances (or higher temperatures), the problem of determining reliable abundances can only deteriorate: for example, in the two near-main-sequence-turnoff dwarfs in the above-mentioned sample there are fewer than 20 non-Fe lines detected in the preferred wavelength range (3700--4700 {\\AA}), of which fewer than half are stronger than 20 m{\\AA}.  To improve the accuracy of the existing abundance analyses we are attempting to obtain high resolution (R = 40,000) data with S/N = 100 of objects with [Fe/H] $<$ --3.5.  The purpose of the present work is to present results for CS~22876--032.  In a subsequent paper we shall present results for CD--24$^{\\rm o}$17504, CD--38$^{\\rm o}$245, CS~22172--002, and CS~22885--096.  CS~22876--032 was discovered by Beers, Preston, \\& Shectman (1985), and with a high-resolution abundance determination of [Fe/H] = --4.3 (Molaro \\& Castelli 1990; Norris, Peterson, \\& Beers 1993) it is currently the most metal-poor dwarf for which such analysis exists.  The abundance determinations for this object are, however, compromised to some extent by the fact that it is a double-lined spectroscopic binary (Nissen 1989) of long but undetermined period.  We set out to determine the nature of the two components and thus improve the reliability of the abundances in the system.  CS~22876-032 has the added attraction that it is the main-sequence dwarf closest in abundance to the material which emerged from the Big Bang. It therefore has a potentially important role in constraining the primordial abundance of Li, especially if the Li abundance in higher metallicity stars has been elevated by Galactic chemical evolution, as we contend elsewhere (Ryan et al. 1999). However, only after its atmospheric parameters have been reliably determined will one be in a position to fully realize that potential.  The outline of the paper is as follows. The observational material upon which our analysis is based is presented in \\S~2, while in \\S~3 we use the available material to present a first estimate of the period and orbital elements of the system.  These results are coupled with the observed {\\it BVRI} colors and available isochrones in \\S~4 to determine the atmospheric parameters T$_{\\rm eff}$, log~$g$, [Fe/H], and ${\\xi}$ of the components, together with element ratios [X/Fe].  Finally, in \\S~5 we discuss the results and revisit the question of the Li abundance of the system, and its role in constraining the primordial value. ",
        "conclusions": "\\subsection{Heavy Element Abundances}  Evidence has mounted over the past decade that below [Fe/H]~=~$-3.0$, Galactic chemical enrichment was a rather patchy business (e.g. Ryan, Norris \\& Bessell 1991; McWilliam et al. 1995; Ryan et al. 1996), leading to models in which such metal-poor objects are formed in the interaction of supernova remnants with material in their immediate vicinity (Ryan et al. 1996; Tsujimoto et al. 1999). In such models, one expects the abundance patterns of the most metal-poor stars to be representative of individual supernova, rather than of a time-averaged ensemble.  Hence, high-resolution spectroscopic analyses of extremely metal-poor stars provide a powerful probe of the nature of the first supernovae in the Milky Way.  Given the paucity of lines in CS~22876--032 we have only Al and the $\\alpha$--elements to discuss.  In the context of star-to-star scatter we shall defer further consideration until we present results on the other four extremely metal-poor stars for which we have data, and which we mentioned in \\S~1 -- except to note here that the value [Al/Fe] = --0.44 for CS~22876--032 lies within $\\sim$~0.20 of the mean line determined for the sample of metal-poor stars discussed in the investigation of Ryan et al. (1996, Figure~3).  In the spirit of associating the yields of particular supernovae models with the observed abundances in extremely metal-poor stars, we conclude by discussing the $\\alpha$--elements.  For Mg, which is produced mainly during hydrostatic carbon and neon burning we find [Mg/Fe] = 0.50.  For Si and Ca we find a weighted mean value, which we designate [$\\langle$Si,Ca$\\rangle$/Fe], of --0.05 (the data were weighted by the inverse square of their errors). These elements are produced by a combination of hydrostatic oxygen shell burning and explosive oxygen burning, which varies from star to star (Weaver \\& Woosley 1993). As discussed by Woosley \\& Weaver (1995), the predicted yields of Si and Ca are expected to vary relative to Mg, depending sensitively on several physical effects.  (These include among others the treatment of convection, the density structure near the iron core, the location of the mass cut, and the amount of material which falls back in the explosion.)  For CS~22876--032 we have [$\\langle$Si,Ca$\\rangle$/Mg] = --0.55. It is then interesting to consider the production factors of the zero-heavy-element model supernovae of Woosley \\& Weaver (1995, Table 17). For those models which produce significant amounts of Mg, Si, and Ca we find the results presented in Figure~8.  If one were to seek to interpret the data for CS~22876--032 in terms of enrichment from a single zero-heavy-element supernova one could identify it, quite reasonably, with the Woosley \\& Weaver (1995) 30 M$_{\\odot}$ model Z30B.  Titanium is produced under more extreme conditions than Si and Ca, as discussed by Woosley \\& Weaver (1995).  In CS~22876--032, as for Si and Ca, [Ti/Fe] is low relative to [Mg/Fe]: we find [Ti/Mg] = --0.39.  This is also best reproduced by Woosley \\& Weaver model Z30B.  Is our result of [Al/Mg] = --0.94 $\\pm$ 0.17 consistent with the hypothesis?  Before comparison with theory one must correct the present LTE abundance for Al for non-LTE effects.  According to Baum\\\"{u}ller \\& Gehren (1997), LTE abundances for Al based on the line we have used underestimate the abundance by $\\sim$~0.65 dex at [Fe/H] $\\sim$~--2.5$-$--3.0. Adopting this correction for CS~22876--032, one therefore has [Al/Mg] $\\sim$~--0.3 $\\pm$ 0.2. For the above-mentioned models of Woosley \\& Weaver (1995), in comparison, one finds [Al/Mg] in the range $\\sim$~--0.65$-$--0.35.  That is, the models do not strongly constrain the situation, and within the observational errors one has reasonable agreement between theory and observation.  \\subsection{Lithium Abundance}  Given the improved atmospheric parameters for the components of CS~22876--032, one may revisit the question of its lithium abundance.  The best currently available observational material is that of Thorburn \\& Beers (1993).  They obtained data with S/N = 150 and reported W$_{\\rm A}$(Li I 6707{\\AA}) = 11 $\\pm$ 1.3 m{\\AA} at JD 2448850.8, when the lines of the two components were well separated.  They did not detect the Li line in the secondary.  Norris et al. (1994) have also observed this object and report W$_{\\rm A}$(Li I 6707{\\AA}) = 15 $\\pm$ 2.1 m{\\AA} from a spectrum with S/N = 70, obtained at JD 2448490.5 when the lines of the primary and secondary would have been well separated, by some 20 km~s$^{-1}$.  As part of our continuing interest in the lithium problem we attempted to remeasure the equivalent width during an observing run on the Anglo-Australian Telescope in 1999 September, using techniques similar to those described by Norris et al. (1994) and Ryan, Norris, \\& Beers (1999). From a spectrum having S/N = 70 per 0.04 {\\AA} increment we obtain W$_{\\rm A}$(Li I 6707{\\AA}) = 14.4 $\\pm$ 2.3 m{\\AA}.  At the time of observation the lines of the primary and secondary were separated by 21 km s$^{-1}$, and the line in the secondary was below our threshold of detectability, which we estimate to be 7 m{\\AA} (3$\\sigma$).  We seek to compare these data with the recent accurate and homogeneous observational material of Ryan et al. (1999).  We first ask if there are systematic equivalent width errors between the various authors.  For 16 stars in common between Ryan et al. (1999) and Thorburn (1994) (who used the same equipment and techniques as those of Thorburn \\& Beers) one finds $\\langle$W$_{\\rm RNB}$--W$_{\\rm T}$$\\rangle$ = --1.9 $\\pm$ 1.1 (s.e.) m{\\AA}, while for 5 stars in common between Ryan et al. and Norris et al. (1994) one finds $\\langle$W$_{\\rm RNB}$--W$_{\\rm NRS}$$\\rangle$ = +2.3 $\\pm$ 1.7 (s.e.) m{\\AA}.  We do not regard these differences as statistically significant, and make no correction to the reported equivalent widths.  We then adopt W$_{\\rm A}$(Li~I 6707{\\AA}) = 12.5 $\\pm$ 1.4 m{\\AA} as the best currently available value for CS~22876--032, having weighted the data by the inverse square of their estimated errors.  (The error is based on small sample statistics following Keeping (1962) without cognisance of the weights.)  At $\\lambda$6707{\\AA} our model shows that to obtain the line strength intrinsic to the primary this should be corrected by the factor f$_{6707,\\rm P}$ = 1.35\\footnote{Based on {\\it B-R} results from the RYI and on model atmosphere synthetic spectra, which give consistent results.}, which leads to W$_{\\rm P}$(Li~I 6707{\\AA}) = 16.9 $\\pm$ 1.9 m{\\AA}.  The recent study of Li by Ryan et al. (1999) used an effective temperature scale different from the RYI one that accompanies the isochrones used in the decomposition of the binary. It would be inappropriate to compare the Li abundance of CS~22876-032 computed using the RYI temperature scale, with a set of data based on an entirely different one. For this reason, we use the decomposed {\\it colors} of the binary, as given in Table~4, to infer the temperature on the same scale as our 1999 Li work.  We obtain T$_{\\rm eff}$~=~6223~$\\pm$~38~K.  The inferred lithium abundance\\footnote{The Li abundance was calculated using exactly the same formalism as in Ryan et al. (1999), i.e. computing synthetic spectra for Bell (1983) models using four $^7$Li components, and measuring the equivalent width of the synthesized line.}  A$_{\\rm P}$(Li) (= log (N$_{\\rm P}$(Li)/N$_{\\rm P}$(H))~+~12.00)~=~2.03~$\\pm$~0.07\\footnote{The perceptive reader will note that in this work we have combined the higher S/N data of Thorburn \\& Beers (1993; S/N = 150) with our lower S/N data (two spectra each with S/N = 70), whereas in a recent work (Ryan et al. 1999) on the spread about the plateau we advocated the importance of homogeneity. We would not advocate the coaddition of data were we studying the spread here, since the statistic used to describe spread -- the standard deviation -- is acutely sensitive to errors, which enter as squared terms. However, the current comparison is with the mean abundance of the stars, and the mean is less affected than the standard deviation by divergent phenomena. Consequently we have chosen to use the superior-S/N data of Thorburn \\& Beers as well. We note that if we instead adopted only the less accurate (S/N-limited) value of $W_{\\rm A}$ = 14.4$\\pm$2.3~m{\\AA} reported in the present work, the Li abundance would be 0.06~dex higher and the uncertainties larger at $A_{\\rm P}$(Li) = 2.09 $\\pm$ 0.08 (1$\\sigma$). In this case the star moves further above the trend, but on account of its much larger error bars, has little impact on the slope.}. (The uncertainty in the quadratic addition of errors corresponding to $\\Delta$f$_{\\rm 6707,P}$~=~0.01, $\\Delta$T$_{\\rm eff}$~=~38~K, and $\\Delta$W$_{\\rm P}$(Li I 6707\\AA)~=~1.9~m\\AA.  We recall that here `P' by convention refers to `Primary' and not necessarily `Primordial'.)  The non-detection of the secondary by Thorburn \\& Beers (1993) and the present work is also important.  If this component has T$_{\\rm eff}$ = 5600~K and the same Li abundance as the primary, one would expect an intrinsic line strength W$_{\\rm S}$(Li~I 6707{\\AA}) = 45 m{\\AA}.  Applying a model correction factor f$_{6707, \\rm S}$ = 3.9 (applicable at $\\lambda$6707 {\\AA} for the secondary), the observed line strength would then be 12 m{\\AA}.  The most likely explanation for the non-detection of Li in the secondary is that with T$_{\\rm eff}$ = 5600~K it lies at the cool edge of the Spite Plateau and has experienced modest depletion to below its primordial value.  The uncertainty in the temperature of the secondary is also greater than that of the primary, as is evident from Figure~6, and could explain part of the problem.  From their investigation of an unbiased sample of 23 near-main-sequence-turnoff stars having 6050~K $<$ T$_{\\rm eff}$ $<$ 6350~K and --3.6 $<$ [Fe/H] $<$ --2.3 Ryan et al. (1999) reported a dependence of A(Li) on [Fe/H].  Their \"best fit\" was A(Li) = 2.447($\\pm$0.066) + 0.118($\\pm$0.023)$\\times$[Fe/H].  They argued that the most likely explanation of the positive trend with metallicity was that the primordial lithium abundance had been increased by early Galactic chemical enrichment involving the interaction of cosmic rays with the interstellar medium (see also Ryan et al. (2000)).  Their material and its \"best fit\" trend, together with the present result for CS~22876--032, are presented in Figure~9.  It is clear that the data for the latter object, which is currently the lowest metallicity dwarf known, are completely consistent with the trend reported by Ryan et al. and support their hypothesis.  Equally clearly, as noted by those authors, more data are required at [Fe/H] $<$ --3.5 to place their result, and its implications for the primordial Li abundance, on a firmer basis.  We also note that to infer the primordial abundance from the observed one, a range of potentially systematic as well as random uncertainties needs to be taken into consideration. Furthermore, the computations of the primordial Li abundance as a function of baryon-to-photon ratio, $\\eta$, or equivalently $\\Omega_{\\rm B}$, are themselves subject to significant uncertainties. The impact of these issues is discussed more fully elsewhere (Ryan et al. 2000)."
    },
    "0004/astro-ph0004399_arXiv.txt": {
        "abstract": "We suggest a new component of the Milky Way that can account for both the optical depth and the event durations implied by microlensing searches targeting the Large Magellanic Cloud.  This component, which represents less than $4\\%$ of the total dark matter halo mass, consists of mainly old white dwarf stars in a distribution that extends beyond the thin and thick disks, but lies well within the dark matter halo. It is consistent with recent evidence for a significant population of faint white dwarfs detected in a proper motion study of the Hubble Deep Field that cannot be accounted for by stars in the disk or spheroid. Further, it evades all of the current observational constraints that restrict a halo population of white dwarfs. ",
        "introduction": "The past few years have yielded exciting new data from observational teams searching for evidence of MACHOs (dark compact objects) in the halo of our galaxy, data which have raised many new questions. The microlensing optical depth toward the Large Magellanic Cloud (LMC) obtained by the MACHO collaboration, $\\tau_{\\rm LMC} = (1.2^{+0.4}_{-0.3})\\times 10^{-7}$ (Cook et al. 2000), is consistent with a significant fraction of the galactic halo ($\\sim 20\\%$) in the form of MACHOs.  The duration of these events also indicates that, under the assumption of a spherical isothermal halo distribution for the MACHOs, the average MACHO mass is $ \\sim 0.5 M_\\odot$ (with large statistical uncertainties).  Such masses suggest several candidates for the lenses including faint halo stars, white dwarfs, and black holes. Direct searches for faint halo stars have placed severe limits on their contribution to halo, requiring it to be less than about $3\\%$ (Flynn, Gould, \\& Bahcall 1996; Graff \\& Freese 1996). Primordial black hole candidates require a fine tuning of the initial density perturbations, and details of QCD phase transition black hole formation remain to be worked out in order to assess their viability and mass function (Jedamzik \\& Niemeyer 1999). Thus in the standard halo model interpretation, white dwarfs appear to be the remaining strong candidate for the lenses.  However, other possible interpretations of the MACHO results have been proposed in order to avoid the difficulties associated with a large halo population of white dwarfs.  The mass estimate for the lenses depends upon the assumed phase space distribution of MACHOs. There are large uncertainties in the halo model parameters, including the distribution and velocity structure of the dark matter, and attempts have been made to exploit these uncertainties in order to obtain mass estimates from the current data that are consistent with lenses in the substellar regime (e.g. brown dwarfs). Previous work by the authors and others have examined a wide range of halo models, including flattened halos, halos with a bulk rotational component to the velocity structure (Gyuk \\& Gates 1998) and halos with anisotropic velocity dispersions (Gyuk, Evans, \\& Gates 1998).  These analyses have shown that for any reasonable (smoothly varying) phase space distribution of the lenses, the implied lens mass is still much larger than the hydrogen burning limit, and thus one cannot appeal to modeling uncertainties in order to invoke brown dwarfs as candidates for the MACHOs.  Other work has explored the possibility that the lenses are not in the halo of the Milky Way.  LMC self-lensing was originally suggested by Sahu (1994), but recent work by Gyuk, Dalal \\& Griest (1999) has argued that this is unlikely to be able to account for the current data. More recently, an unvirialized component of the LMC has been proposed in order to produce sufficient LMC self-lensing (Graff 2000). Zaritsky \\& Lin (1997) and Zhao (1998) have suggested an intervening population of stars toward the LMC (tidal debris or a dwarf galaxy) could be responsible for the microlensing events. Again this suggestion has been subject to much debate, and a recent paper by Gould (1999) argues strongly against such a scenario.  Galactic models in which dark extensions of known populations, such as a heavy spheroid or thick disk, could be the source of the lenses were explored by Gates, Gyuk, Holder \\& Turner (1998).  However, recent results from a proper motion study of the Hubble Deep Field (HDF) (Ibata et al. 1999 ), and from a comparison of the north and south HDF images (Mendez \\& Minniti 1999) have added a new piece to the puzzle. These studies provide further evidence that there may be a previously undetected population of old white dwarfs in the galaxy. Using new models for white dwarf cooling (Hansen 1999), the sources detected by Ibata et al. are consistent with 2-5  old ($> 12 Gyr$), $0.5 M_{\\odot}$ white dwarfs with kinematics roughly consistent with that expected for halo objects. While earlier searches for the local counterparts did not find any such objects (Flynn et al. 2000), recent surveys have turned up several candidates (Hodgkin et al. 2000; Ibata et al. 2000). Overall, these data strengthen the interpretation of the MACHO lenses as white dwarfs in our own galaxy, and thus make alternative scenarios, such as LMC self-lensing, less appealing.  While these new data are still somewhat preliminary, they do raise the intriguing possibility that there is a previously undetected population of white dwarf stars that are not part of the disk or (known) spheroid. Along with the MACHO data and the inability of modeling to significantly change the mass estimates, these new results, seem to be relentlessly pointing to white dwarfs as the lenses.  So, is the halo of our galaxy filled with white dwarfs? A standard halo interpretation of these data would say yes -- a significant fraction of the galactic halo must be in white dwarfs.  However, such a scenario faces serious challenges from many directions. ",
        "conclusions": "We have argued that the microlensing data toward the LMC, combined with observations of white dwarf stars in a proper motion study of the HDF indicate the presence of a new component of the galaxy.  This component can be generally described as an extended distribution that extends at least 2 kpc above the galactic plane, but resides well within the dark matter halo.  It is consistent with all data and observations of the structure and kinematics of the galaxy, and significantly alleviates the considerable problems with a halo population of white dwarf stars that is consistent with microlensing data.  Much work remains to carefully consider the implications of such a component, in particular the formation and evolution of the early population of progenitor stars (and resulting metal enriched gas) that produced this component. However, the significantly lower mass in the progenitor population as compared to that for a halo population of white dwarfs will allow a reasonable fraction of the baryonic mass of the Universe to remain in gas that has not been processed through these very early stars.  Moreover, this component may be a more reasonable distribution for the remains of an early starburst population, in which one would expect a more condensed distribution than that of the halo."
    },
    "0004/astro-ph0004020_arXiv.txt": {
        "abstract": "We present near infrared spectroscopy of the superluminal microquasar \\grs\\ obtained with the first unit of the VLT \\footnote{Based on observations collected at the European Southern Observatory, Chile (ESO No 63.H-0261).} and the ISAAC spectro-imager.  The emission features detected in the VLT data have been identified as He I, Br $\\gamma$, He II and Na I. The detection of Na I is reported here for the first time, while our confirmation of weak He II emission provides support to previous marginal detections of this feature. By comparing the observed spectra with those of massive stars, we find that our results are very consistent with \\grs\\ being a high mass X-ray system with an early type primary, as previously proposed by Chaty et al. (1996) and  Mirabel et al. (1997). The VLT spectra also provide evidence of P Cygni profiles, that turn into blue emission wings when the system is in outburst. This observed line profile evolution implies that \\grs\\ must be surrounded by an expanding envelope, that is partially blown out during the X-ray outbursts. The presence of such circumstellar gaseous material around \\grs\\ is more naturally understood in the context of a massive luminous star than if the system was a low-mass X-ray binary. ",
        "introduction": "\\label{intro}  The X-ray transient \\grs\\ was first discovered  by Castro-Tirado et al. (1992) using the WATCH all-sky monitor on board the Russian satellite GRANAT. Special attention to this source increased enormously a few years later when Mirabel \\& Rodr\\'{\\i}guez (1994) reported that \\grs\\ was producing repeated ejections with apparent superluminal motion.  These events could be first followed at radio wavelengths thanks to the Very Large Array interferometer of NRAO and represented a phenomenon never observed before in the Galaxy. There are only two confirmed galactic superluminal sources known, the other one being GRO J1655$-$40 (Hjellming \\& Rupen 1995; Tingay et al. 1995). Both of them have been classified as `microquasar' systems, i.e., a selected group of radio emitting X-ray binaries whose physics and external appearance is strongly reminiscent of quasars and active galactic nuclei. We refer the reader to the review by Mirabel \\& Rodr\\'{\\i}guez (1999) for an updated account.  \\begin{figure*}[htb] \\mbox{} \\vspace{11.0cm} \\special{psfile=Fig1.ps angle=-90 hoffset=-10 voffset=70 hscale=70 vscale=70} \\caption[]{VLT low resolution spectra of \\grs\\ obtained on 1999 May 29 (before a radio outburst) and 1999 July 04 (during a radio outburst) using the SK filter with ISAAC. The emission line identifications are indicated.} \\label{isaac} \\end{figure*}   Since the discovery of superluminal episodes in \\grs, this object has been under strong surveillance by several observers at different wavelengths. In spite of this work, the nature of the X-ray binary system in \\grs\\ is still a matter of debate. The strong interstellar absorption towards the source, that implies optical extinctions of $A_V \\sim 27$ magnitudes (Chaty et al. 1996), appears as the main cause of this uncertainty. Therefore, it is not surprising that the best approaches to the problem have come from imaging and spectroscopic studies at infrared wavelengths, where the effects of extinction are much less severe ($A_K \\sim 3.0$ mag). Infrared results, however, have not yet reached an unanimous consensus. \\grs\\ has been proposed to be a low-mass X-ray binary (LMXRB) by Castro-Tirado et al. (1996) and an high-mass X-ray binary (HMXRB) by Chaty et al. (1996) and Mirabel et al. (1997), who favour an Oe/Be spectral type. There is also a discussion on the reality of weak emission lines of He II, that have been marginally detected by Castro-Tirado et al. (1996) and Eikenberry et al. (1998a).  The existence, or not, of He II lines is relevant because they are naturally expected in LMXRBs as a result of excitation from the hard photons of the accretion disk. The prototypical LMXRB Sco X-1 is a good example of this statement (Bandyopadhyay et al. 1999). In contrast, He II is seldom observed in HMXRBs. In massive systems, accretion often proceeds directly from the stellar wind and photons from typical companion photospheres would not be energetic enough to ionize helium atoms. However, the He II criterion is not a reliable one to distinguish between high and low mass systems. Indeed, a few objects in the HMXRB group do exhibit He II lines in emission. The best representative is certainly A0538$-$66 (Charles et al. 1983) in the Large Magellanic Cloud. This B2 III-IV system has been observed to produce strong and variable He II 4648 \\AA\\ emission when in outburst. An interesting property of some emission lines in A0538$-$66 (e.g. C IV and He II) is that they alternate between a P Cygni profile and a noticeable blue wing in emission when in quiescence and outburst, respectively. We advance here that the same behavior may be present in \\grs. Another case of massive He II emitter is the classical black hole candidate Cygnus X-1, whose spectrum displays He II 4846 \\AA\\ emission believed to form in the stellar wind above the substellar point of the O9.7 Iab primary (Gies \\& Bolton 1986).  In this context, we present new infrared spectroscopic observations of \\grs\\ obtained with the first 8.2 m unit (Antu) of the Very Large Telescope (VLT) of the European Southern Observatory (ESO). The VLT spectra show, among others ones, possible evidence for He II and a newly detected doublet of Na I, both in emission. When taking all these features into account we find compelling evidence that a HMXRB scenario, with a surrounding envelope, provides a very likely interpretation for \\grs. ",
        "conclusions": "\\begin{enumerate}  \\item We have presented the first VLT infrared spectra of the superluminal source \\grs. The strong emission line of He I at 2.0587 $\\mu$m provides suggestive evidence of a P Cygni profile evolving into a blue wing in emission when the system goes from quiescence to outburst. A similar trend is also likely to occur for Br $\\gamma$. This behavior is consistent with what is observed during the outbursts of other HMXRBs and points towards the existence of an expanding shell or atmosphere around the \\grs\\ companion. We speculate that the appearance of such shell is changing depending on the state of activity of the high energy source.  \\item The \\grs\\ companion is very likely to be a massive star. Not only the presence of an extended envelope is a clue of its HMXRB nature, but also the comparison of the \\grs\\ spectrum with infrared spectral atlas supports this interpretation. The newly discovered line features, specially Na I, as well as those previously detected show a strong resemblance with the spectra of such massive stars. We are also able to propose a tentative spectral classification of the primary as an Of/WN star.  \\item In this context, the long discussed He II line should be excited by the photons of the massive and luminous companion star, although we cannot rule out completely that the accretion disk is also playing a role. In any case additional sensitive observations, in different states of activity of the source, would be very valuable to settle definitively this issue and to provide a more accurate spectral type.  \\end{enumerate}"
    },
    "0004/astro-ph0004216_arXiv.txt": {
        "abstract": "A local luminosity function at 15$\\mu m$ is derived using the bivariate (15$\\mu m$ vs. 60$\\mu m$ luminosity) method, based on the newly published ISOCAM LW3-band (15$\\mu m$) survey of the very deep IRAS 60$\\mu m$ sample in the north ecliptic pole region (NEPR). New IRAS 60$\\mu m$ fluxes are obtained using the SCANPI/SUPERSCANPI software at the new ISOCAM positions of the sources in the sample. It is found to be in excellent agreement with the 15$\\mu m$ local luminosity function published by Xu et al (1998), which is derived from the {\\it predicted} 15$\\mu m$ luminosities of a sample of IRAS 25$\\mu m$ selected galaxies. Model predictions of number counts and redshift distributions based on the local luminosity function and assumptions of its evolution with the redshift are calculated and compared with the data of ISOCAM 15$\\mu m$ surveys. Strong luminosity evolution on the order of $L\\propto (1+z)^{4.5}$ is suggested in these comparisons, while pure density evolution can be ruled out with high confidence. The sharp peak at about 0.4mJy in the Euclidean normalized differential counts at 15$\\mu m$ can be explained by the effects of MIR broadband emission features, eliminating the need for any hypothesis for a 'new population'. It is found that the contribution from the population represented by ISOCAM 15$\\mu m$ sources can account for the entire IR/submm background, leaving little room for any missing 'new population' which can be significant energy sources of the IR/submm sky. ",
        "introduction": "Our understanding of galaxy formation and evolution in the early epochs of the Universe has been vastly improved in the past a few years, thanks mainly to new deep surveys in a wide range of wavebands, ranging from the HST's WFPC2 (UV and optical) and NICMOS (NIR) surveys in the Northern and Southern Hubble Deep Fields (Williams et al. 1996; Williams et al. 1998; Thompson et al. 1999) to the SCUBA submm surveys (see Blain et al. 1999b for a review). In particular, several mid-infrared (MIR) and far-infrared (FIR) deep surveys (see Elbaz et al. 1998b, Puget \\& Lagache 1998 for reviews) were conducted using the Infrared Space Observatory (ISO) (Kessler et al. 1996) during its 29 months mission (Nov. 1995 -- April. 1998).  The new results from these surveys (Aussel et al. 1999; Puget et al. 1999; Dole et al. 1999) indicate significant improvements in sensitivity and accuracy over the earlier published results (Rowan-Robinson et al. 1997; Kawara et al. 1998). Strong cosmic evolution in the population of infrared-emitting galaxies is indicated in these results (Rowan-Robinson et al. 1997; Kawara et al. 1998; Aussel et al. 1999; Puget et al. 1999), consistent with the results of SCUBA surveys (Blain et al. 1999a) and with the scenario hinted at by the newly discovered cosmic IR background (CIB) (Puget et al. 1996; Hauser et al. 1998; Dwek et al. 1998; Fixsen et al. 1998). These results challenge those from the UV/optical surveys (Madau et al. 1998; Pozzetti et al. 1998) in the sense that substantially more (i.e. a factor of 3 -- 5) star formation in the earlier Universe with $z\\gsim 2$ may be hinted at in the IR/submm counts and in the CIB (see, e.g. Rowan-Robinson et al. 1997) compared to that derived from the UV/optical surveys (Madau et al. 1998; Pozzetti et al. 1998). The reason of this discrepancy is attributed to dust extinction which may hide much of the star formation in the early Universe from the UV/optical surveys (see Lonsdale 2000 for a review).  The best observed band in these ISO surveys is the ISOCAM LW3 ($15\\mu$m) band, with 14 surveys covering a wide range of flux density from 0.05 mJy to 50 mJy (Elbaz et al. 1998b).  Compared to the longer wavelength ISOPHOT surveys (e.g. 175$\\mu m$ FIRBACK survey, Puget et al. 1999), ISOCAM LW3 surveys have the advantage of using large detector arrays ($32\\times 32$ compared to the $2\\times 2$ array of the ISOPHOT-C200 camera) and having much better angular resolution which allowed the surveys to go very deep ($\\sim 0.1$ mJy compared to sensitivity limits of $175\\mu m$ surveys of $\\sim 100$ mJy) before reaching the confusion limit.  In addition to the indication of the significant cosmic evolution of IR galaxies, two other very interesting results are emerging from these MIR surveys: \\begin{enumerate} \\item The Euclidean normalized differential number counts of the 15$\\mu m$ band have a sharp peak at about 0.4 mJy (Elbaz et al.  1999) which, Elbaz et al. (1998a) claim, can only be explained by adding to the number count model a new population of objects which emerge (with increasing z) rapidly after $z\\sim 0.4$ and start to dominate the counts below 1 mJy, but contribute negligibly in brighter flux range (Fig.10 of Elbaz et al. 1998a). \\item Subject to the substantial uncertainties in their IR SEDs, the integrated light of the galaxies detected in 15$\\mu m$ surveys may account for most the IR/submm background (Elbaz et al. 1999). \\end{enumerate} Taken at face values these results may have far-reaching impact on the studies of galaxy evolution, suggesting that the objects most responsible for the CIB, which are mostly missed by the UV/optical surveys, are already identified in the 15$\\mu m$ surveys and, perhaps more interestingly, they are from a new population not seen in the local Universe (i.e. not among the IRAS sources which in general have $z\\lsim 0.1$).  However, scrutiny of these interpretations of the 15$\\mu m$ counts is imperative because of the following two complications: (1) the effects of prominent emission features in the wavelength range of 3 -- 13$\\mu m$ (Puget and Leger 1989) which can cause very significant K-corrections in MIR surveys (Xu et al. 1998); and (2) the lack of a local luminosity function (LLF) in the 15$\\mu m$ band which is needed for the quantitative determination of the evolution rate from the number counts (to date, the IRAS 12$\\mu m$ luminosity functions have been mostly used in the interpretation of the ISOCAM 15$\\mu m$ counts, resulting in large uncertainties due to the significant variations of the $f_{15\\mu m}/f_{12\\mu m}$ among galaxies; see Elbaz et al. 1999). Indeed, the source count model of Xu et al. (1998), which takes into account the effect of the MIR emission features, did predict the bumps and dips in the counts similar to what is seen in the 15$\\mu m$ surveys. Accordingly, Xu et al. (1998) made the warning that determinations of the evolution rate based on the slope of source counts will have to treat the effects of the MIR emission features carefully. Xu et al. (1998) also derived a LLF at 15$\\mu m$ from the {\\it predicted} 15$\\mu m$ flux densities of a sample of 1406 IRAS 25$\\mu m$ selected galaxies based on a three component (cirrus/PDR, starburst, AGN) SED model.  Recently, Aussel et al. (2000) published an ISOCAM 15$\\mu m$ survey of the very deep IRAS 60$\\mu m$ sample in the northern ecliptic pole region (NEPR) (Hacking and Houck 1987, hereafter HH87; Hacking 1987). In this paper, we will derive a new 15$\\mu m$ LLF based on the data of this survey, and on the new IRAS SCANPI measurements at 60$\\mu m$ of the same sources using the new ISOCAM positions. This is then compared to the LLF of Xu et al. (1998). With the confidence in the 15$\\mu m$ LLF gained from this new study we explore further, using the model of Xu et al. (1998), the quantitative interpretation of results of the ISOCAM deep surveys as published by Elbaz et al. (1999) and by Aussel et al. (1999). ",
        "conclusions": ""
    },
    "0004/astro-ph0004284_arXiv.txt": {
        "abstract": "Direct imaging of the $\\zeta$ Cnc system has resolved the fourth star in the system, which is in orbit around $\\zeta$ Cnc C. The presence of the fourth star has been inferred for many years from irregularities in the motion of star C, and recently from C's spectroscopic orbit.  However, its mass is close to that of C, making its non-detection puzzling.  Observing at wavelengths of 1.2, 1.7, and 2.2 $\\mu$ with the adaptive-optics system of the CFHT, we have obtained images which very clearly reveal star D and show it to have the color of an M2 star. Its brightness is consonant with its being {\\it two\\/} M stars, which are not resolved in our observations but are likely to be in a short-period orbit, thereby accounting for the large mass and the difficulty of detection at optical wavelengths, where the magnitude difference is much larger. The positions and colors of all four stars in the system are reported and are consistent with the most recent astrometric observations. ",
        "introduction": "$\\zeta$ Cancri is a visual triple stellar system which has been studied as such for over 200 years.  The system consists of three bright stars, all dwarfs of close to solar type with visible magnitudes near 6. The orbital planes of the system are not far from the plane of the sky.  There is a close pair AB separated by about 1 arcsec with an orbital period 58.9 years, and a third star C about 6 arcsec away with an orbital period estimated at 1100 years. Star C itself has proper-motion irregularities that were first remarked upon by Sir John Herschel (1831) and were interpreted by Otto Struve (1874) in terms of orbital motion in a period on the order of 20 years around an unseen companion.  Subsequent observations and discussions have corroborated Struve's interpretation, and have recently been reinforced by a spectroscopic orbit of C by Griffin (2000), who has also described in some detail the interesting observational history of the system.  The situation concerning the three visual orbits in which the $\\zeta$ Cnc stars are involved (those of AB, CD, and AB--CD) have been summarized most recently by Heintz (1996).  All three have appreciable eccentricity, with that of CD being lowest at 0.12.  The principal question that has remained outstanding concerning the $\\zeta$ Cnc system is the nature of the companion to star C, since its orbit indicates a companion of comparable mass and a separation of about 0.4 arcsec.  Griffin (2000) recounts some of the earlier claims regarding the nature of star D, and it is clear that it must be considerably fainter than C at visible wavelengths. Suggestions have consequently been made that it is a white dwarf, or else a close pair of M stars which would be quite faint in the visible.  Griffin concluded that the question remained open, but that a short exposure by a large telescope with adaptive optics might provide the definitive evidence. Accordingly, the system was observed during a run with the Canada France Hawaii telescope by one of us (JBH) in January 2000, and followed up by further data from the same system a month later by another of us (FM). As we describe below, we discovered star D immediately and easily. ",
        "conclusions": ""
    },
    "0004/hep-ph0004191_arXiv.txt": {
        "abstract": "\\tightenlines  We study neutrino conversions in a  recently envisaged source of high energy neutrinos ($E {\\buildrel > \\over {_{\\sim}}} \\, 10^{6}$ GeV), that is, in the vicinity of cosmological Gamma-Ray Burst fireballs (GRB). We consider the effects of flavor and spin-flavor conversions and point out that in both situations, a some what higher than estimated high energy tau neutrino flux from GRBs is expected in new km$^{2}$ surface area under water/ice neutrino telescopes. ",
        "introduction": "Recently, cosmological fireballs are suggested as a possible production site for gamma-ray bursts as well as the high energy ($E {\\buildrel > \\over {_{\\sim}}} \\, 10^{6}$ GeV)  neutrinos \\cite{WB}. Although, the origin of these cosmological Gamma-Ray Burst fireballs (GRB) is not yet understood, the observations suggest that generically a very compact source of  linear scale $\\sim \\, 10^{7}$ cm  through internal or/and external shock propagation produces these gamma-ray bursts (as well as the burst of high energy neutrinos) \\cite{P}. Typically, this compact source is hypothesized to be formed possibly due to the merging of binary neutron stars or due to collapse of a supermassive star.  The main source of high energy tau neutrinos in GRBs is the production and decay of $D^{\\pm}_{S}$. The production of $D^{\\pm}_{S}$ may be through $p\\gamma $ and/or  through $pp$ collisions. In \\cite{PX}, the $\\nu_{e}$ and $\\nu_{\\mu}$  flux is estimated in $pp$ collisions, whereas in \\cite{WB}, the $\\nu_{e}$ and $\\nu_{\\mu}$ flux is estimated in $p\\gamma$ collisions for GRBs. In $pp$ collisions, the flux of $\\nu_{\\tau}$ may be obtained through the main process of $p+p\\rightarrow D^{+}_{S}+X$. The $D^{+}_{S}$ decays into $\\tau^{+}$ lepton and $\\nu_{\\tau}$ with a branching ratio of $\\sim \\, 3\\%$. This $\\tau^{+}$ lepton further decays into $\\nu_{\\tau}$. The cross-section for  $D^{+}_{S}$ production, which is main source of $\\nu_{\\tau}$'s, is $\\sim $ 4 orders of magnitude lower than that of $\\pi^{+}$ and/or $\\pi^{-}$ which subsequently produces $\\nu_{e}$ and $\\nu_{\\mu}$. The branching ratio for $\\nu_{e}$ and/or $\\nu_{\\mu}$ production is higher up to an order of magnitude than that for $\\nu_{\\tau}$ production (through $D^{\\pm}_{S}$). These imply that the $\\nu_{\\tau}$ flux in $pp$ collisions is suppressed up to 5 orders of magnitude relative to corresponding $\\nu_{e}$ and/or $\\nu_{\\mu}$ fluxes. In $p\\gamma$ collisions, the main process for the production of $\\nu_{\\tau}$ may be $p+\\gamma \\rightarrow D^{+}_{S}+ \\Lambda^{0}+\\bar{D}^{0}$ with similar relevant branching ratios and corresponding suppression for cross-section  values. Here the corresponding main source for $\\nu_{e}$ and $\\nu_{\\mu}$ production is $p+\\gamma \\, \\rightarrow \\Delta^{+}\\rightarrow \\pi^{+}+n$. Therefore, in $p\\gamma $ collisions, the $\\nu_{\\tau}$ flux is also suppressed up to 5 orders of magnitude relative to $\\nu_{e}$ and/or $\\nu_{\\mu}$ flux. Thus, in both type of collisions, including the relevant kinematics, the intrinsic $\\nu_{\\tau}(\\bar{\\nu}_{\\tau})$ flux is estimated to be rather small relative to $\\nu_{e} (\\bar{\\nu}_{e})$ and/or $\\nu_{\\mu} (\\bar{\\nu}_{\\mu})$ fluxes from GRBs, typically being, $F^{0}_{\\tau} /F^{0}_{e,\\mu}\\, <10^{-5}$ \\cite{VZA}.  In this paper, we consider the possibility of obtaining higher $\\nu_{\\tau}$ flux, that is, $F^{0}_{\\tau}/F^{0}_{e,\\mu}\\, >\\, 10^{-5}$, from GRBs through neutrino conversions as compared to no conversion situation. The present study is particularly useful as the new under ice/water \\v{C}erenkov light neutrino telescopes namely AMANDA and Baikal (also the NESTOR and ANTARES) will have energy, angle and flavor resolution for high energy neutrinos originating at cosmological distances \\cite{M}. Recently, there are several discussions concerning the signatures of a possible neutrino burst from GRBs correlated in time and angle \\cite{W}. In particular, there is a suggestion of measuring $\\nu_{\\tau}$ flux from cosmologically distant  sources through a double shower (bang) event \\cite{L} or through a small pile up of up ward going $\\mu$-like  events near (10$^{4}-10^{5}$) GeV \\cite{HS}.   The plan of this paper is as follows. In Sect. II, we briefly describe  the matter density and magnetic field in the vicinity of GRBs. In Sect. III, we discuss in some detail, the range of neutrino mixing parameters that may give rise to relatively large precession/conversion probabilities resulting from neutrino flavor/spin-flavor conversions. In Sect. IV, we give  estimates for separable but contained  double shower event rates induced by high energy $\\nu_{\\tau}$'s originating from GRBs without/with conversions for km$^{2}$ surface area under water/ice neutrino telescopes for illustrative purposes and finally in Sect. V, we summarize our results. ",
        "conclusions": "1. Intrinsically, the flux of high energy cosmic tau neutrinos is quite small, relative to non tau neutrino flavor, typically being $F^{0}_{\\tau}/F^{0}_{e, \\mu}\\, <\\, 10^{-5}$ (whereas $F^{0}_{e}/F^{0}_{\\mu} \\sim 1/2$) from cosmologically distant GRBs.   2. Because of neutrino oscillations, this ratio can be greatly enhanced. In the context of three flavor neutrino mixing scheme which can accommodate the oscillation solutions to solar and atmospheric neutrino deficits in terms of oscillations between three active neutrinos, the final down ward going ratio of fluxes of high energy cosmic neutrinos on earth is $F_{e}\\sim F_{\\mu}\\sim F_{\\tau} \\sim F^{0}_{e}$, essentially irrespective of the oscillation solution to solar neutrino problem.  The (vacuum) flavor oscillations leads to an essentially energy independent flux of high energy neutrinos of all flavors originating from cosmologically distant GRBs at the level of electron neutrino flux, whereas spin-flavor precessions/conversions may lead to an energy dependence or/and change in this situation.  The spin-flavor conversions may occur possibly through several mechanisms. We have discussed in some detail mainly the spin-flavor precession/conversion situation induced by a non zero neutrino magnetic moment and by a relatively small VEP as an example to point out the possibility of obtaining some what higher tau neutrino fluxes as compared to no oscillations/conversions scenarios from GRBs.  The matter density in the vicinity of GRB is quite small (up to 4$-$5 orders of magnitude) to induce any resonant flavor/spin-flavor neutrino conversion due to normal matter effects. We have pointed out that a resonant character in the neutrino spin-flavor conversions may nevertheless be obtained due to possible VEP. The corresponding degree of VEP  may be $\\sim \\, (10^{-35}-10^{-25}$) depending on $\\delta m^{2}$ value for vanishing gravity mixing angle.  3. This enhancement in high energy cosmic tau neutrino flux may lead to the possibility of its detection in km$^{2}$ surface area high energy neutrino telescopes. For $2\\times 10^{6}\\leq E$/GeV $\\leq 2\\times 10^{7}$, the down ward going high energy cosmic tau neutrinos may produce a double shower signature because of charged current deep inelastic scattering followed by a subsequent hadronic decay of associated tau lepton   The double shower  event rate for intrinsic (no oscillations/conversions) high energy tau neutrinos originating from GRBs turns out to be small as compared to that due to precession/conversion effects up to a factor of $\\sim \\, 10^{-5}$. Thus, the high energy neutrino telescopes may possibly provide useful upper bounds on intrinsic properties of neutrinos such as mass, mixing and magnetic moment, etc.. The relevant tau neutrino energy range for detection in km$^{2}$ surface area under water/ice neutrino telescopes may be  $2\\times 10^{6} {\\buildrel < \\over {_{\\sim}}}\\, E/\\mbox {GeV}\\, {\\buildrel < \\over {_{\\sim}}} \\, 2\\times 10^{7}$ through characteristic contained but separable double shower  events.  Observationally, the high energy $\\nu_{\\tau}$ burst from a GRB may possibly be {\\em correlated} to the corresponding gamma-ray burst/highest energy cosmic rays (if both have common origin) in time and in direction thus raising the possibility of its detection. If the range of neutrino mixing parameters pointed out in this study is realized terrestially/extraterrestially then a relatively large (energy dependent) $\\nu_{\\tau}$ flux from GRBs is expected as compared to no oscillation/conversion scenario.   \\paragraph*{Acknowledgments.}  The author thanks Japan Society for the Promotion of Science for financial support."
    },
    "0004/astro-ph0004147_arXiv.txt": {
        "abstract": "Far-UV observations in and near the Hubble Deep Fields demonstrate that the Space Telescope Imaging Spectrograph (STIS) can potentially obtain unique and precise measurements of the diffuse far-ultraviolet background.  Although STIS is not the ideal instrument for such measurements, high-resolution images allow Galactic and extragalactic objects to be masked to very faint magnitudes, thus ensuring a measurement of the truly diffuse UV signal. The programs we have analyzed were not designed for this scientific purpose, but would be sufficient to obtain a very sensitive measurement if it were not for a weak but larger-than-expected signal from airglow in the STIS 1450--1900~\\AA\\ bandpass.  Our analysis shows that STIS far-UV crystal quartz observations taken near the limb during orbital day can detect a faint airglow signal, most likely from \\ion{N}{1} $\\lambda 1493$~\\AA, that is comparable to the dark rate and inseparable from the far-UV background.  Discarding all but the night data from these datasets gives a diffuse far-ultraviolet background measurement of $501 \\pm 103$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$, along a line of sight with very low Galactic neutral hydrogen column ($N_{HI} = 1.5\\times10^{20}$ cm$^{-2}$) and extinction ($E(B-V)=0.01$ mag).  This result is in good agreement with earlier measurements of the far-UV background, and should not include any significant contribution from airglow.  We present our findings as a warning to other groups who may use the STIS far-UV camera to observe faint extended targets, and to demonstrate how this measurement may be properly obtained with STIS. ",
        "introduction": "\\label{secintro}  Over the past several decades, the far-ultraviolet background (FUVBG) has been the subject of many studies and the center of considerable controversy (for a review and an example of such debates, see Bowyer 1991\\markcite{B91}; Henry 1991\\markcite{H91}).  The FUVBG is defined as the diffuse astronomical flux incident upon the Earth at wavelengths of 912--2000~\\AA, thus excluding sources such as terrestrial airglow and Galactic stars.  Because the zodiacal light decreases dramatically below 2500~\\AA, and because the diffuse UV background is so faint, O'Connell (1987\\markcite{O87}) has noted that this spectral region offers a window of very low sky background for observations of objects with faint surface brightness.  We have been utilizing this spectral window for UV observations of faint extended sources, but these data also provide insight into the FUVBG itself.  A number of FUVBG measurements have shown a correlation between the FUVBG intensity and the Galactic neutral hydrogen column density ($N_{HI}$), leading to the conclusion that a significant fraction of the FUVBG must be Galactic in origin (see, e.g., the relatively recent data set of Hurwitz, Bowyer, \\& Martin 1991\\markcite{HBM91}). In a 1400--1850~\\AA\\ bandpass, extrapolating to \\medskip  {\\small $^1$Based on observations with the NASA/ESA Hubble Space Telescope obtained at the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS~5-26555.  $^2$Laboratory for Astronomy \\& Solar Physics, Code 681, Goddard Space Flight Center, Greenbelt, MD 20771  $^3$STScI, 3700 San Martin Drive, Baltimore, MD 21218.  ferguson@stsci.edu  $^4$NOAO Research Associate. tbrown@pulsar.gsfc.nasa.gov  $^5$Raytheon ITSS Corp.  $^6$kimble@ccd.gsfc.nasa.gov  $^7$gardner@harmony.gsfc.nasa.gov  $^8$collins@zolo.gsfc.nasa.gov  $^9$bhill@virgil.gsfc.nasa.gov }  \\noindent $N_{HI}=0$ gives an FUVBG intensity of $\\sim$300 ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$; this intensity increases to $\\sim$1200 ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$ with increasing $N_{HI}$, but saturates before high column densities are reached (Hurwitz et al.\\ 1991\\markcite{HBM91}; Bowyer 1991\\markcite{B91}).  Curiously, measurements below Ly-$\\alpha$ may indicate an upper limit of 30--100 ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$ (Henry 1991\\markcite{H91}; Murthy et al.\\ 1999\\markcite{M99}), although this measurement has been contested (Edelstein, Bowyer, \\& Lampton 2000\\markcite{E00}). Herein lies the controversy: a step in the FUVBG at Ly-$\\alpha$ could imply an extragalactic origin for the $N_{HI}=0$ component, such as redshifted Ly-$\\alpha$ emission from an intergalactic medium (IGM) (Henry 1991\\markcite{H91}); demonstration of such an origin of the FUVBG would have extremely important consequences, as canonical values for IGM parameters predict a negligible contribution to the FUVBG (e.g., Jakobsen 1993\\markcite{J93}). On the other hand, others maintain that the dominant contribution to the $N_{HI}=0$ component is mainly dust-scattered Galactic light (Bowyer 1991\\markcite{B91}); under that assumption, the residual dust at $N_{HI}=0$ that is needed to produce the FUVBG corresponds to $E(B-V)=0.015$~mag (Hurwitz et al.\\ 1991\\markcite{HBM91}).  We note that we cannot directly discriminate between an extragalactic and Galactic origin of the FUVBG, as we are measuring the FUVBG longward of Ly-$\\alpha$, but our measurements have all been taken along lines of sight with extremely low $N_{HI}$ column {\\it and} low dust extinction (see Table \\tabobs).  New maps of Galactic extinction (Schlegel, Finkbeiner, \\& Davis 1998\\markcite{SFD98}) are based on a direct measurement of the IR emission from the dust, instead of inferring the extinction from dust-to-hydrogen ratios; these new dust maps can potentially yield new insight into the FUVBG debate.  The Hubble Space Telescope (HST) currently employs one instrument with significant sensitivity in the far-UV: the Space Telescope Imaging Spectrograph (STIS) (Woodgate et al.\\ 1998\\markcite{W98}).  However, STIS is not an ideal instrument for measuring the FUVBG: the throughput in the crystal quartz bandpass  \\newpage  \\parbox{5.5in}{ {\\sc Table \\tabobs:} Far-UV Data  \\begin{tabular}{clccccrcc} \\tableline Field    & Exposure & RA             & Dec                            & $l$    & $b$     & coverage        & $N_{HI}$$^a$ & $E(B-V)^b$ \\\\ ID       & (sec)    & (J2000)        & (J2000)                        & (deg)  & (deg)   & (square arcsec) &  cm$^{-2}$   & (mag) \\\\ \\tableline HDFN-Fol & 124330   & $12^h36^m44^s$ & $\\rm +62^o12\\arcmin 12\\arcsec$ & 125.908 & +54.840 & 3659 & 1.51$\\times 10^{20}$ & 0.01 \\\\ HDFN-Par &  17315   & $12^h35^m41^s$ & $\\rm +62^o10\\arcmin 38\\arcsec$ & 126.124 & +54.853 & 1003 & 1.51$\\times 10^{20}$ & 0.01 \\\\ HDFS-Pri &  52124   & $22^h33^m38^s$ & $\\rm -60^o33\\arcmin 29\\arcsec$ & 328.158 & -49.274 &  878 & 2.39$\\times 10^{20}$ & 0.03 \\\\ \\tableline \\end{tabular}  $^a$ Dickey \\& Lockman (1990)\\markcite{DL90}.\\\\ $^b$ Schlegel et al.\\ (1998)\\markcite{SFD98}. }  \\bigskip  \\noindent is not very high (peaking at 573 cm$^2$ effective area), the pixels and field of view are small (respectively $0.025 \\arcsec$ and $25\\arcsec$), and the dark background is spatially and temporally variable (but very low). Nonetheless, we have found that the deepest imaging programs can yield interesting measurements of the diffuse UV background, because deep, high-resolution, coincident imaging at longer wavelengths can provide a mask for all objects brighter than $\\sim$ 29 mag in the optical. Because we can mask very faint galaxies, we can determine how much of the FUVBG is truly diffuse.  In this paper, we demonstrate the measurements of the FUVBG that could be made using the STIS far-UV camera.  The programs we analyze here, from observations in and near the Hubble Deep Fields (HDFs), were not intended for measurements of the FUVBG, but they indicate that an HST program could be tailored to characterize the FUVBG with excellent statistics. ",
        "conclusions": "\\label{secdis}  Restricting our data to that obtained in orbital night, we measured a diffuse far-UV background of $501 \\pm 103$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$.  Our $5 \\sigma$ detection of the FUVBG, a signal that is $\\sim$10\\% of the STIS dark rate, is in reasonable agreement with earlier measurements at low $N_{HI}$.  However, the unique aspect to STIS measurements of the FUVBG is that STIS can measure a truly diffuse signal, because high-resolution coincident imaging in the far-UV and optical bandpasses allows masking of all objects down to a very faint magnitude.  The night data used for our measurement of the FUVBG were taken in two fields at high Galactic latitude, separated by about 8.5 arcmin.  The measurement in the HDFN-Par field was $296 \\pm 135$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$, and the measurement in the HDFN-Fol field was $712 \\pm 157$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$.  The HDFN-Fol field was observed much closer to the Earth's limb than the HDFN-Par field, but the night exposures should show no measurable contribution from airglow.  Barring systematic errors, the difference between the two fields may reflect a true patchiness in the FUVBG, whether it originates in scattering from Galactic dust or an extragalactic source.  Note that the Galactic $E(B-V)$ along each line of sight is very low (0.01~mag).  \\bigskip  \\hskip -0.2in \\parbox{3.0in}{\\epsfxsize=3.5in \\epsfbox{TMBrown.fig4.eps}} \\vskip 0.1in  \\centerline{\\parbox{3.5in}{\\small {\\sc Fig.~\\figlimb--} Our measurements of the FUVBG, as a function of the average zenith angle during the exposure, using only those exposures taken in complete orbital night.  Individual measurements from the HDFN-Fol data ({\\it open diamonds}) and the HDFN-Par data ({\\it open triangles}) are shown with a representative sample of their associated uncertainties.  The filled symbols show the sum of the measurements near 39$\\rm ^o$, 73$\\rm ^o$, and 90$\\rm ^o$, with their associated uncertainties; there is a weak trend for increasing signal at high zenith angle (near the limb, which is at 114$\\rm ^o$), but the known night airglow in this bandpass should be completely negligible at the HST altitude of 600 km.  The sum of the data taken at a zenith angle less than 80$\\rm ^o$ is $296 \\pm 135$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$, and the sum of the data taken at a zenith angle greater than 80$\\rm ^o$ is $712 \\pm 157$ ph cm$^{-2}$ sec$^{-1}$ ster$^{-1}$ \\AA$^{-1}$. }} \\addtocounter{figure}{1}  \\medskip  To measure the FUVBG with high statistical certainty, future STIS observations should be obtained in TIMETAG mode, to exclude day data, and to ensure that the night data shows no correlation with zenith angle.  The target should be a reasonably empty portion of a field with deep optical imaging, such as one of the fields described herein. If the STIS cooling system to be installed in HST Servicing Mission 3B is as effective as expected in eliminating dark glow, a much larger fraction of the detector area will be available for sensitive measurements.  Off-nominal positioning of the entrance aperture mask could also increase the shadowed detector area available for concurrent monitoring of the dark rate, though HST operating constraints may not permit this.  Prior to our analysis, the crystal quartz filter on the STIS far-UV camera was thought to produce a bandpass that was completely insensitive to both zodiacal emission and airglow.  Because the dark rate in this camera is so low, the low sky enables photometric measurements of very faint extended UV sources (see, e.g., Brown et al.\\ 2000\\markcite{B00}).  However, our analysis showed that while the sky signal is always lower than the dark rate, it may be nonnegligible.  The dark glow can increase the background signal by a factor of 20 in sections of the detector, but the sky can reach levels comparable to the minimum dark rate; the sky can thus increase the background signal by a factor of two in the darkest parts of the detector during near-limb day observations.  Potential STIS observers should take this sky signal into account when planning their S/N estimates, if they will be observing near the limb (e.g., in the HST Continuous View Zone)."
    },
    "0004/astro-ph0004237_arXiv.txt": {
        "abstract": "The simple physics of microlensing provides a well-understood tool with which to probe the atmospheres of distant stars in the Galaxy and Local Group with high magnification and resolution.  Recent results in measuring stellar surface structure through broad band photometry and spectroscopy of high amplification microlensing events are reviewed, with emphasis on the dramatic expectations for future contributions of microlensing to the field of stellar atmospheres. ",
        "introduction": "The physics of microlensing is simple.  For most current applications, the principles of geometric optics combined with one relation (for the deflection angle) from General Relativity is all that is required.  For observed Galactic microlensing events, the distances between source, lens and observer are large compared to intralens distances, so that small angle approximations are valid.  Although it is possible that most lenses may be multiple, $\\sim$90\\% of observed Galactic microlensing light curves can be modeled as being due to a single point lens. Usually, though not always (cf., Albrow et al.~2000), binary lenses can be considered static throughout the duration of the event.  The magnification gradient near caustics is large, producing a sharply peaked lensing ``beam'' that sweeps across the source due to the relative motion between the lens and the sight line to the source (Fig.~1). Furthermore, the combined magnification of the multiple microimages (which are too close to be resolved with current techniques) is a known function of source position that is always greater than unity, so that more flux is received from the source during the lensing event. The net result is a well-understood astrophysical tool that can simultaneously deliver high resolution and high magnification of tiny background sources. In Galactic microlensing, these sources are stars at distances of a few to a few tens of kiloparsecs.  The great potential of microlensing for the study of stellar polarization (Simmons, Willis, \\& Newsam 1995; Simmons, Newsam \\& Willis 1995; Newsam et al.~1998; Gray 2000), stellar spots (Heyrovsk\\'y \\& Sasselov 2000; Bryce \\& Hendry 2000), and motion in circumstellar envelopes (Ignace \\& Hendry 1999) will not be treated here. Instead, the focus will be on how the composition of spherically-symmetric stellar atmospheres can be probed by microlensing. ",
        "conclusions": ""
    },
    "0004/astro-ph0004001_arXiv.txt": {
        "abstract": "The close-in extrasolar giant planets [CEGPs], $\\ltorder$ 0.05 AU from their parent stars, may have a large component of optically reflected light. We present theoretical optical photometric light curves and polarization curves for the CEGP systems, from reflected planetary light. Different particle sizes of three condensates are considered. In the most reflective case, the variability is $\\approx 100$ micromagnitudes, which will be easily detectable by the upcoming satellite missions MOST, COROT, and MONS, and possibly from the ground in the near future. The least reflective case is caused by small, highly absorbing grains such as solid Fe, with variation of much less than one micromagnitude. Polarization for all cases is lower than current detectability limits. We also discuss the temperature-pressure profiles and resulting emergent spectra of the CEGP atmospheres. We discuss the observational results of $\\tau$~Boo b by Cameron et al. (1999) and Charbonneau et al. (1999) in context of our model results. The predictions --- the shape and magnitude of the light curves and polarization curves --- are highly dependent on the size and type of condensates present in the planetary atmosphere. ",
        "introduction": "The discovery of the planet 51 Peg b in 1995 (Mayor \\& Queloz 1995), only 0.051 AU from its parent star, heralded an unexpected new class of planets. Due to gravitational selection effects, several more Jupiter-mass close-in extrasolar giant planets (CEGPs) have been discovered since that time (Butler et al. 1997, 1998; Mayor et al. 1999; Mazeh et al. 2000). To date there are 5 extrasolar giant planets $\\ltorder$ 0.05 AU from their parent stars, and an additional 9 $\\ltorder$ 0.23 AU (see Schneider 2000). Relevant data about the close-in planet-star systems (orbital distance $\\ltorder$ 0.05 AU) are listed in Table 1. Ongoing radial velocity searches will certainly uncover more CEGPs in the near future. The CEGPs are being bombarded by radiation from their parent stars, and could be very bright in the optical. At best the CEGPs could be four to five orders of magnitude fainter than their primary star, compared to Jupiter which is 10 orders of magnitude fainter than the Sun.  The recent transit detection of HD~209458~b by Charbonneau et al. (2000) and Henry et al. (2000b) confirms that the CEGPs are gas giants, gives the planet radius, and fixes the orbital inclination, which removes the $\\sin i$ ambiguity in mass and provides the average planet density. HD~209458 has $R_*=1.2 \\pm 0.1 R_{\\odot}$ and $R_P=1.40 \\pm 0.17 R_J$ (Mazeh et al. 2000), where $R_*$ is the stellar radius and $R_P$ is the planet radius. Transits are definitely ruled out for $\\tau$ Boo b, 51 Peg b, $\\upsilon$ And~b, HD~187123, and $\\rho^1$ Cnc b, whether they are assumed to be gas giants with radius 1.2 $R_J$, or smaller rocky planets with radius $\\sim$0.4 $R_J$ (Henry et al. 2000a, 1997; Baliunas et al. 1997; G. Henry, private communication). Transits are also ruled out for HD~75289 (M. Mayor, private communication). For a transit to be observable, a CEGP must be aligned with the star as seen from Earth with an inclination $i > \\theta_T$, where $\\theta_T={\\cos}^{-1}((R_* + R_P)/D)$. For random orientations, the probability for $i$ to be between $90^{\\circ}$~and~$j^{\\circ}$ is $P(j) = {\\cos}(j)$. With $\\theta_T \\sim 83^{\\circ}$, the CEGPs have transit probabilities of ten percent. By the same $\\theta_T$ criterion, the non-detection of transits puts limits on the orbital inclinations to approximately 83$^{\\circ}$ (for $R_*=1.16$~$R_{\\odot}$, $R_P=1.2$~$R_J$, and $D=0.051$~AU). Several groups (e.g. STARE (PI T. Brown), Vulcan Camera Project (PI W. Borucki), WASP (PI S. Howell)) are monitoring thousands of stars without known planets, searching with high precision photometry for periodic fluctuations indicative of a planetary transit. Follow-up observations by radial velocity techniques (or astrometry in the future) will be needed to fix the orbital radius in order to determine the planet mass. Edge-on CEGP systems are the most promising for reflected light signals.  Several observational approaches to detecting and characterizing CEGP atmospheres have been developed. These include spectral separation, transmission spectra observations during transit, infrared observations, and optical photometric light curve observations.  Charbonneau et al. (1999) and Cameron et al. (1999) have developed a direct detection technique: a spectral separation technique to search for the reflected spectrum in the combined star-planet light. Both groups have observed the $\\tau$ Boo system. $\\tau$ Boo A is one of the brightest (fourth magnitude), hottest (F7V) parent stars, and $\\tau$~Boo~b has one of the smallest semi-major axes; these three properties make $\\tau$ Boo a promising candidate for this technique. From a non-detection, Charbonneau et al. (1999) have put upper limits on the planet-star flux ratio ranging from $5 \\times 10^{-5}$ for $\\sin i \\sim 1$ to $1 \\times 10^{-4}$ for $\\sin i \\sim 0.5$. The result is within the strict assumptions that the light curve is fairly isotropic and that the reflected spectrum is an exact copy of the stellar spectrum from 4668 to 4987 \\AA. Their upper limit on the geometric albedo is 0.3 for $\\sin i \\sim 1$. The same technique for the $\\tau$ Boo system has been used by Cameron et al. (1999) who claim a possible detection at an inclination of 29$^{\\circ}$, and give a planet-star flux ratio of $1.9 \\times10^{-4}$ at $i=90^{\\circ}$. Given $R_P$, the albedo derived from this type of observation can provide a weak constraint on theoretical models.  A second approach is to observe transmission spectra during a planet transit. The stellar flux will pass through the optically thin part of the planet atmosphere. Theoretical predictions show the planetary absorption features will be at the $10^{-4}$ to $10^{-3}$ level (Seager \\& Sasselov 2000). Successful observations will constrain the cloud depth and may give important spectral diagnostics such as the presence of CH$_4$ which is a good temperature indicator for the upper atmosphere layers.  A third technique under development is the use of the Keck infrared interferometer in the differential phase mode to directly detect and spectroscopically characterize the CEGPs. The technique is based on the difference between the very smooth infrared stellar spectrum and the strong water absorption bands and possibly methane bands in the CEGP's infrared spectrum. See Akeson \\& Swain (1999) for more details.  In this paper we present theoretical photometric light curves and polarization curves of the CEGP systems. As the planet orbits the star, the planet changes phase as seen from Earth. The planet and star are too close together for their light to be separated, but this small separation means the stellar flux hitting the planet is large, and the reflected light variation in the combined light of the system from the planet's different phases may be detectable. We focus on the optical where there is a clear signature of reflected light: the planet's dark side has no reflection or emission in optical light. In contrast, there is no large light variation in the infrared where the CEGPs are bright on both the day and night side from reemission of absorbed light. Scattered light is minimal in the infrared, and is difficult to disentangle from the emitted light. Unlike transits, which can only be seen for inclinations $ > \\theta_T$, the reflected light curves of lower inclinations are theoretically visible and may be detectable. This work is motivated by upcoming microsatellites MOST ($\\sim$2002) (Matthews 1997), COROT ($\\sim$2003) (Baglin 1998, 2000), and MONS ($\\sim$2003) (Christensen-Dalsgaard 2000). Initially intended for asteroseismology, these satellites have capabilities to detect $\\mu$mag variability. MOST will observe known stars in a broad visual waveband, one at a time for a period of roughly one month, including one with a known CEGP in its first year. COROT's exoplanet approach will use two CCDs in two colors to observe several fields of $\\sim$6000 stars for a few months each. Because of this wide-field approach the stars with known CEGPs will not be observed by COROT. While COROT's main exoplanet focus is on transits, probability estimates suggest several CEGP light curves from reflection should be detected. Precision of ground-based photometry on the CEGP parent stars is currently at 100~$\\mu$mag, and could reach 50~$\\mu$mag in the near future with dedicated automatic photometric telescopes (Henry et al. 2000a). We also present polarization signatures although they are well under the current limits of detectability which is a few~$\\times 10^{-4}$ in fractional polarization of the system (e.g. Huovelin et al. 1989).  This paper, to our knowledge, is the first to describe photometric light curves and polarization of CEGP systems: gas giants in close orbits around Sun-like stars. Although our own Solar System planets have been well studied in reflected and polarized light, the CEGPs have effective temperatures an order of magnitude higher, so completely different cloud species and atmospheric parameters are expected. If observable, the light curves would roughly constrain the type and size distribution of condensates in the planetary atmosphere. In \\S2 we present definitions and analytical estimates of reflected light from the CEGPs, in \\S3 a description of our model, and in \\S4 results and discussion.  \\section {Analytical Estimate of the Light Curves and Polarization} \\label{sec-analytical} An analytical estimate of the amount of reflected light and polarization of an EGP system is useful for both comparison with simulations and for upper limit predictions. A good idealized case for such estimates is provided by modeling a planet as, for example, a Lambert sphere. The Lambert sphere derives from the law of diffuse reflection proposed by Lambert, postulating a reflecting surface with a reflection coefficient that is constant for all angles of incidence. The reflection coefficient is simply the ratio of the amount of light diffusely reflected in all directions by an element of the surface to the incident amount of light which falls on this element. The general conditions postulated by Lambert, for example angle independence, are satisfied strictly only for an absolute blackbody and an ideal reflecting surface (often called ``absolutely white'', ``ideally matted'', etc.). Thus derives Lambert's definition of albedo, with its inherent ambiguities as discussed at the end of the 19$th$ century by Seeliger, and the ultimate decision by Russell (1916) to endorse Bond's (1861) definition of albedo for use in the Solar System.  The arguments offered by Russell (1916) in favor of the Bond albedo (over other albedo definitions in use at the time) are still relevant for Solar System objects, but not necessarily for EGPs. One important point in favor of the Bond albedo was that it is derivable from observations. The Bond albedo is defined as the ratio of the total amount of reflected light to the total amount of\tincident plane-parallel light integrated over all angles. Note that at the time of Russell --- before multi-wavelength observations of the Solar System planets --- the Bond albedo, $A$, was not defined as an integrated quantity over all wavelengths as it is today. However, the discussion below is still valid with either definition. The Bond albedo $A$ can be separated into two quantities, \\begin{equation} A = pq, \\end{equation} where $p$ is the geometric albedo and $q$ is the phase integral. The geometric albedo is defined as the planet flux divided by the reflected flux from a perfectly diffusing disk of the same radius. The phase integral $q$ is defined as \\begin{equation} q = \\int_0^{\\pi} \\phi(\\alpha) \\sin\\alpha d\\alpha, \\end{equation} where $\\phi(\\alpha)$ is the phase function, or the brightness variation of the planet at different phases. The phase angle, $\\alpha$, is the angle between the star and Earth as seen from the planet; $\\alpha=0$ corresponds to ``opposition'' when the planet is maximally illuminated as seen from Earth. $p$ is measureable for all Solar System planets because it is a geometric and photometric quantity. $q$ is measureable from Earth for Mercury, Venus, Mars, and the Moon; for the outer planets whose phase angle variation is only up to several degrees from Earth satellite mission observations were necessary. Thus the benefit of the Bond albedo: it is a physically meaningful quantity but it can be determined empirically for the Solar System planets.  In contrast, the Bond albedo cannot be determined from observations for extrasolar planets. Because the EGP systems are so distant from Earth, only the CEGPs have prospects for measurement of $p$ in the foreseeable future, and even then only the most reflective CEGPs will be bright enough, and only $\\sim$10\\% of those will have orbital inclinations near $90^{\\circ}$. Charbonneau et al. (1999) and Cameron et al. (1999) have developed a spectral separation detection technique which can put upper limits on --- and in the best case measure --- $p$ in a narrow wavelength region. For CEGPs that are reflective, and at $i\\sim90^{\\circ}$, $q$ should be measureable with the upcoming satellite missions. However if a given CEGP system is at $i<90^{\\circ}$, the full range of phases will not be visible (i.e. $\\alpha$ will not be fully probed), and $p$ and $q$ will not be measureable. As discussed in \\S\\ref{sec-EGPbeyond}, EGPs beyond $D=0.1$~AU will not be detectable in optically reflected light even with the upcoming satellite missions. The EGPs certainly have promise for detection in the infrared where they emit most of their energy. However, most of this energy is reprocessed absorbed energy; it is not possible to measure the Bond albedo with infrared observations. To summarize, the Bond albedo came into standard usage because it was a measureable quantity. This is not possible for almost all of the EGPs because of the distance of the systems and random orbital inclinations.  The goal of this paper is a presentation of light curves at all viewing angles and at different inclinations, instead of a Bond albedo. We begin with the analytical estimate, where in the idealized case we are simply interested in the ratio, $\\epsilon$, of the observed flux at Earth from the EGP at full phase ($\\alpha=0$) to that of the star: $\\epsilon = p(R_P/D)^2$. Here $D$ is the star-planet distance, and $p$ and $R_P$ are as previously defined. For a Lambert sphere the single scattering albedo $\\tilde{\\omega}=1$; no photons are absorbed, and so $A=1$. For a Lambert sphere, all incoming photons are singly, isotropically scattered, and $p = 2/3$. The light variation of the Lambert sphere is only due to phase effects $-$ the phase function, $\\phi$, is simply (Russell 1916): \\begin{equation} \\phi(\\alpha) = \\frac{\\sin(\\alpha) + (\\pi - \\alpha)\\cos(\\alpha)}{\\pi}, \\end{equation} and the phase-dependent flux ratio is ${\\epsilon}\\phi(\\alpha)$. Note that the phase angle, $\\alpha$, is a function of the orbital phase and inclination. We convert this flux ratio to variation in micromagnitudes by \\begin{equation} \\label{eq:dm} \\Delta m = -2.5 {\\rm log}_{10}(1 + \\epsilon \\phi(\\alpha)) 10^6. \\end{equation}  Figure~\\ref{fig:lambert}a shows the photometric light variation at $i=90^{\\circ}$ for a Lambert sphere of $R=1.2 R_J$ at various $D$ corresponding to known CEGP systems. The large variation of 110 and 140 $\\mu$mag for planets with $D$ corresponding to $\\tau$ Boo b and HD 187123 respectively is above current ground-based limits (Henry et al. 2000a). The Lambert sphere light curve is unrealistic, and the point of this paper is to show that the situation is far more complex and almost always conducive to a smaller light variation for a few reasons. Firstly, the single scattering albedo, $\\tilde{\\omega}$, is generally different from one. When optical photons are absorbed by condensates or gas in the CEGP atmospheres, they are re-emitted in the infrared or contribute to the thermal pool. Secondly, multiple scattering gives more of a chance for absorption over single scattering for the same single scattering albedo. Each time a photon scatters, its next encounter has a scattering probability of $\\tilde{\\omega}$, but when a photon is absorbed it cannot contribute to the scattered light. This effect depends on density, i.e. on the mean free path of the photon. Thirdly, when particles are large compared to the wavelength of light, the particle scatters preferentially in the forward direction. In this case, photons that enter the atmosphere are likely to be multiply scattered down into the atmosphere and eventually absorbed, rather than to be backscattered and escape the planetary atmosphere.  Figure~\\ref{fig:lambert}b shows the fractional polarization of the total light of a Lambert sphere at $i=90^{\\circ}$ at various $D$ corresponding to known CEGP systems, and assuming $R_P=1.2R_J$. We assume Rayleigh scattering linear polarization of unpolarized incident light $Pol_{Ray}= {{\\sin^2} \\theta_S}/({1 + {\\cos^2} \\theta_S})$ (Chandrasekhar 1960). Here $\\theta_S$ is the scattering angle: $\\theta_S=0^{\\circ}$ is the forward direction and $\\theta_S=180^{\\circ}$ the backward direction. $Pol_{Ray}$ peaks at a scattering angle of $90^{\\circ}$. For single scattering and $\\tilde{\\omega} = 1$, all scattered light is polarized. The polarization signature plotted in Figure~\\ref{fig:lambert}b does not peak at $\\alpha=90^{\\circ}$ because the polarization is modulated by the reflected light curve; the scattered light is maximally polarized at $\\alpha=90^{\\circ}$, but the amount of scattered light peaks at $\\alpha=0^{\\circ}$. Plotted is $Pol = ({S_{\\bot} - S_{\\|}})/({S + F}) = \\epsilon \\phi(\\alpha) Pol_{Ray}$, where $S$ is the total scattered light, $F$ is the unpolarized stellar flux, $S \\ll F$, and $S_{\\bot}$ and $S_{\\|}$ are the perpendicular and parallel components of the scattered light respectively. In reality polarization is much lower than this best case estimate, for a few reasons. Firstly, the amount of scattered light is expected to be lower as described above. Secondly, for the case of multiple scattering not all of the scattered light is polarized -- multiple scatterings mean the photon loses some of its polarization signature. Thirdly, when the particle is large compared to the wavelength of light, different light paths through the particle and interference effects cause the polarized light to be lower and to have more than one peak, and so a strong single peaked signal such as shown in Figure~\\ref{fig:lambert}b may not be reached.  As shown in Figures~\\ref{fig:lambert}a~and~\\ref{fig:lambert}b, the light curves and polarization are very sensitive to $D$, since $\\epsilon \\sim 1/D^2$. For example, the Lambert sphere at $D=0.042$~AU (corresponding to HD~187123) has a light curve and polarization curve with amplitude twice as high as a Lambert sphere at $D=0.059$~AU (corresponding to $\\upsilon$ And b). The light curve and polarization curve estimates are also sensitive to $R_P$. Both can be scaled --- the light curve approximately and the polarization curve exactly --- for different $R_P$ by the factor $(R_P/1.2 R_J)^2$. For HD 209458~b with $R_P=1.40 \\pm 0.17 R_J$ (Mazeh et al. 2000), this factor is 1.36. ",
        "conclusions": "We have presented photometric light curves and fractional polarization curves for 51 Peg b for 4 mean particle sizes, and discussed the differences with color and with other CEGPs. The light curves are very sensitive to condensate type and size distribution, hence observations will be able to distinguish between extreme scenarios. However, more detailed information such as the exact size distribution and particle type will be more difficult to extract.  The temperature-pressure profiles are also extremely dependent on the condensate assumptions. We have briefly discussed these, along with the emergent spectra. In contrast to T dwarfs which have no observeable clouds, the CEGPs should have clouds closer to the top of their atmospheres because irradiation heats the upper atmosphere to temperatures closer to the equilibrium condensation curve (i.e. the cloud base). The condensates that may not contribute to reflected light because they are sequestered below the top cloud layer will still affect the temperature-pressure profile by heating the lower atmosphere. Thus observations of the light curves which should constrain the general cloud properties will help distinguish between atmosphere models.  The light curves may be very different from sine curves; their shape depends not just on the particle size and type, but also on $\\tilde{\\omega}$ and the atmospheric density. Inclinations other than the narrow angular range possible for transits are theoretically detectable for the CEGPs. Many cases of the light curves are detectable by upcoming space missions, and some of the largest amplitude cases (e.g. from pure MgSiO$_3$ particles with $\\overline{r}=0.1~\\mu$m) might be detectable from the ground in the near future. See Table~2, and equation~(\\ref{eq:dm}) using $\\epsilon = p(R_P/D)^2$. Results from this paper show that the light curve observations that would best constrain atmosphere models are those at different colors in narrow wavelength bands.  Geometric albedos at 5500 \\AA~in this study for a cloud mix of particles of MgSiO$_3$, Fe, and Al$_2$O$_3$ range from $p=0.44$ for particles with $\\overline{r}=10~\\mu$m that have a strong backscattering peak, to $p=0.0013$ for particles with $\\overline{r}=0.01~\\mu$m which include highly absorbing Fe. Clouds of pure MgSiO$_3$ are much more reflective for all particle sizes (see Table~2).  The polarization of the CEGP systems is not detectable with current techniques. Rayleigh scattering polarization peaks at orbital angle 90$^{\\circ}$, but with modulation of the reflected light the fractional polarization has an asymmetric peak at around 70$^{\\circ}$. Polarization from particles that are large compared to the wavelength will have more than one peak due to interference effects from different light ray paths through the particle.  Other CEGP systems in our study give similar light curves to 51 Peg b, but effects from distance from the parent star and density are important. Planets farther than 0.1~AU from their parent stars are too faint in reflected light to be detected photometrically in the foreseeable future.  We emphasize that there are many unknowns in the model atmospheres, and much room for improvement --- mainly more realistic cloud modeling, heat redistribution by winds, photochemistry, non-spherical particles, and other types of condensates. These ingredients will result in different light curves than those shown in this paper. That the predictions are so varied means observations should be able to identify the gross cloud characteristics. In this sense the theory should be seen as an interpretive rather than a predictive tool. Observations by the upcoming satellites MOST, COROT, and MONS and ground-based work will help constrain the CEGP atmosphere models and at best will reveal the nature of their atmospheres directly.   \\bigskip"
    },
    "0004/astro-ph0004223_arXiv.txt": {
        "abstract": "The CCD magnitudes in Johnson $V$ and Cousins $R$ and $I$ photometric passbands are determined for GRB 000301C afterglow starting $\\sim$ 1.5 day after the $\\gamma-$ray burst. In fact we provide the earliest optical observations for this burst. Light curves of the afterglow emissions in $U, B, V, R, I, J $ and $K^{'}$ passbands are obtained by combining the present measurements with the published data. Flux decay shows a very uncommon variation relative to other well observed GRBs. Overall, there is a steepening of the optical and near-infrared flux decay caused by a geometric and sideways expanding jet. This is superimposed by a short term variability especially during early time ($\\Delta t < 8$ days). The cause of variability is not well understood, though it has occurred simultaneously with similar amplitude in all the filters. After removing the superposed flux due to variability, we derive the early and late time flux decay constants using jet model. The late time flux decay is the steepest amongst the GRB OTs observed so far with $\\alpha \\sim 3$. Steepening in the flux decay seems to have started simultaneously around $\\Delta t \\sim 7.6$ day in all passbands. On the other hand no such variations are observed in the quasi-simultaneous broad-band photometric spectral energy distributions of the afterglow. The value of spectral index in the optical-near IR region is $\\sim -1.0$. Redshift determination with $z=2.0335$ indicates cosmological origin of the GRB having a luminosity distance of 16.6 Gpc. Thus it becomes the second farthest amongst the GRBs with known distances. An indirect estimate of the fluence $>$ 20 keV indicates, if isotropic, $\\ge 10^{53}$ ergs of release of energy. The enormous amount of released energy will be reduced, if the radiation is beamed which is the case for this event. Using a jet break time of 7.6 days, we infer a jet opening angle of $\\sim$ 0.15 radian. This means the energy released is reduced by a factor of $\\sim$ 90 relative to the isotropic value. ",
        "introduction": "We have used the published data in combination with the present measurements to study the flux decay of GRB 000301C afterglow. Fig. 3 shows the plot of photometric measurements as a function of time. The X-axis is log ($t-t_0$) where $t$ is the time of observation and $t_0$ is the time of GRB burst which is 2000 March 1.411 UT. All times are measured in unit of day. Before deriving the flux decay constants of the OT, it is mandatory to subtract the contributions from foreground/background galaxies, if there is any. Both ground based  and the early HST images clearly show that any underlying galaxy would have to be fainter than $R >$ 25 mag (Fruchter et al. 2000a). In fact, the late-time HST images taken on 2000 April 3.9 UT by Fruchter et al. (2000b) showed that the $R$ magnitude of the host galaxy is 27.8$\\pm$0.25. We have therefore not applied any correction upto $R <$ 23 mag of the OT for the contamination by host galaxy.  \\begin{figure*} \\begin{center}\\vspace*{-1.5cm}\\hspace*{-1cm} \\psfig{file=fig3.ps,width=20cm,height=20cm} \\end{center}\\vspace*{-1.1cm} \\caption{\\label{light} Light curve of GRB 0000301C afterglow in optical and near-IR photometric passbands. Measurements from UPSO, Nainital have been indicated as asterisk. Suitable offsets have been applied to aviod overlapping in data points of different passbands. Flux decay can not be fitted by a single power-law. Solid line represents the least square non-linear fit to the densely observed data for a jet model while dotted lines are the jet model curves for $\\alpha_1 = 1.2$ and $\\alpha_2 = 3.0$. In all cases the value of $\\beta$ is taken as 5. Short term variability observed in different passbands is shown in the upper right corner box.} \\end{figure*}  \\medskip The flux decay of most of the earlier GRB afterglows is generally well characterized by a single power law $F(t) \\propto (t-t_0)^{-\\alpha}$, where $F(t)$ is the flux of the afterglow at time $t$ and $\\alpha$ is the decay constant. However, optical and near-IR light curves of GRB 000301C (Fig. 3) show erratic behaviour with an overall flux decay. Observers therefore took relatively long time to accept it as an OT of the GRB 000301C.  \\medskip UPSO observation in $R$ filter on 2000 March 2.93 UT is the earliest optical observations published so far. Bhargavi \\& Cowsik (2000) measurements are just after us. Fig. 3 clearly indicates peculiar behaviour of the light curve and perhaps, even shows $\\sim$ 0.5 mag brightening of the $R$ magnitude with $\\alpha = -0.5\\pm1.0$ during $\\Delta t = 1.5 - 1.8$ day. This could be an indication of a rising phase similar to that seen in GRB 970228 (Guarnieri et al. 1997) and GRB 970508 (Castro-Tirado et al. 1998). Contrary to most of the earlier GRB afterglows, light curve of GRB 000301C can not be fitted by a single power-law (see also Masetti et al. 2000, Rhoads \\& Fruchter 2000, Berger et al. 2000; Jensen et al. 2000). Overall the OT flux decay seems to have broken power-law as expected in GRB afterglows having jet-like relativistic ejecta (Sari et al. 1999; Rhoads 1999). This appears to be superimposed with some shorter time flux variability especially during $\\Delta t < 8$ day. Among equally well monitored GRB afterglows, GRB 000301C appears therefore peculiar. Correlated variability can be clearly noticed in $B, R$ and $I$ passbands. The lack of such apparent correlation in the light curves of other passbands is most probably due to non-strict simultaneity of the data points. Broken power-law in these light curves can be empirically fitted by functions of the form (see Rhoads \\& Fruchter 2000)  \\medskip $F(t) = 2F_0/[(t/t_b)^{\\alpha_1\\beta} + (t/t_b)^{\\alpha_2\\beta}]^{1/\\beta}$,  \\medskip where $\\alpha_1$ and $\\alpha_2$ are asymptotic power-law slopes at early and late times with $\\alpha_1 < \\alpha_2$ and $\\beta > 0$. $\\beta$ controls the sharpness of the break, with larger $\\beta$ implying a sharper break. With $\\beta = 1$, this function becomes the same that Stanek et al. (1999) fit the optical light curve of GRB 990510 afterglow. $F_0$ is the flux at the cross-over time $t_b$. The function describes a light curve falling as $t^{-\\alpha_1}$ at $t << t_b$ and $t^{-\\alpha_2}$ at $t >> t_b$. The function can be written in magnitudes as  \\medskip $m=m_b + \\frac{2.5}{\\beta} [log_{10} \\{ (t/t_b)^{\\alpha_1\\beta} +(t/t_b)^{\\alpha_2\\beta}\\} - log_{10}(2)]$,  \\medskip where $m_b$ is the magnitude at time $t_b$. In jet models, an achromatic break in the light curve is expected when the jet makes the transition to sideways expansion after the relativistic Lorentz factor drops below the inverse of the opening angle of the initial beam. Slightly later, the jet begins a lateral expansion which causes a further steepening of the light curve. Before fitting jet model to the light curve to derive accurate flux decay parameters of the afterglow, it is mandatory to deconvolve the short term variability component. Otherwise, it will confuse the determination of $t_b$. Perhaps it is the main reason for having a range of $t_b$ values in the literature. The short term variability component of the light curve is determined as described below.  \\medskip In order to minimize the effects of short term variability on the determination of the parameters of the fireball model, Berger et al. (2000) fitted the entire data set from radio to optical simultaneously and derive $t_b = 7.5\\pm0.5$ days, $\\alpha_1 = 1.28$ for $t < t_b$ and $\\alpha_2 = 2.70$ for $t > t_b$ as global parameters for the jet fireball model. The observed $U, B, V, R, I,$ and $K^{'}$ fluxes are divided by the values obtained from the jet model fit yields, as also is noticed by Berger et al. (2000) that the variability is simultaneous and of similar amplitude in all bands (see upper right corner box in Fig. 3). There is a sharp rise and decline centered on $\\Delta t = 4$ day. Berger et al. (2000) also found similar variability at 250 GHz. All these indicate that variability is the result of a real physical process which produces simultaneously similar level of absolute variation over a large range in frequency. Berger et al. (2000) therefore explain this fluctuation in terms of non-uniform ambient density which varied by about a factor of 3.   \\medskip We use the densely covered observations in $B$, $R$ and $I$ to determine the parameters of jet model using the above function. For this, the short term variability was deconvolved from the observed light curves. It has been noticed that the minimum value of $\\chi^2$ is achieved for $\\beta \\ge 5$. This indicates that the observed break in the light curve is sharp and is unlike the smooth break observed in the optical light curve of GRB 990510 (cf. Stanek et al. 1999; Harrison et al. 1999). In order to avoid a fairly wide range of model parameters for a comparable $\\chi^2$ due to degeneracy between $t_b, \\alpha_2$ and $\\beta$, we have used fixed value of $\\beta = 5$ in our further analyses. The least square best fitted parameters $t_b, m_b, \\alpha_1,$ and $\\alpha_2$ have values 7.51$\\pm$0.63, 22.15$\\pm$0.15, 1.18$\\pm$0.14 and 3.01$\\pm$0.53 respectively in $R$. The corresponding values are 8.27$\\pm$1.11, 23.20$\\pm$0.24, 1.24$\\pm$0.20 and 3.48$\\pm$2.07 respectively in $B$ and 7.27$\\pm$1.04, 21.64$\\pm$0.26, 1.17$\\pm$0.29 and 2.92$\\pm$2.93 respectively in $I$. This indicates that average values of $t_b, \\alpha_1,$ and $\\alpha_2$ are 7.6$\\pm$0.5 day, $1.2\\pm0.1$ and $3.0\\pm0.5$ respectively. The light curves derived with these averaged parameters using the jet model are shown by dotted curves in the $U, V, J$ and $K^{'}$ passbands. This clearly indicates the presence of simultaneous short term variability in all passbands. We therefore conclude in agreement with Masetti et al. (2000) and Rhoads \\& Fruchter (2000) that optical and near-IR flux decays of GRB 000301C afterglow are peculiar in comparison to other such well observed GRB afterglows.  \\subsection{Spectral index of the GRB 000301C afterglow }  The flux distribution of the GRB 000301C afterglow has been studied using the broadband photometric measurements listed in Table 1 along with the published radio, millimeter and ultra-violet observations. We used the reddening map provided by Schlegel, Finkbeiner \\& Davis (1998) for estimating Galactic interstellar extinction towards the burst and found a small value of $E(B-V) = 0.05$ mag. We used the standard Galactic extinction reddening curve given by Mathis (1990) in converting apparent magnitudes into fluxes and used the effective wavelengths and normalisations by Bessell (1979) for $B, V, R$ and $I$ and by Bessell \\& Brett (1988) for $J$ and $K^{'}$. The fluxes thus derived are accurate to $\\sim$ 10\\%. Fig. 4 shows the spectrum of GRB 000301C afterglow from optical to radio region. The fluxes closest to $\\Delta t = 4.8$ day at 1.4 GHz, 4.86 GHz, 8.46 GHz, 15 GHz, 22.5 GHz, 100 GHz, 250 GHz and 350 GHz are taken from Berger et al. (2000). It is observed that as the frequency decreases the flux increases from optical to millimeter wavelengths and then it turns over. The spectrum thus can be described by a single power law in some frequency interval as $F_{\\nu} \\propto  \\nu^{\\beta}$, where $F_{\\nu}$ is the flux at frequency $\\nu$ and $\\beta$ is the spectral index. In the optical to millimeter region, the value of $\\beta $ is $-0.73\\pm$0.06 at $\\Delta t =$ 4.8 day. The optical-near-IR spectrum has not changed significantly (see Table 2 and Fig. 4) and has average value around $-1.0$. This is in agreement with a single value of $\\beta = -1.1$ derived from the low-resolution spectrum taken on 2000 March 3.47 UT by Feng et al. (2000) in the wavelength range of 0.3 to 0.6 $\\mu$m. The HST observations taken around $\\Delta t$ = 33.5 day by Fruchter et al. (2000b) also indicate similar slope. All these, perhaps, indicate no change in the spectral slope of GRB 000301C at later times. There is thus no evidence for a cooling break passing through the optical band on these time scales. This is unlike GRB 980326 ( Bloom et al. 1999) and GRB 970228 (Fruchter et al. 1999; Galama et al. 2000) where spectral index changed to a value of $\\beta \\sim -3.0$ after $\\Delta t >$ 20 days. The spectral slope at radio to millimeter frequencies is generally expected to be $+1/3$ at these early times. However, the observed slope is much larger with a value of $+0.90 \\pm 0.08$. The peak frequency seems to lie in millimeter region. This peak frequency is thus similar to that of GRB 970508 (cf. Galama et al. 1998) but different from that of GRB 990123 (Galama et al. 1999) where the peak is in radio region and that of GRB 971214 for which the peak is in optical/near-infrared waveband (Ramaprakash et al. 1998). From this, one may infer that the synchrotron peak frequency may span a large range in GRB afterglows.  \\begin{figure*} \\begin{center}\\vspace*{-0.5cm}\\hspace*{-1cm} \\psfig{file=fig4.ps,width=8cm,height=8cm} \\end{center}\\vspace*{-0.6cm} \\caption{\\label{spec} The spectral flux distribution of the GRB 000301C afterglow at $\\sim$ 4.8 day after the burst. Fluxes measured at optical, near-IR, radio and millimeter wavelengths closest to the epoch are plotted. The least square linear relations derived using fluxes at optical, near-IR, 350 GHz and 250 GHz and at 4.86 GHz, 8.46 GHz, 15 GHz, 22.5 GHz and 100 GHz are shown by solid lines. Inside the box, optical-near-IR spectral slopes derived at different epochs (see Table 2) are shown as a function of $\\Delta t$.} \\end{figure*} ",
        "conclusions": ""
    },
    "0004/astro-ph0004153_arXiv.txt": {
        "abstract": "We suggest a solution to an important problem of observational helioseismology of the separation of lines of solar acoustic (p) modes of low angular degree in oscillation power spectra by constructing optimal masks for Doppler images of the Sun. Accurate measurements of oscillation frequencies of low-degree modes are essential for the determination of the structure and rotation of the solar core. However, these measurements for a particular mode are often affected by leakage of other p modes arising when the Doppler images are projected on to spherical-harmonics masks. The leakage results in overlaping peaks corresponding to different oscillation modes in the power spectra. In this paper we present a method for calculating optimal masks for a given (target) mode by minimizing the signals of other modes appearing in its vicinity. We apply this method to time series of 2 years obtained from Michelson Doppler Imager (MDI) instrument on board SOHO space mission and demonstrate its ability to reduce efficiently the mode leakage. ",
        "introduction": "Accurate measurements of frequencies and other physical properties of solar oscillations are crucial for understanding the internal structure and dynamics of the Sun. Up to now the determination of low-degree p-mode parameters like cyclic frequency $\\nu$, linewidth or frequency splitting is based on fitting the so-called $m$-$\\nu$ diagrams obtained by projecting time series of solar images (usually Dopplergrams) on to a basis of spherical harmonics of angular degree $l$ and azimuthal order $m$. Because only one hemisphere of the Sun is observed the projection on to spherical harmonics is not orthogonal. This results in leakage of modes which makes the line fitting quite complicated especially when modes overlap with the target mode. We have therefore decided to apply a method originally developed by Kosovichev (1986) based on a singular value decomposition (SVD) technique in order to minimize the leakage.  We discuss in  Sec. 2 the principle of the method, and introduce the concepts of `global' and `local' optimal masks. The global optimal masks are designed to minimize the contribution of all modes except the target one. The local masks minimize only modes within a narrow frequency interval centered at the target mode. In Sec. 3 we show the theoretical efficiency of the method compared to the usual spherical harmonics projection.  In Sec. 4 we illustrate this method with a 2-year time series of MDI LOI-proxy data, obtained by rebinning of the original 1024 by 1024 CCD pixels into 180 bins (Hoeksema et al., 1998; Scherrer et al., 1995) according to the shape of the LOI instrument (Appourchaux et al. 1997).  \\begin{figure}[t] \\epsscale{1} \\plotone{fig1a.eps} \\caption{An example of MDI LOI-proxy Dopplergrams. The full-disk 1024x1024-pixel MDI Dopplergrams are rebinned on board SOHO into 180 bins with the boundaries along heliospheric meridians and parallels. The variation from dark to light colors is due to solar rotation.\\label{figure1}} \\end{figure}  \\section {The optimal-mask method}  The aim of the optimal masks (Gough, 1984 and Christensen-Dalsgaard, 1984) is to find an optimal linear combination of $N$ signals arising from $N$ bins of the solar image in order to produce a time series containing mainly the contribution of a mode with given $l$ and $m$, the so-called target mode, and minimizing the contribution of others modes. It is particularly important to minimize the contribution of modes which are close to the target mode in the power spectrum. These are mainly modes of the same rotationally split multiplet as the target mode. At frequencies about 2.5 mHz and higher the linewidth of p modes becomes larger than the separation between the peaks. Therefore, it is important to separate the overlapping modal lines. We describe this method in the application to the MDI LOI-proxy data. However, this method is quite general and can be applied to other spatially resolved helioseismic data.  Because we observe only one hemisphere of the Sun, the commonly used spherical-harmonic masks are not able to remove all the components of the target mode multiplet.  Instead of using the spherical harmonics masks we apply a SVD method to derive the optimal masks.  To apply this method we first need to model the signal produced in each image bin by a mode of given angular degree $l$, azimuthal order $m$ and radial order $n$. For simplicity, in this paper, the horizontal component of velocity is neglected because we are only interested in low-degree p modes. The inclusion of the horizontal component is straightforward, and generally improves the quality of the optimal masks in the low-frequency part of the solar oscillation spectrum where this component becomes significant. Hence, to the first order, the signal of a normal mode in each pixel $j$ of the CCD detector is: \\begin{equation} v^{(j)}_{nlm}(t) = \\Re\\left\\{V_{nlm}e^{i\\omega_{nlm}t}s^{(j)}_{lm}(t)\\right\\}, \\end{equation} where $V_{nlm}$ is the mode amplitude, $\\omega_{nlm}$ is the mode frequency, and $s^{(j)}_{lm}(t) = Y_{lm}[\\theta_{j}^{*}(t),\\phi_{j}^{*}(t)]\\mu_{j}$ is the sensitivity function which depends on the spatial eigenfunction, spherical harmonic $Y_{lm}$, and the factor $\\mu_{j}$ arising from the projection of the radial velocity onto the line of sight. Both $Y_{lm}$ and $\\mu_{j}$ are expressed in terms of the heliographic coordinates, $\\theta_{j}^{*}(t)$ and $\\phi_{j}^{*}(t)$, of the pixels, which are generally functions of time $t$.   We have to take into account the fact that the orientation of the solar image, which is usually described by two angles of the rotation axes with respect to the ecliptics, $P_0$ and $B_0$, and the image size on the detector are possibly changing with time. The time dependences of the $P_0$ and $B_0$ angles should therefore be taken into account by rotating the spherical coordinates accordingly. In our case the variations of the image size were not taken into account because the data were combined into fixed 180 bins on board spacecraft prior to any mask application. An example of the rebinned Dopplergrams (`MDI LOI-proxy') is shown in Fig. 1.   The k$^{th}$ signal of each of the $N$ bins is obtained by averaging the signals coming from $N_{k}$ CCD pixels withing this bin as defined by the mapping between the MDI pixels and the LOI-proxy bins: \\begin{equation} v^{(k)}_{nlm}(t) = \\Re\\left\\{V_{nlm}e^{i\\omega_{nlm}t}S^{(k)}_{lm}(t)\\right\\}, \\end{equation} where $S^{(k)}_{lm}(t) = \\frac{1}{N_{k}}\\sum_{j=1}^{N_{k}} s^{(j)}_{lm}(t)$ is  $S^{(k)}_{lm}$ is the sensitivity of bin $k$ to a normal mode of angular degree $l$ and order $m$. To extract the signal of a target mode of $n_0$, $l_0$ and $m_0$ we combine the $N$ signals with $N$ weights $w_{n_0l_0m_0}^{(1)}$,...,$w_{n_0l_0m_0}^{(N)}$: \\begin{eqnarray} v_{n_0l_0m_0}(t) = \\sum_{k=1}^N  w_{n_0l_0m_0}^{(k)} v^{(k)}_{nlm}(t)=\\nonumber\\\\ \\Re\\left\\{V_{nlm}e^{i\\omega_{nlm}t}\\left[\\sum_{k=1}^N w_{n_0l_0m_0}^{(k)}S^{(k)}_{lm}(t)\\right]\\right\\}. \\end{eqnarray} Ideally, if  the sum in the square brackets is equal to the product of two Kroneker symbols, $\\delta_{ll_{0}}\\delta_{mm_{0}}$, then $v_{n_0l_0m_0}(t)$ consists only of the signal of the target mode. However, in practice, this sum does not represent $\\delta_{ll_{0}}\\delta_{mm_{0}}$ and contains not only the target mode but also contributions from other modes.  Vector $\\vecw_{n_0l_0m_0}\\equiv\\{w_{n_0l_0m_0}^{(1)},...,w_{n_0l_0m_0}^{(N)}\\}$ defines an optimal mask if this vector is chosen to maximize the signal for the target mode and minimizes the signals for other, say $M$, modes located close to the target mode in the Fourier domain. These conditions can be written as a minimization of the quadratic form: \\begin{equation} r=||\\vecS\\cdot\\vecw - \\vecI||^2, \\end{equation} where $\\vecS$ is a $(M+1)\\times N$ matrix with elements $S^{(k)}_{nlm}(t)$ and $\\vecI$ is a $(M+1)$-element vector with elements $\\delta_{ll_{0}}\\delta_{mm_{0}}$, $l_0$ and $m_0$ being the angular degree and order of the target mode.  Finding $\\vecw$ is equivalent to finding a least-squares solution to the system of linear equations, $\\vecS\\cdot\\vecw = \\vecI$. Depending on $M$, the number of modes we want to minimize, we have a system which is over-determined (N $>$ M+1) or under-determined (N $<$ M+1). For both cases the problem is efficiently solved using the singular value decomposition (SVD) of the matrix $\\vecS$ (e.g. Press et al. 1992). This method optimally solves the linear system and its solution $\\vecw$ gives the optimal mask. This  mask provides an efficient filtering of the selected $M$ modes. Of course, the minimization is less efficient when $M$ is large. There is therefore a trade-off between the number of modes to minimize and the efficiency of this minimization.  Matrix $\\vecS$ describes the sensitivity of each detector bin to the mode of oscillation. For the special case of an under-determined system (N $<$ M+1) this technique is equivalent to the technique of diagonalization of the leakage matrix (Appourchaux et al., 1998). The SVD method allows us to consider general cases. With this technique we can specifically decide which modes must be filtered and to what extend their contribution to the power spectra should be reduced, making the technique extremely flexible. In practice we filter as much as possible the modes having frequencies falling in a given frequency window around the target making a kind of local cleaning of the spectra.  We therefore call the corresponding masks `local optimal masks' to distinguish from the `global optimal masks' obtained when minimizing the contribution of all the modes not only those around the target mode.  The larger the window around the target the more modes have to be minimized, and thus the less efficient is the filtering. In principle we are able to remove efficiently as many modes as the number of the binned signals, 180 in our case. Above that number the SVD acts as a least square minimization and the modes cannot be fully filtered. In practice, the efficiency of the filtering depends on the selected mode set and the noise level.  The solar noise obviously will be less filtered than modes by the masks because its spatial structure is incoherent.  Thus the signal-to-noise ratio of the target mode depends on the mask. More efficient masks typically require higher weights $w_{nlm}^{(k)}$, and, therefore amplify the noise. This leads to the trade-off between the mask efficiency and the noise level. To adjust this trade-off we slightly modify Eq.~(4) and add the noise contribution through a regularization term. Then the quadratic form to minimize is \\begin{equation} r=||\\vecS\\cdot\\vecw - \\vecI||^2 + \\alpha ||\\vecn \\cdot \\vecw ||^2, \\end{equation} where $\\vecn=\\{n^{(1)},...,n^{(N)}\\}$ is the vector of the standard deviation of the noise in each of the $N$ bins, and $\\alpha$ is the regularization parameter. The SVD technique is still applicable here. Eq. (4) corresponds to the case of $\\alpha$=0 which means no regularization.  The noise level is estimated by taking the standard deviation of the signal in each bin, assuming the main contributor is the solar noise. Generally, the noise level is the lowest at the center of the solar disk and increases toward the limb. Table 1 gives a sample of the relative noise level of the MDI measurements as a function of the angular distance from the disk center. Obviously, the absolute noise level is not required for Eq.(5). \\begin{table} \\caption{Relative noise level, $n$, in the MDI data as a function of the angular distance, $d$, from the disk center.} \\begin{center} \\begin{tabular}{crrrrrr} \\hline\\hline $d$, deg.  & 1.72 & 11.0 & 20.5 & 30.7 & 42.1 & 56.1\\nl $n$  & 1.00 & 1.15 & 1.40 & 1.73 & 1.99 &  2.27 \\nl \\hline \\end{tabular} \\end{center} \\end{table}  The regularization parameter $\\alpha$ is chosen such as the amplitudes of the filtered modes is reduced to the level of noise. It turns out that the signal-to-noise ratio obtained with the local optimal masks is similar to what is obtained with spherical-harmonic masks, though it may be somewhat smaller (up to 50\\%). Only for $l$=0 the difference in the signal-to-noise ratio between the two masks can reach a factor of two or more.  Because of solar noise we also avoid using the limb bins. These bins improve very little  the efficiency of the masks but contribute significantly to the noise. Removing not only limb bins but also the next round of bins helps to increase the signal-to-noise ratio at low frequencies for $l=0$ modes, but this is not effective for modes of higher degrees.  The power of this technique is also that the optimal masks for any $P_0$ and $B_0$ angles are easily obtained applying the spherical harmonic rotation matrix to the optimal masks with $P_0$=0 and $B_0$=0. Therefore it is possible as for MDI to apply the optimal masks for each day according to the value of $B_0$ which is changing with time without recomputing the SVD problem which is a heavy computational task . ",
        "conclusions": "Mode leakage which results in overlapping mode peaks in solar oscillation power spectra is one of the most significant problems in observational helioseismology. It may lead to significant errors in the determination of mode frequencies and other properties. We have shown that this problem can be efficiently solved by using the optimal masks which are applied to series of solar images and reduce the signals of solar modes in a narrow frequency interval around a target mode to the level of noise. The optimal mask, however, may increase the noise level. Therefore, the filtering efficiency of the masks has to be balanced with the noise level. This is achieved by a regularization technique. We have demonstrated in the case of low-degree modes that the optimal mask technique allows us efficiently isolate individual modes in rotationally split multiplets, which are unresolved in the commonly used disk-integrated data, and also isolate modes in the case of overlapped multiplets of different angular degree. The optimal masks are quite efficient for filtering out unwanted modes as long as their number is less or of the order of the number of solar image bins.  We used the MDI LOI-proxy data to calculate the optimal masks. However, this method is quite general and can be used for any spatially resolved helioseismic data.  The method has been applied for the determination of the central low-degree frequencies by Toutain et al (1999). In a future paper, we will present the results for rotational splitting of low-degree modes."
    },
    "0004/astro-ph0004403_arXiv.txt": {
        "abstract": "The sequence of phase transitions during the hot history of the universe is followed within a phenomenological framework. Particular emphasis is put on the QCD confinement transition, which is at reach under earth laboratory conditions. A tepid inflationary scenario on the GUT scale with bubble growth at moderate supercooling is discussed. ",
        "introduction": "The standard cosmology teaches that the universe is expanding. This implies a steady change of the state of matter. Particularly interesting are phase transitions where the structure and the relevant degrees of freedom change. Probably the most drastic transition was the formation of the early universe. There is a number of ideas how did this happen, including the creation from a quantum state (cf.\\ \\cite{Kiefer} and further references there). Since the extrapolation of our quantum field theoretical models to high temperature and density points to a maximum temperature of $T \\sim 10^{16}$ GeV, where thermal equilibrium between various particles can be maintained, one can guess that the thermal history starts at such a temperature scale and corresponding world age of $10^{-37}$ sec. Immediately afterwards one conjectures that a symmetry breaking phase transition happens. At a scale of $T \\sim 10^{16}$ GeV, or slightly less, the extrapolated coupling strengths of the strong, weak and electromagnetic interactions merge, and one generally believes that the behavior of matter is to be described within the framework of a grand unified theory (GUT). The very structure of GUT is still matter of debate, and also the relation to supersymmetric and string theories is not yet settled down. Nevertheless, inflation, i.e. a stage of accelerated expansion and a huge blowing off of the space is an important scenario, which allows to understand some of the presently observed properties of the universe.  The subsequent evolution appears quite unspectacular for a long time span: expansion means cooling, and particles with masses in the order of 40\\% of the corresponding temperature disappear due to annihilation.  The next interesting stage is the symmetry breaking on the electroweak scale, i.e. $T \\sim 100$ GeV. Cooling further, at $T \\sim 160$ MeV the quarks and gluons, roaming up to this stage freely, become bound in hadrons. This is the confinement transition, which might be called the hadrosynthesis stage.  Once the temperature falls below 1 MeV the nucleosynthesis starts, creating the light primordial elements. In nuclear matter, a liquid-gas phase transition is conjectured since a long time (cf.\\ \\cite{our_liquid_gas}), and experimental hints have been verified \\cite{liquid_gas}. However, in the cosmic evolution this transition does not play a role since it is related to large baryon density.  From such a point of view the history of the early universe can be considered as sequence of phase transitions. Indeed, based on quite general models, one can develop a scheme to combine the dynamics of phase transitions with the cosmic expansion, as envisaged in \\cite{Boyanowski}. Here we follow, however, a different avenue and consider the specific features of individual transitions.  Among the above mentioned phase transitions the confinement transition from quarks and gluons to hadrons has been elaborated by far in most details. This is because it is believed to be reproduceable in the laboratory. Indeed, for colliding nuclei at sufficiently high energies, the relevant degrees of freedom should be represented by partons, i.e. point-like structures. The presently terminating first round of heavy-ion experiments at CERN accumulated such a wealth of data which give rise to the interpretation that in central nuclear collisions a thermalized state of deliberated quarks and gluons has been created. Indeed, on February 9, 2000 a press release of CERN announced ``compelling evidence for the existence of a new state of matter in which quarks, instead of being bound up into more complex particles such as protons and neutrons, are liberated to roam freely, $\\cdots$ this state must have existed at about 10 microseconds after the Big Bang'' \\cite{CERN}.  In accordance with this new information we will focus here on the confinement transition after the Big Bang. We will compare schematically the confinement dynamics of the Big Bang with that of the Little Bang (i.e. relativistic heavy-ion collisions). By using the known details of the transition we demonstrate the interplay of local matter properties and global evolution dynamics. More specifically, we analyze how drives the local matter state the global expansion, and vice versa, how the expansion causes cooling and transformation of matter forms into each other.  There is some apparent similarity in cosmology and heavy-ion physics,  which we try to visualize in Fig.~\\ref{expanding_matter}. From present observations of remote objects in the universe, we look somewhat into the past. Combining observational facts, like the distribution and the red shift of galaxies, as seen in the left panel of Fig.~\\ref{expanding_matter}, one can develop a picture of the early stages of the universe. The situation in heavy-ion physics is very similar to this. Observing the created hadrons in a very late stage, one tries to extrapolate back to the hottest and densest stages. The artist's view of a semi-peripheral collision of heavy-ions, displayed in the right panel of Fig.~\\ref{expanding_matter}, shows a stage where the matter is converted into hadrons which still interact for a while before disassembling. However, there is a chance in heavy-ion collisions to receive a direct signal of the very early stage: the mean free path of real and virtual photons \\cite{direct_probes} is large in comparison with the system size. For such direct probes the problem, however, arises to disentangle the primordial signal from a huge background \\cite{Gallmeister_2}.  \\begin{figure} \\psfig{file=netscape7.eps,width=5.9cm,angle=-0} \\hfill \\psfig{file=aaa1.eps,width=3.9cm,angle=90} \\caption{Expanding matter: The left panel displays a view on a group of remote galaxies as seen by the Hubble telescope (source of the underlying picture: NASA picture gallery in {\\tt www} \\cite{NASA_1}), while the right panel shows the distribution of hadrons in a semi-peripheral collision of heavy nuclei a few fm/c after the first touch (source: CERN press release in {\\tt www} \\cite{CERN}).} \\label{expanding_matter} \\end{figure}  The present survey is restricted to a phenomenological framework. That means, we do not go into a microscopic theory of phase transitions, rather describing them only in terms of the equation of state. We restrict ourselves to the standard model and the grand unifying theory models of the strong and electroweak interactions. Only in a few exceptional cases we quote interesting effects occurring beyond the standard model of particle physics within supersymmetric or string-theoretical models.  The paper is organized as follows. In section 2 the phenomenological framework is presented resulting in Friedmann's equations for the Big Bang and the corresponding equations of matter in the Little Bang. Section 3 deals with quantum chromodynamics (QCD), quarks and gluons and their equation of state. Section 4 exercises a simple way to construct first-order phase transitions. The basic issues of nucleation theory for handling the phase transformation dynamics are presented in section 5. Section 6 is devoted to details of the cosmic confinement transition. A short remark concerning the present status of knowledge on the electroweak transition is made in section 7. In section 8 we deal with an attempt of a phenomenological realization of an inflationary scenario. The summary can be found in section 9.  The material represented is based on {\\tt www} scans till end of April 2000. ",
        "conclusions": "We should be aware that the evolution of the universe is accompanied by a permanent change of the states of matter, which can be in some interesting cases be considered as phase transitions. There is a change of certain order parameters, like a hypothetical scalar field $\\phi$ on the GUT scale, the Higgs field on the electroweak scale, the chiral condensate on the QCD scale, or even now the quintessence. This causes partially drastic changes of matter like the transition from freely roaming quarks and gluons to hadrons, or the binding of nucleons into light nuclei during nucleosynthesis, or simply the annihilation of particles like $e^+ e^- \\to 2 \\gamma$. These changes induce also changes of the expansion dynamics of the universe, most drastic during inflation, e.g.. Furthermore in the transition from a radiation dominated universe to a matter dominated one, the time-evolution of the scale factor $R(t)$ changes from $ R \\propto t^{1/2}$ to $R(t) \\propto t^{2/3}$. Some uncertainty is introduced by the evidence recently found of a dominating vacuum energy density, the quintessence, because its back-extrapolation needs more definite knowledge on its nature.  Assuming that in early stages quintessence is not dominating the energy density, we considered in some detail the confinement transition in the Big Bang and, very briefly, contrasted it to the Little Bang, presently investigated in heavy-ion collision experiments. In the Little Bang one observes direct electromagnetic radiation like real photons and virtual photons via $\\gamma^* \\to e^+ e^-$. This points to temperature scales of 170 MeV and larger, i.e., above the deconfinement transition. In contrast to this, there do not seem to be specific definite relics from the confinement era of the Big Bang, apart the protons and neutrons created in hadrosynthesis. We summarized possible relics, whose existence is hypothetical and the predictions of which are sometimes related to poorly known details of the confinement dynamics.  Whereas the results of the nucleosynthesis processes specifically depend on the interplay of neutrino decoupling, neutron life time, and cosmic cooling rate, the hadrosynthesis at confinement seems to be rather unspecific. Nucleons emerge as other hadrons do as well, but the unstable hadrons decay shortly after confinement and the proton-to-neutron ratio is determined by $\\beta$ equilibrium. Hence, hadrosynthesis does not cause a specific hadron composition due to the cosmic evolution since chemical equilibrium means memory loss.  Strictly speaking, the success of the primordial nucleosynthesis makes us strongly believe in a maximum temperature of $T_R > 1$ MeV, while the prediction of a flat universe and of an appropriate fluctuation spectrum for structure formation by inflationary scenarios makes us believe in $T_R > 10^{15}$ GeV. Within these scales, a consistent picture of the evolution of the universe seems to emerge, but it leaves a consistent treatment of the ``initial state'' and the reheating (or, better, ``pre-heating'') to $T_R$ as the most challenging problem probably to be solved within quantum gravity.  \\vspace*{0.25cm} \\baselineskip=10pt{\\small \\noindent Acknowledgments: K. Gallmeister is grateful for preparing the figures. I thank F.W. Hehl and A. Gross for carefully reading the manuscript. Some of the presented material is an up-grade of \\cite{our_booklet}. The author thanks his colleagues for the fruitful collaboration. The work is supported in part by BMBF grant 06DR829/1.}"
    },
    "0004/astro-ph0004290_arXiv.txt": {
        "abstract": "The complex ten-component gravitational lens system B1933+503 has been monitored with the VLA during the period February to June 1998 with a view to measuring the time delay between the four compact components and hence to determine the Hubble parameter $H_0$. Here we present the results of an `A' configuration 8.4-GHz monitoring campaign which consists of 37 epochs with an average spacing of 2.8 days. The data have yielded light curves for the four flat-spectrum radio components (components 1, 3, 4 and 6). We observe only small flux density changes in the four flat-spectrum components which we do not believe are predominantly intrinsic to the source. Therefore the variations do not allow us to determine the independent time delays in this system. However, the data do allow us to accurately determine the flux density ratios between the four flat-spectrum components. These will prove important as modelling constraints and could prove crucial in future monitoring observations should these data show only a monotonic increase or decrease in the flux densities of the flat-spectrum components. ",
        "introduction": "B1933+503 is the most complex arcsec-scale gravitational lens system known and the most noteworthy system discovered in the Cosmic Lens All-Sky Survey (CLASS; Browne et al. 1998), a survey of approximately 15,000 flat-spectrum radio sources observed with the VLA at 8.4 GHz with a resolution of 200 milliarcsec (mas). Sykes et al. (1998) report the discovery of this lens and present VLA and MERLIN maps that reveal up to ten components, four of which are compact and have flat spectra while the rest are more extended and have steep spectra. The ten components are believed to be the multiple images of a background source that consists of a flat-spectrum core (quadruply imaged) and two compact lobes symmetrically situated on opposite sides of the core (one quadruply imaged and the other doubly imaged). A {\\em Hubble Space Telescope (HST)}/WFPC2 image ($I$-band) of B1933+503 (Sykes et al. 1998) shows a faint galaxy with a compact core (the lensing galaxy) but none of the images of the background object are detected. Spectra taken with the Keck telescope have given a redshift for the lensing galaxy of 0.755 (Sykes et al. 1998). Sub-mm observations of B1933+503 (Chapman et al. 1998) pointed to the fact that the source of the lens system should be lying above a redshift of 2 and that has been confirmed recently by observations made with the United Kingdom Infra-red Telescope (UKIRT) that have yielded a redshift of 2.62 for the background source (Norbury et al., in preparation).  {\\em HST} NICMOS F160W-band observations (Marlow et al. 1998) of B1933+503 have uncovered the infra-red counterparts to two of the four compact components. These same two components were also detected in a VLBA 1.7-GHz map with a resolution of 6 mas. Marlow et al. suggest that the failure to detect the other two flat-spectrum components in the infra-red can be due to either rapid variability on a time-scale less than the time delay or due to extinction in the interstellar medium (ISM) of the lensing galaxy. The absence of these same two components in the radio could arise due to scatter-broadening of the images when propagating through the ISM of the lens galaxy as well as from time variability.  B1933+503 offers a unique opportunity to construct an accurate mass model due to the wealth of observational constraints that are available. As well as the large number of images, yielding both positional and flux density ratio constraints, some of these are extended. Perhaps more crucial is that three epochs of VLA 15-GHz data (Sykes et al. 1998) suggest that the flat-spectrum components are variable, by as much as 33~per~cent over the period July to September 1995. This means that it should be possible to determine the relative time delays from flux density monitoring. Taken together this makes B1933+503 a promising system for a determination of the Hubble constant, $H_0$ (Refsdal 1964). The observational constraints given by Sykes et al. (1998) have been used as a starting point for the model presented by Nair (1998). The time delays of the flat-spectrum components with respect to component 1 derived from this model are approximately 8~d for component 3, 7~d for component 4 and 9~d for component 6.  In this paper we present the results of a four-month monitoring campaign using the VLA in its `A' configuration at a frequency of 8.4~GHz. This frequency was chosen over 15~GHz (where the magnitude of flux density variations are likely to be greater) for a number of reasons. As well as the sensitivity of the VLA being greater at 8.4~GHz, the phase stability of the array is much poorer at 15~GHz. In addition, VLA 15-GHz monitoring of the lens system B0218+357 (Biggs et al. 1999) showed that at this frequency the data are subject to calibration errors of a few per~cent due to the gains of the antennas varying with elevation. In Section~2 we describe the observations and observing strategy as well as the reduction methods that we used on the data. In Section~3 we show the light curves that were obtained from the data and we follow this with a discussion of the results. ",
        "conclusions": "We have presented the results of a four month program of monitoring the ten-component lens system B1933+503 with the VLA which have not allowed us to measure the independent time delays in this otherwise promising system for determining $H_0$. The component light curves show fluctuations of 0.7--2.5 per cent which we do not believe represent source variability. Instead they are more likely to result from a combination of thermal noise, deconvolution and model fitting errors. Since the system has already shown indications of variability (Sykes et al. 1998), we think that we have caught this system in a quiescent phase. For this reason, we have embarked on a second monitoring campaign that commenced in June 1999, also with the VLA at 8.4~GHz. An advantage that we have now over the previous attempt is that the flux density ratios between the components are known to high accuracy. Since we did not detect any variability during the 4 month monitoring campaign, which is a period significantly longer than the predicted time delay of 7--9~d, we have determined the intrinsic flux density ratios (Table~\\ref{fluxtab}). Therefore, should this second set of observations again not show reliable features that can be identified in each light curve, but do show a monotonic increase or decrease in flux density, a time delay may still be determined. This has previously been demonstrated by Koopmans et al. (1999) for the CLASS lens B1600+434."
    },
    "0004/astro-ph0004081.txt": {
        "abstract": "In this paper we derive stellar parameters for the Wolf-Rayet star in the $\\gamma$~Velorum binary system (WR11), from a detailed non-LTE model of its optical and infrared spectra. Compared to the study of Schaerer et al., the parameters of the WC8 star are revised to a hotter effective temperature ($T_{\\rm eff} \\sim 57$~kK), a higher luminosity ($\\log(L$/L$_{\\odot})$ = 5.00), and a lower mass loss rate (log($\\dot{M}$ / M$_{\\odot}$/yr) = --5.0 - using a 10\\% clumping filling factor. These changes lead to a significant decrease in wind efficiency number, from 144 to 7, so that the driving mechanism of the wind of this WR star may be simply radiation pressure on lines. The derived spectroscopic luminosity is found to be 40\\% lower than that derived by De Marco \\& Schmutz through the mass-luminosity relationship for WR stars ($\\log(L$/L$_{\\odot})$ = 5.2).  The paper furthermore presents a comparison of the independently-developed modelling programs, \\cmf\\ and \\isa . Overall, there seems to be very reasonable agreement between the derived parameters for WR11, except for the carbon content, which is 2 times higher for \\cmf\\ (C/He=0.15 vs. 0.06, by number). The comparison also confirms a disparity in the predicted flux at $\\lambda$$<$400 \\AA , found by Crowther et al., which will have effects on several nebular line strengths.  The paper also presents the first independent check of the photon loss mechanism proposed by Schmutz. We conclude that, not only is it important to include very many lines to realistically model line blanketing, but in particular those ones that critically interact with strong resonance lines (e.g. \\heii\\ \\la 303.78). The inclusion of these latter lines may significantly alter the wind ionization structure. ",
        "introduction": "\\label{sec:intro}  The double-lined spectroscopic binary $\\gamma$~Velorum (WC8 + O7.5) can be used to derive the mass of the Wolf-Rayet star, a quantity which cannot be directly determined in the case of single objects, but which provides an important constraint in computations of the late stages of massive star evolution. From it, and the mass-luminosity relationship for WR stars (Schaerer \\& Maeder 1992) a value for the WR luminosity can be derived and compared to that derived from spectroscopic analyses.  Orbital parameters together with the mass ratio were determined by Schmutz et al. (1997), while the 7.3$\\sigma$ {\\sc hipparcos} distance allowed Schaerer et al. (1997) to derive the system's total mass ($\\cal M$(O+WR)=29.5$\\pm$15.9~\\msun ), the absolute visual magnitude ($M_V$(O+WR)=$-5.39$) and to carry out a preliminary spectroscopic analysis of both components. Schaerer et al. (1997), however, modelled only the WR \\hei\\ \\la 4471 / \\heii\\ \\la 4541 line strength ratio, {\\it assuming} the mass-loss rate and the carbon abundance. Additionally, line blanketing was not accounted for by their model and their calculation used simple helium and carbon model atoms. Their stellar parameter determination should therefore be considered preliminary.  De Marco \\& Schmutz (1999; hereafter Paper~I) carried out a full spectroscopic analysis of the O star through which its parameters were determined {\\it simultaneously} with the light ratio between the WR and O stellar components. From a solution of the hydrodynamic wind equations, the mass of the O star was also derived and found to be consistent with the evolutionary O star mass. From it and the mass ratio of Schmutz et al. (1997), the mass of the WR star could be calculated (M=9.0$\\pm$1.0~M$_\\odot$). The combined mass of the system was therefore derived to be 39~M$_\\odot$, higher than the value of Schaerer et al. (1997), but consistent within their uncertainty. From the mass of the WR star and the theoretical mass-luminosity relation of Schaerer \\& Maeder (1992), a value for the WR bolometric luminosity was derived (log($L$/L$_\\odot$) = 5.2), significantly higher than the spectroscopic luminosity derived by Schaerer et al. (1997; log($L$/L$_\\odot$) = 4.8). Finally, using the synthetic O star spectrum, the WR11 spectrum was de-convolved from the O star component allowing the WR emission lines to be measured with greater accuracy.  In this paper we determine the stellar parameters of the WR star using the clumped, line-blanketed stellar atmosphere model \\cmf\\ (Hillier 1987, Hillier \\& Miller 1998). Ultimately, higher  spectroscopic luminosities would lower the wind performance number, indicating that a smaller number of photon scatterings is needed to achieve the radiative driving of the wind. This would facilitate the derivation of the wind velocity law. $\\gamma$~Velorum provides us with a new tool to test the model luminosity with an independent method in such a way as to impose a further check on the wind momentun number.  A second aim of this paper is to continue the comparison of the non-LTE stellar atmosphere code \\cmf\\ and the independently developed Sobolev approximation code \\isa\\ (de~Koter et al. 1993, 1997), which is used in combination with the line-blanketing Monte Carlo code of Schmutz (1991). This study follows from the comparison carried out by Crowther et al. (1999) for the WN8 star WR124.  Finally, Schmutz (1997) showed that line-line interaction between \\heii\\ Ly$\\alpha$ and metal lines at similar wavelengths can lead to a larger spectral luminosity. The inclusion of this interaction in the models, an effect that he called ``photon loss\", showed that WR6 might have a larger luminosity than previously thought, and allowed him to calculate a velocity law for its outer wind. Since line-line interaction is included in the \\cmf\\ model code, we can test whether the approximation of Schmutz (1997) is valid and quantify the importance of the photon loss effect in the case of WR11.   In Sect.~\\ref{sec:obs} we summarise the observations, while in Sect.~\\ref{sec:red} we discuss the light ratio between the O and WR stars and the reddening towards the system. In Sect.~\\ref{sec:mod} we describe the model atmosphere codes used in the determination of the stellar parameters. We present our results in Sect.~\\ref{sec:res}, together with the code comparison, while in Sect.~\\ref{sec:pholos} the effect of photon loss is studied. We finally draw our conclusions in Sect.~\\ref{sec:conc}.  % =============================================================== ",
        "conclusions": "\\end{center} \\end{table}    \\begin{figure} \\vspace{9cm} \\special{psfile=\"H1986_f3.eps\"  voffset=-270 hoffset=-150} \\caption{Comparison of the combined O+WR synthetic model (dotted lines) and the IUE/LORES and ISO/SWS observations (solid lines) for reddenings of $E(b-v)$ = 0.00 and 0.06 mag, with a light ratio $L_{\\lambda 4700}$(O)/$L_{\\lambda 4700}$(WR) = 3.8} \\label{fig:red} \\end{figure}  % =================================================================  \\label{sec:conc}  In the following sections we present a critical summary of the work carried out in this paper.  \\subsection{Stellar parameters and bolometric luminosity}  We have presented a detailed spectroscopic analysis of the WR star in the $\\gamma$~Velorum WR+O binary system. Using the wavelength-dependent light ratio and the synthetic O star spectrum from Paper~I we have generated a corrected WR optical and IR spectrum which we have fitted with a synthetic spectrum using the \\cmf\\ non-LTE model atmosphere code. Results are summarised in Table~\\ref{tab:wr11_summary}. Our derived mass-loss is in excellent agreement with the mass-loss derived by Stevens et al. (1996) from X-ray {\\sc asca} observations, once their value is scaled to the {\\sc hipparcos} distance. Although the parameters of temperature, mass-loss and carbon abundance are well constrained, the oxygen abundance is assumed. The discrepancy between the \\heii\\ lines at 4340, 4859 and 6560~\\AA\\ and other \\heii\\ lines remains to be explained.  \\begin{table*} \\begin{center} \\caption{Summary of the stellar parameters obtained for WR11 using {\\sc cmfgen} and {\\sc isa}-wind (blanketed with photon loss activated), compared with those derived by Schaerer et al. (1997) and De Marco et al. (1999a) with a non-line blanketed model with no photon loss implemented. The results obtained by Dessart et al. (2000) for WR135 are also included} \\label{tab:wr11_summary} \\begin{tabular}{cccccc} \\hline Parameters & {\\sc cmfgen} & {\\sc isa}-wind & De Marco et al. & Schaerer et al. & WR135: Dessart et al. \\\\ Filling factor, $f$ &  10\\% &  100\\%  & 100\\% & 100\\% & 10\\% \\\\ Photon Loss & -- &  yes &  no & no & -- \\\\ \\hline                                   % 124Z2 $T_{\\ast}$ (kK)                & 57.1     & 56.0   & 75.9   & 51.0   & 63.0 \\\\ log($\\mdot/\\sqrt{f}$) $M_{\\odot}$yr$^{-1}$& --4.53 &--4.48  & --4.33 & --4.2 & --4.9 \\\\ C/He(by number)                & 0.15     & 0.06  & 0.143  & 0.25   & 0.13 \\\\ O/He(by number)                & 0.03     & 0.01   & --     & --     & 0.03 \\\\ R (R$_{\\odot}$)                & 3.2      & 3.2   & 1.84   & 3.3    & 3.4 \\\\ log($L_{\\ast}/L_{\\odot}$)      & 5.00     & 4.95   & 5.01   &  4.8   & 5.20 \\\\ %M$_V$ (mag)               & --3.72   & --3.72 & --3.67 &--3.86$^*$& $-$4,3\\\\ BC (mag)                       & --4.00   & --3.89 & --4.1  & --3.5  & --4.0 \\\\ $v_\\infty$ (km~s$^{-1}$)       & 1550     & 1550   & 1300   & 1450   & 1400 \\\\ $\\eta$/$\\sqrt{f}$              & 21       & 28     & 29     &  144   & 5.4   \\\\ \\hline \\multicolumn{4}{l}{$^*$For $\\Delta M_V$ = 1.36.}\\\\ \\end{tabular} \\label{tab:wr11_summary} \\end{center} \\end{table*}  Previous analyses of WR11 were carried out by Schaerer et al. (1997) and by De Marco et al. (1999a - using the stellar atmosphere code of Koesterke \\& Hamann, 1995). Schaerer et al. determined the stellar effective temperature by fitting the \\hei\\ \\la 4471 / \\heii\\ \\la 4541 ratio, assuming a (higher) mass-loss and carbon abundance. De Marco et al. (1999a) derived larger mass-loss and temperature, and a lower wind terminal velocity. This can be attributed to the fact that in the parameter space of $\\gamma$~Velorum the helium spectrum can be reproduced {\\it approximately} with a range of temperature and mass-losses. On the other hand, when the lines are analysed in detail, it is clear that for temperatures larger than about 57~kK and log($\\dot{M}$/M$_\\odot$~yr$^{-1}$)$>$--4.45, predicted line shapes deteriorate, becoming wider. This effect was partly compensated for by De Marco et al. (1999a) by adopting a lower wind velocity, although it is clear from their fits that most lines are still too wide. Further, their neglect of line-blanketing and therefore of photon loss (no line-line interaction between the key lines), might have contributed to the different parameters.  With respect to the analysis of Schaerer et al. (1997), the star becomes twice as luminous, 10\\% hotter and shows a mass-loss lower by a factor of two. The carbon abundance is lower by 40\\% (although we remind the reader that Schaerer et al. {\\it assumed} both the mass-loss and the carbon abundance). This carries implications for studies of wind-wind interaction. From this study and the results presented in Paper~I for the O star, we determine:  \\[ {(\\dot{M}\\ v_\\infty)_{WR} \\over (\\dot{M}\\ v_\\infty)_O} \\simeq %{1.44 \\times 10^{-2} \\over 4.48 \\times 10^{-4}} = 33 \\]  \\noindent compared to a value of 208 derived by Schaerer et al. (1997). Adopting the relationships of Eichler \\& Usov (1993) we determine that the region of stellar wind collision is approximately 9-30~R$_\\odot$ from the O star (0.15 of the distance that separates the two stars (63-200~R$_\\odot$ - Schmutz et al. 1997) vs. 0.06 for the study of Schaerer et al. (1997)). Overall the stellar parameters derived are not dissimilar from those derived for the single WC8 star WR135 by Dessart et al. (2000; see Table~\\ref{tab:wr11_summary}). This indicates that it is unlikely that substantial levels of emission are produced by the collision region or that an anomalous ionization is induced by the O star ionizing flux.  The derived spectroscopic luminosity ($\\log L/L_{\\odot} = 5.0\\pm0.1$) is increased with respect to that found in previous studies, partly due to line-line interaction, and is close to the luminosity derived in Paper~I via the mass-luminosity relationship for WR stars ($\\log L/L_{\\odot} = 5.2\\pm0.1$). An important implication of the higher luminosity and of a 10\\% clumping factor, is that it helps to reduce the performance number by a factor of $\\sim$20, from 144 to 7, assuming a 10\\% clumping filling factor. Lucy \\& Abbott (1993) showed that radiation pressure on spectral lines should in principle be able to drive winds with performance numbers $\\simeq$ 10. So, the strongly reduced $\\eta$ suggests that line driving is the mechanism responsible for the dense wind of Wolf-Rayet component in $\\gamma$~Velorum.  \\subsection{\\cmf\\ vs. \\isa}  Parameters derived with the fast Sobolev approximation code \\isa\\ are similar to those derived with \\cmf , apart from the derived carbon mass fraction, with \\cmf\\ resulting in a carbon abundance which is a factor of two larger than \\isa . This discrepancy is not totally understood although we should note that the co-moving frame code of Koesterke \\& Hamann (1995) compares favourably with \\cmf\\ for cases in common (e.g. for WR135 - Dessart et al. 2000).  A secondary difference is in the fluxes at $\\lambda$$<$400~\\AA . Crowther et al. (1999) tested the far UV model fluxes of \\cmf\\ and \\isa\\ by modelling the H~{\\sc ii} region associated with the WN8 star WR124, showing that, for that system, the \\isa\\ flux was too hard. Also in the current analysis we find that the \\isa\\ flux is harder for \\la $<$400~\\AA : we suspect that this discrepancy is due to the lack of photon branching in the Monte Carlo code, which would lead to distributing energetic photons to longer wavelengths. Although the ionising properties of the two synthetic atmospheres (the values of $Q_0$ and $Q_1$) are similar, the difference in the extreme UV can carry implications when model WR fluxes are used to interpret nebular line emission from extra-galactic stellar populations (Leitherer et al. 1999).  We stress that the comparison between the two codes is reasonable only if photon loss is included in \\isa\\ via the approximation detailed in Sect.\\ref{ssec:pholos}. If photon loss were not accounted for, the determined \\isa\\ effective temperature would be about 8000~K lower, comparable to the difference found in Sect.\\ref{sec:pholos} when reducing the interaction between metal lines and \\heii Ly$\\alpha$.  \\subsection{The photon loss effect}  The photon loss effect demonstrated by Schmutz (1997) as the cause of a lower wind ionization, was tested here using the \\cmf\\ code which treats line-line interaction as part of the blanketing. For this parameter space the interaction between the \\heii\\ Ly$\\alpha$ and nearby \\oiii\\ and \\ciii\\ lines causes a shift to lower ionization equilibrium, equivalent to $\\sim$10\\,000~K or a $\\sim$20\\% increase in the luminosity\\footnote{The quantification of the luminosity shift resulting from a different temperature depends on the location of the $R_\\ast$, since this can change from model to model, it is not possible to simply derive the change in luminosity from the change in temperature and the Boltzman relationship.}. By reducing the turbulent velocity from 50 to 5~km~s$^{-1}$ for lines residing near 303~\\AA , we have shown that the wind ionization is enhanced in a similar way, demonstrating that line-line interaction, and not continuum blanketing, is at the origin of the ionization shift observed when eliminating metal lines near the \\heii\\ Ly$\\alpha$. If line-line interaction between strong resonance lines and weaker but more numerous metal lines leads to a shift in the wind ionization, a detailed treatment of even the weakest lines is important. We remind the reader that our test remains incomplete because, although \\cmf\\ treats line-line interaction not all lines are accounted for. Amongst the ones that are {\\it not} treated there might be some that could play an important role in this mechanism. Additionally, our conclusion is only valid within the parameter space appropriate for WR11, a WC8 star, since for different temperature and mass-loss the strength of the \\heii\\ and other strong lines will be different (cf. Crowther et al. (1999) who found that photon-loss did not play a role in the atmosphere of their WN8 star).  One should conclude from the presence of the photon loss mechanism that a proper treatment of the metal lines not only requires one to include very many lines (to realistically model the blanketing), but also the very critical ones, such as those near the strong \\heii\\ Ly$\\alpha$ and possibly at other resonance lines such as \\civ\\ at 1550~\\AA . Furthermore, the presence of photon loss shows that applying mean blanketing factors (Pauldrach et al. 1996) - averaged over 20--50~\\AA\\ intervals - may miss out some interaction with an overall effect on the derived stellar parameters for WR stars. Finally, it is clearly pointed out that macro-turbulence is connected with photon loss, as the amount of turbulence essentially dictates the size and strength of the set of lines involved in the line-line interactions."
    },
    "0004/astro-ph0004123_arXiv.txt": {
        "abstract": "We present the first detections of the black hole X-ray binary GRS 1915+105 at sub-millimetre wavelengths. We clearly detect the source at 350 GHz on two epochs, with significant variability over the 24 hr between epochs. Quasi-simultaneous radio monitoring indicates an approximately flat spectrum from 2 -- 350 GHz, although there is marginal evidence for a minimum in the spectrum between 15 -- 350 GHz. The flat spectrum and correlated variability imply that the sub-mm emission arises from the same synchrotron source as the radio emission. This source is likely to be a quasi-steady partially self-absorbed jet, in which case these sub-mm observations probe significantly closer to the base of the jet than do radio observations and may be used in future as a valuable diagnostic of the disc:jet connection in this source. ",
        "introduction": "Since its discovery in 1992, simultaneously with the SIGMA instrument on the {\\it Granat} satellite and the BATSE instrument on the {\\it CGRO} satellite (Castro-Tirado, Brandt, Lund, 1992; Harmon, Paciesas, Fishman, 1992), GRS 1915+105 has been one of the most extensively studied sources of recent times.  The popularity of this source is due to its highly complex and variable nature at all wavelengths from gamma rays to radio.  The many unusual features include relativistic jets (Mirabel \\& Rodr\\'{\\i}guez 1994; Fender et al.\\ 1999), X-ray Quasi-periodic Oscillations (QPOs) (Morgan \\& Remillard 1996) and accretion instabilities resulting in jet formation (Pooley \\& Fender 1997; Eikenberry et al. 1998; Mirabel et al. 1998).  The system exhibits a wide variety of radio behaviour on timescales from minutes to weeks (Foster et al. 1996; Pooley \\& Fender 1997). High time resolution observations using the Ryle Telescope (RT) and more continuous coverage on hourly time-scales, revealed new aspects of the radio emission (Pooley \\& Fender 1997 hereafter PF97).  Both 20--40 min period QPOs associated with soft X-ray variations (first reported in the radio by Pooley (1995,1996)) and a change of QPO period, were clearly visible during some individual observations.  The infrared variability from this source is similar. Following observed 1 and 2 mag variations in the {\\it J}, {\\it H} and {\\it K}-band fluxes (Castro-Tirado et al.\\ 1993; Mirabel et al.\\ 1994; Chaty et al.\\ 1996), Fender et al.\\ (1997) reported rapid infrared flares which had amplitudes, rise, decay and recurrence time-scales strikingly similar to those of the radio flares observed with the RT 8 hours later, suggesting infrared synchrotron emission.  Further RT observations taken simultaneously with the William Herschel Telescope (WHT) showed 26-min oscillations in both the radio and infrared emission (Fender \\& Pooley 1998).  From these observations, the authors showed that the radio variations were delayed by 33 (or perhaps 59) minutes relative to the infrared.  Is there a correlation between the radio or infrared emission and the X-rays?  There is clearly some correlation between X-ray states and radio emission ({\\it CGRO} Burst And Transient Source Experiment (BATSE) correlation reported by Foster et al.\\ 1996, {\\it Rossi X-ray Timing Explorer} Proportional Counter Array ({\\it RXTE} PCA) correlation reported by PF97), and PF97 were also the first to observe frequency-dependent delays in the radio emission.  Quasi-periodic flaring in the infrared, reported by Fender \\& Pooley (1998), was also observed by Eikenberry et al.\\ (1998), who had simultaneous observations with {\\it RXTE} PCA.  The source of these flares is consistent with infrared emitting plasmons from {\\it ejected disc material}, following rapid X-ray variability and an X-ray dip.  The confirmation of such a hypothesis was presented by Mirabel et al.\\ (1998) with simultaneous {\\it RXTE} PCA, United Kingdom Infrared Telescope (UKIRT) and Very Large Array (VLA) data.  Several X-ray dips and subsequent infrared and radio flares were observed. ",
        "conclusions": "We have obtained the first sub-mm photometry of GRS 1915+105 at 350 GHz (850-$\\umu$m), together with radio data on two separate epochs during a time when the source was mildly active.  At all frequencies a significant amount of variability was observed, with a roughly flat spectral index, supporting the hypothesis that emission from radio to infrared wavelengths is dominated by synchrotron emission. Scaling with respect to both direct VLBI imaging and radio--mm--infrared time delays indicate that future sub-mm observations have the potential to probe very close to the base of the jet."
    },
    "0004/astro-ph0004409_arXiv.txt": {
        "abstract": "FUSE spectra of the white dwarf stars G191-B2B, GD~394, WD~$2211-495$ and WD~$2331-475$ cover the absorption features out of the ground electronic states of N~I, N~II, N~III, O~I and Ar~I in the far ultraviolet, providing new insights on the origin of the partial ionization of the Local Interstellar Medium (LISM), and for the case of G191-B2B, the interstellar cloud that immediately surrounds the solar system.  Toward these targets the interstellar abundances of Ar~I, and sometimes N~I, are significantly below their cosmic abundances relative to H~I.  In the diffuse interstellar medium, these elements are not likely to be depleted onto dust grains.  Generally, we expect that Ar should be more strongly ionized than H (and also O and N whose ionizations are coupled to that of H via charge exchange reactions) because the cross section for the photoionization of Ar~I is very high. Our finding that Ar~I/H~I is low may help to explain the surprisingly high ionization of He in the LISM found by other investigators.  Our result favors the interpretation that the ionization of the local medium is maintained by a strong EUV flux from nearby stars and hot gases, rather than an incomplete recovery from a past, more highly ionized condition. ",
        "introduction": "Our solar system is traveling through a low-density ($n_H\\sim 0.1-0.3\\,{\\rm cm}^{-3}$), warm ($T\\approx 7000$K), partially ionized gas cloud called the Local Interstellar Cloud (LIC)  (Lallement et al. 1996) with a diameter of about 4~pc  (Redfield \\& Linsky 2000). Near the LIC are similar clouds that give rise to absorption features with different velocities  (Gry 1996; Lallement 1996), known collectively as the Local Interstellar Medium (LISM).  From the available evidence, the physical conditions in the LISM clouds are only sightly different from those of the LIC  (Linsky 1996; Ferlet 1999).  Presently, we do not have a clear understanding of the physical processes that are responsible for high fractional ionization of helium in the LISM.  From observations that $\\langle n({\\rm He~I})/n({\\rm H~I})\\rangle=0.07$  (Dupuis et al. 1995), along with the knowledge that ${\\rm He/H}=0.10$ in all forms, it is evident that He is usually slightly more ionized than H.  Also, helium shows less variability in how strongly it is ionized from one region to the next  (Wolff, Koester, \\& Lallement 1999). While the EUV radiation from white dwarf and other stars can explain the observed fractional ionization of H  (Vallerga 1998), these sources do not produce enough photons with energies above 24.6~eV to explain the ionization of He.  To resolve this problem, two main proposals have been offered.  One is that the LISM is not in a steady-state condition; it is returning from a much more highly ionized state produced by an energetic event in the recent past ($t\\lesssim 10^6\\,{\\rm yr}$), such as the flash from a nearby supernova  (Reynolds 1986; Frisch \\& Slavin 1996) or its shock wave  (Lyu \\& Bruhweiler 1996).  The other possibility is that the ionization of He is maintained in a steady state by the diffuse EUV radiation arising from conductive interfaces between the cloud edges and the surrounding, hot medium at $T\\sim 10^6$K  (Slavin 1989; Slavin \\& Frisch 1998), or, alternatively, by recombination radiation from highly ionized but cooled gases that might surround us  (Breitschwerdt \\& Schmutzler 1994).  Sofia \\& Jenkins  (1998) have proposed that the apparent abundances of the neutral forms of N, O and Ar, elements that are generally very lightly depleted in the interstellar medium (if at all), can help to unravel the mystery of the He ionization.  UV absorption lines in the spectra of white dwarf stars present good opportunities for studying these abundances in the LISM, since most of the brighter objects are well removed from major gas clouds elsewhere and remain immersed within a ``Local Bubble'' (diam. $\\sim 200$~pc) of very hot, low density material surrounding the nearby cloud complex  (Cox \\& Reynolds 1987; Sfeir et al. 1999).  We report here on initial spectra of four white dwarf stars taken with the Far Ultraviolet Spectroscopic Explorer (FUSE) (Moos et al. 2000; Sahnow et al. 2000) taken at a resolution of $20-25\\,{\\rm km~s}^{-1}$ (\\S\\S\\ref{obs}$-$\\ref{results}), and we interpret the findings in the light of the Sofia \\& Jenkins proposal to obtain a better understanding on how the LISM is ionized (\\S\\ref{interpretation}). ",
        "conclusions": ""
    },
    "0004/astro-ph0004159_arXiv.txt": {
        "abstract": "In recent years a handful of systems have been observed to show ``negative'' (nodal) superhumps, with periods slightly {\\it shorter} than the orbital period.  It has been suggested that these modes are a consequence of the slow retrograde precession of the line of nodes in a disk tilted with respect to the orbital plane.   Our simulations confirm and refine this model: they suggest a roughly axisymmetric, retrogradely-precessing, tilted disk that is driven at a period slightly less than {\\it half} the orbital period as the tidal field of the orbiting secondary encounters in turn the two halves of the disk above and below the midplane.  Each of these passings leads to viscous dissipation on one face of an optically-thick disk --- observers on opposite sides of the disk would each observe one brightening per orbit, but $180^\\circ$ out of phase with each other. ",
        "introduction": "Cataclysmic variable (CV) systems in the mass-ratio range $q\\equiv M_2/M_1\\lesssim0.33$ (e.g., Murray 2000) are observed to show the disk oscillations known as superhumps (see Warner 1995 and Patterson 1999 for recent reviews).  First observed in the SU Ursae Majoris stars, {\\it common} (a.k.a. {\\it positive} or {\\it apsidal}) superhumps have periods a few percent in excess of their system orbital periods.  The story and classification scheme, however, have in recent years become more complicated.  First, some old novae and nova-likes show ``permanent'' positive superhumps indicating a high and roughly constant mass transfer rate, e.g., V603 Aquilae (Patterson et al.\\  1993, 1997) V795 Herculis (Patterson \\& Skillman 1994), H 0551-819 (Patterson 1995), MV Lyrae (Skillman, Patterson, \\& Thorstensen 1995), and V1974 Cygni (Retter, Leibowitz, \\& Ofek 1997; Skillman et al.\\ 1997).  Second, a few systems, including V603 Aql, H 0551-819, V503 Cyg (Harvey et al. 1995), V1159 Orionis (Patterson et al.\\ 1995), TT Arietis (Skillman et al.\\  1998; Kraicheva et al.\\ 1999), AM CVn (Skillman et al.\\ 1999) display so-called {\\it negative} (or nodal) superhumps, which have periods a few percent {\\it shorter} than the orbital period.  The favored model explaining positive superhumps is that for systems in the range $0.03 \\lesssim q \\lesssim 0.33$ with high mass transfer rates and viscosities, the outer disk can expand to the radius near the 3:1 eccentric inner Lindblad resonance (just outside the 3:1 co-rotation radius) where it can be driven periodically by the tidal field of the orbiting secondary (Whitehurst 1988, 1994; Hirose \\& Osaki 1990; Lubow 1991a,b; Murray 1996).   Numerical simulations by Simpson \\& Wood (1998) show that over one oscillation period, as viewed face-on, the disk is pulled into a distorted, elongated shape followed by relaxation back to a nearly circular shape, and that during this oscillation, the spiral shock pattern advances $180^\\circ$ in the slowly precessing frame (n.b., spectroscopic time-series of an eclipsing superhumper may reveal this motion).  Simulations further suggest that viscous energy dissipation through the cycle is the dominant source of the observed photometric oscillations and is highest when the disk is maximally distorted and the streamlines are most non-circular.  Because both modes have been observed {\\it simultaneously} in several of the systems listed above (e.g., V603 Aql, V1159 Ori, H 0551-819, and V503 Cyg), the physical origin of the negative superhumps must be distinct from the above mechanism.  The model suggested by Harvey et al.\\  (1995) invokes the retrograde precession of the {\\it line of nodes} in an accretion disk {\\it tilted} with respect to the orbital plane (see also Bonnet-Bidaud, Motch, \\& Mouchet 1985 and Barrett, O'Donoghue, \\& Warner 1988).  Having identified physical processes associated with each superhump mode, we hereafter follow Harvey et al.\\ (1995) and refer to positive and negative superhump modes as apsidal and nodal, respectively, thereby avoiding the awkward term ``negative period excess.''  The linear theory of tilted disks (Papaloizou \\& Pringle 1983; Papaloizou \\& Terquem 1995) suggests that for the disk to precess as a solid body, the sound crossing timescale must be shorter than the precession timescale, a condition that is well satisfied in our simulations. Larwood et al.\\ (1996) studied the general dynamics of tilted disks using the method of smoothed particle hydrodynamics (SPH), and found that for a constant surface density $\\Sigma$ disk tilted at angle $\\delta$ with respect to the orbital plane, the nodal precessional rate is given by \\begin{equation} {\\omega_p\\over\\Omega(R)} = -{15\\over32}q\\beta^3\\cos\\delta \\label{eq: prec} \\end{equation} where $\\beta\\equiv R_d/a$ is the disk radius in units of the orbital separation, and where a ratio of specific heats $\\gamma=5/3$ was assumed.  For typical values, $q=0.25$, $\\beta=0.4$, $\\Omega(R)=3\\Omega_{\\rm orb}$, equation \\ref{eq: prec} gives $\\sim$40 orbits per precession period, which is in factor-of-two agreement with observations and our simulations below.  In this work we report the preliminary results of a large parametric study (Montgomery, Wood, \\& Simpson 2000, in preparation) exploring the dynamics of tilted disks and the origin of the nodal superhump phenomenon.  Our primary result is that the model of a tidally-driven retrogradely-precessing disk is basically confirmed, but that the driving frequency is slightly higher than twice per orbit as each of the two out-of-orbital-plane halves of the disk are perturbed sequentially and out of phase with each other.   Because we receive photons from only one face of an optically-thick disk in superoutburst, we {\\it measure} a photometric period which is just shorter than the orbital period.  This model predicts that the waveform observed from a tilted disk that is optically-thin or viewed from high inclination should be approximately doubly-periodic each orbit. ",
        "conclusions": "Although these are preliminary results, it appears the tilted-disk model neatly explains many of the observational characteristics of nodal (``negative'') superhumps, and also suggests further observational and numerical tests.  Our findings here can be summarized as follows.  \\begin{itemize}  \\item The results are consistent with the nodal regression model for the origin of negative superhumps.  We tilted a disk both before and after the development of common apsidal superhump oscillations, and find power at twice the expected (observed) nodal superhump frequency. In superoutburst the disk is optically thick and so we observe only one face and hence only a single nodal mode brightening event per orbit --- the other occurs on the opposite (hidden) face of the disk and is $180^\\circ$ out of phase.  This model predicts that optically thin and/or edge-on tilted disks should display power at $2\\nu_n$.  \\item In disks oscillating in both nodal and apsidal modes simultaneously, there is also power at the combination frequencies $2\\nu_n-\\nu_a$ and $2\\nu_n+\\nu_a$.  These combination modes have not yet been identified in nature, but if demonstrated would provide additional support for this model.  \\item In our simulations, where we keep the particle number constant, the amplitude of the nodal mode decreases on a timescale of $\\sim$1000 orbits (several weeks if scaled to a physical orbital period).  This is consistent with the observations, where the nodal superhump modes appear to have a much longer decay time than the apsidal modes.  \\item The origin of the disk tilt remains an open question. Simply displacing the $L_1$ point out of the orbial midplane does not appear to be sufficient to generate a tilted disk, but more simulations are required.  \\end{itemize}"
    },
    "0004/astro-ph0004190_arXiv.txt": {
        "abstract": "We have made the first interferomeric measurements of the wavelengths of the important ultraviolet diagnostic lines in the spectra \\ion{C}{4} near 155\\,nm and \\ion{Si}{4} near 139\\,nm with a vacuum ultraviolet Fourier transform spectrometer and high-current discharge sources. The wavelength uncertainties were reduced by one order of magnitude for the \\ion{C}{4} lines and by two orders of magnitude for the \\ion{Si}{4} lines. Our measurements also provide accurate wavelengths for resonance transitions in \\ion{Al}{3}, \\ion{Al}{2}, and \\ion{Si}{2}. \\vspace{5mm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004229_arXiv.txt": {
        "abstract": "We present the results of the classification of spectral activity types for 193 galaxies from a new sample of 49 compact groups of galaxies in the southern hemisphere (SCGs).  The SCGs is a new sample of compact groups selected in automated fashion from a digitized galaxy catalogue, covering an area of $\\sim 5200 ~sq~deg$ around the South Galactic Pole. It is complete up to m$\\sim 14.5$ in $b_j$ for the brightest galaxy of the group. This sample is very well suited for statistical studies of compact groups properties as it is unaffected by the biases introduced by visual selection methods.  The spectral analysis of the SCG galaxies confirms the results previously obtained from the observation of a smaller sample of 17 CGs in Hickson's sample (HCG). We confirm the luminosity--activity and morphology--activity relations previously observed in HCGs: 1) low luminosity AGNs and AGNs are preferentially located in the most luminous galaxies in the groups, while non--emission--line galaxies and star--forming galaxies share a common distribution among the less luminous galaxies; 2) the non--emission--line galaxies, the low luminosity AGNs and AGNs are nearly all in early--type galaxies, while the star--forming galaxies are preferentially located in late--type spirals.  We also verified that the number of ``evolved'' galaxies (early--type non--emission--line galaxies) significantly increases with the number of members in the group. Finally, we confirm that AGNs (including low luminosity or dwarf AGNs) are the most frequent (41\\%) activity type encountered in CGs.  The SCGs contain more star--forming galaxies and less non--emission--line galaxies than HCGs. This difference is possibly related to the different selection criteria of SCGs with respect to HCGs. Being selected in an automated way, the SCGs probe a wider range of physical properties than the HCGs, which are biased towards more easily detectable, denser and therefore more evolved groups.  The star--forming galaxies in SCGs are composed of few starburst galaxies and a higher number of \\ion{H}{2} Nucleus Galaxies, which generally have less intense star formation than starburst galaxies. As a consequence, the star formation activity in SCGs is remarkably low. At the same time, we found clear evidence for nuclear activity in all the star--forming galaxies.  Adding up the percentage of AGNs and star--forming galaxies showing evidence of nuclear activity would result in more than 70\\% of the galaxies in SCGs having an active nucleus.  If these results are further confirmed, CGs will be in the local Universe the best location where to find AGNs. Curiously, however, this characteristic of CGs generally excludes Seyfert 1 galaxies. ",
        "introduction": "Studies on Compact Groups (CGs) of galaxies mainly concentrate on two issues: 1) what is the origin and relative importance of CGs in the Universe? 2) is there a relation between the makeup of these structures and the formation and evolution of galaxies which form them?  Actually, the possibility that the two issues are connected is what makes CGs so interesting in the first place (see the review by Hickson 1997).  Now that we have some unambiguous evidence of galaxy evolution driven by dynamical environment effects, even at relatively modest redshifts (\\eg O'Connell 1999, VanDokkum et al. 1999), CGs with their high spatial densities and small velocity dispersions look more than ever like the ideal objects to investigate how dynamical effects forge galaxy evolution.  Although the first dynamical simulations of CGs predicted that they should evolve within only a few orbital times into a system dominated by one central merger remnant (Barnes 1989), evidence for on--going or post merger events in CGs are scarce (Mardirossian et al.\\ 1982; Zepf \\& Whitmore 1991; Zepf 1993; Moles et al.\\ 1994).  The high number of these systems also suggests longer lifetime for CGs, as compared to what is expected based on the small crossing times of the galaxies. But, these contradictions may only be apparent, due to our still limited knowledge of how mergers really proceed and what kind of remnants they leave behind (see the review by Mihos 1999). This is the point made by several recent numerical simulations, which show that the evolution of CGs may take more than a few Gyrs if they are embedded in massive dark matter halos (Athanassoula et al. 1997) or even longer if the merging galaxies had originally a wide range in masses (Governato et al. 1991).  Our difficulties in understanding CGs may also be a consequence of our lack of knowledge of what these structures really are. Short dynamical lifetimes and a high number of CGs, for instance, are not a problem if these structures are collapsing physical systems continuously forming in loose groups (Diaferio et al.\\ 1994; Governato et al. 1996). This seems to be the converging conclusion of many recent redshift studies of galaxies around compact groups (Garcia 1995; Ribeiro et al.\\ 1998; Barton et al.\\ 1998). In Ribeiro et al.\\ (1998), in particular, it was shown that amongst CGs there is a wide range of different dynamical stages, with different surface density profiles; the core radius increasing from values typical of very compact to loose groups. These authors confirmed that the fraction of spiral galaxies is lower in CGs than in the field and that it decreases as the velocity dispersion increases (Hickson et al.\\ 1988), suggesting that the environment influenced the formation and/or evolution of the galaxies in the groups. In these structures, interactions and mergers are suggested by the decrease of spiral fraction with the decrease of crossing time and increase of X-ray luminosity (Ponman et al.\\ 1996). The linear relation between velocity dispersion and gas temperature also seems to get steeper in CGs than in cluster, which is consistent with effects of dynamical friction (Bird et al. 1995).  The dynamical evidences in favor of environmental effects, as described above, are consistent with our previous analysis of activity types found in HCGs (Ribeiro et al.\\ 1996; Coziol et al.\\ 1998a, 1998b; de Carvalho \\& Coziol 1999). We have finds that a significant fraction of the brightest galaxies in the groups displays some sort of activity, either star formation or AGN. We also discovered the presence of a large number of low luminosity AGNs, making AGN the dominant activity type in CGs (Coziol et al.\\ 1998a). Related to this activity, we find a density--morphology--activity relation, according to which the AGNs (low luminosity AGNs included) and galaxies without emission lines are the most luminous and invariably early-type galaxies and tend to concentrate in the core of HCGs, while the star--forming galaxies, which may also be preferentially late--type spirals, are more numerous in the halo (Coziol et al.\\ 1998b).  Despite looking like normal field galaxies with similar morphologies, galaxies in CGs seem to have stellar metallicity that is relatively high for their luminosity. They also have slightly narrower metal absorption line equivalent widths and relatively strong Balmer absorption lines which suggests that a meager extra population of intermediate--age stars is still present (Coziol et al.  1998b).  This phenomenon may be reminiscent of the ``post--starburst'' phase associated to the Butcher--Oemler effect (Belloni et al.\\ 1995; Caldwell \\& Rose 1997). But, the ``post--starburst'' signs in CG galaxies look weaker than usually observed (Poggianti \\& Barbaro 1996; Zabludoff et al. 1996; Leonardi \\& Rose 1996), which suggests a more advanced stage of evolution (or milder level of past star--forming activity) for the galaxies in the groups (Caon et al.\\ 1994; Coziol et al.\\ 1998b).  Considering the significance of our results for understanding how the environment of galaxies affects their formation and evolution, we believe that it is important to confirm our previous observations on the activity in CGs on a larger and complete sample.  This is the main subject of the present article.  In Section~2, we present the new sample of CG galaxies used for our analysis and give some details on their spectroscopic observations.  In Section~3, we discuss our classification method and present the results. A summary of our study and conclusions can be found in Section~4. ",
        "conclusions": "We have presented the results of the classification of the level of activity of 193 galaxies in a new sample of 49 CGs of galaxies observed in the southern hemisphere. This analysis confirms the results previously obtained from a smaller sample of 17 HCGs.  \\begin{itemize}  \\item We show that AGN (including low luminosity AGN) is the most frequent (41\\%) activity type encountered in CGs. We found however more SFGs and less non--emission--line galaxies in the SCGs than in the HCGs. We suggest that these differences could be explained by the visual biases affecting the HCG sample. The SCGs sample, filling completely the parameter space allowed by the criteria used for the search, contain groups with a broader range of properties than the HCGs sample, and therefore groups whose properties (\\eg spatial density) are less extreme than those of HCGs.  \\item We confirm the luminosity--activity and morphology--activity relation previously observed in HCGs: 1) low luminosity AGNs and AGNs are preferentially located in the most luminous galaxies in the groups, while non--emission--line galaxies and SFGs share a common distribution among the less luminous galaxies; 2) the non--emission--line galaxies, the low luminosity AGNs and AGNs are more common in early--type galaxies while SFGs are preferentially located in late--type spirals.  \\end{itemize}  We also observed some new phenomena:  \\begin{itemize}  \\item We find a strong variation of the dominant activity type with the increase of galaxies members in the group: the number of non--emission--line galaxies increases while the number of SFGs decreases. But what truly characterizes CGs seems to be the association of one or more SFGs with more evolved or active galaxies (low luminosity AGNs, LINERs or Seyfert 2). It is quite rare to find a group composed only of SFGs, which suggests that what we are looking at is really a physical entity, and not just a projection effect within loose groups.  \\item The star--forming galaxies in CGs are mostly \\ion{H}{2} Nucleus Galaxies. As these galaxies are characterized by having low star formation activity (Coziol 1996), we conclude that the star formation level in CGs is generally low. This conclusion is consistent with the recent results on star formation obtained by Severgnini et al. (1999) based on an H$\\alpha + [{\\rm H II}]$ imaging survey in 31 HCGs. Comparing the distribution in H$\\alpha +$ [\\ion{N}{2}] EW measured by these last authors with the distribution presented in Coziol (1996) for the \\ion{H}{2} Nucleus Galaxies (multiplying by 0.66 to 0.5 to take into account the unusual strengths of the nitrogen lines in these galaxies), one can see that they are fully consistent.  \\item We also see in the \\ion{H}{2} Nucleus Galaxies a clear evidence for non stellar activity.  If the transition nature of these galaxies is taken as a sign of AGN activity, then CGs would be exceptionally rich in AGNs. More than 70\\% of the galaxies in the SCGs would have an AGN nature (adding the SFGs to the AGNs and LLAGNs). The AGN nature of these galaxies, however, must be confirmed, as alternatives, like shocks produced by supernovae or excess of excited gas in the disk could also be valid.  \\end{itemize}  To understand which mechanisms in CGs produce the phenomena described above is not an easy task.  As was already observed (Coziol 1998a; 1998b), we cannot relate directly the high density in compact groups of galaxies to an increase (or decrease) of star formation or to an increase of nuclear activity.  What we observe instead is a sort of density--morphology--activity relation.  The reason for this multiple relation may have something to do with the history of galaxies in CGs and the mechanisms by which these systems form.  The dynamical analysis of Ribeiro et al.  (1998) and Barton et al. (1998) reveal that most of the compact groups are some sort of substructure of larger structures.  These substructures are not all at the same dynamical evolutionary stage.  Some may already be in dynamical equilibrium while others are still replenished in spiral galaxies, which fall from the larger structure where they first formed (Diaferio et al.  1994; Governato et al.  1996). This is specially true for SCGs, where the spread in physical properties is bound to be larger than that of the HCGs.  According to the scenario of sub--clustering in loose groups (Diaferio et al. 1994), during the expansion and collapse of gravitating systems subclumps form and merge in a hierarchical manner. The subclumps are often bound and can be identified with compact groups. In their models, Diaferio et al. 1994 have found that most of the subclumps contain 2 or 3 galaxies. The frequency for groups with 4 and 5 galaxies are, at maximum, 34\\% and 12\\%, respectively. These numbers are comparable to what we find: 46\\% and 13\\% SCGs have 4 and 5 members, respectively. However, to be consistent with the simulations done by Diaferio et al. (1994), the loose groups where SCG would have formed would need to be richer in galaxies. Loose groups with 15 and 30 galaxies are necessary to reproduce our observations. In their study of poor groups with diffuse X-ray, Zabludoff \\& Mulchaey (1998) find a number of galaxies between 20 and 50. Barton et al. (1998), on the other hand, have found more heterogeneous environments, the number of galaxies in their sample varying between 5 members up to the population of rich clusters of galaxies. It is not clear which results are representative of the environment of CGs in general, but it seems that the Diaferio et al. scenario of sub--clustering could a good working hypothesis.  However, there could be a serious problem with the sub--clustering scenario, in the relatively short lifetime it predicts for these systems. After $\\sim 1$Gyr the galaxies should all merge. In the Diaferio et al. simulations, CGs are currently forming and, consequently, are not sufficiently evolved to contain a large number of merger remnants. This in contradiction with the advance level of evolution of the galaxies usually found in the core of CGs (Coziol et al. 1998a; 1998b). This fast merger scenario is also in contradiction with our observation of the variation of activity status of the galaxies with the number of group members: the groups which have a high number of galaxies are also the ones which are more evolved and less active. Furthermore, while the Diaferio et al. model correctly predicts that the more massive galaxies form the CGs, it does not explain how these galaxies were allowed to grow in mass. It seems reasonable to assume that mergers must have taken place sometime in the evolution of these structures to produce both the more massive and early--type galaxies necessary to form the core of the groups. But, as we have emphasized before, signs of mergers are generally scarce, and in the more evolved galaxies in the groups star formation activity has probably ceased more than 2 Gyrs ago (Caon et al. 1994; Coziol et al. 1998b). It seems, therefore, that another parameter is missing in our present history of the formation of these structures which would allow galaxies to survive merging.  The replenishment scenario (Governato et al.  1996) offers one simple solution to the fast merger problem. According to this scenario, if CGs are so conspicious it is because they are continuously replenished by galaxies falling into their core. Now, if we assume that the infalling galaxies are preferentially late--type SFGs, we would then expect to find a high number of these galaxies in each group if replenishment is important. Moreover, these infalling galaxies should be associated preferentially with galaxies with more evolved characteristics. We have verified that late--type SFGs are indeed present in 80\\% of the groups. We have also shown that they are preferentially associated to other activity types, in particular LLAGNs and no--emission line galaxies which are the more evolved galaxies in CGs. This supports the idea that replenishment is an important aspect of the compact group phenomenon (Governato et al. 1996).  According to the replenishment scenario, the SFGs would be ``now'' falling into the core of the groups.  By doing so, they will interact with the group as a whole, but possibly also with individual galaxies. For CGs, the models describing these interactions should explain both the small amount of star formation observed in these galaxies and the activation of non--thermal process.  Simulations of low--velocity galaxy--galaxy interaction have shown that the majority of galaxy--galaxy encounters lead to normal--star forming properties and that the triggering of very active starbursts are rare events (Mihos, Richstone \\& Bothun 1991; 1992).  In some of the closest encounters star formation could even be inhibited as a result of mass loss from the disk (by tidal tails) or perturbations that increase the disk gas scale height (Mihos, Richstone \\& Bothun 1991). Only mergers, seem to be able to increase significantly the star formation rate in interacting galaxies (Mihos, Richstone \\& Bothun 1992). At first, these results seem consistent with the low star formation observed in CGs. The small number of starburst galaxies in these systems could also be the result of rare mergers, as it is clearly suggested for some galaxies in HCG 16.  According to the simulations, however, low--velocity galaxy--galaxy interactions would systematically produce morphological evidences (Howard et al.  1993).  Although we do see peculiarities in some of the SFGs, morphological distortions are not always present. According to simulations, it is also very difficult to induce star formation in the nucleus of galaxies by galaxy--galaxy interaction. Even more so to move the gas toward the center to feed a possible AGN.  Considering that gravitational forces are additive, it is probably more realistic to assume that a galaxy which is falling into the core of the group will first interact with the potential of the group as a whole. In the literature, one can find many models explaining how environmental effects may change the evolution of galaxies falling into rich clusters. Some of the mechanisms proposed in these models may also apply to galaxies in CGs. In some aspects, CGs do resemble rich clusters of galaxies. The density of galaxies in the core of CGs, for instance, is comparable to those in clusters. On other aspects, however, CGs are quite distincts. The galaxies in CGs, in particular, have smaller velocity--dispersions and there is much less galaxies in number. The differences should limit the efficiencies of the possible mechanisms in rich clusters and produce some notable differences.  The effect of tidal interaction between cluster gravitational potential and a disk galaxy was studied by Merritt (1982; 1983), Byrd \\& Valtonen 1990, Henriksen \\& Byrd (1996) and more recently Fujita (1998).  According to Merritt (1982; 1983) a galaxy falling into a rich cluster will have its outer parts rapidly stripped away by the tidal field. But, before tidal stripping becomes effective, the perturbation produced by the tidal field can also trigger an inflow of gas toward the center of the galaxies which may either feed an AGN (Byrd \\& Valtonen 1990) or start a burst of star formation (Henriksen \\& Byrd 1996; Fujita 1998).  Using a simple point--mass model, Byrd \\& Valtonen (1990) have shown that infall of matter happens long before tidal stripping in the disk becomes effective. Therefore, gas rich spiral falling into the group will become active long before reaching the center.  The consequences of this model are interesting. Among other things, it predicts that CGs should have a halo of SFGs or even starburst galaxies around them. When the spiral galaxies enter the core of the group, tidal stripping becomes effective and the disks loose part of their gas (Hutchmeieir 1997). Tidal stripping will not touch the molecular gas which have formed in the center of the galaxy (Verdes--Montenegro 1997; Leon et al. 1998), but it will stop the infall of gas and decrease the star formation in the disk. The galaxies may then resemble an \\ion{H}{2} Nucleus Galaxies, with a small residual of star formation and non thermal activity in their nucleus.  If the source of the non thermal activity is the feeding of an AGN in the center of the galaxies, the question which naturally arise is why is the formation of Seyfert 1 inhibited in CGs?  One reason may be because the infall of gas is followed automatically by star formation. The AGN would then have to compete with star formation for the gas (de Carvalho \\& Coziol 1999). When star formation decreases, a low luminosity AGN then becomes apparent.  Tidal triggering added to tidal gas stripping will combine their effects to transform late type spirals into early--type ones. Triggered star formation in the nucleus of the galaxies increase their bulge in mass and light, while gas stripping makes them loose their disk. Although such a scenario seems viable, an important caveat is the fact that the total mass involved in CGs are small compared to those of rich clusters, making the efficiency of these transformations quite low, which could explain why a significant fraction of CGs are still dominated by spiral galaxies.  Other mechanisms could be important in rich cluster of galaxies but probably do not take place in compact groups. These are ram pressure gas stripping and galaxy harassment. Ram pressure gas stripping should be less efficient in CGs, since it increase as $\\sim L_x^{1/2} v_cg^2$. Although CGs may exhibit X-ray luminosities similar to clusters (Ponman et al. 1996), the velocity dispersions are much smaller. Galaxy harassment (Moore, Lake \\& Katz 1998; Lake. Katz \\& Moore 1998), on the other hand, needs frequent collisions at high speeds, which is difficult to achieve in CGs, where galaxies have lower velocity and there are also much fewer galaxies to interact with. Comparison of the evolution of CGs galaxies with those in rich clusters may allow to distinguish the specific signatures of these two particular mechanisms in rich cluster of galaxies."
    },
    "0004/astro-ph0004137_arXiv.txt": {
        "abstract": "We propose a new method to constrain the location of microlensing objects using EAGLE (Extremely Amplified Gravitational LEnsing) events. We have estimated the rate of EAGLE events by taking the finite-source effect into account. We found that the EAGLE event rate for using a 1-m class telescope whose limiting magnitude is $V \\sim 21$ is the same as or higher than that of the ordinary microlensing events which have been found to date. We have also found that the fraction of transit EAGLE events is large enough to detect: between $4 \\sim 80 \\% $ depending on the lens location. Since the lens proper motion can be measured for a transit event, one can distinguish whether the lens is a MACHO (MAssive Compact Halo Object) in our halo or one of the known stars in the Large Magellanic Cloud (LMC) from the proper motion measurement for each transit EAGLE event. Moreover, we show that the fraction of transit EAGLEs in all EAGLE events significantly depends on the lensing locations: the transit EAGLE fraction for the self-lensing case is $2 \\sim 15$ times larger than that for halo MACHOs. Thus, one can constrain the location of lens objects by the statistics of the transit events fraction. We show that we can reasonably expect $0 \\sim 6$ transit events out of 21 EAGLE events in 3 years. We can also constrain the lens population properties at a greater than $99 \\%$  confidence level depending on the number of transit events detected.  We also present the duration of EAGLE events, and show how an hourly observational mode  is more suitable for an EAGLE event search program. ",
        "introduction": "Several groups have been engaged in gravitational microlensing observations toward the Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC) in order to investigate the nature of MAssive Compact Halo Objects (MACHOs) in the Galaxy's halo. Until now, 8 microlensing event candidates have been found toward the LMC (\\cite{alc97}), and 2 candidates, including one possible binary lens event, have been found toward the SMC (\\cite{afo98}; \\cite{afo99}; \\cite{alb99}; \\cite{alc99a}).  The microlensing light curve informs us only of the event duration $\\hat{t}$, in which the mass, velocity and distance of the lensing object are degenerate, and thus it is difficult to determine the mass or the distance of the lensing object for each microlensing event. Therefore, the nature of MACHOs in the Galactic halo still remains unclear, and one cannot rule out the possibility that the microlensing events are not caused by MACHOs in the Galaxy's halo but by unknown populations lying between the Galaxy and the Clouds (e.g., Zaritsky \\& Lin 1997; Zhao 1998), or by normal stars in the Clouds themselves (e.g.,Sahu 1994).  In order to understand the nature of lensing objects, some additional information that partially or fully break the three-fold degeneracy is required. Over the past years several investigations have been made to extract additional information from special types of microlensing events. Examples of which are binary events and parallax events, for which the lens proper motion and lens distance can be constrained (e.g., Schneider \\& Weiss 1984; \\cite{alc95}; Miyamoto \\& Yoshii 1995). A transit event, in which the lens transits the extended surface of the source star, is another example of such types of microlensing events (Gould 1992, 1994; \\cite{nem94}; \\cite{wit94}; \\cite{pen97}). When the lens passes over the surface of the source star, the point source approximation fails and the peak of light curve deviates from that for a point source. From such a light curve, one can derive the radius of the source star scaled by the Einstein ring radius projected onto the source plane. Since the radius of the source star can be estimated from its spectrum, one can obtain the proper motion of the lens for transit events. The proper motion of the lens will be of great use in investigating the lens location, as the expected proper motion distribution of objects in the Galaxy's halo differs significantly from that in the Clouds. Unfortunately however, the probability for transit events is expected to be extremely small ( $\\sim 0.1\\% $) for normal microlensing events (\\cite{gou94}).  The probability for transit events is relatively high for EAGLE  events, in which an invisible star is amplified above the detection limit, because EAGLE events are guaranteed to have a small impact parameter (\\cite{wan96}; \\cite{gou97}; \\cite{nak98}). Since there are many more faint stars below the detection limit than visible stars that are usually used for microlensing search, the EAGLE event rate is expected to be fairly high (\\cite{nak98}). However EAGLE events have rarely been seen in the current observation programs in which the brightness of visible stars are measured based on a DoPHOT-type analysis. Recently however, a new CCD photometry method called {`}image subtraction{'} has been developed  (\\cite{ala98}, 2000; \\cite{alc99b}; I.A. Bond 1999, private communication). This method, which compares an exposure frame directly with a reference frame, is much more powerful for detecting EAGLE events than the current monitoring of visible stars, and one can expect that plenty of EAGLE events will be detected with this method in the near future. If many EAGLE events are detected, one may be able to find several transit events because EAGLE events have small impact parameters. Therefore, a search for EAGLE events will possibly provide a new opportunity to measure the lens proper motion in transit events, and thereby to constrain strongly the lens location. Furthermore, statistics of transit EAGLE events alone may be useful for studying the nature of lensing objects. The fraction of transit events depends on the ratio of the physical radius of the source star $R_*$ to the Einstein radius $r_E$ in source plane, i.e. $R_* / r_E$. In the case where the lenses are the stars in the LMC itself (self-lensing), the Einstein radius $r_E$ in source plane is much smaller than in the case where the lenses are MACHOs in the Galactic halo. Thus the fraction of transit EAGLE events out of all EAGLE events for self-lensing in the LMC will be much larger than that for lensing due to MACHOs in the halo. Therefore the fraction of transit EAGLE events may be used to constrain the location and the nature of lens.  EAGLE events will thus become one of the most powerful tools for investigating the nature of lensing objects. For these reasons, in this paper we investigate the properties of EAGLE events, and present a detailed calculation of the probability for transit EAGLE events. In section $\\ref{sec:basic}$ we briefly summarize the basic equations, and in section $\\ref{sec:event}$ we estimate the event rate of EAGLEs. In section $\\ref{sec:ratio}$ we calculate the fraction of transit events, and in section $\\ref{sec:duration}$ we derive the event duration of EAGLEs. Discussion on the observation strategy for EAGLE search is made in section $\\ref{sec:disc}$. ",
        "conclusions": "\\label{sec:disc}  We have seen that the EAGLE event rate is as high as that for normal events. Since the duration of an EAGLE event is usually short (1 day $\\sim$ 40 days), the nightly observation program currently undertaken by most groups is not adequate. Hourly observation is necessary for finding EAGLE events. Moreover, monitoring with a 1-m class telescope is sufficient to search for EAGLEs and to make follow up observations near the peak, but not enough to follow up the wing of the light curve. This is because the source is visible only at the peak of the light curve but invisible during most of the remainder of the event. To get overall structure of the light curve, one has to observe the event with a larger telescope for a longer period (but we do not need to observe so frequently).  Current microlensing experiments are being carried out toward very dense stellar fields. So, the determination of the absolute flux of the source has been an issue in microlensing because of source star blending. In the case of EAGLE events, the source stars are fainter than the detection limit and thus the blending of brighter stars is significant, making the problem worse. Attempts have been made to correct the blending effects for individual events by introducing an additional parameter, the blended flux, but this method suffers from very large uncertainties in the derived lensing parameters due to degeneracies among the parameters (\\cite{ala97}; \\cite{han97}). One solution to correct for blending problem is to use the  Hubble Space Telescope (HST). The high resolving power of the HST and the color information from ground-based observations enable us to uniquely separate the lensed source star from blended stars (\\cite{han97}). The image subtraction method (\\cite{ala98}, 2000; \\cite{alc99b}; I.A. Bond 1999, private communication) should be used to locate the centroid of the lensed flux in the ground-based event image accurately to find the lensed star in the HST image. Other high resolution telescopes like the VLT are also good candidates for performing follow-up observations. Using a large telescope is also required for spectroscopic observations of the source star, from which one can estimate the source radius. We can do this even after the event is finished.  Using the image subtraction method has another advantage in microlensing experiments. The color change induced by the limb-darkened extended source helps the measurement of the proper motion of the lens and increases the possible number of proper motion measurement events (\\cite{wit95}; \\cite{gou96}). Han (\\cite{han00}) found the color change measurement by using the image subtraction method enables one to obtain the same information about the lens and the source star, but with significantly reduced uncertainties due to the absence of the blending effect. So measuring the color change with differential photometry helps to derive the proper motion of the lens along with the other microlensing parameters, and at least, inform us whether a lens transited the surface of the source or not without requiring the absolute baseline flux. This information is useful in understanding the location of the lens object which we discuss below.  In order to perform follow-up monitoring and spectroscopic observations, one has to detect an EAGLE event in real-time and issue an alert. For the real-time detection of EAGLE events, the image subtraction method is more suitable than the DoPHOT analysis or the pixel analysis, since the image subtraction method can detect the luminosity variation at any position of the fields, even where no star was identified previously. Thus, the most reasonable and practical observation strategy is to observe hourly (with a 1-m telescope) and to perform the real-time analysis with the subtraction method. An alert can then be issued, to observers around the world enabling frequent observations around the event peak. After that, high resolution observations with larger telescope should be carried out to determine the baseline flux of the source.  As seen in section \\ref{sec:ratio}, a large fraction of EAGLE events is likely to be transit EAGLE events. If we can issue an alert for such an EAGLE event, observations with the global network will allow us to obtain a precise light curve near the maximum amplification. Because the transit time is $\\approx 13$ hours for self-lensing and $\\approx 2$ hours for halo MACHOs, it is possible to measure precisely the light curve in the transit region. Since the overall structure of the light curve can be determined by follow-up observations with larger telescopes, one can obtain the source radius scaled by Einstein radius $u_*$ from the full light curve and from the color change measurements due to a limb-darkened extended source (\\cite{han00}). If one can obtain $u_*$ and also the radius of the source star $R_*$ from the color of the source star, one can obtain the Einstein radius $r_E$ projected in source plane. Since the event duration $\\hat{t}$ can be also determined from the full light curve, one can measure the proper motion of the lens object for a transit EAGLE event (\\cite{gou94}). The proper motion of the lens object is of great use in determining the lens location. Since the proper motion of objects in the halo is $30 \\sim 40$ times larger than that of stars in the LMC, one can know whether the lensing object exists in the halo or not from the proper motion of transit EAGLE events. Even if we cannot derive an accurate proper motion of lens, the fraction of transit events can be used to constrain the lens location, because that ratio for the self-lensing case is $2 \\sim 15$ times larger than that for the halo MACHOs case.  To demonstrate how we can constrain the nature of microlensing objects, we show in figure $\\ref{fig:poiss}$ the probability distribution of the transit events for both the halo lensing case and the LMC self-lensing case. For the typical parameters $\\alpha = 2.35$, $V_{th} = 20$ and a detection efficiency of 50\\%, one can expect to find $\\sim 21$ EAGLE events after a 3 year observation period of 11 square degrees of the LMC central region (as the MACHO collaboration does). In these 21 EAGLE events, we can expect 1.1 (0.4) and 6.1 (4.6) transit events for halo MACHOs whose mass is $0.1 (1.0)M_\\odot$ and for self-lensing with $d_{max}=300$ (1k) pc, respectively. So if we do not detect any transit events in 3 years, we can reject the possibility of self-lensing at more than $99\\%$ confidence level according to Poisson statistics. In table \\ref{tbl:conf}, we show some other results on how strongly we can discriminate between the two possible locations of the lensing objects based on Poisson statistics. It is possible that we may be able to constrain strongly the lens location based on the 3-year statistics of transit EAGLE events.  In conclusion, the study of EAGLE events is very effective for constraining the nature of lens objects. If an EAGLE search is made for a few years, we can form strong constraints on lens objects in microlensing events."
    },
    "0004/astro-ph0004301_arXiv.txt": {
        "abstract": "Using all available major samples of Seyfert  galaxies and their corresponding control samples of closely matched  non-active galaxies, we find that the bar ellipticities (or axial ratios) in Seyfert galaxies are systematically different from those in non-active galaxies. Overall, there is a deficiency of bars with large ellipticities (i.e., `fat' or `weak' bars) in Seyferts, compared to non-active galaxies. Accompanied with a large dispersion due to small number statistics, this effect is strictly speaking at the $2\\sigma$ level.  To obtain this result, the active galaxy samples of near-infrared surface photometry were matched to those of normal galaxies in type, host galaxy ellipticity, absolute magnitude, and, to some extent, in redshift. We discuss possible theoretical explanations of this phenomenon within the framework of galactic evolution, and, in particular, of radial gas redistribution in barred galaxies. Our conclusions provide further evidence that Seyfert hosts differ systematically from their non-active counterparts on scales of a few kpc. ",
        "introduction": "The relationship between the large-scale morphology of Seyfert (Sy) host galaxies and the central nonstellar activity is a long-standing problem and the focus of an ongoing debate. Shlosman, Frank \\& Begelman (1989) argued that stellar-dynamical processes on scales of a few kpcs, and gas-dynamical processes on smaller scales, combine to drive the gas towards the centers and  fuel the active galactic nuclei (AGNs). Sufficient evidence, observational  and theoretical, supports the idea that nonaxisymmetries in the background  gravitational potential, e.g., stellar bars, induce radial mass redistribution in disk galaxies (e.g., Simkin, Su \\& Schwarz 1980; Balick \\& Heckman 1982; Shlosman, Begelman \\& Frank 1990; Athanassoula 1994; Buta \\& Combes 1996). On  the other  hand, a number of optical surveys claimed no correlation between the large-scale  morphology and the central activity (e.g., Moles, M\\'arquez \\& P\\'erez 1995; Ho, Filippenko \\& Sargent 1997; Mulchaey \\& Regan 1997).  The perennial question, therefore, to be addressed is {\\it whether the AGN host galaxies differ morphologically from `normal' (non-active) galaxies, and on what spatial scales}.  High-resolution near-infrared (NIR) observations are clearly advantageous in determining the mass distribution and hence detecting large-scale bars (McLeod \\& Rieke 1995; Mulchaey \\& Regan 1997; Peletier et al. 1999, hereafter Paper~I). Knapen, Shlosman \\& Peletier (2000, Paper~II) used sub-arcsec resolution imaging in three NIR bands ($J, H, K$) to study the complete CfA sample of Sy's and a {\\it matched} control sample of normal galaxies using objective and stringent criteria for assigning bars. Here we look at a different aspect, and, instead of studying the frequency of bars in galaxies, investigate the bar axial ratios (i.e., bar ellipticities). This is done for {\\it all}   matching samples which are large enough and available in the literature. Bar parameters, such as strength, mass, and pattern speed, are very difficult to estimate from observations of stellar morphology alone. Even for the simplest models, bar strength depends on bar's quadrupole moment and on the radial distribution of axisymmetric mass in the disk, bulge and halo. The optical RC3 catalogue (de Vaucouleurs et al. 1991) recognizes three broad morphological classes, A, X, and B, i.e., nonbarred, intermediate barred and strongly barred. This RC3 classification, however, is subjective and has, to our knowledge, not been properly documented. We state that presently, it is not feasible to estimate the distribution  of bar strengths proper in any statistically significant sample, but bar axial ratios can provide a reasonable alternative (Martin 1995). Here we develop this idea.  In this Letter we report a systematic difference in the distribution of the deprojected ellipticities of large-scale stellar bars between four samples of Sy and normal `control' galaxies, of which three are independent.  We describe the samples used, provide results of our analysis, and discuss their implications for understanding the AGN-host galaxy connection. ",
        "conclusions": "All samples used in the previous section provide a coherent picture of bar ellipticity distribution in barred Sy hosts and matching control samples of normal barred galaxies. The most intriguing and important result is the apparent deficiency of stellar bars with large ellipticities (i.e., large $\\epsilon_{\\rm b}$) in Sy's, compared to those found in their non-active counterparts.  Although this effect is at the level of $\\sim 2\\sigma$, due to small number statistics, it is consistent across all the Sy and matching control samples.  There are two possible explanations for the observed phenomena within the framework of galaxy evolution both of which are dependent on the cold gas component in the disk. First, numerical simulations of pure stellar disks have shown that the bar instability becomes milder, i.e., produces a bar with smaller ellipticities, if the disk is hotter initially, prior to instability (Athanassoula 1983). This can be understood as weaker swing amplification in disks with larger velocity dispersion (Toomre 1981). An additional caveat is that the presence of a cold and clumpy gas component in the disk provides a heating source for stars and acts to weaken {\\it all} dynamical instabilities, including the bar instability (Shlosman \\& Noguchi 1993) and the vertical bending of the bar (Berentzen et al. 1998). Gas gravity is crucial here and even reasonable amounts of cold gas are sufficient to suppress the instabilities altogether. If indeed Sy disks have a larger fraction of cold gas than normal galaxies (e.g., Hunt et al. 1999), which may also be more clumpy, this trend should be explored more fully. The resulting difference(s) between Sy and normal disks will be long-lasting because the stellar component, once heated up, will not be able to cool down easily, as the stellar `fluid' is non-dissipative.  Our understanding of the evolution of barred galaxies points to an alternative and possibly more elegant explanation for the observed difference in bar properties between active and normal galaxies. Numerical simulations of bars revealed their weakening with time in response to a growing mass concentration in the galactic centers (Hasan \\& Norman 1990; Friedli \\& Benz 1993; Hasan, Pfenniger \\& Norman 1993; Berentzen et al. 1998). The radial gas inflow towards the central kpc and further inwards is the prime suspect. All or part of this gas can contribute to the growth of galactic bulges, nuclear rings, disks and bars, and ultimately to central black holes (BHs) by dissolution of the main family of periodic orbits supporting the large-scale stellar bars. These so-called $x_1$ orbits (e.g., Sellwood \\& Wilkinson 1993) are replaced by stochastic orbits when the mass of the central BH exceeds $\\sim 1\\%$ of the total mass. Heller \\& Shlosman (1996) also found that nuclear rings with masses greater than a few $\\times 10^9\\ {\\rm M_\\odot}$ lead to $x_1$ orbit dissolution exterior to the ring,  leaving smaller bar remnants escaping detection inside the central kpc.  In view of the fact that the supermassive BHs appear to be ubiquitous in normal galaxies (Kormendy \\& Richstone 1995) and not only in Sy nuclei, it seems relevant to ask why the Sy hosts are affected more by the bar dissolution processes than normal galaxies. A resolution of this paradox can lie in that the BHs in Sy's are more massive than in their non-active counterparts, but only statistically. If indeed supermassive BHs play a role in the bar dissolution, we anticipate that their instantaneous mass distributions in Sy's and normal hosts peak at different values, but have a large overlap. Moreover, it is plausible that ground-based observations based on stellar-dynamical considerations overestimate the BH masses in normal galaxies due to insufficient spatial resolution. These mass estimates may be lowered by upcoming HST observations.  In summary, we have analyzed all available reasonably-sized samples of Sy host galaxies and independent control samples of normal galaxies  carefully matched in type, disk ellipticity, absolute magnitude, and, to some extent, in redshift. We find that samples of active galaxies are systematically deficient in high-ellipticity  `thin' stellar bars compared to normal galaxies. The associated uncertainty is quite large and the overall effect is at the $2\\sigma$ level.  We discuss the corollaries of this effect. The acceptable alternatives point towards the cold and clumpy gas component in the disk as being responsible for the observed effect. Our result provides an indication  that Sy host galaxies differ systematically on scales of a few kpc from their normal counterparts, and that the gas component may be responsible for this."
    },
    "0004/astro-ph0004071_arXiv.txt": {
        "abstract": "The large number of microlensing events discovered towards the Galactic Bulge bears the promise to reconstruct the stellar mass function. The more interesting issue concerning the mass function is certainly to probe its low mass end, near the region occupied by the brown dwarfs. However due to the source confusion, even if the distribution and the kinematics of the lenses are known, the estimation of the mass function is extremely biased at low masses. The blending due to the source confusion biases the duration of the event, which in turn dramatically biases the estimation of the mass of the lens. To overcome this problem we propose to use differential photometry of the microlensing events obtained using the image subtraction method. Differential photometry is free of any bias due to blending, however the drawback of differential photometry is that the baseline flux is unknown. In this paper we will show that even without knowing the baseline flux, purely differential photometry allow to estimate the mass function without any biases. The basis of the method is that taking the scalar product of the microlensing light  curves with a given function and taking its sum over all the microlensing events is equivalent to project the mass function on another function. This method demonstrates that there is a direct correspondancy between the space of the observations and the space of the mass function. Concerning the function to use in order to project the observations, we show that the principal components of the light curves are an optimal set. We also demonstrate that there is no additional information about the distribution of the scalar products of the data beyond their sum (first order moment). Higher order moments are only linear combination of the first order moment. Thus the sum of the projection on the principal components contains all the information, and translate in an equal number of projection of the mass function with functions associated with the principal components. To illustrate the method we simulate data sets consistent with the microlensing experiments. By using 1000 of these simulations, we show that for instance the exponent of the mass function can be reconstructed without any biases. ",
        "introduction": "The success of the microlensing experiments has been impressive, several collaborations have reported the detection of microlensing amplification of stars, OGLE (Udalski {\\it et al.} 2000), MACHO (Alcock {\\it et al.} 1998), EROS (Derue {\\it et al.} 1999)), DUO (Alard \\& Guibert 1997), VATT (Uglesich  {\\it et al.} 1999). In particular, the recent release by the OGLE II collaboration of more than 200 microlensing events is of great interest. One of the obvious promises of such data sets is the possibility to explore the mass function near the low mass end. Although, the relation between the mass function and the microlensing observations is not straightforward. In the classical scheme of analysis, one must first estimate the duration of the events by fitting the theoretical light curves, then compute the lensing rates, and finally relate these lensing rates to the mass function. The trouble is that fitting the theoretical light curve to the microlensing data to obtain the duration of the event is a highly degenerated process (Alard 1997, Han 1997, Wozniak \\& Paczy\\`nski 1997). Due to the source confusion in crowded fields, it is almost impossible to estimate the flux of the amplified source reliably. The flux of the source is severely biased by blending with neighboring sources. In such case, an over-estimation of the baseline flux will result in an under-estimation of the duration of the event. This bias is severe for unresolved stars and will result in a reduction in the estimation of the duration of a factor of 2,3, or even more. Since the bias on the mass of the lens goes like the square of the bias on the duration, the relevant bias on the mass function will be very large. It was already shown (Alard 1997, Han 1997) that the contribution of these unresolved, highly blended stars to the lensing rated was dominant at short durations. The consequence is that the estimation of the mass function at low mass end will be largely over-estimated, suggesting the existence of a large number of brown dwarf that actually do not exists. It is obvious that as long as the bias due to the blended sources has not been solved, any correct estimation of the mass function is impossible. This paper proposes a solution to this problem. We will see that using differential photometry obtained using the image subtraction method (Alard \\& Lupton 1998, Alard 2000) it is possible to estimate the mass function, even without knowing the baseline flux of the source. This method assumes that the distribution and kinematics of the lenses are known with good accuracy. It is important to emphasize that the uncertainties related to the structure of the Galaxy in the bulge region are several order of magnitude smaller than the uncertainties due to the blending bias. Many good model of the kinematic and structure of the bulge of our Galaxy are already available (Zhao, 1996, Fux 1997, Bissantz {\\it et al.}, 1997, Binney, Gerhard, \\& Spergel 1997) ",
        "conclusions": ""
    },
    "0004/astro-ph0004247_arXiv.txt": {
        "abstract": "The strength of the H$\\beta$ index has been computed for the integrated spectra of model globular clusters from the evolutionary population synthesis. These models take into account, for the first time, the detailed systematic variation of horizontal-branch (HB) morphology with age and metallicity. Our models show that the H$\\beta$ index is significantly affected by the presence of blue HB stars. Because of the contribution from blue HB stars, the H$\\beta$ does not monotonically decrease as metallicity increases at a given age. Instead, it reaches a maximum strength when the distribution of HB stars is centered around 9500 K, the temperature where the H$\\beta$ index becomes strongest. Our models indicate that the strength of the H$\\beta$ index increases as much as 0.75 {\\AA} due to the presence of blue HB stars.  The comparison of the recent Keck observations of the globular cluster system in the Milky Way Galaxy with those in giant elliptical galaxies, NGC 1399 and M87, shows a systematic shift in the H$\\beta$ against metallicity plane. Our models suggest that this systematic difference is understood if the globular cluster systems in giant elliptical galaxies are several billion years older, in the mean, than the Galactic counterpart. Further observations of globular cluster systems in the external galaxies from the large ground-based telescopes and space UV facilities will enable to clarify whether this difference is indeed due to the age difference or other explanations are also possible. ",
        "introduction": "For distant stellar populations, one relies upon the integrated colors or spectra to investigate their ages and metallicities since individual stars are not resolved. In this paper, we will specifically focus on the H$\\beta$ index that is widely used as an age indicator. It has been highlighted because it is (1) the only index that shows the anticorrelation with metallicity and (2) sensitive to temperature by reaching a maximum strength at around 9500 K (\\cite{wor94}).  Most of the previous works (e.g., \\cite{wor94}; \\cite{buz95}), however, have been done on the basis that stars near the main-sequence turnoff (MSTO) region are the most dominant sources for the integrated strength of H$\\beta$. Consequently, without the meticulous consideration for stars beyond the red giant branch (RGB), they claimed that the strength of H$\\beta$ depends on the location of the MSTO, which in turn depends on the age at a given metallicity. Several investigators, however, have cast some doubt upon the sensitivity of the H$\\beta$ index given the presence of other warm stars, especially blue horizontal-branch (HB) stars (see, e.g., \\cite{rab82}; \\cite{chr83}; \\cite{bur84}; \\cite{buz94}; \\cite{def95}; \\cite{fer95}; \\cite{fis95}; \\cite{bre96}; \\cite{lee96}; \\cite{gre97}; \\cite{jor97}).  On the observational side, it was barely possible to obtain low S/N spectra of globular clusters in systems outside the Local Group (e.g., \\cite{mou87}; \\cite{mou90}; \\cite{huc87}; \\cite{bro91}; \\cite{gri94}). These spectra have only been used for kinematic information. With the advent of 10 m-class telescopes, however, Kissler-Patig et al. (1998) and Cohen, Blakeslee, \\& Ryzhov (1998) have successfully obtained relatively high S/N spectra that provide reliable line indices calibration for globular clusters in NGC 1399 and M87, the central giant elliptical galaxies in Fornax and Virgo clusters. Figure 1 shows the distribution of globular clusters in the Milky Way Galaxy and those in NGC 1399 in the H$\\beta$ versus Mg$_{2}$ diagram. At first glance, there seems to be little difference between these two systems especially when considering the large uncertainties in NGC 1399 globular cluster system. Some careful scrutiny of the data, however, indicates that there is a systematic difference in the sense that the metal-rich clusters in NGC 1399 have higher H$\\beta$ compared to Galactic counterparts, while the opposite seems to be the case for the more metal-poor clusters. In fact, a Kolmogorov-Smirnov (K-S) test gives only 2.14\\% probability, excluding the two very metal-rich NGC 1399 clusters, that the two samples are extracted from the same parent population.  This systematic difference has motivated us to explore the sensitivity of the integrated H$\\beta$ strength with detailed population synthesis models that reproduce the systematic variation of HB morphology with age and metallicity. The HB stars are not unusual components of old stellar systems since they are ubiquitously found from the color-magnitude diagrams (CMDs) of Galactic globular clusters and the stellar evolution theories remarkably reproduce them to a certain degree (e.g., Lee, Demarque, \\& Zinn 1990, 1994). In the following section, our population models without and with considering HB stars are presented. Section 3 reinforces the validity of our models with HB stars by comparing them to a sample of Galactic globular clusters and compares our results with the observations of globular cluster systems in the Milky Way Galaxy, M31, NGC 1399, and M87. Finally, section 4 discusses major implications from our works. ",
        "conclusions": "We have found in this work that (1) the integrated H$\\beta$ strengths of old stellar systems, such as globular clusters, are significantly affected by HB stars, and (2) the present observational data for the globular cluster systems are best understood in terms of systematic age differences among them, in the sense that globular clusters in more massive galaxies are older. If our age estimation is confirmed to be correct, this would indicate that the star formation in denser environments has proceeded much more rapidly and efficiently, so that the initial epoch of star formation in more massive (and denser) systems occurred several billion years earlier than that of the Milky Way (see also Lee 1992). This is also consistent with the view that a substantial population of the massive early-type galaxies formed earlier at very high redshift (e.g., \\cite{lar90}; \\cite{mao90}; \\cite{ben96}; \\cite{har98}; \\cite{sta98}).  Further observations of globular cluster systems in nearby galaxies from the large ground-based telescopes are imperative to manifest this possible systematic age difference. In addition, satellite UV photometry could also provide a test for the validity of our results presented in this paper, because it is expected that older globular clusters would reveal the stronger UV fluxes and the bluer UV colors as their blue HB stars should be hotter than those in the Galactic counterparts at given metallicity (see Figure 1 of Park \\& Lee (1997)). These new observations, together with the detailed population models presented here, will undoubtedly help to clarify our understanding of the formation epoch of galaxies."
    },
    "0004/astro-ph0004320_arXiv.txt": {
        "abstract": "We present a model to investigate the formation and evolution of cluster galaxies using cosmological high-resolution $N$-body simulations. The $N$-body simulations are used to construct merging history trees of dark halos. Gas cooling, star formation, supernova feedback and mergers of galaxies within dark halos are included by using simple prescriptions taken from semi-analytic models of galaxy formation. In this paper, we examine the merger-driven bulge formation model and represent the morphology-density relation of cluster galaxies at $z = 0$. We find that this morphological evolution model can explain the distribution of elliptical galaxies in the clusters well and cannot reproduce the distribution of S0 galaxies. This result suggests that the elliptical galaxies are mainly formed by the major mergers, while, in the S0 formation, the processes other than major mergers play important role. ",
        "introduction": "It is well established that galaxy populations vary with the density of neighbouring galaxies in clusters of galaxies \\citep{dre80}. This morphology-density relation (hereafter MDR) indicates that dynamical processes that depend on environment of each galaxy mainly affect the final configuration of the stellar component. This impression seems to be supported by the Hubble Space Telescope images of clusters at intermediate redshifts, which show an abnormally high proportion of spiral and irregular types at $z \\sim 0.5$ and increase of S0 fraction toward the present time \\citep{cou94, dre97}.  Some mechanisms that may transform one morphological type into another have been proposed, for example, ram-pressure stripping of interstellar medium of spirals by a intracluster medium \\citep{gun72, fn99}, galaxy harassment by cumulative tidal interactions in clusters (Moore et al. 1996, 1998), and galaxy mergers  \\citep{too77}. $N$-body simulations have confirmed that merging disk galaxies produce galaxies resembling ellipticals as merger remnants (e.g. Barnes 1996), so that the galaxy merger is a favorite explanation for the predominance of ellipticals and S0s in rich clusters \\citep{kau93, kau96, bau96}.  The purpose of our study is to check this {\\it merger hypotheses} by investigating the MDR for the simulated clusters in the cosmological context. For this purpose, it is difficult in the present situation to use numerical simulations including both gravity and gas dynamics (e.g. Katz et al. 1992;  Evrard et al. 1994), because such simulations are expensive in CPU time and so only a limited range of parameter space can be explored with an insufficient spatial resolution.  Semi-analytic modeling of galaxy formation has already proved to be a powerful technique (e.g. Kauffmann et al. 1993; Cole et al. 1994). However, we cannot identify the position of galaxies by such approaches, because these models follow the collapse and merging histories of dark halos by using a probabilistic method on the mass distribution based upon an extension  of the Press-Schechter formalism \\citep{pre72, bon91, bow91, lac93}.  One approach to identify the position and velocity of each galaxy is to track the merging histories of dark halos by using $N$-body simulations and to combine them with the simple prescriptions of the semi-analytic models \\cite{rou97, kau99, ben00}. Their schemes, however, do not deal with the substructures within dark halos. Since the dynamics within clusters may strongly affects the evolution of cluster galaxies (Okamoto \\& Habe 1999, 2000), we should use a new galaxy tracing method.  In this paper, we adopt the galaxy tracing method provided by Okamoto \\& Habe (1999, 2000), which enables us to trace the individual galactic dark halos within dense environments using high-resolution $N$-body simulations, that is, we can obtain the three-dimensional distribution of the galaxies within clusters. Our formula of modeling gas cooling, star formation, supernova feedback, and galaxy mergers is directly taken from the previous semi-analytic work. We also adopt a merger-driven scheme for the production of galactic bulges and the way of the morphological classification based on bulge-to-disk ratios as earlier studies \\cite{kau93, bau96}.  In \\S 2, we describe a brief outline of the galaxy formation model used here. We show the results about the morphology of the galaxies in \\S 3. These results are discussed in \\S 4. ",
        "conclusions": "In this paper, we represent a new method that combines high-resolution $N$-body simulations with a semi-analytic galaxy formation model for cluster  galaxies. The high-resolution simulations enable us to identify and trace galactic halos within clusters directly. Therefore, we can obtain the three-dimensional distribution of the cluster galaxies and also we can incorporate the dynamical processes, for example merging and tidal stripping of the galactic halos in the clusters, into our modeling of galaxy formation. For the first study of the cluster galaxies using this method, we have shown the MDRs for the simulated clusters of galaxies with two cosmological models based on the merger-driven bulge formation scenario.  The model which reproduce the observed luminosity function well (model C) is successful in explaining the observed elliptical fraction as a function of the local density. However, any models cannot reproduce the observed trend in the distribution of the cluster S0s. In all models, the S0 fractions are too small, especially in the high density regions.  We also examine the case in which the S0 galaxies in the cluster environment have the different $B/D$ range from the S0 galaxies in the field, because Koopmann \\& Kenney (1998) found that the galaxies with low central concentration, which should be identified as spiral galaxies in the field, were often classified into the S0 type in clusters. We have confirmed that the S0 fraction are increased only in the low density regions even if we adopt lower threshold value of $B/D$ between S and S0.  We then investigate how the parameter, $f_{\\rm bulge}$, affects on the MDR. When we choose smaller $f_{\\rm bulge}$, the early-type fraction is increased. In this case, however, the elliptical fraction is mainly increased, the shortage of S0s is then still significant. Even if we classify the low $B/D$ ellipticals as S0s in the model with small $f_{\\rm bulge}$, this change increases the S0 fraction only in the low density regions again.  Above results imply that, under the merger-driven bulge formation model, the galaxies separate into two types, i.e. almost pure bulge and pure disk systems, in the high density regions independent on the choice of the parameters and cosmologies.  The reason is probably considered as follows. In the model adopted here, a S0 galaxy is  only formed by the disk formation after the last major merger. The efficiency of merging is, however, high before cluster formation, and then it rapidly decreases due to the large internal velocity dispersions of the clusters and the reduction of the size of tidally truncated halos as the mass of clusters grows. The accretion onto the galactic halos is also prevented by the strong tidal field after beginning of the cluster formation \\citep{oka99}. Hence, the cluster ellipticals formed by the mergers at high redshifts hardly change their morphologies into S0s by minor mergers or gas accretion, which form additional stellar disks.  Our results suggest that the {\\it merger hypothesis} can give a good explanation of the distribution of cluster ellipticals. On the other hand, to reproduce the observed S0 fraction, the environmental effects that change the morphology of spiral galaxies into S0 galaxies directly should be considered, for example the ram-pressure stripping \\citep{gun72, fn99}, the galaxy harassment (Moore et al. 1996, 1998), and the minor mergers \\citep{wak96}.  The morphological classification based on $B/D$ is the simplest way. There are various ways of classification other than the $B/D$ classification. For instance, Caon \\& Einasto (1995) found another type of morphological segregation of early-type galaxies in the Virgo cluster, based on the configurations of their isophotes. Comparing local projected densities, galaxies with boxy isophotes are located in local density enhancements, while galaxies with disky isophotes lie in regions of lower local densities, independently of their classification as elliptical or lenticular. Although, it is interesting to investigate what physical processes are responsible for this segregation, more numerical studies are needed to include such morphological classification in our model because we need to understand the dynamical states of the internal structures of galaxies to modeling it, and it is then left for further studies.  In the next paper, we will investigate whole properties of cluster galaxies in detail, i.e. colors, metalicities, velocities, and so on, using the method presented in this paper. We will also discuss some environmental effects."
    },
    "0004/astro-ph0004116_arXiv.txt": {
        "abstract": "We present a study of globular clusters in the inner region of the giant elliptical galaxy NGC 4472, % based on the $HST$ WFPC2 archive data. We have found about 1560 globular cluster candidates at the galactocentric radius $r<4$ arcmin, most of which are located within $r=1.1$ arcmin. $V-(V-I)$ diagram of these objects shows a dominant vertical structure which consists obviously of two components: blue globular clusters (BGCs) and red globular clusters (RGCs). The luminosity function of the globular clusters is derived to have a peak at $V$(max) $=23.50\\pm 0.16$ mag and $\\sigma = 1.19\\pm 0.09$ from Gaussian fitting. From this result the distance to NGC 4472 is estimated to be $d=14.7\\pm1.3$ Mpc, for the foreground reddening $E(V-I)=0.03$ and the peak luminosity for the Galactic globular clusters of $M_V$(max) $= -7.4$ mag. The peak luminosity for the RGCs is similar to that for the BGCs, which indicates that the RGCs may be several Gyrs younger than the BGCs. The luminosity function of the globular clusters shows little systematic variation depending on the galactocentric radius at $r<4$ arcmin. However, the mean luminosity of the bright BGCs decreases by 0.2 mag with increasing galactocentric radius over the range of 9 arcmin, while that of the RGCs does not. The observed color distribution of these globular clusters is distinctively bimodal with peaks at $(V-I) = 0.98\\pm0.01$ and $1.23\\pm0.01$, which correspond to the metallicities of [Fe/H] $= -1.41$ and $-0.23$ dex, respectively. The mean observed color of all the globular clusters with $V<23.9$ mag is derived to be $(V-I) =1.11\\pm0.01$ (m.e.), corresponding to the metallicity of [Fe/H] $= -0.79 \\pm 0.05$ dex. These colors are exactly the same as those for the globular clusters in M87. It is found that the relative number of the BGCs to the RGCs is increasing with the increasing galactocentric radius. There are more RGCs than the BGCs  within $r \\sim 2$ arcmin, while there are fewer RGCs than the BGCs in the outer region of NGC 4472. Radial and azimuthal structures of the globular cluster system in NGC 4472 are also investigated. Surface number density profiles of both the BGCs and RGCs get flat in the central region, and the core radii of the globular cluster systems are measured to be $r_c = 1.9$ arcmin for the BGCs, $r_c = 1.2$ arcmin for the RGCs, and  $r_c = 1.3$ arcmin for the total sample, which are much larger than the stellar core of the galaxy. In general the properties of the globular clusters in the inner region of NGC 4472 are consistent with those of the globular clusters in the outer region of NGC 4472. ",
        "introduction": "NGC 4472 (M49) is a giant elliptical galaxy located at 4$^\\circ$ south of the center of the Virgo cluster. NGC 4472 is the brightest member of the Virgo cluster, some 0.2 mag  brighter than the cD galaxy M87. NGC 4472 is an outstanding example of giant elliptical galaxies showing a bimodality in the color distribution of globular clusters (Gebhardt \\& Kissler-Patig 1999, Kundu 1999). Geisler, Lee, \\& Kim (1996) and Lee, Kim, \\& Geisler (1998) have found from the deep wide ($16.4 \\times 16.4$ arcmin$^{2}$) field CCD imaging of  NGC 4472 that the color distribution of the globular clusters in NGC 4472 is remarkably bimodal. The bimodal color distribution of the globular clusters in NGC 4472 has shown that there are two kinds of cluster populations in this galaxy: a metal-poor population with a mean metallicity of [Fe/H] $= -1.3$ dex and a  metal-rich population with a mean metallicity of [Fe/H] $= -0.1$ dex. Interestingly it is found that the metal-rich globular clusters show some properties in common with the galaxy halo stars in their spatial distribution and color profiles, while the metal-poor globular clusters do not show such behavior. Also it is found that the metal-rich globular clusters are more centrally concentrated than the metal-poor globular clusters. This result indicates that there may exist some connection between the metal-rich globular clusters and halo stars in NGC 4472 in the formation and evolution processes. However, these studies based on the ground-based observations could not investigate the properties of the globular clusters close to the center of NGC 4472, because of the severe crowding problem and strong background halo light in ground-based images of the inner region of the galaxy. Hubble Space Telescope ($HST$) data are ideal for the study of the globular clusters close to the center of bright galaxies such as NGC 4472. They are especially useful for the study of the structure of the globular cluster system in the central region of NGC 4472 for which little is known by now.  In this paper we present a study of the globular clusters in the inner region of NGC 4472 using the $HST$ archive data. During the preparation of this paper Puzia \\etal (1999) published a paper based on the analysis of the data which are the same as three sets among the four sets used in this study. However, there are several differences between the two studies, details of which will be described later. Here is given a brief summary of the differences between the two studies. First, we use one more data set in addition to the data set used by Puzia et al., with which the size of the bright globular cluster sample is increased by 29\\% and the complete areal coverage of the bright globular cluster sample is increased significantly (by a factor of three in the range of the galactocentric radius). Secondly, we use different methods of data reduction and analysis from the Puzia et al.'s. As a result, our photometry goes about 1 mag deeper and is more complete than the Puzia et al.'s before incompleteness correction. Third, $V$ and $I$ magnitudes in our photometry are on average about 0.2 mag fainter than the Puzia et al.'s for the common objects, while the colors are almost the same between the two studies. Fourth, our results for the peak luminosity of the globular cluster luminosity functions are significantly different from the Puzia et al.'s, leading to an opposite conclusion on the age of the globular clusters. Fifth, we present for the first time a study of the structure of the central region of the globular cluster systems which was not done by Puzia et al. Sixth, we use both the $HST$ data of the inner region of NGC 4472 and the ground-based data of the outer region of NGC 4472 to investigate any systematic variation over a wide range, while Puzia et al. worked with only the $HST$ data of the inner region.  This paper is composed as follows. Sections 2 describes the data used in this study, and 3 explains the reduction of the data. Section 4 presents the color-magnitude diagram of the measured objects in the NGC 4472 field. Sections 5 and 6 show, respectively, the luminosity functions and color distribution of globular clusters. Section 7 investigates the spatial structure of the globular cluster system in NGC 4472. Primary results are discussed in Section 8, and a summary of this study is given in the final section. ",
        "conclusions": "We have presented $VI$ photometry of the globular cluster candidates in the inner ($r<4$ arcmin) region of NGC 4472, derived from the $HST$ WFPC2 archive data for the four fields of NGC 4472. We have found about 1560 globular cluster candidates from the colors and image classifiers. Primary results and conclusions obtained in this study are summarized as follows.  \\begin{enumerate}  \\item The color-magnitude diagram of the measured point sources shows a dominant vertical structure at the color range of $0.75<(V-I)<1.45$, most of which are globular clusters in NGC 4472. This structure consists obviously of two populations: the blue globular clusters and the red globular clusters.  \\item The luminosity function of the globular clusters is well fit by a Gaussian function with a peak at $V=23.50\\pm0.16$ mag and $\\sigma = 1.19\\pm0.09$. From this the distance to NGC 4472 is derived to be $d=14.7\\pm1.3$ Mpc. The globular cluster luminosity function shows little systematic variation depending on the galactocentric radius. There is little difference in the peak luminosity between the BGCs and RGCs. This indicates that the RGCs may be several Gyrs younger than the BGCs, contrary to the conclusion given by Puzia \\etal (1999).  \\item The color distribution of the 609 bright globular clusters with $V<23.9$ mag is distinctively bimodal with peaks at $(V-I)=0.975$ and $1.233$. The metallicities corresponding to these peaks are [Fe/H] $= -1.41$ and $-0.23$ dex. The mean color of the bright globular clusters with $V<23.9$ mag is derived to be $(V-I) =1.11\\pm0.01$(m.e.), corresponding to the metallicity of [Fe/H] $= -0.79 \\pm 0.05$ dex.  \\item The ratio of the number of the BGCs to that of the RGCs is increasing with increasing galactocentric radius from the center out to about 9 arcmin. Thus the RGCs are more centrally concentrated than the BGCs.  \\item The surface number density profile of the globular clusters in the inner region is approximately fit by a King model with core radius $r_c = 1.3$ arcmin, while that of the globular clusters in the outer region is better fit by de Vaucouleurs law. The core radius of the globular cluster system obtained in this study is much larger than that of the stellar halo.  \\item The mean colors of the BGCs and RGCs change little with increasing radius, and the mean color of the RGCs is very similar to that of the stellar halo of NGC 4472.  \\item The radial gradient of the metallicity of the entire sample of the globular clusters with $V<23.9$ mag is approximately fit by [Fe/H] $ = - 0.390 \\log r$  $+ 0.120$ with $\\sigma = 0.037$, where $r$ is given in the unit of arcmin.  \\item In general the property of the globular clusters in the inner region of NGC 4472 is consistent with that of the globular clusters in the outer region of NGC 4472, except that there are more RGCs than the BGCs in the inner region.  \\end{enumerate}"
    },
    "0004/astro-ph0004266_arXiv.txt": {
        "abstract": "Most models of anomalous X-ray pulsars (AXPs) account for the observed X-ray spectra and pulsations by means of radiation processes that occur on the surfaces of neutron stars. For any such model, general relativistic deflection of light severely suppresses the amplitude of the observed pulsations. We calculate the expected pulsation amplitudes of AXPs according to various models and compare the results with observations. We show that the high ($\\lesssim 70$\\%) pulse amplitudes observed in some AXPs can be accounted for only if the surface emission is localized (spot radius $<40^\\circ$) and strongly beamed ($\\sim \\cos^n\\theta'$ with $n\\gtrsim 2$, where $\\theta'$ is the angle to the normal). These constraints are incompatible with those cooling and magnetar models in which the observed X-rays originate as thermal emission from the neutron-star surface. Accretion models, on the other hand, are compatible with observations for a wide range of parameters. Finally, definitive conclusions cannot be reached on magnetospheric models, since their localization and beaming properties are not well understood. ",
        "introduction": "Among pulsating compact X-ray sources are a small subset with pulsation periods between 6 and 12 seconds, soft spectra, and no identifiable companions. The first suggestion that these objects might form a separate class of neutron stars -- later called the anomalous X-ray pulsars (AXPs) -- was made by Mereghetti \\& Stella (1995), who proposed that they might be powered by accretion from a very low mass companion (this was also earlier hinted at by Hellier 1994). The lack of optical counterparts, however, and the absence of observable Doppler shifts in the frequency of the X-ray pulses led van Paradijs, Taam, \\& van den Heuvel (1995) to favor a different accretion model for AXPs.  According to their suggestion, material from a fossil accretion disk, possibly the debris of a disrupted binary companion after a period of common-envelope evolution, is being accreted by a solitary neutron star. Recently, a similar model has been proposed by Chatterjee et al.\\ (2000), in which the accreting material is supplied by the post-supernova fallback material from the neutron star progenitor itself.  In a different class of models, AXPs are considered to be isolated neutron stars, spinning down by magnetic dipole radiation. Because of their unusually high period derivatives, a simple application of the relationship between the dipole magnetic field strength and the spin-down rate implies a field strength for these objects of $\\sim 10^{14}-10^{15}$~G (see, e.g., Thompson \\& Duncan 1996). Two main types of models that rely upon the presumed high magnetic field of the stars have been proposed.  Thompson \\& Duncan (1996; see also Duncan \\& Thompson 1992) suggested that the released energy may be drawn from the decay of the magnetic field itself and from differential movements in the stellar crust. This model also serves to explain the bursts observed from soft gamma-ray repeaters as being produced by larger-scale magnetospheric phenomena. Because of such models, AXPs are often called ``magnetars'' and are grouped into the same class of sources as the soft gamma-ray repeaters (see, e.g, Thompson \\& Duncan 1995, 1996; also Hurley 2000 for a review of SGR properties and models). In the alternative high magnetic-field model, Heyl \\& Hernquist (1997a, 1997b) suggested that AXPs draw their energy from the residual thermal energy of the star itself.  All these models face a number of difficulties. For example, if AXPs are powered by accretion from a stellar companion, the absence of detectable Doppler shifts in the arrival of X-ray pulses cannot be easily explained(Mereghetti et al.\\ 1998). If accretion is from either a companion or a fossil disk, the optical fluxes one would expect directly from the disk or due to X-ray reprocessing are too high compared to the observed upper limits (see, e.g., Perna, Hernquist, \\& Narayan 2000; Hulleman et al.\\ 2000). On the other hand, in the magnetar model the absence of bright faster AXPs and the observed variations in spin-down rates are hard to account for (see, e.g., Baykal \\& Swank 1996; Marsden et al.\\ 2000; see, however, Heyl \\& Hernquist 1999; Melatos 1999).  In all the above models of AXP, nearly all of the X-ray emission is produced at the surface of the neutron star. It is well known that strong gravitational fields tend to smooth out the variability produced by a spinning compact star, even if the emission is highly localized in bright spots (Pechenick, Ftaclas, \\& Cohen 1983). Indeed, the X-ray pulse amplitudes of the three radio pulsars that show thermal emission from their surfaces are only $\\lesssim 30$\\% (see, e.g., the discussion in Page 1995; Harding \\& Muslimov 1998). This is in contrast to the non-thermal emission from radio pulsars (which is magnetospheric in origin) and from accretion-powered X-ray pulsars (which is from collimated accretion columns) which often show pulse amplitudes as high as $\\sim 90$\\% (see, e.g., Nagase 1989). In this respect, AXPs are similar to accretion-powered pulsars, showing X-ray pulse amplitudes anywhere between $\\sim 10$\\% and $\\sim 70$\\%.  We study in this paper a set of variability diagnostics that may be used in constraining emission models of AXPs. We examine three parameters: the pulse fraction observed at infinity, which is a measure of the overall amplitude of the variations, and the Fourier amplitudes at the first and second harmonics of the neutron-star spin frequency.  We find tight constraints on the properties of magnetar models. Most importantly, we are unable to reproduce the observed variability properties of AXPs with thermal cooling models. ",
        "conclusions": "In this paper we have used the high pulse amplitudes observed from a number of AXPs to constrain the properties of their emission mechanism. We find that the observations can be accounted for only if the surface emission is localized (half-opening angle $<40^\\circ$) and strongly beamed  \\vbox{ \\centerline{ \\psfig{file=fig9.ps,angle=0,width=9.5truecm,height=7.0truecm} } \\figcaption[]{\\footnotesize Contours of maximum pulse fraction for different beaming functions ($\\cos^n\\theta'$; cf.\\ eq.[\\ref{eq:cosn}]) and half-opening angles of each hot spot. The $\\simeq 70$\\% pulse fraction observed from the source 1E~1048.1-5937 restricts any viable model of AXPs to the left of the uppermost contour.}} \\vspace{0.5cm}  \\noindent ($n\\gtrsim 2$ in eq.~[\\ref{eq:cosn}]), as summarized quantitatively in Figure~9. These constraints are a consequence of the compactness of the neutron stars and the resulting strong general relativistic deflection of photon trajectories. Our conclusions are valid for all realistic neutron star masses and radii.  The properties of individual sources offer a number of additional clues. For example, the double-peaked pulse profile of 1E~2259$+$586 requires that the emission is localized around {\\em two\\/} antipodal spots on the neutron-star surface, probably associated with the magnetic poles. Furthermore, the change in the relative strength of the two peaks observed with {\\em GINGA\\/} (Iwasawa et al.\\ 1992) implies that the pulse shape cannot be solely due to geometric effects but should also reflect a flux contrast between the two antipodal spots. Furthermore, this flux contrast should be variable and, therefore, cannot be caused by a non-dipolar magnetic field configuration or a non-uniform fallback of low metallicity material.  Such arguments, together with the constraints presented in Figures~2, 4, 5, and 9 appear to rule out thermal cooling models for AXPs.  They are also inconsistent with those magnetar models in which most of the X-ray flux originates from heating in the deep surface layers of the neutron star. On the other hand, the localized emission and beaming predicted by accretion models seem to be consistent with the observations.  A magnetospheric model in which the neutron-star surface is heated over localized spots by particle bombardment may also be viable, though the beaming properties of such a model are unknown."
    },
    "0004/astro-ph0004099_arXiv.txt": {
        "abstract": "We present the first comprehensive catalogue of variable stars in the Leo II dwarf spheroidal galaxy.  We have identified 148 RR Lyrae type variables, of which 140 were amenable to derivation of variability parameters with our data.  We have also confirmed the existence of four anomalous Cepheids as identified in previous studies.  The average period of the RR Lyrae ab variables (0.62 days), the fraction of c variables (0.24) and the minimum period of the RR Lyrae ab variables (0.51 days) all define Leo II as an ``Oosterhoff intermediate\" galaxy.  We have used the properties of these variables to derive a metallicity for Leo II of approximately [Fe/H]=-1.9.  We attempt to resolve discrepancies between this value and those determined by previous efforts.  The presence of longer period, higher amplitude RR Lyrae variable implies a metallicity distribution that extends to as poor as [Fe/H]=-2.3.  Leo II's location on the period-metallicity relation of clusters, like that of other ``Ootershoff intermediate'' objects, falls between the Oosterhoff Class I and Oosterhoff Class II clusters. The properties of the variable populations of these objects are consistent with the idea that the Oosterhoff ``dichotomy\" is a continuum.  The gap between the classes seems to be explained by the horizontal branch of Galactic globular clusters shifting away from the instability strip at at intermediate metallicities. However, Leo II, as well as other Oosterhoff intermediate objects, has a second parameter effect strong enough to leave horizontal branch stars in the instability strip. ",
        "introduction": "The Leo II dwarf spheroidal galaxy (dSph) was first identified by Harrington and Wilson (1950) on the Palomar Sky Survey.  Detailed information about Leo II did not appear in the literature for over three decades due to its extreme distance and small size (except the starcount study of Hodge 1962 and abstracts on variables in Swope 1967 and Swope 1968) . A wave of interest in the 1980's (Hodge 1982; Demers and Harris 1983, hereafter DH83; Aaronson et.al. 1983; Azzopardi et al. 1985; Suntzeff et al. 1986, S86) revealed Leo II to be a metal-poor, second-parameter horizontal branch object with a handful of carbon stars.  Recently, a series of studies have expanded our knowledge of the Leo II system. Demers \\& Irwin (1993, DI93) with deep BV CCD photometry, derived a metallicity of [Fe/H] = -1.9 and revised Leo II's distance to 215 kpc. This was followed by the VI CCD studies of Lee (1995, hereafter L95) and Mighell \\& Rich (1996, hereafter MR96).  The former derived a metallicity of -1.97 while the latter derived a metallicity of -1.6 and determined that Leo II formed almost all of its stars between 7 and 14 Gyr ago. Both confirmed a distance modulus close to the DI93 value. Vogt et al. (1995) measured radial velocities in Leo II and derived a mass-to-light ratio of 11.1.  One of the gaps in our knowledge of Leo II remains its variable star content.  The seminal work in this field was to be the comprehensive survey of Baade and Swope.  Their collection of over a hundred Palomar 200-inch plates was intended to produce a complete catalogue of Leo II variables.  Unfortunately, while the catalogue exists, it has not been published, nor has a rigorous analysis been applied to the data.  Swope did publish two brief abstracts, one reporting the identification of 152 variables and period measurement for 76 (Swope 1967) and another reporting four anomalous Cepheids (Swope 1968).  A later update by van Agt (1973) reported 64 Bailey (1902) type ab RR Lyrae stars, six type c and an average ab period of 0.59 days.  Neither report was considered complete by the authors and neither included coordinates, finding charts or light curves. DI93 reported the identification of 80 variable candidates, but did not have enough images to fit light curves.  In this paper, we publish the first comprehensive list of positively identified variable stars in Leo II.  We have found 148 RR Lyrae variables, for which we have successfully derived periods and amplitudes for 140.  The distribution of periods is very similar to other dSph galaxies, with a moderate fraction of RRc variables.  All of our observations are consistent with an ``Oosterhoff intermediate'' classification for Leo II and this classification is exhibited on a star-by-star basis.  We have also identified four anomalous Cepheid variables.  The presence and characteristics of these populations confirm that Leo II is metal-poor with a large intermediate age population.  The longest period RRab variables hint at the existence of Leo II stars as metal-poor as [Fe/H]=-2.3. ",
        "conclusions": "\\subsection{Leo II and the Oosterhoff Continuum}  The RR Lyrae variables in Leo II place its dominant metallicity at [Fe/H]=-1.9.  A handful of large $\\Delta log P$ stars appears to imply a lower metallicity population.  This is not conclusive, but warrants a more comprehensive look at the color-magnitude diagram of Leo II, which we will perform in a future contribution.  At present, we can make no contribution regarding Leo II's distance modulus due to the non-photometric conditions during our observations.  Leo II is Oosterhoff intermediate class, like several other objects, including a number of dSph galaxies.  Its violation of the Oosterhoff gap is exhibited in both the bulk and individual properties of its RR Lyrae stars.  The Oosterhoff intermediate objects fill in the Oosterhoff gap, and therefore support the interpretation that metallicity is the dominant parameter in determining average RRab period.  \\subsection{The Fornax-Leo-Sculptor Stream Revisited}  Lynden-Bell (1982) was the first to note that the dSph galaxies appear to lie in two great streams in the sky.  One of these streams appears to include the Fornax, Leo I, Leo II and Sculptor dSph galaxies.  Majewski (1992) noted that Phoenix, Sextans and several second parameter objects also lie along this plane.  Palma et al. (2000) found that the outermost ($R_{GC} > 25$ kpc) second parameter globular clusters have the highest probability of being aligned with one of the two streams. While such alignments could be coincidental, they could also result from a common formation mechanism, in which case one might expect some similiarities in the stellar populations of brethren objects.  Comparisons between the populations of Leo II and Sculptor are possible from the available data.  Both show a strong low metallicity ([Fe/H]=-1.7 in Sculptor), second parameter, Oosterhoff intermediate population. Both also show evidence of a less populous, even more metal-poor population.  In Sculptor, this population does not have a second parameter effect and is spatially dispersed (M99).  The presence of a spatial gradient in the horizontal morphology of Leo II (Da Costa et al. 1996) may indicate a similar abundance-HB morphology pattern to that observed in Sculptor.  Fornax also has multiple populations. While it is dominated by a population at [Fe/H]=-1.5, it shows evidence of populations from -0.7 to -2.2 in both its field star and globular cluster population (c.f. Buonanno et al. 1985; Beauchamp et al. 1995).  The most metal-poor population, like that in Sculptor, is much more extended than the metal-rich populations (Grebel \\& Stetson 1998). One could envision a scenario in which Fornax, Leo and Sculptor all originated in a common [Fe/H]$\\sim$-2.3, first parameter HB progenitor, then, after disociation, each object followed it own star formation history.  Alternatively, perhaps there is something fundamental about the ``threshold\" metallicity [Fe/H]=-2.3 in small stellar systems.  We note that the lowest metallicities in the Milky Way globular cluster system are also around [Fe/H]=-2.3 (Harris 1996).  Any firm conclusions about correlations between Leo II, Sculptor and Fornax must await the measurement of their absolute proper motions. This will be the single most powerful discriminant for or against a common history.  However, the similarities between the most metal-poor populations in each of Leo II, Sculptor and Fornax are intriguing."
    },
    "0004/astro-ph0004050_arXiv.txt": {
        "abstract": "We present the first grating-resolution X-ray spectra of the Seyfert~1 galaxy NGC\\,3783, obtained with the High Energy Transmission Grating Spectrometer on the {\\it Chandra} X-ray Observatory. These spectra reveal many narrow absorption lines from the H-like and He-like ions of O, Ne, Mg, Si, S and Ar, as well as \\ion{Fe}{17}--\\ion{Fe}{21} L-shell lines. We have also identified several weak emission lines, mainly from O and Ne. The absorption lines are blueshifted by a mean velocity of $\\approx 440\\pm 200$~km\\,s$^{-1}$ and are not resolved, indicating a velocity dispersion within the absorbing gas of a few hundred km\\,s$^{-1}$ or less. We measure the lines' equivalent widths and compare them with the predictions of photoionization models. The best-fitting model has a microturbulence velocity of 150~km\\,s$^{-1}$ and a hydrogen column density of $1.3\\times 10^{22}$~cm$^{-2}$. The measured blueshifts and inferred velocity dispersions of the X-ray absorption lines are consistent with those of the strongest UV absorption lines observed in this object. However, simple models that propose to strictly unify the X-ray and UV absorbers have difficulty explaining simultaneously the X-ray and UV absorption line strengths. ",
        "introduction": "X-ray spectra of Active Galactic Nuclei (AGNs) often show evidence for deep \\ion{O}{7} (739~eV) and \\ion{O}{8} (871~eV) absorption edges. The ionized gas component creating these edges is named the `warm absorber' and is seen in most ($\\gtrsim$70\\%) Seyfert~1s and some quasars. Theoretical models for the warm absorber suggest it should also be a source of X-ray emission (e.g., Netzer 1996) and absorption (e.g., Nicastro, Fiore, \\& Matt 1999) lines. While the current X-ray data are consistent with the presence of such lines (e.g., George et~al 1998a), it has been impossible to study them in any detail.  The bright Seyfert~1 galaxy NGC\\,3783 ($V\\approx13.5$ mag, $z=0.0097$) has one of the strongest X-ray warm absorbers known with a column density of ionized gas of $N_{\\rm H}\\approx$~(1--2)$\\times 10^{22}$~cm$^{-2}$. {\\it ASCA} spectra also show warm absorber variability and excess flux around 600~eV which has been interpreted as emission lines from the warm absorber, particularly \\ion{O}{7}~568~eV (e.g., George et~al. 1998a, 1998b). UV spectra of NGC\\,3783 show intrinsic absorption features due to \\ion{C}{4}, \\ion{N}{5} and \\ion{H}{1}. Two kinematic absorption components are known with radial velocities of $\\approx -560$~km\\,s$^{-1}$ and $\\approx -1420$~km\\,s$^{-1}$ relative to the optical redshift from [\\ion{O}{3}]. The \\ion{C}{4} absorption shows extreme variability (e.g., Crenshaw et~al. 1999 and references therein). It is not clear whether the UV and X-ray absorbing gas are the same, although some relation seems likely (e.g., Shields \\& Hamann 1997).  This letter describes the discovery of a large number of X-ray absorption lines in the {\\it Chandra} spectrum of NGC 3783. In \\S~2 we describe the observation and data analysis, and \\S~3 gives the preliminary interpretation of the results. \\vspace{1cm} ",
        "conclusions": "\\subsection{Absorption Lines, Column Density and Ionization Parameter} \\label{ionsec}  The strongest observed absorption lines in the X-ray spectrum of NGC\\,3783 are from the H-like and He-like ions of O, Ne, Mg, Si, S and Ar, as well as \\ion{Fe}{17}--\\ion{Fe}{21} L-shell lines (see Table~1 but note that some features represent blends, and the values given are the summed EWs).  As explained above, the lines are not resolved, and the only firm conclusion is that, given a Gaussian profile, their FWHMs do not exceed a few hundred km\\,s$^{-1}$. Given this uncertainty, we have taken an approach that relies solely on the measured EWs of the lines. We assumed a certain microturbulence velocity in a single line-of-sight cloud and calculated a theoretical spectrum with the photoionization code ION99, \\ the 1999 \\ version of the  \\footnotesize \\begin{center} {\\sc TABLE 1\\\\ Absorption and Emission Lines From NGC\\,3783} \\vskip 4pt \\begin{tabular}{lcccc} \\hline \\hline {} & {Predicted} & {Velocity} & {Predicted} & {Measured} \\\\ {} & {$\\lambda$} & {Shift} & {EW} & {EW} \\\\ {Ion} & {(\\AA)} & {(km\\,s$^{-1}$)} & {(m\\AA)} & {(m\\AA)} \\\\ \\hline \\multicolumn{5}{c}{Absorption Lines} \\\\ \\hline \\ion{S }{16}                  &  4.728 & $-$571$\\pm$317          &   0.7 &   8.0$\\pm$2.6 \\\\ \\ion{S }{15}                  &  5.039 & $-$535$\\pm$297          &   8.2 &   7.5$\\pm$2.5 \\\\ \\ion{Si}{14}                  &  5.217 & $-$920$\\pm$287          &   1.4 &  10.1$\\pm$3.5\\phn \\\\ \\ion{Si}{14}                  &  6.181 & $-$534$\\pm$242          &   7.7 &  17.3$\\pm$4.6\\phn \\\\ \\ion{Si}{13}                  &  6.648 & $-$586$\\pm$225          &  18.4 &  14.7$\\pm$4.8\\phn \\\\ \\ion{Mg}{12}                  &  7.106 & $-$379$\\pm$211          &   6.6 &   6.2$\\pm$2.2 \\\\ \\ion{Mg}{11}                  &  7.473 & $-$280$\\pm$200          &   8.6 &   8.0$\\pm$3.6 \\\\ \\ion{Mg}{11}                  &  7.851 & $-$382$\\pm$191          &  15.6 &   9.7$\\pm$3.1 \\\\ \\ion{Mg}{12}\\tablenotemark{a} &  8.420 & $-$463$\\pm$178          &  $>$19.9 &  25.6$\\pm$6.6\\phn \\\\ \\ion{Mg}{11}\\tablenotemark{a} &  9.170 & $-$261$\\pm$163          &  $>$25.9 &  34.5$\\pm$10.2 \\\\ \\ion{Ne}{10}                  &  9.362 & $-$384$\\pm$320          &   6.2 &  16.4$\\pm$4.5\\phn \\\\ \\ion{Ne}{10}\\tablenotemark{b} &  9.481 & $-$474$\\pm$158          &  12.0 &  18.6$\\pm$4.9\\phn \\\\ \\ion{Ne}{10}\\tablenotemark{a} &  9.709 & $-$216$\\pm$154          &  $>$15.7 &  29.7$\\pm$8.2\\phn \\\\ \\ion{Ne}{10}\\tablenotemark{a} & 10.239 & $-$410$\\pm$146          &  $>$24.7 &  19.8$\\pm$5.5\\phn \\\\ \\ion{Ne}{ 9}\\tablenotemark{a} & 11.001 & $-$518$\\pm$190          &  $>$20.1 &  50.4$\\pm$12.2 \\\\ \\ion{Fe}{17}                  & 11.250 & $-$186$\\pm$133          &  23.3 &  15.7$\\pm$4.3\\phn \\\\ \\ion{Fe}{18}\\tablenotemark{a} & 11.435 & $-$629$\\pm$131          &  $>$20.5 &  21.0$\\pm$14.2 \\\\ \\ion{Ne}{ 9}\\tablenotemark{a} & 11.547 & $-$701$\\pm$129          &  $>$49.3 &  27.2$\\pm$7.4\\phn \\\\ \\ion{Ne}{10}\\tablenotemark{b} & 12.132 & $-$667$\\pm$123          &  69.8 &  49.2$\\pm$19.8 \\\\ \\ion{Fe}{17}\\tablenotemark{b} & 12.263 & $-$391$\\pm$122          &  43.1 &  51.9$\\pm$12.8 \\\\ \\ion{O }{ 8}                  & 14.635 & \\phn$-$20$\\pm$102       &  34.3 &  24.8$\\pm$12.2 \\\\ \\ion{O }{ 8}                  & 14.821 & $-$182$\\pm$101          &  38.9 &  27.8$\\pm$12.8 \\\\ \\ion{O }{ 8}                  & 15.176 & \\phn$-$79$\\pm$197       &  42.8 &  35.8$\\pm$11.2 \\\\ \\ion{Fe}{17}                  & 15.264 & $-$471$\\pm$196          &  42.7 &  30.6$\\pm$8.4\\phn \\\\ \\ion{Fe}{18}                  & 15.628 & $-$624$\\pm$249          &  18.2 &  27.7$\\pm$7.9\\phn \\\\ \\ion{O }{ 8}\\tablenotemark{a} & 16.007 & $-$346$\\pm$121          &  $>$52.3 &  43.9$\\pm$14.2 \\\\ \\ion{O }{ 7}                  & 18.628 & $-$515$\\pm$322          &  52.1 &  41.9$\\pm$14.8 \\\\ \\ion{N }{ 7}                  & 19.120 & $-$510$\\pm$109          &  11.2 &  18.8$\\pm$6.8\\phn \\\\ \\hline \\multicolumn{5}{c}{Emission Lines} \\\\ \\hline \\ion{Mg}{12}                  &  9.300 & \\phm{$-$}306$\\pm$225    &  15 &  11.0$\\pm$6.8\\phn \\\\ \\ion{Ne}{10}                  & 12.132 & \\phm{$-$}444$\\pm$123    &  36 &  44$\\pm$24 \\\\ \\ion{Ne}{ 9}                  & 13.447 & \\phm{$-$}\\phn64$\\pm$111 &  19 &  99$\\pm$88 \\\\ \\ion{Ne}{ 9}                  & 13.553 & \\phm{$-$}464$\\pm$110    &  13 & 163$\\pm$140 \\\\ \\ion{Ne}{ 9}                  & 13.699 & \\phm{$-1$}26$\\pm$175     &  41 & 245$\\pm$215 \\\\ \\ion{O }{ 8}                  & 18.967 & \\phm{$-$}415$\\pm$126    & 460 & 335$\\pm$310 \\\\ \\ion{N }{ 7}                  & 20.912 & \\phm{$-$}161$\\pm$100    &  28 & 118$\\pm$92\\phn \\\\ \\ion{O }{ 7}                  & 21.602 & \\phm{$-1$}54$\\pm$69\\phn &  99 & 197$\\pm$169 \\\\ \\ion{O }{ 7}                  & 21.807 & \\phm{$-$}229$\\pm$68\\phn &  52 & 135$\\pm$109 \\\\ \\ion{O }{ 7}                  & 22.102 & \\phm{$-$}154$\\pm$67\\phn & 201 & 631$\\pm$583 \\\\ \\hline \\end{tabular} \\vskip 2pt \\parbox{3.485in}{ \\small\\baselineskip 9pt \\footnotesize \\indent $\\rm ^a${Blended\\,with\\,Fe\\,absorption lines which\\,are\\,not\\,taken into account in the model.} \\\\ $\\rm ^b${Blended with Fe absorption lines which are taken into account in the model.} } \\end{center} \\setcounter{table}{1} \\normalsize  \\noindent code ION (Netzer 1996). We then compared the results with the observations. The most important ingredients of the modelare the microturbulence velocity, the column density, the spectral energy distribution (SED), the ionization parameter, and the gas composition.  The atomic data required for calculating line optical depths are taken from various sources. The oscillator strengths for the H-like and He-like transitions of all metals are well known (e.g., Porquet \\& Dubau 2000 and references therein). This is not the case for the various iron L-shell lines where there is no standard published data set. In this paper we rely on various publications (e.g., Mason et al. 1979; Bhatia \\& Mason 1980; Cornille et al. 1992, 1994; Phillips et al. 1996; Saba et al. 1999) as well as on the National Institute of Standards and Technology (NIST) data set. These were combined with the Phillips et~al. (1999) line lists, which are based on solar X-ray spectra and served as our major source for the wavelengths of the iron lines. Only 3--5 of the strongest iron lines, per ion, have been used in the analysis.  Our models are based on those in George et~al. (1998b) that gave satisfactory fits to the 1993 and 1996 {\\it ASCA} observations of the source. They include a weak UV-bump AGN SED (the weak IR case in Netzer 1996) with $\\Gamma_{0.1-50~{\\rm keV}}$=1.75, a single constant density cloud with $N_{\\rm H}=1.3\\times 10^{22}$~cm$^{-2}$, solar com- \\centerline{\\epsfxsize=8.5cm\\epsfbox{ewcomp.eps}} \\figcaption{Measured versus modeled absorption line EWs. Filled circles are model EWs that are a sum of all lines contributing to the blend. Right-pointing triangles are lower limits for model EWs (see \\S~\\ref{obser} for details). The agreement between the model and the measured points is noticeable (Spearman rank-order correlation probability of $\\approx 10^{-6}$). A line with a slope of unity is drawn to guide the eye. \\label{ewcomp}} \\centerline{}   \\footnotesize \\begin{center} {\\sc TABLE 2\\\\ Best Model Column Densities} \\vskip 4pt \\begin{tabular}{lclc} \\hline \\hline {} & {Column} & {} & {Column} \\\\ {Ion} & {($10^{17}$ cm$^{-2}$)} & {Ion} & {($10^{17}$ cm$^{-2}$)} \\\\ \\hline \\ion{O}{7} &    5.0 & \\ion{S}{16} &  0.05   \\\\ \\ion{O}{8}&     38 & \\ion{Ar}{17} &  0.06    \\\\ \\ion{Ne}{9}&  4.0  & \\ion{Ar}{18} & 0.003    \\\\ \\ion{Ne}{10}& 6.9  & \\ion{Fe}{17} & 1.2     \\\\ \\ion{Mg}{11}&  2.5  & \\ion{Fe}{18} & 1.5    \\\\ \\ion{Mg}{12} &  1.5 & \\ion{Fe}{19} & 1.0   \\\\ \\ion{Si}{13} &  2.2 & \\ion{Fe}{20} & 0.4   \\\\ \\ion{Si}{14} &  0.5 & \\ion{Fe}{21} & 0.07  \\\\ \\ion{S}{15} & 0.5  & \\ion{Fe}{22} & 0.02   \\\\ \\ion{C}{4} & 2.4$\\times 10^{-5}$ &  \\ion{N}{5} & 4.7$\\times 10^{-4}$  \\\\ \\hline \\end{tabular} \\end{center} \\setcounter{table}{2} \\normalsize \\centerline{}  \\noindent position, and a 0.1--10 keV ionization parameter of $U_{\\rm x}=0.13$. The standard ($E>13.6$~eV) ionization parameter for this continuum is 19.3 and the mean temperature is 1.9$\\times 10^5$~K. We have assumed Doppler profiles with $V^2=V^2_{\\rm thermal}+V^2_{\\rm turbulence}$ and calculated absorption EWs for $100<V_{\\rm turbulence}<600$~km\\,s$^{-1}$. The best agreement with observations is for $V_{\\rm turbulence}$=150~km\\,s$^{-1}$. The model EWs are tabulated in Table~1. For a few lines we list only the model's EW lower-limit since the measurement includes blends for which we do not yet have a modeled EW. In most cases, the blends are expected to increase the modeled EW by $\\approx$~10--20\\%. The comparison between the observed and calculated EWs is shown in Fig.~\\ref{ewcomp}, and the deduced column densities, for the ions producing the strongest lines, are given in Table~2. All of the strong lines predicted by the model are observed in the spectra. The microturbulence velocity of our best fitted model is consistent with the velocity dispersion of the UV absorption lines (161$\\pm$22~km\\,s$^{-1}$; Crenshaw et~al. 1999).  We have carried out basic tests to verify that the deduced line widths are in the right range. If the `true' widths are much larger than assumed, the resonance lines in the H-like and He-like series will be optically thin. The observed EWs show that this is not the case. On the other hand, if the `true' line widths are much smaller than 150~km\\,s$^{-1}$, several of these lines will be saturated, giving very similar EWs to all. This is not observed either. While these are not strong conclusions given the S/N of the observations, it is encouraging to see the good agreement between the observed and calculated EWs, suggesting that the line optical depths are indeed in the intermediate range. According to the model, the largest optical depth for a line with $E>1$~keV is $\\tau$(\\ion{Ne}{10}~12.132~\\AA)=35.  \\subsection{Dynamics of the Absorbing Gas} \\label{dinamic}  All the identified absorption lines in the HETG spectrum of NGC\\,3783 are blueshifted by a few hundred~km\\,s$^{-1}$ with respect to the systemic velocity. The unweighted mean blueshift is 440$\\pm$200~km\\,s$^{-1}$. Kaastra et~al. (2000) have recently observed the bright Seyfert~1 NGC\\,5548 using the Low Energy Transmission Grating Spectrometer (LETGS) on {\\it Chandra}. The low-energy X-ray lines reported in that paper indicate a similar level of ionization, and the absorption lines are blueshifted by similar amounts. Judging by these first two results, it seems that the warm X-ray material in Seyfert~1 galaxies is outflowing with typical velocities of several hundred km\\,s$^{-1}$. This has important implications for the dynamics of the absorbing gas as well as for the X-ray-UV connection.  The acceleration of X-ray ionized gas by the strong central radiation source in AGN has not been studied in detail so far (for some discussion see Reynolds \\& Fabian 1995). The recent work by Chelouche \\& Netzer (2000) is the first detailed study of this kind. This study discusses the motion of discrete, X-ray ionized clouds in hydrostatic equilibrium, under a large range of conditions. It shows that acceleration to velocities of a few hundred to a few thousand km\\,s$^{-1}$ is a natural consequence of the AGN environment. The terminal velocity is of the same order as the escape velocity in the region where the flow originates. Using these results, and assuming that the conditions in the nucleus of NGC\\,3783 are similar to those considered by Chelouche \\& Netzer (2000), we conclude that the origin of the warm absorbing gas in this source is outside the Broad Line Region, where the escape velocity drops below $\\approx 1000$~km\\,s$^{-1}$. Better line profile analysis will enable us to constrain this location more accurately.  Another important comparison is with the UV absorption lines observed from this source. Several papers (e.g., Mathur, Elvis, \\& Wilkes 1995; Shields \\& Hamann 1997 and references therein) suggest that the UV and X-ray absorbers may originate in the same component. The Crenshaw et al. (1999) observations reveal two UV absorption systems, the strongest of which has a velocity consistent with the mean observed velocity of the X-ray gas and a velocity dispersion consistent with our best fitted model. Our model, which depends on the assumed SED, predicts column densities for \\ion{C}{4} and \\ion{N}{5} (see Table~2) which are 37 and 5 times smaller, respectively, than what has been calculated by Crenshaw et al. (1999) using the 1993--1995 data. Hence the UV absorption lines appear to be inconsistent with our X-ray absorber model and are likely coming from a separate component.  \\subsection{Emission Lines and Covering Factor}  The HETG spectra of NGC\\,3783 revel several weak emission lines (see Table~1) some of which show hints of P Cygni profiles. These have a marginally significant redshift with an unweighted mean of 230$\\pm$170~km\\,s$^{-1}$ relative to the systemic velocity measured from the [\\ion{O}{3}] line. Kaastra et~al. (2000) found similar redshifts for the emission lines in the LETG spectrum of NGC\\,5548. While the EWs of some of the lines are well within the MEG and HEG capability, the low S/N and bumpy continuum do not allow reliable EW measurements. The values given in Table~1 are therefore highly uncertain. However, even these rough measurements are good enough to test the idea that the central source in NGC\\,3783 is surrounded by gas clouds with properties similar to those deduced for the line-of-sight absorber. This is done by comparing the EWs calculated by ION with those observed, for various assumed covering factors.  The calculated emission-line EWs, for an assumed geometry of full covering, are listed in Table~1. Given these values, we find that the mean deduced covering factor is close to unity, in agreement with previous estimates (e.g., George et al. 1998b). This is the {\\it first direct constraint} on the covering factor of a warm absorber in any AGN. Unfortunately, the uncertainty on the covering factor is large, due to the uncertain EWs. Further discussion of this issue, as well as detailed analysis of the continuum, emission features, and absorption features, are deferred to a future paper."
    },
    "0004/hep-ph0004109_arXiv.txt": {
        "abstract": "We discuss a problem concerning the ultrahigh energy neutrino propagation through the Earth. We present  calculation of the neutrino-nucleon cross section at high energies, based on the unifed evolution equation at small $x$. We also show the solution of the transport equation for different neutrino fluxes originating from active galactic nuclei, gamma ray bursts and top-down model. ",
        "introduction": "The ultrahigh energy neutrino physics is an interdisciplinary subject which can address many vital problems in astrophysics, particle physics and geophysics, for a review see \\cite{GQRS2,GQRS1}. Neutrinos can traverse large distances without being disturbed and thus give us important information about distant astronomical objects like active galactic nuclei or gamma ray bursts. The ultrahigh energy neutrinos can have energies up to $10^{12} \\; \\rm GeV$, much larger than currently accessible at present colliders, like HERA. At these ultrahigh energies the structure of nucleon is probed at very small values of Bjorken $x$. It means that we are possibly entering a region in which partons have large density. Therefore detailed knowledge of parton distribution function at very small $x$ is vital to estimate the neutrino-nucleon cross section. This cross section rises strongly with  energy and therefore neutrino propagation through the Earth can be affected by the increased interaction with matter at these ultrahigh energies. Thus it is necessary to consider the effect of attenuation of high energy neutrinos while traversing the Earth. This phenomenon could be used as a possible method of the \"Earth tomography\" by neutrinos. All these topics will be studied at large neutrino telescopes like AMANDA, NESTOR and ANTARES (for review see \\cite{HALZEN}).  In this paper we present a method of calculation the neutrino-nucleon cross section using gluon and quark distributions obtained from a unifed evolution equation for small $x$ \\cite{KMS}. This equation embodies both DGLAP and BFKL evolutions on equal footing as well as important subleading effects via consistency constraint \\cite{AGS}. This constraint limits the virtualities of the exchanged gluon momenta in the ladder to their transverse components. We then extrapolate the results to ultrahigh energies and compare the predictions with the other, based on standard global parton analysis \\cite{GQRS2},\\cite{GKR}.  We use the resulting cross section as an input to the transport equation \\cite{NIK} for the neutrino flux penetrating the Earth. We solve this equation for different incident angles as well as different input neutrino fluxes originating from active galactic nuclei, gamma ray bursts and a sample top-down model. \\\\  The details of the  calculation has been presented elsewhere \\cite{KMS2}. \\\\ ",
        "conclusions": ""
    },
    "0004/astro-ph0004044_arXiv.txt": {
        "abstract": "We present detailed three-dimensional modelling of the stellar luminosity distribution for the disks of 31 relatively nearby ($\\le$ 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 $\\approx 10\\%$. 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. ",
        "introduction": "Global parameters of galactic disks can be used to constrain the formation process as well as the evolution of disk galaxies. Recently it has become possible to deduce parameters such as scalelength or central surface density from numerical or semi-analytical self-consistent galaxy formation models (Syer et al. \\cite{syer}, Mo et al. \\cite{mo}, Dalcanton et al. \\cite{dal}) and compare the results with observed values. Observationally the Hubble Deep Field gives the opportunity to study morphological features and simple structural parameters for galaxies even at high redshifts (Takamiya \\cite{taka}, Marleau \\& Simard \\cite{hdfmorph}, Fasano \\& Filippi \\cite{fasan}). \\\\ Former statistical studies providing sets of homogeneously derived parameters for nearby galaxy samples are those of de Jong (\\cite{dj}) using a two-dimensional and of Courteau (\\cite{cour}) with a one-dimensional decomposition technique. \\\\ However, so far only a few statistical studies based on high quality CCD data (de Grijs \\cite{dgmn}, Barteldrees \\& Dettmar \\cite{bd}, hereafter Paper I) have addressed the actual three-dimensional structure of disk galaxies with regard to the stellar distribution perpendicular to the plane (z-direction) and taking into account an outer truncation {\\it (cut-off radius)} for the disk, first introduced by van der Kruit \\& Searle (\\cite{vdk81a}). \\\\ While the stellar distribution perpendicular to the plane (z-profile) could result from various \"heating\" processes during the galactic evolution (Toth \\& Ostriker \\cite{toth}, Hernquist \\& Mihos  \\cite{hern}), the cut-off radius can be used to constrain either the angular momentum of a protocloud (van der Kruit \\cite{vdk87}) or possible starformation thresholds in the gaseous disk (Kennicutt \\cite{kenni}). \\\\ In the following, we have compiled parameters of galactic disks for a total of 31 edge-on galaxies in different optical filters of which 17 objects have been already discussed in Paper I. One goal of our detailed comparison of several independent fitting procedures is to study the influence of the applied techniques on resulting disk parameters. It also provides the data base for statistical analysis addressing some of the above mentioned questions (e.g., Pohlen et al. \\cite{pol}). For 14 objects surface photometric data are given for the first time in Appendix A. ",
        "conclusions": ""
    },
    "0004/astro-ph0004272_arXiv.txt": {
        "abstract": "The Multibeam survey at the Parkes Radio Telescope has provided a wealth of 21 cm HI data. We have mosaicked ten standard HIPASS cubes to produce a datacube approximately 25\\deg$\\times$25\\deg\\ in size, centred on NGC1399, the optical centre of the Fornax Cluster. Some properties of the initial $\\sim$80 galaxies identified in HI are described.  In the optical, Ferguson's Fornax Cluster Catalogue (FCC) (Ferguson 1989) identified 340 likely member galaxies in the central 40 deg$^2$. In the radio we have detected a more uniform sheet of about 80 galaxies at the cluster velocity extending up to 15\\deg\\  from the cluster centre. At the cluster distance of about 15 Mpc, this corresponds to an elongated structure more than 7 Mpc in extent.  Galaxies were detected to a lower mass limit of $\\sim$$1 \\times 10{^8}$ M$_\\odot$ and fewer than 25 of these were within the central 40 deg$^2$, \\hspace{2pt} suggesting considerable HI depletion of galaxies in the centre of the  cluster. Further, these results strongly indicate that HI surveys do {\\it{not}} readily identify galaxy clusters. ",
        "introduction": "The Fornax Cluster is amongst the closest and most well studied clusters in the southern sky, providing a rich nearby field for the study of galaxy populations, dynamics and evolution in the cluster environment.   Covering an area of nearly 40 deg$^2$, Ferguson's optical catalogue (Ferguson 1989) identified 340 ``likely cluster members'', only 85 of which had redshifts. Cluster membership was largely assigned on the basis of morphology and surface brightness (SB), suggesting a bias whereby high SB compact galaxies, paricularly unusual dwarf galaxies, may have been rejected as ``stars'' and any extended low surface brightness (LSB) galaxies incorrectly assigned a ``background'' status (Disney 1976).  Using the Tully-Fisher relation in Fornax, Bureau et al.\\ (1996) determined the cluster distance to be \\mbox{15.4 $\\pm$ 2.3 Mpc}. Recessional velocities of Fornax cluster galaxies range from about 700 \\kms\\ to 2200 \\kms\\  (Barnes et al.\\ 1997). Very  few galaxies have been detected in front of the cluster or in the void behind the cluster in the region \\mbox{2200--4500 \\kms}.  Galaxies beyond about \\mbox{{\\it{cz}} =  3000 \\kms}  would not be classified as likely cluster members (Drinkwater \\& Gregg 1998).   The fields covered by several surveys of the Fornax cluster are illustrated in Figure 1, including the unbiased spectroscopic Two Degree Field (2dF) Survey which is surveying  {\\it{all}} objects in the magnitude limits 16.5 $\\leq$ $B_J$ $\\leq$ 19.7 in a {\\mbox{$\\sim$12 deg$^2$}} region (Drinkwater et al.\\ 2000a).  The positions of the 10 HIPASS cubes used to compile our mosaic are shown.  The deep neutral hydrogen survey covering 10\\deg$\\times$10\\deg\\ in a ``basket-weave'' pattern has recently been completed  and new detections not found in HIPASS have already been identified (see Drinkwater et al.\\ in these proceedings).  Also, a   spectrographic survey centred on Fornax using the FLAIR instrument on the UK Schmidt Telescope (not shown here) is identifying  bright compact objects over a 6\\deg$\\times$6\\deg\\ field in this region (Drinkwater \\& Gregg 1998).   \\begin{figure} \\plotone{waugh1.eps} \\caption{Relative positions and approximate sizes of various surveys in the Fornax region.} \\end{figure} ",
        "conclusions": "A mosaic of 10 standard HIPASS cubes, spanning about 25\\deg$\\times$25\\deg\\ and centred on the Fornax Cluster, has been created and efficiently searched using an automated galaxy finder.  Fewer than about 30\\% of the detections are within the central 40 deg$^2$ of the cluster, suggesting considerable HI depletion of galaxies in the cluster. Further, our results show that HI surveys do not readily detect  galaxy clusters.  Down to a lower M$_{\\rm HI}$ limit of $10^8$  M$_\\odot$, many low HI mass galaxies were detected. Initial results indicate that generally only low {\\mbox{M$_{\\rm HI}$}} galaxies have high HI mass-to-light ratios ({\\mbox{M$_{\\rm HI}$}/L}). New velocities of at least 8 objects have been measured. Interestingly, several HI detections do not have optical counterparts.  Our initial $\\sim$80 confirmed HI detections have revealed that the Fornax Cluster is imbedded in a large-scale structure  more than 7 Mpc in extent and evidence of a velocity gradient across the LSS is seen."
    },
    "0004/astro-ph0004102_arXiv.txt": {
        "abstract": "We present a  novel approach for investigating lens  phenomenology of cosmic strings  in order  to elaborate  detection strategies  in galaxy  deep field images. To account  for the complexity of the  projected energy distribution of string networks we assume their lens  effects to be similar to those of a straight string carrying a {\\em  random} lineic energy distribution. In such a model we show that, unlike the case of uniform strings, critical phenomena naturally  appear.  We  explore the  properties  of the  critical lines  and caustics. In particular, assuming that the energy coherence length along the string is much smaller than the observation scale, we succeeded in computing the total length of critical lines per unit string length and found it to be $4/\\sqrt{3}\\,{\\bf E}(3/4)$.  The length  of the associated caustic lines can also  be computed to  be $16/(\\pi\\,\\sqrt{3})\\,{\\bf  E}(3/4)$.  The  picture we obtain here for  the phenomenology of cosmic string  detection is clearly at variance with common lore. ",
        "introduction": "Although the current results on large-scale structure formation favor models with initial aiabatic scalar perturbation\\cite{Inflation}, the formation of cosmic     strings     is     general     enough    to     merit     further investigations\\cite{Strings}. It is in  particular not excluded that strings may  form  after  an  inflationary phase\\cite{InflationetStrings}.   We  are interested here in  the phenomenological aspects of cosmic  strings for lens distortion  effects whereas most  of the  previous investigations  have been focused on  string detection  from multiple quasar  images\\cite{Hogan,Hu} or Cosmic Microwave Background fluctuations\\cite{KS}.   With the advance of new generation of large  CCD cameras the best direct  evidence for cosmic string relics is however  likely to be obtained from distortion effects they induce on background objects such as galaxies.  In a companion paper we insist on the generic properties expected for cosmic strings  as far  as  lens effects  are  concerned: string  lens effects  are equivalent to those induced by  a lineic energy distribution. This effective energy  distribution should  obviously  take into  account  the lens  energy density, its tension, as well as kinetic energy that might result from rapid movements or energy currents along the string\\cite{peter94b}.  The explicit  computation of  string induced lens  effects has been  done in various  cases, for  long strings  \\cite{Strings}, circular  loops  in plane transverse  to the  line  of  sight\\cite{delaix96}, or  small  loops with  a multipole  expansion  approach\\cite{Hogan}.    However,  much  more  complex situations   have  been   explored   so  far   with  numerical   experiments only\\cite{delaix97}. In  particular the study  presented in~\\cite{delaix97b} suggests a quite  different lens phenomenology for the  images of background galaxies but lacks analytical insights. Having in mind such phenomenological effects,  we  try  in  this  paper  to  get  more  quantitative  results  by introducing a  simplified description of strings.  In  Sect. \\ref{SModel} we present  the   model  of  ``Poisson''   energy  distribution  we   use.   In Sect. \\ref{ElPhenomenology} general  phenomenological properties observed in such  a model  are  presented.   Sect.  \\ref{StatProp}  is  devoted to  more precise calculations  on the properties  of the critical lines  and caustics for such a model of strings and finally we evoque observational aspects. ",
        "conclusions": "\\label{Discuss}   This  investigation  provides  a  new description  of  the  phenomenological properties  of lens  physics for  cosmic  strings.  The  model of  ``Poisson string'' we  present is based on  general results which state  that the lens effects of string  are those obtained from lineic  energy density.  We think such a model must have captured  most of the generic properties expected for string lens phenomenology  as the resemblance with the  results of numerical experiments (see Fig.  3 of \\cite{delaix97b}) strongly suggests  it.  We are however aware that  the validity of the description we  adopted has not been demonstrated  (numerical  simulations of  string  networks  anyway lack  the necessary resolution for describing  correctly the small scale structures of strings).  It is clear for instance  that the projected position of a string on  the sky  does not  follow a  straight line.   We have  checked  that the results we found are not strongly affected if a non-zero curvature radius is introduced,  as long  as this  radius is  larger than  the  other quantities involved in this problem, i.e. $\\xi_0$ and $(s_0\\xi_0^2)^{1/3}$.   The results we have obtained suggest novel strategies for detecting strings: multiple images are still present, but images can have undergone significant distortions;  small images  of  background  objects can  be  found in  large numbers...  The most  adapted strategy depends however on  the resolution of the  imaging  apparatus  and  on   the  energy  scale  of  the  string.   On Fig. \\ref{HSTLens}  we show an  example of an  HST deformed image.   In this case the energy scale of the string expressed in units of $\\xi_0$ is similar to the angular scale of a galaxy.   For GUT strings, $\\xi_0$ would be of the order  of a  few  arcsecs which  is indeed  a  typical galaxy  size in  deep surveys. The probability of observing such an event depends obviously on the survey coverage. The expected angular length of strings to be present within redshift unity  has been estimated in  \\cite{hindmarsh} and is  such that on average one  string is  expected to  cross a 100  deg$^2$ survey.   But even though such  surveys are  within reach in  the near future,  the probability with which one can detect a  string (or conversely the rejection one can put on  the existence  of  strings with  a  given energy  scale) still  deserves further investigations in  particular on the effects of  multi source planes (the amplitude  of the lens  effects depends on  the position of  the source plane) that may partly blur the searched effects."
    },
    "0004/astro-ph0004178_arXiv.txt": {
        "abstract": "s{We apply the Pseudo-$C_l$ formalism to obtain an unbiased, approximate method for efficient simultaneous estimation of several cosmological parameters from large, almost full--sky cosmic microwave background data sets.} ",
        "introduction": "Within the standard model of cosmology there are about 10 parameters which characterise the properties of our Universe. It is one of the key goals of future CMB experiments such as MAP and Planck to determine these cosmological parameters to high precision. This undertaking faces the challenge that realistic CMB data is necessarily incomplete and noisy. The Galaxy obscures roughly a third of the sky and because of the smallness of the anisotropy signal, detector noise is not negligible in the analysis. This leads to the computational challenge which was expertly described at this meeting  in the contribution by Borrill\\cite{borrillMoriond}.  In this talk we apply the pseudo-$C_l$ formalism\\cite{WandeltHivonGorski} to this problem and show that it can be used to develop an approximate form of the likelihood which has several useful properties: it is Gaussian and hence easy to apply; it does not suffer from the usual disadvantages of Gaussian approximations such as obtaining negative estimates of positive definite quantities; it is computationally efficient with memory usage of $N_{pix}$ and number of operations scaling as $N_{pix}^{\\frac32}$ per likelihood evaluation with a very small pre--factor leading to thousands of likelihood evaluations per CPU hour. ",
        "conclusions": ""
    },
    "0004/astro-ph0004387_arXiv.txt": {
        "abstract": "Anomalous X-ray pulsars, compact non-pulsing X-ray sources in supernova remnants, and X-ray bursters are three distinct types of sources for which there are viable models that attribute their X-ray emission to thermal emission from the surface of a neutron star. Inferring the surface area of the emitting regions in such systems is crucial in assessing the viability of different models and in providing bounds on the radii of neutron stars. We show that the spectroscopically inferred areas of the emitting regions may be over- or under-estimated by a factor of $\\lesssim 2$, because of the three-dimensional geometry of the system and general relativistic light deflection, combined with the effects of phase averaging. Such effects make the determination of neutron-star radii uncertain, especially when compared to the $\\sim 5$\\% level required for constraining the equation of state of neutron-star matter. We also note that, for a given spectral shape, the inferred source luminosities and pulse fractions are anticorrelated because they depend on the same properties of the emitting regions, namely their sizes and orientations. As a result, brighter sources have on average weaker pulsation amplitudes than fainter sources. We argue that this property can be used as a diagnostic tool in distinguishing between different spectral models. As an example, we show that the high inferred pulse fraction and brightness of the pulsar RXS~J1708$-$40 are inconsistent with isotropic thermal emission from a neutron-star surface. Finally, we discuss the implication of our results for surveys in the soft X-rays for young, cooling neutron stars in supernova remnants and show that the absence of detectable pulsations from the compact source at the center of Cas~A (at a level of $\\gtrsim 30$\\%) is not a strong argument against its identification with a spinning neutron star. ",
        "introduction": "Neutron stars in nature appear in various different flavors. They were first unambiguously discovered in radio wavelengths as rotation-powered pulsars and later on, in X-rays, as accretion-powered pulsars and bursters. Such systems, however, constitute only a small fraction of the expected number of neutron stars in the galaxy, as inferred from estimates of supernova rates (e.g., Kaspi 2000). As a result, ongoing searches exist for detecting neutron stars in different manifestations, e.g., as isolated stars accreting from the interstellar medium (e.g., Belloni et al.\\ 1997) or simply cooling by thermal emission (e.g., Pavlov et al.\\ 1996). Different emission mechanisms are thought to operate in different types of neutron stars and possibly even between different wavelength bands in a given system. For at least three distinct classes of neutron stars, which we discuss below, viable models exist in which radiation emerging directly from their surfaces is responsible for their high-energy spectra.  A number of slow X-ray pulsars, often called the anomalous X-ray pulsars (AXPs; Mereghetti \\& Stella 1995) form a distinct class of neutron stars with soft X-ray spectra but no radio or optical counterparts. Because of their high spin-down rates and association with supernova remnants (SNRs) they are thought to be young and related to the soft $\\gamma$-ray repeaters, in their quiescent states (see Hurley 2000). It is still an open question whether such objects are neutron stars accreting from a fossil disk (e.g., van Paradijs et al.\\ 1996; Chaterjee, Narayan, \\& Hernquist 2000), are powered by magnetospheric emission (e.g., Thompson \\& Duncan 1995, 1996), or even emitting thermally (Heyl \\& Hernquist 1998).  Another class of objects, potentially related to the AXPs are the compact, non-pulsing, soft X-ray sources that are being discovered within SNRs. The most recent of such compact objects is the central source in Cas A discovered with the {\\em Chandra X-ray Observatory\\/} (Tananbaum 1999; see also Pavlov et al.\\ 2000; Chakrabarty et al.\\ 2000).  Their spectra and luminosities are typical of what is expected for cooling, young neutron stars (see also Pavlov et al.\\ 1996; Zavlin et al.\\ 1998, 1999). However, the absence of coherent pulsations in their X-ray brightness allows for the possibility that they are black holes accreting from fallback material (see, e.g., Pavlov et al.\\ 2000; Chakrabarty et al.\\ 2000).  Finally, weakly-magnetic ($\\lesssim 10^{10}$~G) accreting neutron stars often show thermal emission from their surfaces during thermonuclear flashes, the so-called Type I X-ray bursts (see Lewin, van Paradijs, \\& Taam 1996). In six bursters, relatively coherent oscillations are detected during the bursts, at frequencies $\\sim 300$~Hz (see, e.g., Strohmayer et al.\\ 1996). The coherence and stability of the oscillations suggests that they occur at the neutron-star spin frequencies and are caused by the non-uniform pattern of burning in their surface layers.  Understanding the properties of the X-ray emission from such thermally emitting neutron stars, especially in connection to the presence or absence of detectable coherent pulsations, is crucial in assessing the viability of different models. Moreover, comparing model spectra to observations offers the possibility of measuring the radii of neutron stars and hence constraining the properties of neutron-star matter (see Lattimer \\& Prakash 2000 for a recent discussion). Recently, there has been significant progress in calculating spectra emerging from neutron-star atmospheres with various compositions (see, e.g., Pavlov et al. 1996; Rajagopal, Romani, \\& Miller 1997). Constraints on the amplitudes of oscillations from compact stars with non-uniform surface emission have also been studied in connection to the observed oscillation amplitudes in cooling radio pulsars (Page 1995), bursters (see, e.g., Miller, \\& Lamb 1998; Weinberg, Miller, \\& Lamb 2000), and AXPs (DeDeo, Psaltis, \\& Narayan 2000).  In this article, we address a number of issues related to the emission from a spinning neutron star with a non-uniform surface brightness. In particular, in \\S2 we study the effects of the three-dimensional geometry of the systems, phase averaging, and the general relativistic deflection of light. In \\S3, we demonstrate that, when the radiation emerges from a localized surface area on a rotating star, the radiation flux reaching an observer at infinity may be substantially different compared to the case of isotropic, spherically symmetric emission from a compact star with the same surface area (see also Zavlin, Shibanov, \\& Pavlov 1995; Leahy \\& Li 1995). As a result, when such effects are not taken explicitly into account, the hot-spot sizes, inferred from the observed fluxes and temperatures, may be significantly over- or under-estimated. We show that general-relativistic deflection typically reduces this discrepancy, depending on the compactness of the neutron star.  We also note that the surface area on the neutron star, from which the localized emission emerges, determines both the flux that reaches the observer at infinity and the amplitude of pulsations at the stellar spin frequency. In fact, for stars with the same local surface temperature and emerging spectrum, an anticorrelation is expected between the luminosity and pulsation amplitude.  Therefore, simultaneous consideration of these two properties for a given system can offer strong constraints on the operating emission mechanisms. In \\S4, we investigate this property and its implications for the observations of compact X-ray sources with and without detectable pulsations. ",
        "conclusions": "In this paper we have shown that the inferred properties of thermally emitting neutron stars, i.e., their emitting surface areas and pulse fractions, are significantly affected by the three-dimensional geometry of the systems, phase-averaging effects, and general relativistic light bending. The uncertainties introduced by these effects can be significant (with a spread of $\\lesssim 2$), especially at the limit of small emitting surface areas. However, even when most of the neutron star surface is emitting, the uncertainties in estimating the neutron-star radius are of the same order as the $\\sim 5$\\% level required (Lattimer \\& Prakash 2000) for constraining the equation of state of neutron-star matter.  We have also argued that for a given local radiation spectrum emerging from a bright spot on the stellar surface, faint sources can give rise to larger pulse fractions than brighter sources. The maximum pulse fraction as a function of the brightness of the neutron star, measured by $l$ (cf.\\ eq.~[\\ref{eq:Loo}]), provides a diagnostic tool for distinguishing between different emission models.  These derived constraints depend both on the local radiation spectrum (e.g., thermal versus non-thermal emission) and its beaming and are, therefore, different for different classes of models. Moreover, both quantities are measurable from spectral and timing observations and can be directly compared to the calculated constraints. The high observed pulse fractions of AXPs (see, e.g., Chakrabarty et al.\\ 2000) and X-ray bursters (Strohmayer et al.\\ 1998) have been shown to strongly constrain emission models and such constraints can become only tighter when the apparent luminosities of these systems are also taken into account. As an example, we consider the source RXS~J1708$-$40, which has been identified as an AXP (Sugizaki et al.\\ 1997). Fitting a blackbody spectrum to {\\em ASCA\\/} observations of this source gives a temperature of $T_\\infty=0.41\\pm 0.03$~keV and a blackbody flux of $4\\pi d^2 F_\\infty=3.16\\times 10^{35}$~erg~s$^{-1}$ (Sugizaki et al.\\ 1997; Chakrabarty et al.\\ 2000). The inferred apparent brightness of this source is, therefore, $l=0.8$ and combined with the observed $\\simeq 50$\\% pulse fraction (Sugizaki et al.\\ 1997) is inconsistent with isotropic thermal emission from the neutron-star surface (see Fig.~5).  Finally, our results also have a number of important implications for soft X-ray surveys, in supernova remnants, for young, cooling neutron stars that emit thermally.  For a given spectral shape, the brightest sources, which would be more easily detectable, correspond to smaller pulse fractions, while strong pulsations are only expected from dimmer sources. This effect is shown in Figure~6, where the fraction of systems $N_{\\rm obs}/N_{\\rm tot}$ with a pulse fraction at infinity larger than a threshold $PF_{\\rm 0}$ is plotted against the fractional emitting surface area. For this purpose, we assume a random orientation of the magnetic inclination and the inclination to the observer for a sample of systems and define \\begin{equation} \\frac{N_{\\rm obs}}{N_{\\rm tot}}(PF_0) = \\int_{\\alpha=0}^{\\pi/2}\\int_{\\beta=0}^{\\pi/2} X[PF(\\alpha,\\beta); PF_0] \\sin\\alpha d\\alpha\\; \\sin\\beta d\\beta \\;, \\end{equation} where the step function $X(PF;PF_0)$ is defined such that it is unity when $PF>PF_0$ and zero at all other values of the pulse fraction. For realistic neutron star masses and radii ($p\\sim 2-3$), only a very small fraction of sources shows pulsations that are detectable at a significant level ($\\gtrsim 10-20$\\%) and this fraction drops rapidly with increasing apparent luminosity.  If the central source in the remnant Cas A is a young, cooling neutron star, its surface brightness distribution cannot be uniform, as inferred from fitting thermal models to the observed countrate spectra (Pavlov et al.\\ 2000; Chakrabarty et al.\\ 2000). However, as Figure~6 shows, this property is not inconsistent with the $\\simeq 30$\\% upper limit on its pulse fraction (Chakrabarty et al.\\ 2000). For the assumed isotropic beaming of the emerging radiation and the very small fractional surface area inferred for the central source in Cas~A (Pavlov et al.\\ 2000; Chakrabarty et al.\\ 2000), less than half of the systems would have been detected with pulse fractions higher than the detection threshold, if $p=4$. For more realistic neutron-star properties ($p=2,3$), no system would show such a large pulse fraction. Therefore, the absence of detectable pulsations from this source is not a strong argument against its identification with a spinning neutron star."
    },
    "0004/astro-ph0004208_arXiv.txt": {
        "abstract": "Binary clusters account for more than 10\\% of the cluster population in the Magellanic Clouds.  Statistically fewer than 50\\% of the found pairs are expected to be chance superpositions. We estimated the cluster encounter rate and suggest that tidal capture is an unlikely formation scenario for the formation of binary clusters. Thus, most {\\it true} binary clusters can be expected to have formed together.  Here we present a study of three binary cluster candidates which are located in the bar of the LMC. NGC 1971 \\& NGC 1972 are situated in the association LH\\,59 in the eastern part of the bar. A third star cluster, NGC 1969, is close enough to this pair that all three objects may constitute a triple system. We present the first age determination that is based on CMDs for these star clusters. Our findings suggest that all three clusters are young (40-70 Myr) and may have been formed in the same GMC. It cannot clearly decided whether the clusters are physically interacting or not.  NGC\\,1894 \\& SL\\,341 are located at the south-western rim of the LMC bar. This pair is studied in detail for the first time: The isopleths of both clusters reveal an elliptical shape. Whether this might be interpreted as a sign of interaction or is a peculiarity which is shared with a large amount of LMC star clusters which show higher ellipticities than their counterparts in the Milky Way remains unclear. From our age determination we find that both clusters are coeval with an age of $55\\pm5$ Myr. This makes a formation from the same GMC a likely scenario.  SL\\,385 \\& SL\\,387 are a close pair in the western part of the LMC bar. We derived ages of $170\\pm30$ Myr for SL\\,385 and $\\ge250$ for SL\\,387. The large age difference makes it unlikely that these two clusters formed in the same GMC. ",
        "introduction": "\\label{intro}  The Magellanic Clouds offer the unique possibility to study star clusters in general and binary clusters in particular. These two companion galaxies are close enough to resolve single stars, but distant enough to make the detection of close pairs of star clusters an easy task. Bhatia \\& Hatzidimitriou (\\cite{bh}), Hatzidimitriou \\& Bhatia (\\cite{hb}), and Bhatia et al. (\\cite{brht}) have surveyed the Magellanic Clouds in order to catalogue the binary cluster candidates. The maximum projected centre-to-centre separation of the components of a pair chosen for inclusion in the list of candidates was 18 pc, which corresponds to $\\approx 1\\farcm3$ in the Large Magellanic Cloud (LMC). A binary cluster with larger separation may become detached by the external tidal forces while shorter separations may lead to mergers (Sugimoto \\& Makino \\cite{sm}, Bhatia \\cite{bhatia}). Two clusters may appear to be a binary cluster due to chance line-up while in fact being at different distances within the Magellanic Clouds and not gravitationally bound to each other. An estimation of the number of such chance-pairs of objects (Page \\cite{page}) revealed that less than half of all pairs found are expected statistically. As considerably more double clusters were found, this strongly suggests that at least a subset of them must be true binary clusters, i.e., clusters that are formed together and/or may interact or even be gravitationally bound.  Star clusters form in giant molecular clouds (GMCs), but the details of cluster formation are not yet totally understood (Elme\\-green et al. \\cite{eepz}). Fujimoto \\& Kumai (\\cite{fk}) suggest that star clusters form through strong collisions between massive gas clouds in high-velocity random motion. Oblique cloud\\--cloud collisions result in compressed sub-clouds which revolve around each other and form binary or multiple clusters. Binary star clusters are expected to form more easily in galaxies like the Magellanic Clouds, whereas in the Milky Way the required large-scale high-velocity random motions are lacking. Indeed only few binary clusters are known in our own Galaxy.  The probability of tidal capture of one cluster by another one is small (Bhatia et al. \\cite{brht}), but becomes more probable in dwarf galaxies like the Magellanic Clouds with small velocity dispersion of the cluster systems (van den Bergh \\cite{vdbergh}). In that case the clusters would be gravitationally bound, but age differences are likely. Tidal interactions between clusters can be traced using isodensity maps (de\\,Oliveira et al. \\cite{dodb}, Leon et al. \\cite{lbv}). However, Vallenari et al. (\\cite{vbc}) estimated a cluster encounter rate of $dN/dt\\sim1\\cdot(10^{9} \\mbox{yr})^{-1}$. This makes tidal capture of young clusters very unlikely. Interaction between LMC and SMC might have led to a higher cluster formation rate and, due to the dynamical perturbance, a higher encounter rate might be possible (see Vallenari et al. \\cite{vbc} and references therein).  The formation of low-mass star clusters tends to proceed hierarchically in large molecular complexes over several $10^7$ years (e.g., Efremov \\& Elmegreen \\cite{ee}). Leon et al. (\\cite{lbv}) suggest that in these groups the cluster encounter rate is higher and thus tidal capture is more likely: Binary clusters are not born together as a pair with similar ages but are formed later through gravitational capture. An observational test of this scenario would require the detection of evidence of tidal interactions between clusters, whose age differences need to be compatible with the survival times of GMCs.  Another binary cluster formation scenario is introduced by Theis et al. (\\cite{teph}) and Ehlerov\\'a et al. (\\cite{epth}): Exploding supernovae close to the centre of a giant molecular cloud sweep up the outer cloud material within a few Myrs and accumulate it in the shell. The large amount of matter makes the shell prone to gravitational fragmentation and finally to the formation of many open cluster-like associations. In case of a dense ambient medium outside the cloud or a very massive original molecular cloud the shell can be strongly decelerated resulting in a recollapse. A galactic tidal field acting on this recollapsing shell can split it into two or more large clusters, thus forming coeval twin clusters which are, however, not gravitationally bound (Theis \\cite{theis}). The fragments may stay together for a long time, though they are gravitationally unbound. The evolution of their spatial separation mainly depends on the shape of the shock front at the time of fragmentation.  Depending on their masses and separations binary clusters will eventually merge or become detached. A merged binary cluster will have one single but elliptical core (Bhatia \\& McGillivray \\cite{bm}) which can also be found in the young blue populous star clusters of the Magellanic Clouds. Can mergers of former binary star clusters be responsible for at least some of the blue populous clusters in the LMC?  We are studying binary cluster candidates in the Magellanic Clouds to investigate whether the cluster pairs are of common origin and if they show evidence for interaction. While it is not possible to measure true, deprojected distances between apparent binary clusters an analysis of their properties can give clues to a possible common origin.  In this paper, we investigate three binary cluster candidates, namely NGC\\,1894 (also known as SL\\,344, Shapley \\& Lindsay \\cite{sl}, or BRHT\\,8a, see Bhatia et al. \\cite{brht}) \\& SL\\,341 (or BRHT\\,8b), NGC\\,1971 \\& NGC\\,1972 (or SL\\,481 (BRHT\\,12a) \\& SL\\,480 (BRHT\\,12b), and SL\\,387 \\& SL\\,389 (or BRHT\\,35a \\& BRHT\\,35b). All three cluster pairs are located in the bar of the LMC.  NGC\\,1894 and SL\\,341  constitute a close pair at the south-western rim of the LMC bar. Only the bigger cluster NGC\\,1894 appears in the list of Bica et al. (\\cite{bcdsp}) as an SWB type II cluster (Searl et al. \\cite{swb}). Here we present the first detailed study of this candidate binary cluster.  NGC\\,1971 and NGC\\,1972 are situated in the association (or stellar cloud, Lucke \\& Hodge \\cite{lh}) LH\\,59 in the eastern part of the LMC bar, close to the geometrical centre of this galaxy (see Hodge \\& Wright \\cite{hw}).  A third cluster, NGC\\,1969 (or SL\\,479), is located within $1\\farcm5$ ($\\approx22$ pc). Bica et al. (\\cite{bcdsp}) classified the cluster pair to be of SWB type II while the third cluster belongs to the next older class SWB type III. So far, no accurate ages based on colour magnitude diagrams (CMDs) have been obtained for this potential triple cluster system.  The third cluster pair, SL\\,385 and SL\\,387, is also located close to the geometrical centre, but in the western part of the LMC bar. According to ages for a large sample of star clusters, derived on the base of integrated colours, Bica et al. (\\cite{bcdsp}) propose an age gradient of the LMC bar. Younger clusters are predominantly found in the eastern part, while older clusters of SWB type III and higher are concentrated to its western end. From integrated colours Bica et al. (\\cite{bcdsp}) find these two clusters to be coeval, while Vallenari et al.\\ (1998) determine a large CMD-based age difference of $\\sim 350$ Myr between the clusters. Based on subgiant density profiles Leon et al. (\\cite{lbv}) suggest that the two clusters may constitute a true binary cluster with physical interaction.  This paper is organized as follows. In Sect. \\ref{phot} we describe the photometric data in general. Stellar density maps are presented in Sect. \\ref{stardens}. The following section describes the CMDs for the components of each cluster pair. Ages for each cluster are derived and compared with previous studies. In Sect. \\ref{tidalcapture} we estimate the probability of cluster encounters in the LMC's bar and disk. In Sect. \\ref{summary} we give a summary and discuss the results. ",
        "conclusions": "\\label{summary}  We investigated the stellar densities in and around the triple cluster NGC 1971, NGC 1972, and NGC 1969. The isopleths do not show connections between the clusters, but reveal the distortion of NGC\\,1971 towards the other two clusters. NGC\\,1969 is much more extended than can be seen from short exposed optical images. Furthermore we see a clumpy structure on the stellar density map. Whether this is due to hierarchical stellar clustering (see text below) or due to background fluctuations cannot be clearly said.  Comparing isochrones based on the Geneva models (Schae\\-rer et al. \\cite{smms}) with the CMDs of each star cluster we find the following ages: $63\\pm10$ Myr for NGC 1971, $40\\pm10$ Myr for NGC 1972, and $65\\pm15$ Myr for NGC 1969.  Within the errors all three clusters are similar in age and may have been formed in the same GMC. The slight age differences suggest that cluster formation occurred in several episodes, with NGC 1969 as the oldest of the triple cluster which might have triggered the formation of the other two clusters, and NGC 1972 as the smallest and youngest object. Similar evidence for multiple episodes of cluster formation separated by a few tens of Myr have also been observed in 30 Doradus (Grebel \\& Chu \\cite{gc}). The triple cluster can be seen as an example of hierarchical stellar clustering (Elmegreen et al. \\cite{eepz} and references therein): All three clusters are embedded in the larger association  LH 59. The youngest field star population is part of LH 59 and has an age of 40 to 80 Myr which corresponds to the ages for the three clusters. The clumpy structure of NGC 1969 reveals stellar clustering on smaller scales. Whether the star clusters are physically connected or not cannot be clearly decided. Both NGC 1971 and 1969 show somewhat distorted isopleths which might indicate a possible interaction, but this is highly uncertain. A thorough radial velocity study for binary and multiple cluster candidates would be helpful.  Both NGC\\,1894 and SL\\,341 reveal an elliptical shape on the stellar density map. A significant enhancement between the clusters was not found. A density gradient can be seen in Fig. \\ref{n1894dens} in direction of the LMC bar. Whether the elliptical shape of the two clusters is due to interaction or is a peculiarity that is found for a large amount of single clusters in the LMC (Goodwin \\cite{goodwin}) is unclear.  Fitting isochrones to the CMDs we find that the components of this cluster pair are coeval within the measurement accuracy with an age of $55\\pm5$ Myr. It is likely that both clusters have formed from the same GMC. Whether they have or had some interaction cannot be decided, since most LMC clusters show larger ellipticities than clusters of the Milky Way.  The models of both Fujimoto \\& Kumai (\\cite{fk}) and Theis (\\cite{theis}) (see Sect. \\ref{intro}) provide possible formation scenarios for the triple cluster NGC\\,1971, NGC\\,1972 \\& NGC\\,1969 and the cluster pair NGC\\,1894 \\& SL\\,341.  The star density map of SL\\,385 \\& SL\\,387 shows an enhancement between the two clusters which connects the components of this binary cluster candidate. The enhancement is $2 \\sigma$ above the mean of the surrounding field. However, both clusters have a projected distance of $45\\farcs6$ corresponding to 11.1 pc and an extension of $\\approx27\\farcs2$ or 9.2 pc radius. It is likely that the isopleths of the clusters are overlapping even if the clusters may have different distances along the line of sight, e.g. both clusters may appear as a pair only due to chance superposition. Based on deeper data, Leon et al. (\\cite{lbv}) suggest the presence of two partially overlapping tidal tails distinguishable by their different stellar populations, which are indicative of a possible encounter between these two clusters of different ages.  Based on isochrone fitting to the CMDs we derived an age of $170\\pm30$ Myr for SL\\,385, while only a lower age limit of 250 Myr could be found for SL\\,387. If we adopt the ages estimated by Vallenari et al. (\\cite{vbc}), then the age difference is $\\approx 350$ Myr, far above the maximum life time of protocluster gas clouds suggested by Fujimoto \\& Kumai (\\cite{fk}).  We calculated cluster encounter rates for the bar and the disk of the LMC. We found the cluster encounter rate to be increased by a factor of 10 in the dense, cluster-rich LMC bar as compared to the disk of the LMC, though the rate is still very low (2 encounters per Gyr!). Since additional conditions such as sufficient angular momentum loss must be fulfilled for a tidal capture of one cluster by another during an encounter, this seems an unlikely formation scenario for the formation of binary clusters. Obtaining radial velocities for a large number of stars in cluster pairs, in their potential tidal tails, and in the surrounding field population would help to evaluate cluster internal and external dynamics and to rule out or confirm the occurrence of interactions."
    },
    "0004/astro-ph0004068_arXiv.txt": {
        "abstract": "We present a deep X-ray observation of the low-luminosity active galactic nucleus in NGC~4258 (M~106) using the {\\it Advanced Satellite for Cosmology and Astrophysics} ({\\it ASCA}).  Confirming previous results, we find that the X-ray spectrum of this source possesses several components.  The soft X-ray spectrum ($<2\\keV$) is dominated by thermal emission from optically-thin plasma with $kT\\sim 0.5\\keV$.  The hard X-ray emission is clearly due to a power-law component with photon index $\\Gamma\\approx 1.8$ absorbed by a column density of $N_{\\rm H}\\approx 8\\times 10^{22}\\pcmsq$.  The power-law is readily identified with primary X-ray emission from the AGN central engine.  Underlying both of these spectral components is an additional continuum which is possibly due to thermal emission of a very hot gaseous component in the anomalous arms and/or the integrated hard emission of X-ray binaries in the host galaxy.  We also clearly detect a narrow iron K$\\alpha$ emission line at $\\sim 6.4\\keV$.  No broad component is detected.  We suggest that the bulk of this narrow line comes from the accretion disk and, furthermore, that the power-law X-ray source which excites this line emission (which is typically identified with a disk corona) must be at least $\\sim 100\\,GM/c^2$ in extent.  {\\it This is in stark contrast to many higher-luminosity Seyfert galaxies which display a broad iron line indicating a small ($\\sim 10\\,GM/c^2$) X-ray emitting region.}  It must be stressed that this study constrains the size of the X-ray emitting corona rather than the presence/absence of a radiatively efficient accretion disk in the innermost regions.  If, instead, a substantial fraction of the observed narrow line originates from material not associated with the accretion disk, limits can be placed on the parameter space of possible allowed relativistically broad iron lines.  We include a discussion of various aspects of iron line limb-darkening for highly inclined sources, including the effect of gravitational light-bending on the apparent limb-darkening law.  By comparing our data with previous {\\it ASCA} observations, we find marginal evidence for a change in absorbing column density through to the central engine, and good evidence for a change in the AGN flux.  We conclude with a brief discussion of two serendipitous sources in our field of view; the QSO Q1218+472 and a putative $z\\sim 0.3$ cluster of galaxies. ",
        "introduction": "The nearby low-luminosity active galactic nucleus (AGN) in NGC~4258 (M~106) has become crucially important in our understanding of accreting supermassive black holes.  Exquisite position and velocity measurements of the H$_2$O megamasers in NGC~4258 reveal that the masing material resides in a very thin and slightly warped disk that is in almost perfect Keplerian motion about a central black hole with a mass of $3.5\\times 10^7\\Msun$ (Nakai, Inoue \\& Miyoshi 1993; Greenhill et al. 1995a; Miyoshi et al. 1995).  Detailed studies of the maser proper motions and centripetal accelerations confirm this interpretation (Greenhill et al. 1995b) and allow us to measure the distance to this source independently of the Hubble constant ($d=7.2\\Mpc$; Herrnstein et al. 1999).  Sensitive X-ray observations provide a powerful means of probing both large scale and small scale structures within NGC~4258.  X-ray emission was first detected in a short {\\it Einstein} observatory high resolution imager (HRI) observation (Fabbiano et al. 1992).  More sensitive soft X-ray data collected by the {\\it ROSAT} position sensitive proportional counters (PSPC) and HRI found an extended halo of hot ($4\\times 10^6\\K$) gas around NGC~4258 (Pietsch et al. 1994; Vogler \\& Pietch 1999) as well as X-ray emission associated with the well known helically twisted jets (the anomalous arms; Pietsch et al. 1994; Cecil, Wilson \\& De Pree 1995; Vogler \\& Pietch 1999).  None of these soft X-ray observations penetrated the column of absorbing gas that obscures the AGN itself. This had to await {\\it ASCA} observations (Makishima et al. 1994; hereafter M94). M94 found that the central X-ray source was well described by a power-law with photon index $\\Gamma\\approx 1.8$, and was absorbed by a column density of $N_{\\rm H}\\sim 1.5\\times 10^{23}\\pcmsq$. The soft X-ray spectrum was found to be complex with components arising from thermal plasma emission and, possibly, contributions from the underlying X-ray binary population in NGC~4258.  A marginal detection of an iron line was claimed with an equivalent width of $250\\pm 100\\eV$.  Of great importance is the opportunity that NGC~4258 gives us to study accretion physics when the mass accretion rate is potentially very low. The significance of this issue is highlighted when it is realized that most of the (quiescent) supermassive black holes in the universe are inferred to accrete matter at rates which are comparable to, or less than, that in NGC~4258.  It was realized by several authors that when the accretion rate is low (relative to the Eddington rate), an accretion disk may switch into a hot, radiatively-inefficient mode (Ichimaru 1977; Rees 1982; Narayan \\& Yi 1994; Narayan \\& Yi 1995).  In essence, the plasma becomes so tenuous that the timescale for energy transfer from the protons to the electrons (via Coulomb interactions) becomes longer than the inflow timescale.  The energy remains as thermal energy in the protons (which are very poor radiators) and gets advected through the event horizon of the black hole.  These are the so-called Advection Dominated Accretion Flows (ADAFs).  ADAFs are to be contrasted with `standard' radiatively-efficient accretion disks in which the disk remains cool and geometrically thin all of the way down to the black hole (Shakura \\& Sunyaev 1973; Novikov \\& Thorne 1974).  X-ray observations of broad iron $K\\alpha$ lines have shown that higher luminosity systems do indeed accrete in the radiatively-efficient mode (Tanaka et al. 1995; Fabian et al. 1995).  However, the basic nature of the accretion disk when the mass accretion rate is low is still far from clear.  The existence of the ADAF solution is at the mercy of poorly known physics such as the strength of the electron-ion coupling and the fraction of the viscous energy that is deposited into the electrons (Quataert \\& Gruzinov 1999).  Also, Blandford \\& Begelman (1999) have suggested that the ADAF solutions discussed by Narayan \\& Yi (1994, 1995) are physically inconsistent and necessarily drive powerful outflows (producing the so-called Adiabatic Inflow-Outflow Solutions; ADIOS).  Even if ADAF-type solutions exist, it remains an open question as to whether real disks can make a transition to this mode when the outer regions of the disk are cold and geometrically-thin (as they clearly are in systems such as NGC~4258).  NGC~4258 provides a laboratory in which we can examine all of these issues.  This dichotomy in possible accretion disk physics has produced two models for the central regions of NGC~4258.  By modeling the observed water maser emission in this system, Neufeld \\& Maloney (1995) conclude that the accretion disk in NGC~4258 has a high efficiency ($\\approxgt 10\\%$), and a low accretion rate ($\\dot M/\\alpha \\sim 10^{-4}\\Msunpyr$, where $\\alpha$ is the standard viscosity parameter of accretion disk theory).  Furthermore, the observed maser emission in part traces out a warp in the disk.  Modeling this warp as being driven by radiation pressure from the central X-ray source (Pringle 1996) also implies a radiative efficiency $\\sim 10\\%$ (Maloney, Begelman, \\& Pringle 1996). On the other hand, Lasota et al. (1996) use ADAF models of the 2-10 keV X-ray power law slope and continuum radio data to argue that the system has low efficiency ($\\sim 0.1\\%$) and high accretion rate ($\\dot M/\\alpha \\sim 10^{-2}\\Msunpyr$, $\\alpha \\approxgt 0.3$).  While the original model of Lasota et al. (1996) postulated a large ADAF region ($r\\sim 10^5~GM/c^2$), recent radio data constrain the ADAF region to be smaller than $r\\sim 100~GM/c^2$ (Herrnstein et al. 1998).  In this paper, we report a deep ($\\sim 200$\\,ksec of good exposure time) X-ray observation of NGC~4258 with {\\it ASCA} (Ohasi et al. 1996; Makishima et al. 1996; Yamashita 1997).  We also obtained simultaneous hard X-ray data with the {\\it RXTE} satellite but, due to a gain change prior to our observation, there are currently no robust response matrices or background models for these particular {\\it RXTE} data.  A presentation of these {\\it RXTE} data must await these developments in calibration.  In section 2, we describe the data reduction before discussing the X-ray spectrum of NGC~4258 in Section 3.  Section 4 focuses on the properties of the observed iron line and the implications of this line for the nature of the accretion disk.  Section 5 presents a brief discussion of the other interesting objects in the field of view of the {\\it ASCA} Gas Imaging Spectrometer (GIS).  Our results are summarized in Section 6. ",
        "conclusions": "The X-ray spectrum of NGC~4258 shows 4 distinct components: \\begin{enumerate} \\item a power-law component with a 0.5--10\\,keV luminosity of $7.9\\times 10^{40}\\ergps$ and photon index $\\Gamma=1.79^{+0.31}_{-0.11}$.  This component is readily identified as the primary X-ray emission from the AGN central engine itself.  This emission is absorbed by an intrinsic column density of $N_{\\rm H}=8.2\\pm 0.9\\times 10^{22}\\pcmsq$.  Note that this absorbing column is a factor of 1.4 smaller than that obtained by M94.  Also note that our value of $N_{\\rm H}$ is insensitive to exactly which of the acceptable spectral models we use (cf. Gammie, Narayan \\& Blandford 1999; Chary et al. 2000). \\item thermal plasma emission from optically-thin gas at a temperature of $kT\\sim 0.5\\keV$.  The total luminosity of this component is $\\sim 2\\times 10^{40}\\ergps$, approximately 10--25\\% of the power-law luminosity (depending on how low in energy the power-law extends). \\item a bremstrahlung component with $kT>5\\keV$.   This component may represent thermal emission from a very hot gaseous component in the anomalous arms and/or the integrated hard X-ray emission of the X-ray binary population. \\item a narrow fluroescent K$\\alpha$ emission line of cold iron with an equivalent width of $W_{K\\alpha}=107^{+42}_{-37}\\eV$. \\end{enumerate}  We suggest that the bulk of the narrow iron emission line originates from the accretion disk.  With this assumption, we have shown that the most of the line emission must originate at relatively large radii ($r\\approxgt 100GM/c^2$) in order to produce the small line width. This, in turn, implies a large X-ray source (with a size $D\\approxgt 100GM/c^2$), in contrast to the situation normally found in higher-luminosity Seyfert galaxies (e.g. Nandra et al. 1997).  Note that a model in which the X-ray source is small ($D\\sim 10GM/c^2$) but the inner region of the disk are too hot/ionized to produce line emission has difficulty producing the observed equivalent width of the line. Thus, under the assumption that the bulk of the observed iron line originates from the accretion disk, {\\it there appears to be a clear difference in the size/structure of the X-ray source between this low-luminosity source and higher-luminosity Seyfert galaxies.}  However, we cannot rule out the possibility that the narrow iron line originates from some previously undetected distant matter not associated with the accretion disk.  If the observed line {\\it does} originate from a non-disk source, our data are consistent with but do not require the presence of a ``Seyfert-like'' relativistic broad iron line.  Even in this case, interesting constraints can be placed on the parameter space of possible relativistic broad iron lines.  This AGN is ripe for study with the new generation of X-ray observatories.  The high spatial resolution of Chandra will allow us to study large scale structure in this source directly.  We will be able to pinpoint the location and nature of the thermal component seen in our {\\it ASCA} data.  We will also be able to spatially separate emission from the anomalous arms, the X-ray binary population, and the AGN.  High signal-to-noise spectroscopy with {\\it XMM} (and, in the longer term, {\\it Constellation-X}) will allow the inner accretion disk to be probed."
    },
    "0004/astro-ph0004297_arXiv.txt": {
        "abstract": "We report phase resolved spectroscopy of the transient accreting pulsar, \\srcnm.  For the first time, more than two cylotron resonance scattering features are detected in the spectrum of an \\xray\\ pulsar. The shape of the fundamental line appears to be complex, and this is in agreement with predictions of Monte-Carlo models.  As in other pulsars, the line energies and optical depths are strong functions of pulse phase.  One possible model for this is an offset of the dipole of the neutron star magnetic field. ",
        "introduction": "\\srcnm\\ is a transient accreting \\xray\\ pulsar in an eccentric 24 day orbit \\cite{bil97} with an O9e star \\cite{ung98}. A cyclotron resonance scattering feature (CRSF) was first noted near 20\\,keV by Wheaton, \\etal\\ (1979) \\nocite{whe79} with the UCSD/MIT hard \\xray\\ (10-100\\,keV) experiment aboard \\heaoone. White, Swank \\& Holt (1983) \\nocite{wsh83} analyzed concurrent data from the lower energy (2-50\\,keV) \\heaoone/A2 experiment and found an additional feature at \\aprx12\\,keV, making \\srcnm\\ the first pulsar with two cyclotron line harmonics.  We discuss here phase-resolved spectra derived from an observation of the 1999 March outburst \\cite{wil99,hei99} obtained with the \\emph{Rossi X-Ray Timing Explorer} (\\rxte).  First results of this work are detailed in Heindl \\etal\\ (1999)\\nocite{hei99}.  \\sax\\ has also made detailed observations of this outburst \\cite{San99a,San99b}. ",
        "conclusions": "Table~\\ref{tab:fits} lists the best fit cyclotron line parameters. Two interesting phenomena are apparent from these results.  First, within the individual phases, the lines are not harmonically spaced. And, second, the line energies vary significantly with pulse phase.  A harmonic relation (with small modifications due to relativistic effects) between line energies is expected in simple Landau theory. However, models of cyclotron line formation (e.g. Isenberg, Lamb, \\& Wang 1998\\nocite{ise98}) predict that the fundamental line shape can be quite complex, even having wings resembling \\emph{emission} features.  On the other hand, the higher harmonics should have relatively simple line shapes.  Thus, it may be that our simple Gaussian absorption does not give an accurate measure of the fundamental line energy.  Evidence for this appears in the residuals near 10--15\\,keV (see Fig.~\\ref{fig:delchi}), which are the most significant remaining deviations in our best fit.  Furthermore the third and higher harmonics' energies are consistent with integer multiples of \\emph{half} of the second harmonic energy.  Thus it seems likely that \\emph{half} of the second harmonic energy gives a more accurate measure of the magnetic field.  In the case of phase A, this is $\\rm B = 1.3\\times10^{12}$\\,G, assuming a gravitational redshift of 0.3 to the neutron star surface. \\begin{table} \\caption{\\label{tab:fits} Best fit cyclotron line parameters from the two phase bins shown in Fig.~\\ref{fig:flc}.} \\begin{minipage}{\\linewidth} \\renewcommand{\\thefootnote}{\\thempfootnote} \\begin{center} \\begin{tabular}{lccclccc} \\hline \\hline & \\multicolumn{3}{c}{Phase A} & & \\multicolumn{3}{c}{Phase B}\\\\ \\cline{2-4} \\cline{6-8} Harmonic & Energy & Width\\protect\\footnote{One standard deviation of the Gaussian optical depth profile} & Optical Depth && Energy & Width\\protect\\footnotemark & Optical Depth\\\\ & (keV)  & (keV) & & & (keV)  & (keV)\\\\ \\hline 1\\protect\\footnote{also called the ``fundamental''.} & $\\rm 13.35^{+0.08}_{-0.06}$ & $\\rm  3.29^{+0.13}_{-0.07} $ & $\\rm 1.17^{+0.06}_{-0.04} $ && $\\rm 12.40^{+0.65}_{-0.35}$ & $\\rm  3.3^{+0.19}_{-0.4} $ & $\\rm 0.72^{+0.10}_{-0.17}$ \\\\ 2 & $\\rm 23.7{\\pm0.1} $ &$\\rm 5.43^{+0.18}_{-0.27} $ & $\\rm 2.68^{+0.05}_{-0.07}$ && $\\rm 21.45^{+0.25}_{-0.38} $ &$\\rm 4.5^{+0.7}_{-0.9} $ & $\\rm 1.24^{+0.04}_{-0.06} $\\\\ 3 & $\\rm 36.4^{+0.4}_{-0.5} $ & $\\rm 4.3^{+0.4}_{-0.6}$ & $\\rm 2.41^{+0.11}_{-0.13} $ && $\\rm 33.56^{+0.70}_{-0.90} $ & $\\rm 3.8^{+1.5}_{-0.9}$ & $\\rm 1.01^{+0.13}_{-0.14} $\\\\ 4 & $\\rm 47.8^{+0.4}_{-0.7} $ & $\\rm 5^{+\\infty}_{-1.2}$ & $\\rm 2.3{\\pm 0.2} $ && -- & -- & -- \\\\ 5 & $\\rm 61.7{\\pm 1.1} $ & $\\rm 5 \\emph{fixed} $ & $\\rm 1.8{\\pm 0.3} $ && -- & -- & -- \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table}  Figure~\\ref{fig:linevar} shows the HEXTE flux and the \\aprx20\\,keV line energy and optical depth as a function of pulse phase.  These parameters were determined from the HEXTE data alone in 20 phase bins. The second harmonic line was not required in the HEXTE data \\emph{alone} at phases greater than 0.7.  Both the line energy and optical depth are maximal not at the pulse peak, but on the falling edge of the main pulse.  The line energy varies by \\aprx20\\%.  This behavior is also seen in Cen~X-3 \\cite{hei99b}.  In Cen~X-3, Burderi et al. (2000) \\nocite{Bur00} have modeled the variation of the line energy as due to an offset of the magnetic dipole moment from the center of the neutron star, and this may also be the case in \\srcnm. \\begin{figure} \\centerline{\\includegraphics[width=3.5in,bb=116 82 529 522]{linevar.ps}} \\caption{\\label{fig:linevar}As a function of pulse phase, the energy (top) and depth (bottom) of the \\aprx20\\,keV cyclotron line.  Also plotted is the flux in the HEXTE band.} \\end{figure}     {"
    },
    "0004/astro-ph0004012_arXiv.txt": {
        "abstract": "$BVI_{\\rm C}$ photometry and classification spectroscopy is presented for the young open cluster NGC~6604. Additional Echelle spectroscopy of the brightest members is used to check the reddening against the interstellar NaI and KI absorption lines, to measure the cluster radial velocity and to derive the individual rotational velocities. We obtain 1.7 kpc as the cluster distance, an age of $\\sim 5\\times 10^{6}$ years and a reddening $E_{B-V}$= 1.02 ($\\pm 0.01$ from three independent methods). Pre-ZAMS objects are apparently not detected over the $\\bigtriangleup m$ = 8.5 mag explored range. The cluster radial velocity is in agreement with the Hron (1987) model for the Galaxy disk rotation. \\keywords {Open Clusters: general -- Open Clusters: individual (NGC 6604)} ",
        "introduction": "NGC~6604 ($\\alpha_{2000} =18^{h}~18^{m}.1$, $\\delta_{2000}=-12^{\\circ}~14^{\\prime}$, $l~=~18^{\\circ}.3$, $b~=~+1^{\\circ}.7$) is a fairly compact open cluster belonging to class I3p (Ruprecht 1966) and lies at the core of the HII region S54 (Georgelin et al. 1973). The fairly large extinction to the cluster ($A_{V} \\sim 3$) is surely linked to the partnership with an HII region. The cluster contains several OB type stars (Stephenson and Sanduleak 1971), which suggest the cluster to be fairly young (Forbes and Dupuy 1978). Its distance is rather uncertain, with estimates ranging from 0.7 to 4.4 kpc (Alter et al. 1970). Moffat and Vogt (1975, hereafter MV) presented photoelectric photometry for a dozen stars in the cluster field, deriving a 1.6 kpc distance and a $E_{B-V}$=1.01 reddening. Forbes and DuPuy (1978, hereafter FD) performed photoelectric and photographic photometry, which yielded a 2.1 kpc distance, a $E_{B-V}$=0.96 reddening and a 4x10$^6$ years age. They pointed out the possible presence of pre-ZAMS members which demanded confirmation.  In this paper we report about our CCD $BVI_{\\rm C}$ photometry of the cluster central area, medium resolution spectroscopy of a field larger than the cluster and high resolution Echelle spectroscopy of the brightest cluster members.   \\section {Photometry}   Observations of NGC 6604 have been carried out with the CCD camera mounted on the 1.0 m telescope of the South Africa Astronomical Observatory (SAAO) at Sutherland on July 2, 1992. The journal of observations is given in Table~1 and a finding chart is presented in Figure~1. The surveyed area corresponds to $2.1\\times3.3$arcmin centered on the Of star HD 167971. The reader is referred to Munari and Carraro (1995) for the observing technique employed.  As usual, we observed several E-regions standard stars to determine the extinction coefficients ($K_{B}=0.77\\pm0.04$, $K_{V}=0.52\\pm0.01$ and $K_{I}=0.42\\pm0.02$) and the color equations:  \\begin{eqnarray} B \\ - \\ b \\ &=& \\  -0.002(B-V) \\ - 4.204 \\\\ V \\ - \\ v \\ &=& \\  -0.025(B-V) \\ - 4.387 \\\\ I \\ - \\ i \\ &=& \\  -0.055(V-I) \\ - 4.954 \\\\ B \\ - \\ V \\ &=& \\  +1.024(b-v) \\ + 0.187 \\\\ V \\ - \\ I \\ &=& \\  +1.019(v-i) \\ + 0.584 \\end{eqnarray} \\noindent The resulting $BVI_C$ magnitudes are listed in Table~2 together with their internal errors.  These ones average at $\\sigma \\sim$ 0.0015 mag and increase only for stars fainter than $V = 14$ mag for which photon statistics becomes the dominant source of noise.  Comparison between our CCD profile photometry (DAOPHOT) and the photoelectric observations of MV yields for the seven stars in common: \\begin{eqnarray} V \\ - \\ V_{MV} \\ = \\ 0.034  \\ \\ (\\sigma=0.055) \\\\ (B-V)\\ - \\ (B-V)_{MV} \\ = \\ 0.033  \\ \\ (\\sigma=0.020) \\end{eqnarray}  \\begin{table} \\caption{Journal of observations. The last column is the FWHM of stellar images as measured on the CCD frames.} \\begin{flushleft} \\begin{tabular}{cccc} \\hline \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{Filter} & \\multicolumn{1}{c}{Exp. time} & \\multicolumn{1}{c}{Seeing} \\\\ & &  (sec) & ($\\prime\\prime$) \\\\ \\hline July 9, 1992 & V &  30 &  2.1\\\\ & V & 180 &  2.2\\\\ & B & 300 &  2.4\\\\ & B &  60 &  2.3\\\\ & I &   3 &  2.1\\\\ & I &  30 &  1.9\\\\ \\hline \\end{tabular} \\end{flushleft} \\end{table}  \\begin{figure} \\centerline{\\psfig{file=Barbon_fig1.ps,width=9cm}} \\caption[]{Identification map for the stars with $BVI_{\\rm C}$ photometry in Table~2. The imaged field covers $2.1\\times3.3$ arcmin. North is up, East to the left.} \\end{figure}  The comparison with the photoelectric observations of FD gives for the seven stars in common: \\begin{eqnarray} V \\ - \\ V_{FD} \\ = \\ 0.039 \\ \\ (\\sigma=0.047) \\\\ (B-V)\\ - \\ (B-V)_{FD} \\ = \\ 0.092 \\ \\ (\\sigma=0.028) \\end{eqnarray} Our CCD profile photometry appears slightly fainter and redder than older aperture photometry. It is our opinion that fainter and therefore generally redder stars (quite abundant in a crowded field like that of NGC~6604) have entered and contaminated the older aperture photometry, performed with fixed diaphragms. The different standard stars used (Landolt' equatorial stars for older photometric works, Cousin's E-regions for us) may add a little to the differences in the photometry. The excellent agreement in the interstellar reddening found below via three independent methods gives some support to such arguments.  \\begin{table*} \\caption{$BVI_{\\rm C}$ photometry of NGC~6604. Columns gives our identification number and those used by MV and FD, the X and Y positions on the chip, the corresponding J2000.0 equatorial coordinates (linked to the GSC~I system), and the BVI magnitudes and colours with associated internal errors as provided by the DAOPHOT package.} \\begin{flushleft} \\begin{tabular}{rrrrrrrrrrrrr} \\hline \\multicolumn{1}{c}{ID} & \\multicolumn{1}{c}{MV} & \\multicolumn{1}{c}{FD} & \\multicolumn{1}{c}{X}  & \\multicolumn{1}{c}{Y}  & \\multicolumn{1}{c}{$\\alpha_{J2000}$}  & \\multicolumn{1}{c}{$\\delta_{J2000}$}  & \\multicolumn{1}{c}{$V$}& \\multicolumn{1}{c}{$\\sigma_V$} & \\multicolumn{1}{c}{\\sl B--V}   & \\multicolumn{1}{c}{$\\sigma_{B-V}$} & \\multicolumn{1}{c}{\\sl V--I} & \\multicolumn{1}{c}{$\\sigma_{V-I}$}  \\\\ \\hline 1&    &   D&     276.51&      63.89&  18 18  11.26 &--12 13  54.71 &12.101&      0.001&   0.763&      0.001&    1.025&    0.002 \\\\ 2&    &   M&     128.33&     144.03&  18 18  09.12 &--12 14  52.47 &12.201&      0.001&   0.800&      0.001&    1.065&    0.002 \\\\ 3&  11&    &     142.21&     293.03&  18 18  05.15 &--12 14  46.96 &12.993&      0.001&   0.840&      0.002&    1.271&    0.013 \\\\ 4&  10&    &     171.07&     310.88&  18 18  04.68 &--12 14  35.69 &12.996&      0.002&   0.807&      0.003&    1.108&    0.004 \\\\ 5&    &  51&      69.04&     321.60&  18 18  04.39 &--12 15  15.48 &14.683&      0.005&   0.943&      0.009&    1.175&    0.004 \\\\ 6&    &    &      40.66&     331.84&  18 18  04.12 &--12 15  26.55 &12.873&      0.001&   0.839&      0.002&    1.138&    0.003 \\\\ 7&    &   J&      21.64&     373.07&  18 18  03.02 &--12 15  33.94 &12.227&      0.001&   0.807&      0.001&    1.094&    0.002 \\\\ 8&    &    &     154.74&     378.57&  18 18  02.88 &--12 14  42.02 &13.025&      0.002&   0.767&      0.003&    1.051&    0.004 \\\\ 9&    &    &     160.40&     388.27&  18 18  02.62 &--12 14  39.80 &12.411&      0.001&   0.751&      0.001&    1.052&    0.003 \\\\ 10&    &   2&     283.59&     417.69&  18 18  01.84 &--12 13  51.72 &15.894&      0.013&   1.020&      0.026&    2.586&    0.025 \\\\ 11&    &  72&      94.91&     453.76&  18 18  00.87 &--12 15  05.31 &15.987&      0.013&   1.026&      0.026&    1.394&    0.035 \\\\ 12&   9&   Z&     254.86&     469.18&  18 18  00.47 &--12 14  02.90 &11.658&      0.001&   0.772&      0.001&    1.057&    0.001 \\\\ 13&    &  52&       5.97&     148.61&  18 18  08.99 &--12 15  40.20 &15.187&      0.008&   0.943&      0.014&    1.236&    0.026 \\\\ 14&    &  49&     121.83&     185.94&  18 18  08.00 &--12 14  54.98 &14.971&      0.006&   1.176&      0.013&    1.505&    0.014 \\\\ 15&    &  31&     308.73&     294.42&  18 18  05.12 &--12 13  42.00 &14.478&      0.004&   0.989&      0.007&    1.185&    0.011 \\\\ 16&    &  71&     135.02&     424.02&  18 18  01.67 &--12 14  49.68 &15.282&      0.007&   0.909&      0.013&    1.200&    0.021 \\\\ 17&    &  76&      50.75&     484.57&  18 18  00.05 &--12 15  22.52 &14.259&      0.003&   0.872&      0.005&    1.239&    0.009 \\\\ 18&    &    &      46.77&     322.57&  18 18  04.36 &--12 15  24.17 &14.083&      0.003&   0.869&      0.006&    1.185&    0.009 \\\\ 19&    &    &     143.02&     271.89&  18 18  05.72 &--12 14  46.66 &14.265&      0.006&   0.892&      0.008&    1.271&    0.013 \\\\ 20&    &  50&      89.39&     319.79&  18 18  04.44 &--12 15  07.55 &14.790&      0.009&   0.963&      0.013&    1.240&    0.016 \\\\ 21&   2&   H&     253.65&     244.38&  18 18  06.45 &--12 14  03.52 &10.127&      0.001&   0.707&      0.001&    0.921&    0.001 \\\\ 22&    &    &     278.91&     250.70&  18 18  06.29 &--12 13  53.66 &15.454&      0.011&   1.037&      0.019&    1.238&    0.027 \\\\ 23&   3&    &     175.74&     369.74&  18 18  03.11 &--12 14  33.83 & 9.681&      0.001&   0.729&      0.001&    0.972&    0.001 \\\\ 24&   4&   Y&     228.36&     444.68&  18 18  01.12 &--12 14  13.25 &10.426&      0.001&   0.735&      0.001&    0.984&    0.001 \\\\ 25&    &    &     208.44&     445.42&  18 18  01.10 &--12 14  21.02 &15.112&      0.010&   0.853&      0.016&    1.233&    0.019 \\\\ 26&   1&   X&     179.19&     265.92&  18 18  05.88 &--12 14  32.55 & 7.524&      0.001&   0.812&      0.001&    1.116&    0.001 \\\\ 27&    &    &     195.26&      90.86&  18 18  10.54 &--12 14  26.39 &13.716&      0.017&        &           &    2.890&    0.017 \\\\ \\hline \\end{tabular} \\end{flushleft} \\end{table*} ",
        "conclusions": ""
    },
    "0004/astro-ph0004224_arXiv.txt": {
        "abstract": "New near-infrared galaxy counts in the $J$ and $K$ bands are presented over a total area of 0.70 and 0.97 degrees$^{2}$, respectively.  The limiting magnitudes of the deepest regions are 19.5 in $J$ and 18.0 in $K$.  At $J>16$ and $K>15$ our $J$ and $K$-band counts number counts agree well with existing surveys {\\em provided} all data are corrected to a common magnitude scale.  There are real differences from field to field, and the ELAIS N1 and N2 fields show an overdensity of $J<16$, $K<15$ galaxies.  The slopes of $\\log N(m) / dm$ are $\\sim$ 0.40 to 0.45 at $15 < K < 18$ and $16 < J < 19.5$.  Our counts favor galaxy models with a high normalization of the local luminosity function and without strong evolution. ",
        "introduction": "Counting galaxies as a function of magnitude has established itself as a powerful way to study galaxy evolution.  This is especially true in the near-infrared, where K-corrections are smaller and dust extinction of less importance than in the visible.  There have been numerous $K$-band galaxy surveys for over a decade now. These have shown the pitfalls of deriving any definite evolutionary conclusions from a single wavelength: e.g.\\ the excess of faint galaxies seen especially in blue bands (see Ellis 1997 for a comprehensive review) is not nearly as severe in $K$-band.  Another near-infrared waveband, the 1.25$\\mu$m $J$-band, is now added to the available data.  Bershady et al.\\ (1998) and Saracco \\etal\\ (1999) have published deep $J$-counts, and there are preliminary wide field DENIS- (Mamon 1998) and 2MASS-counts (Skrutskie 1999) available at $J<16$.  Ultra-deep HST NICMOS counts are also available in $J$ and $H$ bands (Thompson et al.\\ 1999, Yan et al.\\ 1998, Teplitz et al.\\ 1998). The only intermediate magnitude $J$-counts to date are those of Teplitz \\etal\\ (1999); our survey covers a slightly shallower magnitude range but has a sky coverage 15 times as large.  The range of magnitudes, $J$=12.5--19.5, in the present survey is too bright to be of direct interest to cosmology and high redshift galaxy evolution studies.  However, modeling of the faintest counts to study the evolution of distant galaxy populations heavily relies on accurate normalization of the models at brighter magnitudes. Indeed, the normalization is often left as a free parameter, signifying that the local galaxy population is still quite uncertain.  The bright and medium counts are thus instrumental in the studies of both the local galaxy population and the evolving properties of galaxies with large look-back times.  Our survey fields are within the European Large Area ISO-Survey (ELAIS) regions (see Oliver et al.\\ 2000).  ELAIS fields will be heavily studied in wavelengths ranging from X-rays to radio bands, with much of the scientific interest lying in high-redshift, dust-obscured star formation, AGN, and ultraluminous infrared sources.  Our observations were done as part of NIR followup for the ELAIS project.  This paper presents our galaxy counts and compares them to some available data and models.  We also address the question of whether this field, selected because of its low 100~$\\mu$m emission, is a representative sample of the extragalactic sky.  The discussion of individual objects relating to ISO-sources will be presented elsewhere. ",
        "conclusions": "Our $J$ and $K$-band galaxy counts in two ELAIS fields (N1,N2) represent the largest areas to date in the ranges $15 < J < 19.5$ and $16 < K < 18$. The $J$-band counts are the first wide field galaxy counts at this magnitude range. For $J>16$, $K>15$, the data are consistent with existing surveys {\\em provided} significant magnitude scale corrections and large-scale structure effects are taken into account. In particular, the N2 region has a 10\\% overdensity of galaxies due to rich clusters.  The fields of the large Hawaii survey (Huang et al.\\ 1997) seem to have a systematic underdensity of about a factor of 1.3, which explains their counts without the need for a local void.  At $J<16$ and $K<15$, the counts in the ELAIS fields are higher than previous results by up to a factor of two.  This is probably due to a real overdensity of field galaxies in the regions, although nearby large galaxy clusters outside the fields may affect the very brightest magnitude bins.  This overdensity needs to be taken into account when interpreting results of surveys in other wavelengths in these regions.  Our galaxy counts favor a high normalization of the local LF. The slope of the counts, $d \\log N / dm \\approx 0.40 - 0.45$, together with the amplitude of the counts at $K$=15--18 and $J$=16--19 are best fit by minimally evolving galaxy models and luminosity functions having relatively large numbers of faint galaxies."
    },
    "0004/astro-ph0004154_arXiv.txt": {
        "abstract": "We present moderate resolution data for 39 quasi-stellar objects (QSOs) at $z \\approx 2$ obtained at the Multiple Mirror Telescope.  These data are combined with spectra of comparable resolution of 60 QSOs with redshifts greater than $1.7$ found in the literature to investigate the distribution of Ly$\\alpha$ forest lines in redshift and equivalent width. We find a value for $\\gamma$, the parameter describing the number distribution of Ly$\\alpha$ forest lines in redshift, of $1.88\\pm0.22$ for lines stronger than a rest equivalent width of 0.32 $\\AA$, in good agreement with some previous studies.  The Kolmogorov-Smirnov test was applied to the data and it is found that this single power law is a good fit over the relevant redshift ranges. Simulations of the Lyman $\\alpha$ forest were performed to determine the completeness of the line lists and to test how well the analysis recovers the underlying line statistics, given this level of completeness. ",
        "introduction": "The spectra of quasi-stellar objects (QSOs) blueward of Ly$\\alpha$ emission show a large number of absorption lines primarily due to Ly$\\alpha$ absorption by intervening neutral hydrogen along the line of sight to the QSO (Lynds, 1971; Sargent et al. 1980; Weymann, Carswell, \\& Smith, 1981). The evolution of the number density of these lines can be described by a power law in $1+z$ with index $\\gamma$.  It has been debated whether the number densities of absorbers of different column densities, Lyman $\\alpha$ forest, Lyman limit systems, and damped Lyman $\\alpha$ systems, evolve differently with redshift which would indicate that these absorbing structures represent different populations of objects (Lanzetta 1991, Storrie-Lombardi 1994,1995, Giallongo et al. 1996).  Several authors have carried out the analysis of the statistics of the Lyman $\\alpha$ forest on other data sets (Lu et al. 1991, hereafter LWT; Bechtold 1994, hereafter B94; Williger et al. 1994; Cristiani et al. 1995; Giallongo et al. 1996; Kim et al. 1997; Weymann et al. 1998.) At high redshift, LWT find $\\gamma= 2.75 \\pm 0.29$; and B94, whose resolution class M sample is a subsample of our total sample, finds $\\gamma=1.85 \\pm 0.27$. These results were obtained using a fixed equivalent width threshold of 0.36 $\\AA$ and using only Ly$\\alpha$ lines outside the immediate vicinity of the QSO. Work with high resolution spectra has yielded results similar to those of both LWT and B94. Cristiani et al. (1995) find $\\gamma= 1.86\\pm 0.21$ for Lyman $\\alpha$ forest lines with log(N$_{HI}) \\geq 13.8$ from R $\\sim$ 22,000 spectra of 3 QSOs; and Giallongo et al. (1996) find $\\gamma= 2.7$ for lines with log(N$_{HI}) \\gtrsim 14$ from R $\\sim$ 25,000 spectra of 10 QSOs. Similarly, Kim et al. (1997)  report $\\gamma= 2.78$ from a sample of R $\\sim$ 36,000 spectra of 5 QSOs taken using the HIRES Spectrograph on the Keck Telescope.  The evolution is significantly flatter at redshifts less than 1.7. The first Hubble Space Telescope Quasar Absorption Line Key Project paper (Bahcall et al. 1993) reported $\\gamma = 0.50 \\pm 0.77$ for lines from 13 objects with W $>$ 0.32 $\\AA$.  The addition of four more quasars at z $<$ 1.3 yielded $\\gamma= 0.58 \\pm 0.50$ for lines with W $>$ 0.24 $\\AA$ (Bahcall et al. 1996). Weymann et al. (1998) have analyzed the complete HST Key Project Ly$\\alpha$-absorption line sample presented by Jannuzi et al. (1998). They find $\\gamma= 0.15 \\pm 0.23$ for a fixed equivalent width threshold of 0.24 $\\AA$.    From a large sample of archival HST/Faint Object Spectrograph (FOS) spectra which includes the public data from the Key Project, Dobrzycki \\& Bechtold (1997) find $\\gamma = 0.87 \\pm 0.27$ for lines with W $>$ 0.36 $\\AA$ at z $<$ 1.7, and notes that this value flattens to $\\gamma = 0.14 \\pm 0.45$ for Ly$\\alpha$ lines with equivalent widths between 0.24 and 0.36 $\\AA$.  This is supported by Weymann et al. (1998) who also find a trend of increasing $\\gamma$ with increasing line strength.  In this paper, a homogeneous sample of moderate resolution QSO spectra is used to investigate the number density evolution of Ly$\\alpha$ systems and how this changes with redshift and with varying equivalent width thresholds. Specifically, the Ly$\\alpha$ forest in the redshift range between 1.6 and 2.0  was targeted because few lines of sight in the literature cover this range, as it extends down to wavelengths of $\\sim$3200 $\\AA$. Improvements in CCD technology allowed us to obtain data in this spectral region. We present new data for 39 objects and supplement this sample with 60 objects from the literature. Metal line systems in these spectra have been identified and removed from the final analysis of the Ly$\\alpha$ forest (Murdoch et al. 1986, hereafter MHPB). The resulting Ly$\\alpha$ absorption line sample is comprised of 2079 lines in the range 1.6 $<$ z $<$ 4.1 when a variable equivalent width threshold is used, or 1131 lines using a fixed rest equivalent width threshold of 0.32 $\\AA$. Future papers will extend this work to the local universe using a sample of low-redshift spectra from the HST/Faint Object Spectrograph (FOS) archives. Much recent work has promoted treating the Ly$\\alpha$ forest as density fluctuations in the continuous field of the intergalactic medium (Bi 1993; Reisenegger \\& Miralda-Escud\\'{e} 1995; Hernquist et al. 1996; Miralda-Escud\\'{e} et al. 1996; Bi \\& Davidsen 1997; Croft et al. 1998,1999; Weinberg et al. 1999). However, in this paper, we will continue to interpret the Ly$\\alpha$ forest as a series of discrete lines for comparison to previous work. The spectra presented in this data set will be of particular use to those interested in using the Ly$\\alpha$ forest to recover the power spectrum of mass fluctuations in the universe (Croft et al. 1998,1999) and are available on the web. For the spectra used in this paper, see the authors' homepages, which are accessible from either http://www.as.arizona.edu or http://cfa-www.harvard.edu.  In Section~\\ref{sec-absdata}, the details of the observations are outlined and the absorption line sample is presented. Notes on individual sample QSOs and their corresponding absorption line systems are given in Section~\\ref{sec-notes}. In Section~\\ref{sec-results} the statistics of the Ly$\\alpha$ forest are derived and discussed. Lyman $\\alpha$ forest simulations based on our sample of QSOs were performed to test the ability of our analysis to recover the underlying line statistics in our moderate resolution data. These are also presented and discussed in Section~\\ref{sec-results}. In a subsequent paper, Paper II in this series (Scott, Bechtold, Dobrzycki, \\& Kulkarni 2000), we will investigate the measurement of the mean  intensity of the ionizing background radiation using this Ly$\\alpha$ forest sample. ",
        "conclusions": "\\label{sec-results}  The number of Ly$\\alpha$ lines per unit redshift per unit equivalent width can be parametrized as follows: \\begin{equation} \\frac{\\partial^{2} {\\cal N}}{\\partial z \\partial W}= \\frac{A_{0}}{W^{*}}(1+z) ^{\\gamma}\\exp\\left(-\\frac{W}{W_*}\\right). \\end{equation} Integrating this equation over equivalent width with a constant threshold equivalent width throughout each spectrum gives \\begin{equation} \\frac{d{\\cal N}}{dz}= {\\cal A}_{0}(1+z)^{\\gamma} \\label{eq:dndz} . \\end{equation} To solve for the parameters $\\gamma$ and $W^{*}$, we use a maximum likelihood technique which allows for an equivalent width threshold that varies with wavelength.  We also derive these parameters using various fixed threshold values; and in this case, the procedure reduces to the method described in the Appendix of MHPB, using corrected expressions for their equations (A8) and (A2a). However, the variable threshold information is still used in the fixed threshold case, as regions of the spectrum for which the threshold lies above the fixed value, ie. where not all significant lines could be detected even if they were present, are excluded.  The solutions for the statistics $\\gamma$ and $W^{*}$ are listed in Table~\\ref{table-paramtab}. Each sample excludes regions of the spectra within $\\Delta$z of 0.15 of the QSO emission redshift, chosen to eliminate any effects on the line density due to proximity to the QSO. A variable equivalent width threshold gives a value of $1.23 \\pm 0.16$ for $\\gamma$.  This is lower than the value of $2.75 \\pm 0.29$ found by LWT for for a fixed equivalent width threshold of 0.36 $\\AA$ over the range $1.7 < z < 3.8$, and the value of $1.89 \\pm 0.28$ found by B94 for a fixed threshold of 0.32 $\\AA$ over the range $1.6 < z < 4.1$.  Using a fixed threshold of 0.32 $\\AA$, the value of $\\gamma$ derived from our data is $1.88 \\pm 0.22$, in good agreement with that of B94. In Table~\\ref{table-paramtab}, no error is quoted for ${\\cal A}_{0}$ because it is strongly correlated with the error in $\\gamma$.  We calculate the Kolmogorov-Smirnov (KS) probability that a power law number density distribution given by Equ.~\\ref{eq:dndz} for each of these values of $\\gamma$ is a good representation of the data (cf. Appendix of MHPB). A high probability (P$_{KS}$) that the maximum deviation from the cumulative number distribution could occur by chance if the data set is drawn from an assumed parent distribution indicates that the choice of parent distribution is justified. These results are included in Table~\\ref{table-paramtab}. The total sample and each of the subsamples is described well by a single power law, as illustrated by the high KS probabilities obtained. The KS probability obtained from our data set with a fixed equivalent width threshold of 0.32 $\\AA$ and the LWT $\\gamma$ value of 2.75 is 0.0020, while the B94 value of 1.85 gives 0.97, as it is in good agreement with our maximum likelihood result.  The errors in $\\gamma$ and $W^{*}$ are calculated by our software by fitting a parabola to the peak of the logarithm of the likelihood function, using the fact that the likelihood function itself should be distributed as a Gaussian  in $\\gamma$ and $W^{*}$ near its maximum value.   In order to avoid any assumptions about the distribution of the statistics of interest, a resampling technique was used to independently estimate the distribution. Jackknife samples (Babu \\& Feigelson 1996, Efron 1982) of our original data set were constructed, 100 in all, each with one QSO from the original sample removed.  We used the same program to calculate $\\gamma$ and $W^{*}$ for each jackknife sample, for the case of $W_{thr}$=0.32 $\\AA$.   The goal is to understand how the values of these statistics found by our software vary with random variations in the data. The weighted mean of all the jackknife values for $\\gamma$ is 1.91 and for $W^{*}$ it is 0.309 $\\AA$. Since we cannot treat each of the 100 values of these statistics as independent measurements of $\\gamma$ and $W^{*}$, the jackknife errors show how well the error calculated by the software estimates the true distribution of the statistics calculated.  The jackknife results for $\\sigma_{\\gamma}$ and $\\sigma_{W^{*}}$ are 0.26  and 0.011  respectively. The fact that the jackknife errors are $\\sim$20\\% larger than the error calculated by our software may reflect the fact that the jackknife estimate of the variance tends to be conservative (Efron 1982) or it may indicate the the presence of additional sources of random error. In any case, the jackknife results do agree with the total data set result to well within the errors.  The two questions we now ask are whether the number densities of strong and weak lines evolve differently with redshift and whether there is a difference in $\\gamma$ for low and high redshift subsamples, ie. does $\\gamma$ evolve over the history of the universe after the observed break at z $\\sim$ 1.7? In this context, strong lines will refer to lines with rest equivalent widths greater than 0.32 $\\AA$ and weak lines will refer to those with rest equivalent widths between 0.16 $\\AA$ and 0.32 $\\AA$. The total absorption line sample was divided into low and high redshift subsamples  at an absorption redshift of 2.5, giving 1084 and 995 lines in each subset, respectively. For the remainder of this paper, the low and high redshift subsamples will refer to Lyman $\\alpha$ forest absorption lines with redshifts above and below 2.5, respectively.  Figure~\\ref{fig:dndz} is a set of plots of log$(d{\\cal N}/dz)$ versus log$(1+z)$ for the various subsamples of our data set which are binned solely for display purposes.  The straight lines are derived from the parameters given in Table~\\ref{table-paramtab}. Figure 4a shows the low and high redshift subsamples and the solutions for each along with the solution for the total sample.  Each of these are generated with a fixed equivalent width threshold of 0.32 $\\AA$.  Figure 4b shows the results for strong (W $>$ 0.32  $\\AA$) and weak lines (0.16 $<$ W $<$ 0.32 $\\AA$) considered separately. Column 8 of Table~\\ref{table-paramtab} lists the KS probabilities for each case considered.  No log$(d{\\cal N}/dz)$ versus log$(1+z)$ plots are shown and no KS probabilities are quoted for any case in which a variable threshold was used.  This is because the distribution in redshift is now related to the equivalent width of each line.   The separation of these two distributions, which is possible in the case of a constant threshold, is not possible; and the formalism of MHPB can no longer be applied.  Nevertheless, since the implementation of a variable threshold allows the most efficient use of the data, we consider these values of $\\gamma$ to be reliable, especially in light of the reasonable KS probabilities in the constant threshold cases.  Considering the moderate resolution and signal-to-noise of our data, it is worth investigating how well we are recovering the true parameters describing the line distribution.  Recall from the discussion of the simulations in Section~\\ref{sec-absdata} that our 5$\\sigma$ line lists are 55\\% complete due to blending. To address this point, we generated more sets of artificial spectra based on the 56 objects in our data set for which we have detailed spectral information in the way described in Section~\\ref{sec-lineid}. The redshift of each QSO in these sets is equal to that of one of the 39 new MMT spectra presented in this paper or to that of one of the 17 spectra presented in Dobrzycki and Bechtold (1996). In order to investigate how signal-to-noise impacts this analysis, we created three sets of these 56 artificial spectra with the resolution of the data, $\\sim$1 $\\AA$, one set having signal-to-noise ratios half that of the data (median S/N $\\sim$ 5), another having signal-to-noise ratios equal to that of the data (median S/N $\\sim$ 10), and another having twice the signal-to-noise of the data (median S/N $\\sim$ 20). The input parameters used were $\\gamma$=1.88, $\\beta$=1.46, N$_{lower}$=1x10$^{13}$ cm$^{-2}$, N$_{upper}$=1x10$^{16}$ cm$^{-2}$, $<$b$>$=28 km s$^{-1}$, and $\\sigma_{b}$= 10 km s$^{-1}$.  In the low S/N simulation, FINDSL spuriously identified one simulated line, out of 1722 lines above threshold, as two separate lines, both of 5$\\sigma$ significance or greater.  This did not occur in either the data S/N simulation, or in the high S/N simulation, so we remain confident that the Lyman $\\alpha$ lines in our line lists are real absorption features.  We also generated set of synthetic spectra with higher resolution than the data. Two sets were made with resolution $\\Delta\\lambda\\sim$ 0.7 $\\AA$, one with the same signal-to-noise as the data, and another with median S/N $\\sim$ 20. Finally, a Keck/HIRES data set was simulated by generating spectra with $\\Delta\\lambda\\sim$ 0.2 $\\AA$ and median S/N $\\sim$ 40.  The simulation line lists were analyzed in the same way as the data to determine the value of $\\gamma$ input into the FINDSL analysis. This $\\gamma$ is not necessarily equal to the simulation input $\\gamma$, 1.88, because, in generating the artificial spectra, the simulation software does not fix the redshift and equivalent width distributions by the input parameters, but rather draws line redshifts and equivalent widths from a distribution given by Equation~\\ref{eq:dndz}. FINDSL line lists were then generated and $\\gamma$ was calculated again using these line lists. This was done for both the variable threshold and the case of an equivalent width threshold of 0.32 $\\AA$ for all redshifts, and at high and low redshifts separately. The two values of $\\gamma$ for each case are compared with each other in order to determine how well the redshift distribution in the FINDSL line lists reflects the distribution output by the simulations. The results are listed in Table~\\ref{table-sim}. The simulation resolution and median signal-to-noise ratio are given in the first two columns; the redshift range and the threshold used for the $\\gamma$ solution are given in columns (3) and (4); and the values of $\\gamma$ derived from the simulation line lists ($\\gamma_{\\rm simulation}$) and from the FINDSL line lists ($\\gamma_{\\rm FINDSL}$) are given in columns (5) and (6), respectively. DB96 discuss this simulation software in detail and use it to investigate the column density distribution of Lyman $\\alpha$ lines.  Their data set, a subset of ours, encompassed a limited redshift path and was therefore insensitive to a determination of $\\gamma$ from the simulations. Presumably, if we ran the large number of simulations for which this software was designed, we would recover $\\gamma$=1.88 in column (5) of Table~\\ref{table-sim}; but since we are merely trying to determine the reliability of our methods for identifying significant lines, we will leave this for future work. These Monte-Carlo simulations create line lists by distributing lines according to the input value of $\\gamma$, which is independent of redshift and equivalent width.  It is for this reason, and because we have created a relatively small number of synthetic spectra in order to simulate our data set, that we do not take the values of $\\gamma$ derived either from the simulation line lists or from the FINDSL line lists to truly reflect the redshift distribution of Lyman $\\alpha$ lines.  We use these simulations only to investigate how well our techniques for identifying significant lines and calculating $\\gamma$ recovers the value input into the FINDSL analysis.  Figure~\\ref{fig:gamma} also demonstrates these results.  It shows the number of sigma difference between the output (FINDSL line lists) and input (simulation line lists) values of $\\gamma$, (a)-(c) for the variable threshold case and (d)-(f) for the W$_{thr}$=0.32 $\\AA$ case. The square points and solid lines indicate the results for the simulations  at the resolution of the data in this paper, $\\Delta\\lambda\\sim$ 1 $\\AA$.  The open triangles and dotted lines show the results for the simulations at higher resolution,  $\\Delta\\lambda\\sim$ 0.7 $\\AA$; and the filled triangle shows the result for the Keck/HIRES simulation, $\\Delta\\lambda\\sim$ 0.2 $\\AA$.  Histograms of the line distributions used in the input (simulation line lists) and output (FINDSL line lists) $\\gamma$ solutions are shown in Figure~\\ref{fig:hist}(a-d). Also, plots of log$(d{\\cal N}/dz)$ versus log$(1+z)$ analogous to those in Figure~\\ref{fig:dndz} for the data resolution, data S/N, constant threshold simulations are shown in Figure~\\ref{fig:dndzsim}.  As in Figures~\\ref{fig:hist}(a-d), the solid lines correspond to the maximum likelihood solution for $\\gamma$ and ${\\cal A}_{0}$ for the simulation line lists and the dashed lines correspond to the solution for the FINDSL line lists.  These figures demonstrate that the process of simulation lines above threshold being $``$blended out\" with other features in the final FINDSL line lists dominates over lines below threshold being $``$blended in\" by blending with other features below threshold in all cases. Overall, therefore, the FINDSL line lists suffer from a net loss of lines due to the blending out of significant features.  However, this blending has not significantly affected the value of $\\gamma$. The only case for which the simulation and subsequent FINDSL solutions for $\\gamma$ differ by more than 1.5$\\sigma$, indicated by the dashed-dotted lines in Figure~\\ref{fig:gamma}, is the constant threshold solution for the lowest S/N simulation at low redshift, the leftmost point in Figure~\\ref{fig:gamma}(b).  It should be noted that some visual inspection of the simulation spectra was necessary to achieve this overall agreement between the simulation and the FINDSL $\\gamma$'s. This examination was commensurate with that done on the data, especially during the course of the metal line identifications, so no significant bias is introduced into the simulation analysis by doing this. The FINDSL program tended to miss some weak lines in the high redshift spectra due to crowding of features.  Some lines were also missed  by FINDSL at low redshift, where the signal-to-noise is lowest.  The equivalent width thresholds used in the solution for $\\gamma$ required that the weakest lines at low S/N be left out of the simulation line list solution, so missing them with FINDSL had little effect.  However, in some cases, FINDSL either failed to find lines above threshold at low S/N or failed to fit them with the proper equivalent width. These omissions did adversely affect the agreement between the $\\gamma$ solutions, as these lines were included in the solution using simulation line lists. Upon inspection of the simulated spectra, all of these lines were identified and the simulation and FINDSL line list solutions for $\\gamma$ were brought into agreement.  For the total sample and the high redshift subsample, including weak lines in the maximum likelihood solution tends to make $\\gamma$ more shallow. Both our data and  high resolution work (Cristiani et al. 1995, Giallongo et al. 1996) indicate that the tendency for $\\gamma$ to change in either direction when weaker lines are included is not a significant one. Decreasing the column density cutoff from log(N$_{HI}$)=13.8 to 13.3 at z$\\sim$3, Cristiani et al. (1995) find that $\\gamma$ increases from 1.86 to 2.17; but this is a change of less than 1$\\sigma$.  Giallongo et al. (1996) find that decreasing the column density cutoff from log(N$_{HI}$)=14 to 13.3, again at z$\\sim$3, decreases $\\gamma$ from 2.7 to 2.49, $\\sim$1$\\sigma$. However, using only weak lines for our total sample gives a $\\gamma$ of $0.26\\pm0.33$, a value consistent with no evolution for $q_{0}= 0.5$; while using all lines with rest equivalent width greater than 0.32 $\\AA$ gives a value 4$\\sigma$ larger, 1.88.  In the case of the high redshift subsample, this difference is 2.6$\\sigma$.  Weak lines being blended out in the crowded, high redshift regions of the spectra is undoubtedly contributing to this effect. In our simulations, lines with rest equivalent widths between 0.16 $\\AA$ and 0.32 $\\AA$ yield a $\\gamma$ of 2.25$\\pm$0.40 for the simulation output lines, while the FINDSL line lists give a significantly lower value of 1.30$\\pm$0.49. The plots of log$(d{\\cal N}/dz)$ versus log$(1+z)$ analogous to Figure~\\ref{fig:dndzsim} for these weak lines are shown in Figure~\\ref{fig:dndzsimweak}; and this solution for all redshifts is shown in panel (a). Recall that the input simulation redshift distribution is independent of the line width.  By contrast, the simulation line list and FINDSL line list values of $\\gamma$ for lines with equivalent widths greater than 0.32 $\\AA$  are 1.62$\\pm$0.27 and 1.70$\\pm$0.30, respectively. This indicates that though we can be confident that we are recovering the true $\\gamma$ for lines with equivalent widths greater than 0.32 $\\AA$, weak lines blended out at high redshift in our data may indeed produce this flattening of $\\gamma$ seen when weak lines are included in the solution.  For the  low redshift subsample, the  weak lines give a steeper $\\gamma$, but this difference is not statistically significant. The Weymann et al. (1998) results at z $<$ 1.7 suggest the opposite, that lines of higher rest equivalent width yield larger values of $\\gamma$. These authors find a difference in the evolution rates for Ly$\\alpha$ absorbers with and without identified associated metal lines. Their interpretation of this is that it can be attributed to a difference in the rate of evolution of lines of different strengths. This scenario is supported by the higher redshift results of Kim et al. (1997). Their high resolution data suggest that there is a break in the column density distribution of Ly$\\alpha$ lines at log(N$_{HI}$) $\\geq$ 14.8 and z $\\sim$ 3.3 and that this break occurs at lower column densities and becomes more pronounced as redshift decreases.  These results imply that weak lines should show a flatter $\\gamma$ at all redshifts and that the difference in the rate of evolution between strong and weak lines should be more significant at redshifts less than 2.5 than at redshifts greater than 2.5.  The $\\gamma$'s derived from the simulation and FINDSL line lists for the spectra generated at the data resolution and signal-to-noise listed in Table~\\ref{table-sim} are generally in good agreement with one another for strong and weak lines at low and high redshift, noting however, the large uncertainties for the weak line $\\gamma$'s. The FINDSL $\\gamma$'s for weak lines for all redshifts and at low redshift are systematically lower than the simulation $\\gamma$'s, due to blending out of weak features preferentially at high redshift. The high redshift solution does not suffer from this as lines are evenly blended out at all redshifts greater than 2.5, as demonstrated in Figure~\\ref{fig:dndzsimweak}(a-c). In any case, this comparison indicates that there is no tendency for blending to work to artificially produce the trend noted above, namely that the $\\gamma$ for weak lines is steeper than the $\\gamma$ for strong lines at low redshift, contrary to the results of other authors.  For a variable threshold at high redshifts, $\\gamma$ flattens by 1.5$\\sigma$ compared to the value found for low redshift lines, to $0.64 \\pm 0.47$ for z $>$ 2.5 from $1.57 \\pm 0.42$ at $z < 2.5$. Again, the difference is not statistically significant; but a trend exists in that the maximum likelihood values of $\\gamma$  found for the low redshift subsample are larger than those found for the high redshift subsample in all cases in which weak lines are included, while for strong lines, $\\gamma$ increases from low to high redshift. The agreement between the $\\gamma$'s derived from the simulation line lists and the FINDSL line lists indicates that, at the resolution and signal-to-noise of the data, this trend is not artificially imposed by blending.  Equivalently, one can investigate the distribution in equivalent width as a function of redshift.  The value of the parameter $W^{*}$ increases from low to high redshift from 0.282 $\\AA$ to 0.330 $\\AA$ in the case of a constant 0.32 $\\AA$ threshold, a difference of $\\sim$3$\\sigma$ in the sense that the distribution is more shallow at high redshift.  Both of these results imply that there exist more weak lines relative to strong ones at low redshift than at high redshift. Given the discussion above, it is likely that at least some of this difference can be attributed to increased blending of weak lines at high redshifts. Nevertheless, the Kim et al. (1997) analysis supports this interpretation, as do the results of the hydrodynamic simulations of Dav\\'{e} et al. 1999. These authors find that $W^{*}$ does indeed increase with redshift from z=0 to z=3 due to the onset of structure formation. The values of $W^{*}$ they derive from their simulated spectra at high resolution are smaller than those measured in this paper or at low redshift by Weymann et al. (1998). They find, however, that the effects of blending in even low redshift, moderate resolution spectra, comparable to the FOS data, can raise the measured values to those found by Weymann et al. (1998).  This effect is demonstrated by Figure~\\ref{fig:ewhist}, a histogram of the rest equivalent width distribution of lines in the simulation and FINDSL line lists for the data resolution ($\\sim$1 $\\AA$) simulations with median signal-to-noise ratios of 5, 10, and 20 in the variable threshold case. As expected, the number of lines blended out is largest at low equivalent width, flattening out the overall distribution and in turn raising the value of $W^{*}$ derived.  If a fixed equivalent width threshold of 0.32 $\\AA$ is used (rows 9, 11, 15, and 17 in Table~\\ref{table-paramtab}), weak lines are thrown out and the distribution in redshift is flatter at low redshift than at high redshift, though not significantly so: $\\gamma=1.30 \\pm 0.60$ for z $<$ 2.5, versus $1.69 \\pm 0.60$ for z $>$ 2.5, a difference of less than 0.5$\\sigma$. Interestingly, Stengler-Larrea et al. (1995) find $\\gamma= 1.50 \\pm 0.39$ for Lyman limit absorbers between z=0.32 and z=4.11, in reasonable agreement with our values of $\\gamma$ using W$_{thr}$=0.32 $\\AA$ for both the low and high redshift subsamples. The total sample of lines with W $>$ 0.32 $\\AA$ gives a somewhat larger value of $\\gamma$, 1.88 $\\pm$ 0.22, but including the low redshift data of Bahcall et al. (1993) yields a value of 1.70 $\\pm$ 0.19, consistent with the result for Lyman limit systems. It has been proposed that Ly$\\alpha$ absorbers with log(N$_{HI}$) $\\gtrsim$ 14, the value of the break in the column density distribution, are associated with the outer halos of galaxies responsible for Lyman limit systems and damped Lyman $\\alpha$ systems (Giallongo et al. 1996, Lanzetta et al. 1995, 1996, Chen et al. 1998). This column density is approximately equivalent to the equivalent width threshold of 0.32 $\\AA$ used in this study; and the agreement between our values of $\\gamma$ and that for Lyman limit systems lends some credence to this scenario."
    },
    "0004/astro-ph0004362_arXiv.txt": {
        "abstract": "Because we know little about the Galactic force-field away from the plane, the Galactic mass distribution is very ill-determined. I show that a microlensing survey of galaxies closer than $50\\Mpc$ would enable us to map in three dimensions the Galactic density of stellar mass, which should be strictly less than the total mass density. A lower limit can be placed on the stellar mass needed at $R<R_0$ to generate the measured optical depth towards sources in the bulge. If the Galaxy is barred, this limit is lower by a factor of up to two than in the axisymmetric case.  Even our limited knowledge of the Galactic force field suffices to rule out the presence of the amount of mass an axisymmetric Galaxy needs to generate the measured optical depth. Several lines of argument imply that the Galaxy is strongly barred only at $R<4\\kpc$, and if this is the case, even barred Galaxy models cannot generate the measured optical depth without violating some constraint on the Galactic force-field. Galactic mass models that are based on the assumption that light traces mass, for which there is significant support in the inner Galaxy, yield microlensing optical depths that are smaller than the measured value by a factor of more than 2.5. ",
        "introduction": "Notwithstanding the difficulties to which heavy obscuration by dust gives rise, we now have a reasonable idea of how luminosity is distributed in the Galaxy. We know we live in a galaxy that has a bar, $3-4\\kpc$ long, whose nearer end lies at positive longitudes. Within the bar there is a moderately flattened component in which the luminosity density rises roughly as $r^{-1.8}$ with decreasing galactocentric radius $r$. Around the bar there is a roughly exponential disk with scale length $\\sim3\\kpc$, and a moderately flattened metal-poor halo in which the luminosity density varies as $\\sim r^{-3.5}$.  The situation regarding the Galactic distribution of mass is very much less satisfactory. Observations of external galaxies and cosmological theory have convinced us that light is not a good tracer of mass. Hence we cannot straightforwardly translate our models of the luminosity distribution into models of the mass distribution; in principle, we should start afresh and derive the mass distribution by tracing out the Galaxy's gravitational field $\\b F(\\b r)$ and then applying the divergence operator: $\\nabla\\cdot\\b F=-4\\pi G\\rho$.  From observations of gas that flows in the Galactic plane we have a fair idea of what $\\b F(\\b r)$ looks like at points $\\b r$ within the plane, but we have very little secure knowledge of $\\b F$ out of the plane. This deficiency is serious, because until we know $\\b F$ away from the plane, we cannot apply $\\nabla$ {\\em anywhere}, and can say {\\em nothing\\/} with security about $\\rho$ {\\em anywhere\\/}!  What we currently know about $\\b F$ out of the plane relates to points near the Sun. Cr\\'ez\\'e et al.\\ (1998) and Holmberg \\& Flynn (2000) have used proper motions of Hipparcos stars to estimate the gradient of $F_z$ near the plane and hence infer that the local mass density is $0.076\\pm0.015\\msun\\pc^{-3}$ and $0.10\\pm0.01\\msun\\pc^{-3}$, respectively, while Kuijken \\& Gilmore (1991) have used the radial velocities of stars at the SGP to estimate $F_z$ at $z=-1.1\\kpc$ and thus infer that the surface density of the Galaxy within $1.1\\kpc$ of the plane is $71\\pm6\\msun\\pc^{-3}$. To constrain $\\b F$ at points that lie far from the Sun, we have to study objects that move to such points.  High-velocity stars are the obvious tracers to use, because they are so numerous. They strongly constrain models of $\\b F$ because stars on essentially the same orbit can be studied both locally and in situ (Binney, 1994; Dehnen \\& Binney, 1996). Unfortunately, the potential of this approach has yet to be systematically exploited. Recently, there has been considerable interest in using tidal streamers associated with disrupted satellites as tracers of $\\b F$ (Johnston et al.\\ 1999, Helmi et al.\\ 1999). My own view is that high-velocity stars have greater potential because (a) they are vastly more numerous and (b) they do not require approximations of the level of the (manifestly false) assumption that all elements of a streamer are on the same orbit.  Moreover, extracting useful information from streamers requires space-based astrometry, the requisite observations high-velocity stars are available now, and all that's lacking is machinery for modelling them.  Gravitational microlensing directly probes the Galaxy's mass distribution, but in a very different way from classical studies of gas and stars. In fact, microlensing does not measure the smooth Galactic force-field $\\b F$ but  graininess in the mass distribution. Hence, it is insensitive to the contribution to the latter from elementary particles and gas, and is therefore complementary to the traditional approach to the determination of $\\rho$ from $\\b F$, rather than competitive with it. ",
        "conclusions": ""
    },
    "0004/hep-th0004094_arXiv.txt": {
        "abstract": "{ We study inflation and entropy generation in a recently proposed pre-big-bang model universe produced in a collision of gravitational and dilaton waves. It is shown that enough inflation occurs provided the incoming waves are sufficiently weak. We also find that entropy  in this model is dynamically generated as the result of the nonlinear interaction of the incoming waves, before the universe enters the phase of dilaton driven inflation. In particular, we give the scaling of the entropy produced in the collision in terms of the focusing lengths of the incoming waves.} ",
        "introduction": "Pre-big-bang cosmology (PBB) \\cite{gasp} has evolved in recent years into one of the main trends in string cosmology. In spite of its many nice features there is a current discussion in the literature as to what extent the resolution of the flatness and homogeneity problems within the PBB picture requires a certain amount of fine tuning \\cite{ft}.  In a recent paper \\cite{fkvm} we have proposed a cosmological scenario in which the universe starts in a trivial state characterised by a bath of plane gravitational and dilatonic waves which, upon collision, generate PBB bubble universes. These bubbles act as possible seeds for a Friedman-Robertson-Walker (FRW) universe in the spirit of PBB cosmology. The proposed scenario \\cite{fkvm} should be seen as a realisation of the picture of Bounanno, Damour and Veneziano (BDV) \\cite{bdv}.  The aim of this Letter is to study some relevant phenomenological implications of the wave collision induced PBB cosmology. To this end we will be  focusing our attention on two issues: the conditions for successful inflation and the production of entropy in the scenario.  Though it was shown in \\cite{fkvm} that there exists  a dense region in parameter space,  for which the PBB inflation takes place, it is important to clarify further as to whether such favourable initial conditions would lead to sufficient inflation. In particular, we will see that inflationary requirements leading to consider large black holes as the seeds of PBB bubbles in the original BDV picture here imply the weakness of the incoming waves. Therefore, these waves can be regarded as emerging from perturbations of flat space-time. With respect to entropy production, unlike in the original BDV scenario where the origin of entropy remains rather  vague, we will see that entropy in our model is generated dynamically   as a result of the wave collision.  In the wave collision induced model universe, one starts in the remote past with non-interacting plane gravi-dilatonic waves. It was argued in \\cite{fkvm} that such a state should be identified as one with a  minimum entropy, for both gravitational and matter sectors. This assumption is based on the following reasons:  \\begin{itemize} \\item[i)]  Simplicity: Plane waves have no nonvanishing curvature invariants, and thus are, apart from Minkowski spacetime, the simplest geometries one may think of.  \\item[ii)]  Symmetry requirements: With the advent of the string theory and its alphabetic generalisations (M-theory and F-theory) our view on the initial state of the universe has changed considerably. In particular one is led to think that the primordial universe should be described by some exact string background which at later time evolves into a FRW-like universe. In this light, Penrose's proposal \\cite{penrose} of assigning zero gravitational entropy to backgrounds with vanishing Weyl curvature (a particular case of which is the FRW metric) seems to be too restrictive. This is especially evident in the context of PBB cosmology where the universe originates in a highly perturbative state and only evolves into a FRW phase after a graceful exit phase. Plane gravitational waves, being exact string backgrounds \\cite {horstief}, are ideal candidates to represent this PBB primordial state \\cite{gr}.  \\item[iii)]  Absence of particle creation: The fact that in the vicinity of plane waves there is no quantum creation of particles enforces the idea that plane waves are zero gravitational entropy states. Even using Penrose's original suggestion \\cite{penrose,no} where the ``number of gravitons'' contained in the gravitational field could be adopted as a measure of entropy, plane waves are likely to be associated with minimum gravitational entropy. \\end{itemize}   \\begin{figure}[tbp] \\centerline{\\epsfig{file=ent_fig1.eps, width=2.9in}} \\caption{Interaction region in wave collision induced PBB cosmology.} \\end{figure}   We now briefly summarise some features of the model universe described in \\cite{fkvm} (the reader is referred for details to the original paper). The universe starts in a distant past with weak, strictly {\\it plane} waves propagating on a flat background. Eventually these waves collide breaking the plane symmetry in the interaction region (Fig. 1). Just after the collision takes place at some $t_{i}<0$, physics is dominated by the superposition of two noninteracting null fluids so the geometry is described by a Doroshkevich-Zeldovich-Novikov (DZN) line element \\cite{dzn}. This phase is followed by an intermediate region where the nonlinearities of gravity become important and the evolution then becomes nonadiabatic. Finally, as we approach the caustic singularity the universe enters a Kasner phase (for $t_{K}<t<0$) where matter is effectively described in terms of a stiff fluid and the evolution is adiabatic. In this region the geometry is given by a generalised Kasner line element whose exponents depend on one spatial coordinate alone. In Einstein frame this metric can be written as \\begin{eqnarray} ds^{2}=-dt^{2}+\\left({-t\\over -t_{K}}\\right)^{2\\alpha _{1}(z)}dx^{2} +\\left({-t\\over -t_{K}}\\right)^{2\\alpha _{2}(z)}dy^{2}+\\left({-t\\over -t_{K}}\\right)^{2\\alpha _{3}(z)}dz^{2}, \\label{met} \\end{eqnarray} where we have normalized the scale factors to their values at the beginning of the Kasner epoch, $t=t_{K}$.  The generalised Kasner exponents $\\alpha_{i}(z)$ are determined by the initial data on the null boundaries of the interaction region. In terms of the gravitational and scalar source functions $\\epsilon (z)$ and $\\varphi (z)$ given, in turn, by quite complicated line integrals of incoming data \\cite{fkvm}, the Kasner exponents can be written as \\begin{eqnarray} \\alpha _{1}(z) &=&{\\frac{1+\\epsilon (z)}{a(z)+2}},  \\nonumber \\\\ \\alpha _{2}(z) &=&{\\frac{1-\\epsilon (z)}{a(z)+2}},  \\nonumber \\\\ \\alpha _{3}(z) &=&{\\frac{a(z)}{a(z)+2}},  \\label{ke} \\end{eqnarray} where $a(z)\\equiv {\\frac{1}{2}}[\\epsilon (z)^{2}+\\varphi (z)^{2}-1]$. They satisfy the following relations \\begin{eqnarray} \\sum_{i=1}^{3}\\alpha _{i}(z)=1,\\hspace{1cm}\\sum_{i=1}^{3}\\alpha _{i}(z)^{2}=1-{\\frac{2\\varphi (z)^{2}}{[a(z)+2]^{2}}}. \\label{kas} \\end{eqnarray} On the other hand the dilaton field near the singularity behaves like \\begin{eqnarray} \\phi (t,z)=\\beta (z)\\log {\\left( \\frac{-t}{-t_{K}}\\right) }+\\phi _{0} \\label{dil} \\end{eqnarray} with \\begin{eqnarray} \\beta (z)={\\frac{2\\varphi (z)}{a(z)+2}} \\label{beta} \\end{eqnarray} and $\\phi _{0}$ being some constant zero mode that determines the string coupling constant at the beginning of dilaton driven inflation (DDI), $g_{st}=e^{{1\\over 2}\\phi_{0}}$. Since we have normalized the metric and the dilaton at the beginning of DDI the natural scale in Einstein frame in this regime will be the Planck scale at $t_{K}$, $\\ell_{Pl}=e^{{1\\over 2}\\phi_{0}}\\ell_{st}$. We will assume in the following that the dilaton field does not vary very much from the initial boundaries to the onset of the Kasner regime, so $\\phi_{0}\\sim \\phi(t\\rightarrow -\\infty)$. In this case asymptotic past triviality implies that $\\phi_{0}\\ll -1$. ",
        "conclusions": "In this Letter we have investigated some phenomenological aspects of the picture for PBB cosmology proposed in Ref. \\cite{fkvm}. In particular, we have focused our attention on two special aspects of the model: the conditions for successful inflation and the entropy generation.  As to the requirement of successful inflation we have found that in order to meet the necessary number of $e$-foldings to solve the homogeneity and flatness problems the focal lengths of the incoming waves have to be exponentially large with respect to the string scale. Actually it is interesting that the bound on the product $L_{1}L_{2}$ is independent of the physical mechanism terminating inflation, $\\alpha'$ corrections or strong coupling effects. This provides quite a natural picture where the waves might be seen as small perturbations on a flat space-time producing upon collision the PBB bubbles. One of the alleged ``fine-tunings'' of the BDV picture was the necessity of starting with very large black holes in order to produce sufficient inflation during the PBB phase. In our picture, this requirement is translated into a rather ``natural'' condition of having extremely weak incoming waves.  With respect to the production of entropy we have seen, that entropy is generated dynamically as the result of the nonlinear interaction of the incoming waves. Thus, the following picture emerges. As argued in \\cite{fkvm} the universe starts as a bath of gravitational and dilatonic plane waves with zero total (matter+gravitational) entropy. The waves eventually interact and, before producing a singularity, the universe passes through a stage of nonadiabatic evolution where entropy is created (Fig. 1). The process of entropy generation  ends when the universe enters the regime of adiabatic (perfect fluid) evolution. This happens at the onset of the Kasner phase or earlier, depending whether velocities dominate over spatial gradients at or before entering the Kasner regime. We have also seen that the total entropy inside a Hubble volume at the beginning of DDI is related to the focal lengths of the incoming waves by $S=\\kappa L_{1}L_{2}/\\ell_{Pl}^{2}$, with $\\kappa$ a numerical constant depending on the initial conditions of the collision and the thermodynamics of the effective imperfect fluid describing the dilaton condensate in the nonadiabatic intermediate regime.  An immediate question to be asked is whether the analytical relation between the initial data and the structure of the singularity \\cite{fkvm} in this scenario allows a quantitative evaluation of the entropy produced. This, unfortunately, cannot be done since any analysis of this type would require some input about the thermodynamical properties of the effective fluid in the {\\it whole} interaction region. This is easy to understand on physical grounds, since the initial data only determines the dynamics of the background fields (the metric and the dilaton in this case), but not the thermodynamics of the dilaton condensate. In order to specify from first principles the phenomenological parameters required for the description of the fluid in the nonadiabatic regime (which would determine the amount of entropy generated and therefore the value of $\\sigma\\ell_{Pl}^2$) one would need to invoke a microscopical description of the matter content. This situation is analog to the case of standard FRW cosmology, where the gravitational dynamics of a radiation dominated universe does not determine the value of the Stefan-Boltzmann constant.  We have checked that the Bekenstein entropy bound is satisfied close to the singularity and can only be violated at earlier times. However, the conditions favourable for a long lasting inflation seem to comply with the Bekenstein entropy bound. On the other hand we have seen that the Hubble entropy bound will be satisfied depending on the value of the Kasner exponents (and as a consequence on the initial data on the null boundaries) as well as on the amount of entropy generated during the intermediate regime.  To close, the results of our analysis suggest a picture in which some collisions occur within a primordial ``cold'' bath of weak gravi-dilatonic waves. Before entering into the DDI stage the universe undergoes a nonadiabatic transient epoch where entropy is produced. If the waves are weak enough (as one would expect if they result from small fluctuations of the backgrond) the resulting PBB bubble can become the seed of a FRW universe of the kind we see at present. In order to study  the scenario analytically we have confined the initial data to strictly plane incoming waves. Nonetheless, it would be great to understand the problem in terms of more general $pp$ waves which have finite tranverse size. Note, however, that unlike plane waves, $pp$ waves do polarise the vacuum, and assigning initial zero entropy to those would be somehow less natural."
    },
    "0004/hep-ph0004019_arXiv.txt": {
        "abstract": "A gamma-ray burst fireball is likely to contain an admixture of neutrons, in addition to protons, in essentially all progenitor scenarios. Inelastic collisions between differentially streaming protons and neutrons in the fireball produce $\\num(\\barnum)$ of $\\sim$10 GeV as well as $\\nue(\\barnue)$ of $\\sim 5$ GeV, which could produce $\\sim 7$ events/year in km$^3$ detectors, if the neutron abundance is comparable to that of protons. Photons of $\\sim 10$ GeV from $\\pi^0$ decay and $\\sim 100$ MeV $\\barnue$ from neutron decay are also produced, but will be difficult to detect. Photons with energies $\\siml 1$ MeV from shocks following neutron decay produce a characteristic signal which may be distinguishable from the proton-related MeV photons. ",
        "introduction": "Gamma-ray burst (GRB) sources are distributed throughout the universe and their energy output is measured to be a substantial fraction of a solar rest mass equivalent \\cite{fm95}.  A variety of observations support the interpretation that these events are caused by cataclysmic stellar collapse or compact mergers, producing a fireball with bulk expansion Lorentz factor $\\Gamma \\sim 10^2-10^3$. In the standard GRB model a fireball made up of $\\gamma,\\epm$ and magnetic fields with an admixture of baryons is produced by the release of a large amount of energy $E\\simg 10^{53}$ ergs in a region $r_o \\sim 10^7 \\ro7$ cm (e.g. \\cite{pm99}). The observations indicate that typical fireballs are characterized by a luminosity $L\\sim 10^{52}\\L52$ erg s$^{-1}$ and durations $t_w = 10 t_{w1}$ s in the observer frame, with a large spread in both quantities. The outflow is controlled by the value of the dimensionless entropy  $\\eta=(L/{\\dot M} c^2)$ injected at $r_o$. Previous discussions of fireball models have generally focused on the charged particle components, since they determine directly the photon signal. However, consideration of a neutron component introduces qualitatively new effects \\cite{dkk99}.  In a $p,n$ fireball, for values of $\\eta\\simg 400$, the neutrons and protons acquire a relative drift velocity causing inelastic $n,p$ collisions and creating neutrinos.  We investigate here the neutrino and photon signals from $n,p$ collisions following decoupling in GRB. The $\\sim 10$ GeV neutrinos from this mechanism depend upon the presence of neutrons in the original explosion, but the neutrinos are created in simple physical processes occurring in the later stages of the fireball. On the other hand, the $10^5$ GeV neutrinos discussed in refs. \\cite{wb97} require the acceleration in shock waves of ultra-high energy protons interacting with photons. Thus the 10 GeV and the $10^5$ GeV neutrinos reflect very different astrophysical processes and uncertainties. Other processes, e.g. neutrinos from $p,p$ collisions\\cite{px94} also require shocks but have lower efficiencies, while 10-30 MeV neutrinos from the original explosion\\cite{pk99} are much harder to detect due to the lower cross sections.  We show (\\S \\ref{sec:nudet}) that the 10 GeV neutrinos could be detectable by future km$^3$ size detectors. The associated $\\sim 10$ GeV $\\gamma$-ray fluences are compatible with current detection rates, and may be detectable with future space missions.  The dependence of these signals on the neutron/proton ratio $\\xi$ provides a new tool to investigate the nature of the GRB progenitor systems. Moreover the predicted neutrino event rate depends on the asymptotic bulk Lorentz factor of the neutrons, which is linked to that of the protons.  The latter affects all the electromagnetic observables from the GRB fireball, including the photospheric and shock radii, as well as the particle acceleration and non-thermal photon production. ",
        "conclusions": "For characteristic parameters, GRB outflows produce 5-10 GeV $\\num\\barnum$ and $\\nue\\barnue$ from internal inelastic $p,n$ collisions that create pions. The $\\nu\\barnu$ energy output $E_{\\nu\\barnu}\\sim 5\\times 10^{51}\\E53 (2\\xi/[1+\\xi])(2/[1+z]) (\\etspi/\\eta)^{4/3}$ ergs depends on the total energy $E$ of the GRB and on the neutron fraction $\\xi$ as well as on the dimensionless entropy $\\eta$. For a km$^3$ detector, approximately 5 to 10 neutrino events above 10 GeV are predicted per year, for a neutron/proton ratio $\\xi=1$. These events will be coincident with GRB electromagnetic flashes in direction and in time (to an accuracy of $\\sim 10$ s), which can enable their separation from the atmospheric neutrino background. Underground water detectors of the type being planned by BAIKAL \\cite{bai97}, NESTOR \\cite{nes98}, ANTARES \\cite{ant98} and the Antartic detector ICECUBE \\cite{hal99} could potentially detect these relatively low energy neutrino events if a sufficiently high density of phototubes were used. About 80\\% of these neutrinos are $\\num$ and $\\barnum$ (in approximately equal numbers) and the remainder are $\\nue$ and $\\barnue$. These 5-10 Gev $\\nu\\barnu$ are followed by $\\sim$ 120 MeV $\\barnue$ from neutron decay, but the event rate from neutron-decay neutrinos is very low. The higher energy neutrinos are produced for neutron/proton ratios $\\xi >0$ when the dimensionless entropy $\\eta= L/{\\dot M}c^2$ exceeds $\\etspi \\simeq 4\\times 10^2 \\L52^{1/4}\\ro7^{-1/4}$ $(2/[1+\\xi])^{1/4}$, and are accompanied by $\\sim 10$ GeV photons which may be detectable in low redshift cases with GLAST \\cite{glast99}. For a typical GRB at redshift $z\\sim 1$ the number fluences in 10 GeV neutrinos are $N(\\barnue+\\nue)\\sim 0.5 N(\\barnum+\\num)$ $\\sim 10^{-4}$ cm$^{-2}$, and one order of magnitude less for GeV photons.  In all bursts where $\\xi >0$ the lower energy ($\\sim 120$ MeV) neutrinos are produced, and neutron decay occurs on an observer timescale $t_n\\sim 3 \\L52^{-1/4}\\ro7^{1/4} [(1+\\xi )/2]^{1/4} (\\eta/\\etspi)^{1/3}$ s. For outflows of duration $t_w$, these decays will be associated with MeV electromagnetic pulses of duration $\\min [t_n,t_w]$, which are additional to the MeV pulses expected from shocks in the original proton component. For short bursts with $t_w \\siml 3$ s, the proton electromagnetic pulse appears first and is separated from a subsequent neutron electromagnetic pulse, the latter having a slow rise-fast decay envelope and a softer spectrum, which may be detectable with the Swift satellite\\cite{swiftpage}. A systematic study of the time histories of GRB  emission would be useful to search for evidence of delayed pulses that might be caused by neutron decay.  The detection of 5-10 GeV $\\nu\\barnu$ in coincidence with GRB photon flashes will not be easy, but would provide unique astrophysical information.  Constraints on the neutron fraction could provide information about the progenitor stellar system giving rise to GRB. For instance, core collapse of massive stars would lead to an outflow from an $Fe$-rich core with $\\xi\\sim 2/3-1$, while neutron star mergers would imply $\\xi \\geq 1$. Photodissociation during collapse or merger, as well as $n,p$ decoupling and inelastic collisions, would both drive $\\xi$ toward unity, although this equalization process is likely to remain incomplete. For low $\\eta\\siml \\etspi$, inelastic collisions are not expected and the 5-10 GeV $\\nu\\barnu$ are absent, producing only the harder to detect $\\sim 100$ MeV $\\barnue$ from neutron decay. An initially non-baryonic outflow of, e.g. $e^\\pm$ and magnetic fields, would acquire a baryonic load by entrainment from the progenitor environment, with $\\xi \\ll 1$ from massive stellar envelopes, but $\\xi\\simg 1$ for, e.g., compact mergers. Thus, lower values of $\\xi$, leading to lower ratios of 5-10 GeV $\\nu\\barnu$ and a lower ratio of neutron decay MeV photons to total fluences would be expected from massive progenitors than from compact mergers.  {\\it Acknowledgements}\\\\ Partial support was received by JNB from NSF PHY95-13835 and by PM from NASA NAG5-2857, the Guggenheim Foundation and the Institute for Advanced Study. We acknowledge valuable conversations with G. Fishman, V. Fitch, Vl. Kocharovsky, P. Kumar, R. Nemiroff, J. Norris, M.J. Rees, P. Vogel and E. Waxman."
    },
    "0004/astro-ph0004376_arXiv.txt": {
        "abstract": "We derive and solve exactly the Dyer-Roeder equation in a Friedman Robertson Walker cosmological model with non zero cosmological constant. To take into account non homogeneous distribution of matter we use the phenomenological clumpiness parameter $\\tilde{\\alpha}$. We propose also a general form of an approximate solution which is simple enough to be useful in practical applications and sufficiently accurate in the interesting range of redshifts. ",
        "introduction": "Recent determination of distances to high redshift type Ia supernovae revealed that the universe is now expanding faster  than expected. This accelerated expansion is linked with a non zero cosmological constant. Since gravitational lensing of cosmological objects involves also high redshift galaxies it is interesting to extend the cosmological lens equation to cosmological models with non zero cosmological constant.\\\\Furthermore the mass distribution in the universe at the relevant range of redshifts is clearly non homogeneous. In our considerations we will take into account the fact that at the scale up to several hundred Mpc matter distribution is non homogeneous. Moreover, with the advent of the Next Generation Space Telescope (NGT), which will allow imaging of objects at $z\\geq10$ (\\cite{Bark}),  it will be possible to check better and better the influence of  non homogeneities .\\\\Using the general equations describing propagation of light in a given spacetime we derive an equation for the angular diameter distance and we extend the notion of angular diameter distance to two objects at arbitrary redshifts. The general equation for the angular diameter distance, the so called Dyer-Roeder equation, can be exactly solved also in the case when the cosmological constant is different from zero. However, the exact solution is very complicated  and therefore not useful for practical applications. We propose a simple analytic form of an approximate solution. The approximate form of the angular diameter distance depends on four arbitrary parameters. We  fix the values of these parameters by fitting this form to the exact solution.\\\\The paper is organized as follows: In Section $2$ we present the derivation of the angular diameter distance in an arbitrary Friedman- Robertson- Walker cosmological model. To take into account the non homogeneous distribution of matter, following the standard practice, we introduce a constant phenomenological parameter $\\tilde{\\alpha}$ and rewrite the final equation in a few different forms. In Section $3$ we present exact analytic solutions for the angular diameter distance, first in a flat universe and then in the general case of arbitrary spatial curvature. Section $4$ is devoted to the approximate solution and to  discussions of the fitting procedure used to determine the values of arbitrary parameters. Discussion of our results is presented in Section $5$. ",
        "conclusions": "In this paper we have considered the cosmological lens equation in the case when the cosmological constant $\\Lambda\\not=0$ and the curvature of space is different for zero ($k\\not=0$). We have included the effects of non homogeneous distribution of matter which are described by a phenomenological parameter $\\tilde{\\alpha}$. Unfortunately at the moment there are no generally accepted  models that  describe the distribution of baryonic and dark matter at high redshifts and therefore the influence of non homogeneities of matter distribution can be included only at this approximate level.  Following the standard procedure (\\cite{SEF}) we use the Dyer-Roeder distance to find the distance between two objects with redshifts $z_{1}$ and $z_{2}$.  To find the general solution we slightly enlarged the parameter space by considering the non zero cosmological constant, so  the cosmological model is described by the following parameters $\\Omega_{\\Lambda}\\not=0$, $\\Omega_{m}\\not=0$, $\\tilde{\\alpha}\\in [0.3,1]$ (not completely clumpy Universe), $p=0$ (no radiation). The solution that we have found is a functional combination of the solution given in SEF for the case when $\\Omega_{\\Lambda}=0$, $\\Omega_{k}=0$ and two hypergeometric functions. Fig. 1 shows that in the flat universe the Dyer-Roeder distance increases for increasing values of the cosmological constant. In the next step we allow the curvature of space to be different from zero. This further complicates the Dyer-Roeder equation which becomes of a Fuchsian type with 4 regular singular points and one regular singular point at infinity. The general form of this kind of ordinary differential equations is given in terms of the P-Riemann symbol.  The obtained exact solution is so complicated that it is practically useless in practical applications. Therefore we have looked for an approximate analytic solution simple enough to be used in many applications and at the same time sufficiently accurate at least in the interesting range of redshifts. The proposed form of the approximate solution of the Dyer-Roeder equation depends on four arbitrary parameters. To fix values of these parameters we fit the approximate solution to the exact one with the help of a non linear regression method (see {\\it Mathematica~~ 4.0}). Of course, the parameters depend on the density parameters, but their values for different cosmological parameters can be tabulated. Following SEF we have also found the function $\\chi$ which appears in the expression for time delays as well as in the lens equation. We have also proposed approximate analytical form of the function $\\chi$ which contains three arbitrary parameters. To find values of these parameters we use the same method as above.  We have also studied how the redshift corresponding  to the maximal angular diameter distance depends on the basic parameters determining the cosmological model.   Let us now consider the following three combinations of the DR distance: \\begin{eqnarray} \\circ&&\\frac{D_{LS}}{D_{OS}} =\\frac{H_0}{c}\\frac{r_{LS}}{r_{OS}},\\nonumber\\\\ \\circ&&\\frac{D_{OL}D_{LS}}{D_{OS}}=\\frac{H_0}{c}\\frac{r_{OL}r_{LS}}{r_{OS}},\\nonumber\\\\ \\circ&&\\frac{D_{OL}D_{OS}}{D_{LS}}=\\frac{H_0}{c}\\frac{r_{OL}r_{OS}}{r_{LS}}.\\nonumber \\end{eqnarray} We have selected these combinations because of the role they play, respectively, in bending of light, lensing statistics, and time delay. Asada (\\cite{AS97}, \\cite{AS98a}, \\cite{AS98b}) have found that: \\begin{eqnarray} \\circ&&\\frac{D_{LS}}{D_{OS}} (\\alpha_1)<\\frac{D_{LS}}{D_{OS}}(\\alpha_2),\\,\\,\\, {\\rm{for}} \\,\\,\\,\\alpha_1< \\alpha_2, \\nonumber\\\\ \\circ&&\\frac{D_{OL}D_{LS}}{D_{OS}} (\\alpha_1)< \\frac{D_{OL}D_{LS}}{D_{OS}}(\\alpha_2),\\,\\,\\,\\, {\\rm{for}}\\,\\,\\,\\, \\alpha_1<\\alpha_2,\\nonumber\\\\ \\circ&&\\frac{D_{OL}D_{OS}}{D_{LS}} (\\alpha_1)> \\frac{D_{OL}D_{OS}}{D_{LS}}(\\alpha_2),\\,\\,\\,\\, {\\rm{for}}\\,\\,\\,\\, \\alpha_1 <\\alpha_2. \\nonumber \\end{eqnarray} Using our approximate relations for $r$ and $\\chi$ we obtain that: \\begin{eqnarray} \\circ \\,\\,\\,\\,\\,\\ \\frac{D_{LS}}{D_{OS}} &=& (1+z_L) r(z_L) (\\chi(z_L)-\\chi(z_S))=\\nonumber\\\\ &=&(1+z_L) \\frac{1}{\\sqrt{d_1 + d_2(d_3 + d_4/z_L + z_L)^2)}}(\\frac{1}{(e_1 + e_2  z_L + e_3 z_L^2)}-\\nonumber\\\\&&\\frac{1}{(e_2 + e_2 z_S + e_3z_S^2 )}),\\nonumber \\\\&&\\nonumber\\\\ \\circ\\,\\,\\,\\,\\,\\frac{D_{OL}D_{LS}}{D_{OS}}&=& (1+z_L) r^2(z_L) (\\chi(z_L)-\\chi(z_S)) = \\nonumber\\\\&=&(1+z_L) \\frac{1}{(d_1 + d_2(d_3 + d_4/z_L + z_L)^2)}(\\frac{1}{(e_1 + e_2 z_L + e_3 z_L^2)}-\\nonumber\\\\&&\\frac{1}{(e_1 + e_2 z_S + e_3 z_S^2)}),\\nonumber \\\\&&\\nonumber\\\\ \\circ\\,\\,\\,\\, \\frac{D_{OL}D_{OS}}{D_{LS}} &=&\\frac{1}{(1+z_L)(\\chi(z_L)-\\chi(z_S))} =\\nonumber\\\\&=&\\frac{1}{(1+z_L)}\\frac{1}{(\\frac{1}{(e_1 + e_2  z_L + e_3 z_L^2)}- \\frac{1}{(e_1 +e_2  z_S + e_3 z^2)})}\\nonumber, \\end{eqnarray} Using the fitted values for the parameters $(d_{1}, d_{2}, d_{3}, d_{4})$ and $(e_{1}, e_{2}, e_{3})$ we confirm results of Asada for $\\tilde{\\alpha}$ in the considered range $1\\geq {\\tilde{\\alpha}} \\geq 0.3$. \\\\ Finally we like to stress that from our analysis it follows that variations in the angular diameter distance caused by the presence of cosmological constant are quite similar to variations due to changes in the value of $\\tilde{\\alpha}$.   \\begin{figure}[ht] \\centering \\epsfig{figure=f9.eps,height=6cm,width=0.8\\textwidth,clip=} \\caption{r(z) for $\\Omega_{\\Lambda}=0$ and $\\tilde{\\alpha}=0.6$ (upper curve) and $\\Omega_{\\Lambda}=0.7$ and $\\tilde{\\alpha}=1$ (bottom curve). This figure shows that non homogeneous distribution of matter can mimic the effect of non zero cosmological constant.} \\end{figure}  Actually, in Fig.9,  we consider two plots corresponding to  the values $\\Omega_\\Lambda=0$ and $\\tilde{\\alpha}\\ne 0$ and $\\Omega_\\Lambda\\ne 0$ and $\\tilde{\\alpha}=1$ and we see that the $\\tilde{\\alpha}$ parameter could mimic the effect of a non zero cosmological constant. This is an important observation in view of the recent observational results concerning the non zero value of the cosmological constant (\\cite{Koch96}, \\cite{Koch98}, \\cite{Koch96b}, \\cite{Perl97}). In the future we would like to extend our work and include also the radiation density parameter.    ~\\\\{\\it Acknowledgments}.  It is a pleasure to thank V.F. Cardone , G. Covone. C. Rubano and P. Scudellaro for discussions we had on the manuscript. RdR and AAM are financially sustained by the M.U.R.S.T. grant PRIN97 ``SIN.TE.SI`` , MD by the grant 2-P03D-014-17 of the Polish State Committee for Scientific Research, and EP by the Social European Committee."
    },
    "0004/astro-ph0004189_arXiv.txt": {
        "abstract": "I show how the existing observational data on Local Group dwarf galaxies can be used to estimate the average star formation law during the first $3\\dim{Gyr}$ of the history of the universe. I find that the observational data are consistent with the orthodox Schmidt law with a star formation efficiency of about 4 percent if the star formation is continuous (during the first $3\\dim{Gyr}$). The efficiency is proportionally higher if most of the gas in the dwarfs was consumed (and never replenished) in a short time interval well before the universe turned $3\\dim{Gyr}$. ",
        "introduction": "Direct observational measurements of the star formation rate in the very first objects in the universe is exceptionally difficult, and even with the recent observational and technological advances, is well beyond reach. Nevertheless, it is possible to reconstruct the general features of the star formation at high redshift in much the same way as paleontologists reconstruct the prehistoric life from fossils.  In the currently widely accepted hierarchical clustering paradigm the characteristic mass scale of cosmological objects increases with time: low mass objects form first, and more massive objects form later. Thus, a considerable fraction of the dwarf galaxies may be very old, and may represent the fossils of the very first star formation in the universe. Although these dwarf galaxies are small and faint, and thus hard to observe, a vast amount of detailed information exists on the Local Group dwarf galaxies, as summarized in several recent reviews (Grebel 1998; Mateo 1998; van den Bergh 1999). The first two reviews are particularly useful for my purpose, as they contain the star formation histories of a large fraction of all Local Group dwarfs. A careful investigation of the available data reveals a striking feature common to the vast majority (more than 90\\%) of all dwarf galaxies with known star formation histories: a sharp decline in the star formation rate about $10\\dim{Gyr}$ ago.  There are at least five possible explanations for such a feature in the star formation history. First, it may be attributed to the tidal effects from the Galaxy or Andromeda. However, in the hierarchically clustering universe neither the Galaxy nor Andromeda existed at that time as entities, but rather as an interacting collection of proto-galactic clumps (see, for example, Contardo, Steinmetz, \\& Fritze-von Alvensleben 1998), and thus they were not able to have a large impact on the star formation histories of {\\it almost all\\/} dwarf galaxies.  A second possible explanation is the effect of supernovae, which could have expelled the gas from the low potential wells and quenched the star formation rate. However, it is hard to believe that the supernova activity was so well synchronized in 90\\% of all dwarf galaxies as to expel the gas almost simultaneously (this is elaborated further in the Conclusions).  Since the drops in the star formation rates in dwarf galaxies are essentially simultaneous, it is plausible that this was triggered by a process that affected the whole universe. One such process is reionization. It is well established that photoheating of the intergalactic medium during cosmological reionization sharply reduces the gas fractions in the low mass objects (Thoul \\& Weinberg 1996; Quinn, Katz, \\& Efstathiou 1996; Weinberg, Hernquist, \\& Katz 1997; Navarro \\& Steinmetz 1997; Gnedin 2000). This, in turn, will lead to the drop in the star formation rates in dwarf galaxies, including the Local Group dwarfs. This proposal however encounters a difficulty because an age of $10\\dim{Gyr}$ corresponds to about $z=2$ (in a $13\\dim{Gyr}$ old universe with $h\\Omega_0^{1/2}=0.4$), which is too late for cosmological reionization. However, if the drop in the star formation rate in dwarf galaxies occurred some twelve billion years ago, the reionization explanation could work. This requires that the measured stellar ages are off by about $2\\dim{Gyr}$, but the question of whether this is plausible or not is outside the scope of this paper.  The fourth explanation is based on the observational evidence (which is admittedly not completely compelling) that the universe experienced a second epoch of reheating at about $z=3$ (Ricotti, Gnedin, \\& Shull 2000; Schaye et al.\\ 2000). This second reheating (plausibly attributed to the reionization of helium) has a similar effect on dwarf galaxies as hydrogen reionization at $z\\sim7-10$, except that it takes place somewhere between $z=2$ and $z=3$.  It is also possible that photoionization associated with the burst of star formation that made the bulges of the Galaxy and Andromeda (``local reionization'') inhibited the star formation in the surrounding dwarfs (van den Bergh 1994). While possible in principle, it is not clear whether this explanation is plausible within the hierarchical clustering paradigm, or whether it agrees with the observed metallicity distribution of the Galactic bulge (McWilliam \\& Rich 1994).  Whatever the explanation of the observed simultaneous drop in the star formation rate in 90\\% of Local Group dwarfs is, it is not essential to the measurement described in this paper. The only thing that matters is that there exists a subset of the Local Group dwarfs that had little star formation in the last $10\\dim{Gyr}$, so that their stellar content provides a fossil record of the early universe. ",
        "conclusions": "I use a sample of Local Group dwarf galaxies selected by the criterion that they form most of their stars more than about $10\\dim{Gyr}$ ago to measure the star formation law in the first $3\\dim{Gyr}$ of the life of the universe (assuming the age of the universe is $13\\dim{Gyr}$). I find that the star formation rate on scales comparable to the sizes of dwarf galaxies ($100-600\\dim{pc}$) is a power-law function of the gas density, with the power-law index $1.54\\pm0.11$. If the index is fixed at $n=3/2$, the star formation rate is given by equations(\\ref{rhoeos}) with an efficiency of about 4\\% if the gas is consumed continuously. The efficiency cannot be smaller than 2.2\\% but can be higher if the gas was consumed well before the universe turned $3\\dim{Gyr}$.  The efficiency is not a strong function of scale for scales in the range from about 100 to 600 parsecs (if it were, the straight power-law would not be a good fit in Fig.\\ \\ref{figLI}).  Admittedly, the time of the drop in the star formation rate in the Local Group dwarfs is not very well determined, and is, perhaps, uncertain to about one or two billion years.  Thus, the ``simultaneous'' drop in star formation rates in the majority of the dwarfs is perhaps not sufficiently simultaneous to preclude the possibility that it was supernova driven winds rather than reheating of the universe (or local reionization), which was responsible for the observed drop. However, in this case it would be difficult to explain the observed correlation, since the central regions of the dwarfs galaxies can be expected to be affected most by the supernova activity, thus spoiling the beautiful correlation shown in Fig.\\ \\ref{figLI}. Therefore, the present result indicates that the supernova driven winds play only a modest role in the early evolution of low-mass galaxies.  Curiously, equation \\ref{lsicbest}) can be used, in principle, as a novel distance indicator, since only the $x$-axis depends on the distance. However, it is likely to not be competitive with other indicators given the amount of information needed to measure $\\bar{L}_V$ and $\\bar{I}_c$.  The result presented here can be incorrect if the structural form of the Local Group dwarfs changes monotonically with their luminosity as to mimic the observed relationship (\\ref{bestfit}). I would like to emphasize that, while this possibility cannot be completely excluded, there is currently no observational evidence that supports it, and this assumption also contradicts the modern theoretical view on the structure of galaxies."
    },
    "0004/astro-ph0004140_arXiv.txt": {
        "abstract": "Inhomogeneities in wave propagation conditions near and below the solar surface have been detected by means of time-distance helioseismology.  Here we calculate the effect of temperature inhomogeneities on the travel times of sound waves. A temperature increase, e.g. in active regions, not only increases the sound speed but also lengthens the path along which the wave travels because the expansion of the heated layers shifts the upper turning of the waves upward. Using a ray tracing approximation we find that in many cases the net effect of a temperature enhancement is an {\\it increase} of the travel times.  We argue that the reduced travel times that are observed are caused by a combination of magnetic fields in the active region and {\\it reduced} subsurface temperatures. Such a reduction may be related to the increased radiative energy loss from small magnetic flux tubes. ",
        "introduction": "In recent years it has become possible to measure the travel times of acoustic waves travelling through the outer layers of the Sun through `time-distance helioseismology'. These travel times are used to infer information about the sub-surface structure of the Sun and have revealed inhomogeneities in the wave propagation conditions. (Duvall et al. 1993, Duvall 1995, D'Silva \\& Duvall 1995, Duvall 1997, Duvall et al. 1998). These inhomogeneities indicate shorter sound-wave travel times and have been associated with active regions.\\\\  Slight shifts in global p-mode frequencies in the course of a solar cycle have also been detected (Woodard \\& Libbrecht 1993, Dziembowski et al. 2000). As in the case of the time-distance measurements, they have been found to be closely related to magnetic activity (Dziembowski et al. 2000), and indicate shorter travel times in the outermost layers of the convection zone. A common interpretation for both effects is that the sound speed in active regions is somewhat higher because of an increase in temperature. However, it was shown by Goldreich et al. (1991) that, though intuitively appealing, such a temperature increase is unlikely to be the cause of the mode frequency shifts. Contrary to expectation, a higher temperature causes {\\em longer} travel times, because a temperature increase causes a slight expansion of the envelope. The path length increase caused by this expansion dominates over the increase in sound speed since the expansion is proportional to the temperature $T$, while the sound speed increases only as $T^{1/2}$. Instead of a temperature increase, Goldreich et al. propose that most of the effect is due to the photospheric magnetic field. A magnetic field increases the stiffness of the gas as experienced by pressure waves.\\\\  The expansion argument would also affect the interpretation of travel-time anomalies discovered by time-distance helioseismology. The purpose of this paper is to verify to what extent the expansion of the envelope associated with a temperature rise increases the wave travel times. In Sec. 2 we show that in most cases a temperature enhancement does in fact lead to {\\em longer} travel times. To explain shorter travel times one might then invoke magnetic fields. Alternatively, the subsurface temperatures could be slightly {\\em lower} in the active regions, the opposite of what is concluded in most analyses of time-distance measurements.  These possibilities are discussed briefly in Sec. 3. ",
        "conclusions": "Time-distance helioseismology indicates the presence of subsurface inhomogeneities. In order to map these inhomogeneities quantitatively, a propagation model is needed. In most analyses it is assumed that shorter travel times correspond to higher propagation speeds. While this is correct if the inhomogeneities are due to a vertical magnetic field, the most obvious possibility, i.e. a change in temperature, requires careful treatment. This is because a temperature change has two side effects in addition to a change in propagation speed. Since hot gas is less dense, the resulting positive buoyancy causes the envelope to expand vertically in regions of higher temperature. Secondly, the resulting horizontal imbalance also sets up a circulation flow. Both have effects on acoustic wave propagation.\\\\  We have analysed here the effects of a temperature increase in a model in which horizontal imbalance is compensated geostrophically by a circulation (i.e. by Coriolis forces acting on a horizontal flow, as in high- and low-pressure systems in the Earth's atmosphere). We find that if temperatures are increased, but vertical expansion is ignored, the changes in travel time as seen in time-distance measurements can be of either sign, depending on the wavenumber and the inter-skip distance. This is because the increased propagation speed is offset in part by the fact that the upper turning point of the waves is higher in a hotter model.\\\\  This effect is enhanced if vertical expansion of the perturbed model due to vertical pressure balance is taken into account. We have calculated this in a model in which the entropy has been increased locally (corresponding to an increase in the polytropic constant $K$). The results show that the travel times are {\\it increased} by a temperature enhancement, as long as the horizontal extent of the inhomogeneity is not too large. We found it difficult to obtain time delays for inhomogeneities that have widths of more than about 10 Mm without making unrealistic assumptions.  If one takes into account the effect of turbulence onto the sound speed inside the inhomogeneity, one finds that the sign of the time delay can change. This implies that predictions become sensitive to parametrizations of poorly known turbulent flows and therefore less robust (see Sec. \\ref{circs}).\\\\   The effects of temperature enhancements have been considered before by Kuhn and Stein (1996), whose conclusions differ from ours. By a numerical convection simulation, these authors calculate the effect of a temperature increase applied at a depth of 2.6 Mm on a part of the lower boundary. This increase causes both a circulation and a vertical expansion of the model above the hotter boundary region. The temperature change as a function of depth in their figure 2 mimics the {\\it vertical gradient} of the unperturbed model. The vertical shift implied by the figure is about 10 km at the photosphere. This can be compared with a value of roughly 12 km expected from approximate vertical hydrostatic balance for an entropy change as applied by the authors. The authors then perform a ray tracing calculation similar to ours, and find {\\it reduced} travel times.\\\\  The cause of the disagreement can be traced to the treatment by Kuhn and Stein of the layers below 2.6 Mm, i.e. at depths not covered by the numerical simulation. In their ray tracing calculations, they assume higher temperatures, where $\\delta T/T$ declines linearly from 0.006 at 2.6 Mm to 0.003 at a depth of 50 Mm. The vertical expansion of the model, however, is based only on the 2.6 Mm layer included in the simulation. The expansion is approximately proportional to the depth over which the increased temperatures extend (see Eqn. \\ref{pol}). On the one hand, the convection simulation by Kuhn and Stein demonstrates the vertical expansion effect, but on the other hand their ray tracing calculation underestimates its effect on travel times by a large factor.\\\\  Kuhn and Stein do not specify the cause of their assumed temperature enhancement at a depth of 50 Mm, but suggest a source in those layers where the magnetic field of the solar cycle is produced, near the base of the convection zone ($z\\sim 200$ Mm) (see also Kuhn, Libbrecht and Dicke 1988). If this were the case, for a given temperature increase at the surface, the vertical expansion effect would be four times larger than for an assumed depth of 50 Mm. The travel time increase by the vertical expansion would then certainly dominate over the reduction due to the higher sound speed for all time-distance measurements published so far (which reach depths of the order of 20 Mm).  \\subsection{Lower temperatures in active regions?}  The shorter sound travel times in active regions found by time-distance seismology are consistent with the increase of mode frequencies with magnetic activity found by Woodard et al. (1993) and Dziembowski et al. (2000). As noted by Goldreich et al. (1991), the increased mode frequencies are not consistent with increased temperatures, for the same reason as in our time-distance calculations. Instead, Goldreich et al. suggest that the magnetic field of active regions causes the increase in propagation speeds. For the vertical magnetic fields seen near the surface, there would be no associated vertical expansion of the envelope. While this explanation is consistent with the data available then, it is no longer compatible with recent data. This is because a magnetic change in propagation speed is confined to a thin layer near the surface (unless very large magnetic fields are assumed at a depth of 10--20 Mm). The mode frequency changes measured by Woodard et al. and  very accurately by Dziembowski et al. (2000) (with MDI) show that while most of the effect is concentrated near the surface, there are also significant changes at depths of 10 Mm. This would require field strengths of the order 20 kG covering large fractions of the surface.\\\\  Instead of getting involved in a discussion about the difficulties that such large unobserved magnetic fields would cause, we suggest here a more radical and simple solution. One could infer that the subsurface temperatures in active regions, in spite of the increased emission at the surface, are in fact {\\it reduced}.\\\\  If the effect of the small scale magnetic fields were an increase in convective efficiency, or some other effect that increases the radiation losses at the surface, then the increased cooling would cause the intergranular downdrafts to be cooler than average (since by this assumption the granules would have lost more heat). These lower temperatures would be carried down with the downdrafts, and cause horizontal average temperatures to be reduced below active regions. The depth dependence of the effect would depend on the details of the rate of spreading of the downflows by entrainment.\\\\  An effect that would cause just such an increased cooling assciated with active region fields was proposed by Spruit (1977). There it was shown that the `dimples' in the photosphere caused by the reduced opacity in the magnetic elements allow more radiation to escape. The magnetic elements are effectively small leaks through which more heat escapes than from the normal photosphere.  This increased radiation at the same time implies a larger average cooling rate in active regions. The downflows inferred from time-distance helioseismology (Duvall et al. 1998) are consistent with this interpretation, but are hard to understand in models with increased sub-surface temperatures."
    },
    "0004/astro-ph0004230_arXiv.txt": {
        "abstract": "{ The Basel Stellar Library (BaSeL models) is constituted of the merging of various synthetic stellar spectra libraries, with the purpose of giving the most comprehensive coverage of stellar parameters. It has been corrected for systematic deviations detected in respect to UBVRIJHKLM photometry at solar metallicity, and can then be considered as the state-of-the-art knowledge of the broad band content of stellar spectra. In this paper, we consider a sample of 9 F-type stars with detailed spectroscopic analysis to investigate the Basel Stellar Library in two photometric systems simultaneously, Johnson (B$-$V, U$-$B) and Str\\\"omgren (b$-$y, m$_1$, and c$_1$). The sample corresponds to potential targets of the central seismology programme of the COROT space experiment, which have been recently observed at OHP. The atmospheric parameters T$_{\\rm eff}$, [Fe/H], and log g obtained from the BaSeL models are compared with spectroscopic determinations as well as with results of other photometric calibrations. For a careful interpretation of the BaSeL solutions, we computed confidence regions around the best $\\chi$$^2$-estimates and projected them on T$_{\\rm eff}$-[Fe/H], T$_{\\rm eff}$-log g, and log g-[Fe/H] diagrams. We first derive the 3 atmospheric parameters from the full photometric information available (Johnson and Str\\\"omgren data). The BaSeL library has only been calibrated for broad band UBVRIJHKLM photometry, and it presents therefore intrinsic limitations in respect to other photometric systems, especially with different bandwidth. Thus, using this combination (Johnson and Str\\\"omgren), BaSeL temperatures are systematically lower ($\\sim$ 130K), and the discrepancy for the gravity and the metallicity can be quite large in comparison to the other methods. To disentangle these unexpected discrepancies, we fixed the temperatures at their spectroscopic values and studied the relative influence of each colour index in the log g-[Fe/H] diagrams. We find inconsistent results between Str\\\"omgren and Johnson synthetic colours. While the Johnson colours give impressively good log g-[Fe/H] solutions, the combination of the m$_1$ and c$_1$ Str\\\"omgren synthetic indices does not provide reliable results. This is not, of course, due to any intrinsic superiority of the UBV over the  Str\\\"omgren system, but to the properties of the BaSeL models.\\\\ Finally, in order to simultaneously and accurately determine the stellar parameters T$_{\\rm eff}$, [Fe/H] and log g, we suggest to use the combination of the synthetic BaSeL indices B$-$V, U$-$B and b$-$y rather than B$-$V, U$-$B, b$-$y, m$_1$ and c$_1$. ",
        "introduction": "The Basel Stellar Library (BaSeL) is a library of theoretical spectra corrected to provide synthetic colours consistent with empirical colour-temperature calibrations at all wavelengths from the near-UV to the far-IR (see Cuisinier {\\al} 1996 for the correction procedure, and Lejeune {\\al} 1998 and references therein for a complete description). These model spectra cover a large range of fundamental parameters (2000 $\\leq$ T$_{\\rm eff}$ $\\leq$ 50,000 K, $-$5 $\\leq$ [Fe/H] $\\leq$ 1 and $-$1.02 $\\leq$ log g $\\leq$ 5.5) and their photometric calibrations are regularly updated (for instance, the calibrations in the very low metallicity range is being improved  by comparisons with globular cluster populations, Westera et al. 1999) and extended to an even larger set of parameters (see Lejeune {\\al} 1997, Lejeune {\\al} 1998, and Lejeune {\\al} 2000). \\\\ The BaSeL library spectra have been calibrated directly for standard dwarf and giant sequences at solar abundances and using UBVRIJHKLM broad-band photometry, and are hence expected to provide excellent results in these photometric bands. Since they are based on synthetic spectra, they can in principle be used in many other photometric systems taken either individually or simultaneously, and this is another major advantage of these models. For these reasons, they are currently used in an increasing number of astrophysical studies (e.g. Bruzual {\\al} 1997, Weiss \\& Salaris 1999, Kurth {\\al} 1999). \\\\ Nevertheless, little attention has been paid until now to validate the possibilities offered by the BaSeL models in several photometric bands simultaneously. The fact that the library spectra have been calibrated using broad-band observations (UBVRIJHKLM) does not necessarily imply that similarly reliable results can be expected from their application to intermediate- and/or narrow-band photometry data. On the contrary, the extent to which the {\\it coarse} UBVRIJHKLM calibration provides useful results even at higher resolution requires separate investigation. In this paper, we investigate the BaSeL models using Johnson and Str\\\"omgren photometric systems simultaneously. BaSeL Str\\\"omgren photometry has been previously tested with success against stars belonging to double-lined eclipsing binaries for which gravities are accurately known (Lastennet {\\al} 1999a). Here we use a sample of 9 F-type stars for which detailed spectroscopic analysis has been performed. These stars are potential candidates of the COROT (COnvection and ROTation) space experiment main programme (see for instance Catala {\\al} 1995, Mi\\-chel {\\al} 1998, and Baglin {\\al} 1998 for details), and has been observed at the 193cm telescope at Observatoire de Haute-Provence (OHP, France) as part of the target selection process. High signal to noise ratio (S/N $\\simeq$ 150) spectra have been obtained using {\\it Elodie} echelle spectrograph and analysed by comparison with theoretical spectra (Ligni\\`eres {\\al} 1999). This set of new spectroscopic data together with both Johnson and Str\\\"omgren photometry from the literature provides a unique test of the predicted results of the BaSeL library. In addition, other calibration methods, namely the Templogg programme (based on Napiwotzki {\\al} 1993, and K\\\"unzli {\\al} 1997) and the Marsakov \\& Shevelev (1995) catalogue, provide further comparisons. \\\\ The paper is organized as follows: Sect. 2 deals with the description of our working sample of stars and the extinction issue, Sect. 3 presents the various methods used to derive estimates of T$_{\\rm eff}$, [Fe/H], and log g, and Sect. 4 is devoted to the presentation and the discussion of the results, with particular emphasis put on the intrinsic properties of the BaSeL library spectra in Sect. 5. Finally, Sect. 6 draws our general conclusions. ",
        "conclusions": "Several methods of determination of the fundamental stellar parameters T$_{\\rm eff}$, log g and [Fe/H] are compared for nine single F stars. Particular attention has been paid to the simultaneous predictions of the BaSeL models in two photometric systems, Johnson and Str\\\"omgren. We show that using all photometric data is not the best strategy to obtain reliable simultaneous determinations of T$_{\\rm eff}$, log g and [Fe/H] with the BaSeL models in the temperature range relevant for this paper, because of intrinsic limitations of the BaSeL library. As a matter of fact, if one uses all five available photometric indices (B$-$V, U$-$B, b$-$y, m$_1$ and c$_1$), only the agreement with spectroscopic determination for the effective temperature is good, although the BaSeL-derived temperatures are slightly but systematically lower. The discrepancies in gravity and metallicity are however rather high, with the BaSeL predictions being too low. Alternatively, we show that the best results are obtained by using B$-$V, U$-$B, b$-$y in combination. This BaSeL combination is the best because on the one hand the b$-$y synthetic index gives reliable and accurate estimates of the effective temperature and, on the other hand, B$-$V and U$-$B give good estimates of [Fe/H] and the surface gravity. We wish to emphasize that, if Str\\\"omgren photometry is shown here to be less performing than UBV as a device for determining fundamental stellar parameters (and, most particularly, abundances and surface gravities), this is due to the fact that the BaSeL library spectra have been calibrated by colour-calibration in UBVRIJHKLM; it is not, of course, due to an intrinsic superiority of the UBV system over the Str\\\"omgren system. Thus, because the BaSeL library has only been calibrated in the UBVRIJHKLM colours, the best parameters derivations come from this system. Whilst parameters derivations in equivalent bandwidth systems such as Washington seem to be as good with the BaSeL libraries than with empirical methods (Lejeune 1997), parameter derivations from the BaSeL library in narrow band systems such as Str\\\"omgren photometry seem to be of poorer quality. \\\\ This work thus strongly suggests that the BaSeL library should be calibrated in various photometrical systems, in various bandwidth (maybe using higher resolution synthetic stellar spectra libraries) in order to represent in the best way the variety of the knowledge of stellar spectral energy distributions that exists today in various photometric systems. \\\\ We found that the best combination to determine stellar parameters within the BaSeL library, even with its limitation, was to use a combination of two Johnson colours, U$-$B and B$-$V, and a Str\\\"omgren colour, b$-$y.\\\\ We also note that the agreement between Templogg and BaSeL for the hottest stars of the sample could be especially useful in view of the well known difficulty of spectroscopic determinations for fast rotating stars. \\\\ As far as the astrophysical applications are concerned, the BaSeL synthetic colours are of particular interest in evolutionary synthesis and colour-magnitude diagram studies of stellar populations, such as open clusters and young associations, where (i) F-type stars are highly common and, of course, (ii) a vast abundance of data are available in B$-$V, U$-$B and b$-$y colours. Concerning the determination of the atmospheric parameters of the COROT potential targets, a result of the present analysis is that all the methods give consistent solutions. In the context of the further $\\sim$1000 potential targets of the COROT exploratory programme, it will be interesting to compare the results of the BaSeL models with those of automated spectral analysis methods (e.g. Katz {\\al} 1998, and Bailer-Jones 2000)."
    },
    "0004/astro-ph0004006_arXiv.txt": {
        "abstract": "Advection-dominated accretion flows (ADAFs) are known to be convectively unstable for low values of the viscosity parameter $\\alpha$.  Two-dimensional axisymmetric hydrodynamic simulations of such flows reveal a radial density profile which is significantly flatter than the $\\rho\\propto r^{-3/2}$ expected for a canonical ADAF. The modified density profile is the result of inward transport of angular momentum by axisymmetric convective eddies.  We present three-dimensional hydrodynamic simulations of convective ADAFs which are free of the assumption of axisymmetry.  We find that the results are qualitatively and quantitatively similar to those obtained with two-dimensional simulations.  In particular, the convective eddies are nearly axisymmetric and transport angular momentum inward. ",
        "introduction": "Analytical and numerical calculations have shown that advection-dominated accretion flows (ADAFs) experience strong convection, especially for low values of the viscosity parameter $\\alpha$ (Narayan \\& Yi 1994, 1995; Igumenshchev, Chen \\& Abramowicz 1996; Igumenshchev \\& Abramowicz 1999, 2000; Stone, Pringle \\& Begelman 1999).  The numerical work reveals that convection significantly modifies the structure of ADAFs (Stone et al. 1999; Narayan, Igumenshchev \\& Abramowicz 2000, hereafter NIA).  The change in the structure is understood to be a consequence of inward transport of angular momentum by convective eddies (NIA; Quataert \\& Gruzinov 2000, hereafter QG; Ryu \\& Goodman 1992; Stone \\& Balbus 1996).  All the numerical simulations done so far have corresponded to axisymmetric two-dimensional flows.  NIA and QG emphasized that the inward transport of angular momentum seen in these simulations is a consequence of axisymmetry (cf. Stone \\& Balbus 1996) and that the results could in principle be qualitatively very different in three dimensions.  We present in this paper the first fully three-dimensional numerical simulations of ADAFs.  The details of our numerical technique are given in \\S2 and the results are described in \\S3.  The paper concludes with a discussion in \\S4. ",
        "conclusions": "We have used numerical simulations to study three-dimensional hydrodynamic accretion flows with inefficient radiative cooling (ADAFs).  The simulated flows have nearly Keplerian angular momentum on the outside, relatively weak viscosity, $\\alpha=0.01$, $0.03$, and an adiabatic index $\\gamma$ equal to either $5/3$ or $4/3$.  We find that the convective motions are nearly axisymmetric (Fig. 3).  As a result, the basic properties of the flows are very similar to those found in previous 2D axisymmetric simulations (Igumenshchev et al. 1996; Igumenshchev \\& Abramowicz 1999, 2000; Stone et al. 1999; NIA).  Our most important result is that convection transports angular momentum inward (Fig. 5).  This property of convection was identified by Ryu \\& Goodman (1992) in analytical work on thin accretion disks. Recently, NIA showed that convection in 2D axisymmetric simulations of ADAFs also moves angular momentum inward.  They and QG argued, however, on the basis of previous work by Stone \\& Balbus (1996), that this result was a consequence of enforcing axisymmetry on the accreting gas.  The present work shows that even in 3D, where the flow is not constrained to be axisymmetric, convective motions do transport angular momentum inward.  In fact, we find that 2D and 3D simulations show good quantitative agreement in the estimate of the ($r\\varphi$)-component of the Reynolds stress tensor, and in the radial dependences of various fluid variables.  In the models, the turbulent pressure due to convection, $P_{turb}=\\langle\\rho v'^2\\rangle /3$, where $v'$ is the velocity fluctuation, is $\\la 10^{-3}$ of the gas pressure $P$.  Therefore, convection does not have a direct dynamical effect on the flow. However, as argued by NIA and QG, convection can have a very strong indirect effect on the structure of the flow if it causes inward transport of angular momentum.  The accretion flow can then achieve a critically balanced state in which the convective transport almost completely cancels the outward transport of angular momentum by viscosity.  This leads to a flat radial density profile $\\rho\\propto r^{-1/2}$.  Since the amount of transport possible with convection is limited, the viscosity parameter $\\alpha$ must be less than a critical value $\\alpha_{crit}\\la0.1$ for the new structure to be realized (NIA).  The value of $\\alpha$ in an ADAF is not known at present. Model-fitting to the observed light curve and spectra of soft X-ray transients suggests that $\\alpha\\sim0.2-0.3$ (Esin, McClintock \\& Narayan 1997).  Numerical MHD simulations in 3D should provide an independent estimate of $\\alpha$ (assuming that viscosity in ADAFs is produced by magnetic stresses).  If such computations confirm that $\\alpha$ in ADAFs is large, then convection may turn out to have only limited dynamical effect on these flows.  Conversely, if $\\alpha<\\alpha_{crit}$, one may expect ADAFs to resemble the analytical solutions described by NIA and QG.  Our results may be relevant also to the physics of rapidly rotating convective stars. In such stars convection may reduce or even suppress the outward transport of angular momentum by magnetic stresses (see Bisnovatyi-Kogan et al. 1979).  An important property of convective accretion flows is the outward transport of energy by convection.  We find that 2D and 3D simulations give similar results; the effective outward luminosity due to convection is $\\sim 10^{-3}-10^{-2}\\dot{M}c^2$ for the values of $\\alpha$ and $\\gamma$ considered here, where $\\dot{M}$ is the net mass accretion rate of the black hole.  Very likely the energy will be radiated on the outside of the accretion flow.  If so, the results have important implications for the spectra and luminosities of accreting black holes and neutron stars (Igumenshchev \\& Abramowicz 2000).  Convection could also introduce variability on different time-scales (Igumenshchev \\& Abramowicz 1999).  The simulated flows presented in this paper show some interesting dependence on the adiabatic index $\\gamma$.  Models with $\\gamma=4/3$ have a steeper density profile (Fig.~2) than models with $\\gamma=5/3$. In addition, Model~D ($\\alpha=0.01$, $\\gamma=4/3$) shows a quasi-stationary meridional circulation (Fig.~4), whereas other models have chaotic convective motions.  The reason for these differences is unclear."
    },
    "0004/astro-ph0004283_arXiv.txt": {
        "abstract": "s{The majority of present efforts to constrain cosmological parameters with cosmic microwave background (CMB) anisotropy data employ approximate likelihood functions, the time consuming nature of a complete analysis being a major obstacle. We have performed a full (unapproximated) likelihood analysis on several experiments that allows us to examine the various assumptions made in these approximate methods and to evaluate their performance.  Our results indicate that care should be taken when using such approaches.  With an improved approximate method, we present some constraints on cosmological parameters using the entire present CMB data set.} ",
        "introduction": "Since their first detection by COBE, the CMB temperature fluctuations have become an essential tool for constraining cosmological parameters. The number of experiments measuring anisotropies at different scales is now close to 20. With time, the observations become more accurate and their angular resolution improves, increasing the number of pixels. This evolution towards higher precision measurements of CMB fluctuations inhibits complete (unapproximated) likelihood analyses over large sets of models (or free parameters)\\footnote{note that the computation time grows like $\\Npix ^3$}. Approximate methods have thus been proposed in order to explore as much parameter space as possible with the complete CMB data set. Easy to use and fast, $\\chi^2$ minimisation has been widely applied. The approximate nature of this method resides in two assumptions: (i) the one dimensional bandpower distribution is Gaussian; (ii) all pertinant information in a map is contained in the flat--band power. The latter is also assumed by other approximations, such as proposed by Bond et al. (1998), Wandelt et al. (1998) and Bartlett et al. (1999). Focussing on a small number of experiments (COBE, SASKATOON and MAX), we have performed a complete likelihood analysis over a reasonably representative set of (Inflationary) models.  This allows us to compare our results with those from a $\\chi^2$ minimisation and to quantitatively examine these two assumptions. ",
        "conclusions": "We have seen how in some cases a Gaussian is a bad representation of the one--dimensional in--band power likelihood function. To improve on $\\chi^2$ methods, we therefore recommend the use of other, more appropriate approximations, as proposed in the literature, or direct interpolation of the exact likelihood curve, when available. Even this may not be entirely satisfactory because flat--band power estimates do not always sufficiently represent a CMB data set.  We have demonstrated that some of MAX and Saskatoon data (bins) prefer non--zero spectral slopes and different in--band powers.  In these cases, the incorporation of additional information, for example, the in--band spectral slope, should lead to a better reproduction of the actual constraints. Taking these remarques into account, have also seen how the first generation experiments favor the simple Inflationary scenario predictions: curvature close to zero and spectral index close to one."
    },
    "0004/astro-ph0004161_arXiv.txt": {
        "abstract": "We present the first detection of a gravitational depletion signal at near-infrared wavelengths, based on deep panoramic images of the cluster Abell 2219 ($z$=0.22) taken with the Cambridge Infrared Survey Instrument (CIRSI) at the prime focus of the 4.2m William Herschel Telescope. Infrared studies of gravitational depletion offer a number of advantages over similar techniques applied at optical wavelengths, and can provide reliable total masses for intermediate redshift clusters. Using the maximum likelihood technique developed by Schneider, King \\& Erben (1999), we detect the gravitational depletion at the $3\\sigma$ confidence level.  By modeling the mass distribution as a singular isothermal sphere and ignoring the uncertainty in the unlensed number counts, we find an Einstein radius of $\\theta_{\\rm E} \\simeq 13.7^{+3.9}_{-4.2}$ arcsec (66\\% confidence limit). This corresponds to a projected velocity dispersion of $\\sigma_v \\sim 800$ km s$^{-1}$, in agreement with constraints from strongly-lensed features. For a Navarro, Frenk and White mass model, the radial dependence observed indicates a best-fitting halo scale length of \\mbox{125 h$^{-1}$ kpc}. We investigate the uncertainties arising from the observed fluctuations in the unlensed number counts, and show that clustering is the dominant source of error. We extend the maximum likelihood method to include the effect of incompleteness, and discuss the prospects of further systematic studies of lensing in the near-infrared band. ",
        "introduction": "Gravitational lensing is now recognised as a valuable probe of the mass distribution in intermediate redshift ($z < 0.4$) galaxy clusters independent of any assumptions about the nature of the lensing material (see reviews by Fort \\& Mellier 1995, Narayan \\& Bartelmann 1997, Mellier 1999, Bartelmann \\& Schneider 1999). For lenses at such redshifts, uncertainties arising from the redshift distribution of background sources are minimized and the angular scales of both weak and strongly-lensed features are well-suited for precise studies.  On small angular scales in super-critical systems, multiply-lensed arcs can provide useful absolute mass estimates provided spectroscopic redshifts and geometric constraints on the location of the various critical lines are available (e.g. Kneib \\etal 1994, 1996). The most promising progress in constraining the mass on large scales has come from weak shear measurements (e.g. Fischer 1999; Hoekstra \\etal 1998; Clowe \\etal 1998) for which sophisticated inversion techniques have been developed (e.g. Kaiser \\& Squires 1993; Kaiser 1995).  However, cluster mass determinations based on weak shear signals are not without limitations.  Firstly, the mass reconstruction from weak shear is non-trivial because of boundary effects due to the finite field of the data.  To counter this effect, several finite-field methods have been proposed (e.g. Seitz \\& Schneider 1996).  Secondly, since the shear arises from the {\\em gradient} of the gravitational potential, the mass reconstruction is only known to within an additive constant.  Accordingly, masses based on shear measurements are subject to a possible upward correction arising from a `mass sheet degeneracy'.  With sufficiently wide-field data it is possible to make the assumption that the surface mass density will approach zero at large distances from the cluster.  However, with independent knowledge of the magnification of the lens, the mass-sheet degeneracy can be broken regardless of the field of view of the data.  Two methods have been proposed to make use of this magnification information and calibrate the absolute scale of the mass distribution, either through the change of image size at fixed surface brightness (Bartelmann \\& Narayan 1995) or source counts (Broadhurst, Taylor \\& Peacock 1995).  This paper is concerned with exploring the role that infrared imaging offers in the gravitational depletion (or `convergence') method for estimating the total masses of clusters. The depletion method was first suggested by Broadhurst, Taylor \\& Peacock (1995) who predicted the diminution in background galaxy surface number density as a function of radius expected behind a lensing cluster. Here we are concerned with extending the original test to near-infrared wavelengths where, in principle, there are significant advantages, namely the flatter number-count slope and a more accurate colour discrimination between foreground and background populations.  The unique wide-field capabilities of the panoramic near-infrared Cambridge Infrared Survey Instrument (CIRSI, Beckett \\etal 1998) allow us to test the method on the rich cluster Abell 2219 ($z$=0.22).  A plan of the paper follows. In $\\S$2 we review the gravitational depletion method illustrating the difficulties associated with its implementation at optical wavelengths and the potential gains of repeating the experiment at near-infrared wavelengths. In $\\S$3 we present new observations of Abell 2219 made at the prime focus of the 4.2m William Herschel telescope and discuss the techniques used to reduce the data as well as the methods used to create a sample of background galaxies. $\\S$4 discusses the depletion signal observed in the context of various mass models and reviews the uncertainties involved. In $\\S$5 we discuss the prospects of routinely estimating cluster masses using this method both with CIRSI and with the upcoming suite of wide field infrared survey telescopes. ",
        "conclusions": "We present a study of the depletion effect around Abell 2219, the first done in the infrared, using the newly-available panoramic IR camera (CIRSI). We show (see Appendix A) that the sample can be effectively exploited beyond the completeness limit, as long as the lensed field and the background field obey the same incompleteness functions and have minimal contamination by false objects. This allows us to detect a clear dip in the radial number density profile of background galaxies at small distances from the cluster centre.  The optical-infrared colours enable us to select a population of red background galaxies with a flat number-count slope in order to optimise the lensing signal and reduce foreground-background confusion.  For a population of red background galaxies with an extremely flat slope ($\\alpha=0.185$), we employ maximum likelihood methods and a SIS model to derive an estimate of the Einstein radius $\\theta_E=13.7^{+3.9}_{-4.2}$ arcsec (66\\% confidence limit when uncertainties is $n_0$ are ignored), resulting velocity dispersion $\\sigma_v=814^{+112}_{-139}$ km s$^{-1}$.  These values are consistent with the location of the redder of the two giant arcs and the estimate $\\sigma_v \\sim 930$ km s$^{-1}$ of Smail \\etal (1995a).  We examine the uncertainty in the number counts, and derive a fractional error of 8.6\\% on the normalisation of the background number density (consistent with clustering on these scales).  When this error is incorporated into the maximum-likelihood analysis the error bars become too large to make a precise statement about the magnitude of the lensing (although our previous measurement is not ruled out).  This demonstrates the crucial importance of the background density $n_0$ for an accurate depletion analysis as discussed in Schneider, King \\& Erben (2000).  Finally, while we cannot at present use our current data to distinguish between alternative models for the cluster mass distribution, we demonstrate the application of the method to the data using a second mass model.  We show the typical depletion curve induced by a cluster with an NFW density profile, and using this profile as a second estimate of the cluster mass we derive a scale length of 125$h^{-1}$ kpc for the cluster.  The wide field IR coverage available using CIRSI is an ideal match to the current generation of wide-field optical imagers.  The availability of multicolour data taken simultaneously but at large radii from the cluster centre provides offset information about the field population and eliminates the need to normalise the depletion curve to the values at the edge of the chip, and the infrared sample provides a cleaner discriminant between foreground and background and a flat number-count slope.  The depletion method is an elegant and relatively simple way to derive mass estimates of clusters of galaxies.  Deeper IR surveys will provide more accurate knowledge of $\\omega(\\theta)$ on arcminute scales, allowing us to better understand the degree to which clustering affects the method.  Future studies of clusters with wide-field optical-infrared data (e.g. with the VISTA telescope, \\mbox{\\tt http://www-star.qmw.ac.uk/$\\sim$jpe/vista}) covering a wide wavelength range could provide more accurately selected background populations via photometric redshifts and allow us to add another sample of independent mass profiles to be compared with those derived from velocity dispersions, X-ray measurements, and strong and weak lensing.  The problem posed by the uncertainty in the background number counts can be overcome by selecting similar clusters (e.g. by their X-ray temperatures) and stacking the depletion signal accordingly to obtain an average cluster mass profile; a feasible project for future IR survey telescopes."
    },
    "0004/astro-ph0004357_arXiv.txt": {
        "abstract": "This paper presents the results from the Jodrell Bank~--~IAC two-element 33~GHz interferometer operated with an element separation of 32.9 wavelengths and hence sensitive to 1\\dg-scale structure on the sky. The level of CMB fluctuations, assuming a flat CMB spatial power spectrum over the range of multipoles $ \\ell = 208 \\pm 18 $, was found using a likelihood analysis to be $\\Delta T_\\ell=63^{+7}_{-6}{\\rm\\,\\mu K}$ at the 68\\%\\, confidence limit, after the subtraction of the contribution of monitored point sources. Other possible foreground contributions have been assessed and are expected to have negligible impact on this result. ",
        "introduction": "\\label{intro}  Observations of the angular power spectrum of the CMB temperature fluctuations are a powerful probe of the fundamental parameters of our universe. The amplitude and spatial distribution of these fluctuations can discriminate between competing cosmological models. Most inflationary models predict more power on scales of 0\\fdg2~--~2$^\\circ$, in the form of a series of peaks. These are due to acoustic oscillations in the photon-baryon fluid, which are frozen into the CMB at recombination, with the peaks corresponding to regions of maximum compression and troughs regions of maximum rarefaction. Hence, the position of the first acoustic peak is a strong test for the geometry of the universe, since it corresponds to a fixed physical scale at the time of recombination projected onto the sky.  The previous result from the Jodrell~Bank~--~IAC 33~GHz interferometer of $\\Delta T_\\ell=43^{+13}_{-12}{\\rm\\,\\mu K}$, reported in Dicker \\et (1999), corresponds to an angular spherical harmonic $\\ell \\sim 110$, equivalent to $\\sim$ 2\\dg structure. To investigate smaller angular scales the baseline was doubled; in this paper we analyze the data from this wide spacing configuration which corresponds to an angular spherical harmonic $\\ell \\sim 210$. The data presented here were taken at the Teide Observatory, Tenerife, between 27 May 1998 and 9 March 1999. The paper is organised as follows. The instrumental configuration is summarised in Section~\\ref{int33}; a full description can be found in Melhuish \\et (1999). The basic data processing is outlined in Section~\\ref{bdpcal}; for a more complete discussion see Dicker \\et (1999). The calibration method is also discussed in Section~\\ref{bdpcal} and the data analysis in Section~\\ref{like}. A derivation of the fluctuation amplitude, after an estimate of the contribution of possible foregrounds, is given in Section~\\ref{fore}. ",
        "conclusions": "\\label{conclusions}  In this paper we describe the results taken with the Jodrell Bank~--~IAC~33~GHz interferometer in its wide spacing configuration, which is sensitive to structure at $\\ell = 208 \\pm 18$. In the final result of $ \\Delta T = 63.0^{+7.0}_{-6.0}\\,{\\rm \\mu K} $ possible foreground contributors have been considered, the most significant of which are point sources. In Section~\\ref{pts} the contribution of strong, monitored sources are removed from the data. To allow for the contribution of unresolved sources the lower limit on \\DT has been increased. In Sections~\\ref{spin}~and~\\ref{dge}, the contributions from dust and diffuse Galactic emission are found to be negligible. We believe our result represents the intrinsic CMB fluctuation signal at $\\ell = 208 \\pm 18 $. The quoted error in our result is dominated by sample variance, resulting from the finite number of beam areas observed.\\\\  Table~\\ref{cmba} and Fig.~\\ref{clplot} show our result alongside others published in the last year from experiments covering similar angular scales. These experiments have been made at a range of different frequencies and regions of the sky. Our result is in good agreement with the published data around $\\ell \\sim 200$. The results at $\\ell \\sim 200$ appear to be converging on a value of $\\Delta T = 60-70 \\,{\\rm \\mu K}$. However, over the wider-$\\ell$ range of 50~--~300 discrepancies appear in the data sets significantly greater than the quoted errors. There appears to be evidence for unknown systematic effects and possible foreground contamination remaining in the data sets.  Our interferometer results show a rise in the amplitude of the power spectrum between $\\ell \\sim 100 $ and $\\ell \\sim 200$. This is intrinsic to the CMB since the parameters of the interferometer system, particularly the calibration, remain the same except for the spacing. The data in Fig.~\\ref{clplot}, despite the discrepancies referred to above, are strongly indicative of a peak in the power spectrum at  $\\ell \\sim 200$.  The interferometer is currently in its narrow spacing configuration ($\\ell \\sim 110 $), observing declinations spaced by 1\\fdg2 from Dec +37.4 to +43.4; these data will significantly reduce the sample variance of the result published by Dicker \\et (1999) to the order of 5 per cent."
    },
    "0004/astro-ph0004027_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Hubble Deep Fields (North and South, or HDF--N and HDF--S) currently offer the deepest optical images of the distant universe, and their exquisite angular resolution provides the opportunity to study the morphologies of galaxies in detail and explore how the galaxy population has transformed with cosmic time.  When considering the properties of HDF galaxies, it is important to keep some facts/limitations of the data set in mind.  The co--moving volume probed by the central, deepest WFPC2 fields (i.e., neglecting the shallower but wider flanking fields) is quite small. For the currently popular ``supernovae + Cepheids'' cosmology, i.e., $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, $H_0 = 70$~km/s/Mpc, which I will adopt here, the comoving volumes out to $z = 1$ and 2 are approximately 5000 and 20000 Mpc$^3$, respectively.  Multiplying these by the normalization of the local galaxy luminosity function (using $\\phi^\\ast = 0.0055h_{70}^3$~Mpc$^{-3}$, from \\cite{Gardner 1997}) gives a rough estimate of the number of $L^\\ast$ galaxies (or their progenitors) expected within the HDF volume.  This is only $\\sim 30$ at $z < 1$.  Thus, small number statistics alone limit the utility of the HDF for a reliable census of bright galaxy properties at $z < 1$, and given real galaxy clustering the possible uncertainties are greater still.  At $1 < z < 2$ there is room enough for $\\sim 80$ $L^\\ast$ galaxies, still small but much better for statistical purposes.  However, at $z > 1$, the optical rest frame, where we are most familiar with local galaxy properties, redshifts into the near--infrared.  WFPC2 imaging therefore measures the rest frame ultraviolet properties of many (probably most) faint HDF galaxies, and thus primarily traces the light of hot, short--lived stars, modulated by the possible effects of extinction.  Systematically comparison of $z > 1$ HDF galaxies to those at $z < 1$ therefore requires deep, near--infrared data.  The HDF--N was observed in the near--infrared from the ground in several different programs (\\cite{Hogg 1997}, \\cite{Barger 1998}, \\cite{Dickinson 1998}). The depth and angular resolution (typically $\\sim 1\\arcsec$) of these data are a poor match to that of the optical WFPC2 HDF images.  Two programs therefore targeted the HDF--N with NICMOS on board HST, providing much deeper images with high angular resolution.   The NICMOS GTOs \\cite{Thompson 1999} imaged one NIC3 field ($\\sim 51\\arcsec\\times 51\\arcsec$) for 49 orbits each at F110W (1.1$\\mu$m) and F160W (1.6$\\mu$m).  We, on the other hand, took the ``wide field'' approach, mosaicing the complete HDF--N with a mean exposure time of 12600s per filter in F110W and F160W. Sensitivity varies over the field of view, but the mean depth is $AB \\approx 26.1$ at $S/N=10$ in an $0\\secspt7$ diameter aperture. The drizzled PSF has FWHM~=~$0\\secspt22$, primarily limited by the NIC3 pixel scale.  The NICMOS data, then, offer the opportunity to study the photometric and morphological properties of HDF galaxies at rest frame optical wavelengths out to $z \\approx 3$.  The $H_{160}$ bandpass\\footnote{I will use AB magnitudes here throughout, and notate the six WFPC2+NICMOS bandpasses by \\U300, \\B450, \\V606, \\I814, \\J110 and \\H160.} samples the rest frame $I$, $V$, and $B$ bands at $z \\approx 1$, 2, and 2.8, respectively.  The combined WFPC2+NICMOS data set offers several attractive new options for studying distant galaxies.  Given the redshift (or a photometric estimate thereof) for an object, we can measure magnitudes and colors at fixed {\\it rest} frame wavelengths, and even generate fixed rest frame {\\it images} to study galaxy morphologies in a common manner over a very broad range of redshifts. I will try to demonstrate here how this can be quite an educational way to study distant galaxies.  For this review, I will discuss the morphologies, colors, and space densities of giant galaxies in the HDF, on and off the Hubble sequence, at $0 < z \\simlt 2$, where the NICMOS images probe rest frame optical wavelengths at the rest frame $V$ band or redder.  The spectroscopic redshift sample (primarily from \\cite{Cohen 1996} and \\cite{Phillips 1997}) becomes thin (nil, in fact) at $1.4 < z < 2$, and I will rely on photometric redshifts fit to our 7--band WFPC2+NICMOS+$K_s$ data by Tamas Budav\\'ari and collaborators (\\cite{Budavari 1999}, \\cite{Csabai 1999}). The near--infrared photometry should substantially improve the reliability of photometric redshifts (\\cite{Connolly 1997}, \\cite{Fernandez-Soto 1999}), but it is nevertheless important to remember that there are, as yet, no spectroscopic calibrators for the $z_{phot}$ estimates in this redshift range.   I will restrict the $z_{phot}$ sample to $\\H160 < 26$, where we believe the photometry, completeness and reliability to be good. ",
        "conclusions": "Deep, high resolution optical--infrared imaging, together with spectroscopy and (where needed) photometric redshifts, offers the means for studying galaxy properties at common rest frame wavelengths over a broad redshift baseline.  At present, the HDF--N is the only place where all the ingredients are available for a single field.  Even there, our NICMOS map is neither as deep (to minimize surface brightness dimming losses) nor as sharp (to match WFPC2 angular resolution) as we would like.  Moreover, with only a single field covering a small cosmic volume, it is risky to generalize HDF--N results to the universe as a whole.  Nevertheless, several trends are apparent.  We find giant disk galaxies with prominent spiral structure, red bulges, and bars out to $z \\sim 1.25$, and red, apparently mature giant ellipticals out to $z_{phot} \\approx 1.8$. The latter probably formed the bulk of their stars at much higher redshift.  However, bluer early type galaxies are also found at $0.5 \\simlt z \\simlt 1.4$, suggesting that some field ellipticals had extended star formation histories, in contrast with what has been observed in rich clusters.  The morphological peculiarities of most irregular HDF galaxies persist in the NICMOS images, suggesting that they are genuinely disturbed or immature objects.  Finally, there seem to be far fewer high luminosity galaxies of all types at $1.4 \\simlt z < 2$ compared to lower redshifts, although this result must be treated with caution given the small volume of the HDF, its susceptibility to clustering variations, and the reliance on photometric redshifts.  With a revived NICMOS in Cycle 10 we can extend such work to other fields.  Better still, the infrared channel of WFC3, scheduled for installation on HST in 2003, will have a substantially wider field of view and smaller pixel scale, providing improved angular resolution compared to the undersampled NICMOS Camera~3.  Deeper and wider surveys will become enormously more efficient, permitting a much better census of the distant universe in the near infrared, and helping to pave the way for NGST."
    },
    "0004/astro-ph0004211_arXiv.txt": {
        "abstract": "We report the development of a new tool to determine the origin of quasi-periodic oscillations (QPOs) in accretion disk systems. The technique uses the source energy spectrum and the energy dependence of the QPO fractional amplitude to restrict the location of the emission region of the modulated photons, which are assumed to originate in the inner accretion disk. Both Monte Carlo and semi-analytical methods are presented. We assume the accretion disk is enshrouded by a slab atmosphere of hot electrons in which unsaturated Compton scattering produces the high-energy spectrum. Properties of the atmosphere, in particular the electron temperature, are assumed functions of radius from the central compact object. We show that our model reproduces the observed energy dependence of the fractional amplitude of the 67\\,Hz QPO in \\source if the QPO is assumed to originate within five gravitational radii from the innermost stable circular orbit. ",
        "introduction": "\\label{sect:intro}  Quasi-periodic oscillations (QPOs) have been seen in the power spectra of X-ray binaries since work with the EXOSAT and \\textsl{Ginga} satellites (van der Klis 1995)\\markcite{V}. With the launch of the Rossi X-Ray Timing Explorer (RXTE) satellite in 1995, advances in instrumentation have allowed observers to measure QPOs at such high frequencies that their origin has been associated with the inner regions of the accretion disks which are expected to exist in binary systems. Evidence on the exact location of the QPO generation, however, remains unclear. Observed properties of the QPOs, such as centroid frequency, width, amplitude, phase lag, and coherence, have been investigated for insight into the physical processes which produce these oscillations.  Another characteristic of QPOs, the energy dependence of their fractional root-mean-square (rms) amplitude, remains as a relatively unexplored tool to diagnose their origin.  In this Letter we report the development of a tool which utilizes the observed energy dependence of the QPO amplitude to locate the accretion disk region in which the modulation may occur.  Analyses of QPO mechanisms must be commensurate with a model for the overall energy spectrum of the source. In particular, a frequent component of X-ray binary spectra is a high-energy power-law, usually interpreted to arise from unsaturated inverse Comptonization by hot electrons. These electrons are typically associated with a high-temperature atmosphere or corona for which the geometry, production, and physical properties are uncertain (Dove, Wilms, \\& Begelman 1997, and references therein)\\markcite{DWB}. Here we demonstrate a method for localizing the QPO emission site using the observed energy dependence of the QPO amplitude and the observed source energy spectrum. We assume a particular form of a Comptonizing atmosphere which adequately fits the observed source spectrum. Ideally one would like to constrain the atmosphere's true structure by employing additional observed properties of the QPO, but this is beyond the scope of this Letter. ",
        "conclusions": ""
    },
    "0004/astro-ph0004094_arXiv.txt": {
        "abstract": "We investigate shear and buoyancy instabilities in radially stratified, magnetized, cylindrical flows, for application to magnetocentrifugally driven winds - such as those from protostars - and to magnetized accretion disks. Our motivation is to characterize the susceptibility of cold MHD disk winds to growing internal perturbations, and to understand the relation of wind instabilities to known accretion disk instabilities. Using four different linear analysis techniques, we identify and study nine principal types of unstable or overstable disturbances, providing numerical and analytic solutions for growth rates for a wide range of parameters.  When magnetic fields are predominantly {\\it toroidal}, as in protostellar winds far from their source, we find the system is susceptible to growth of five different kinds of perturbations: axisymmetric fundamental and toroidal resonance modes, axisymmetric and non-axisymmetric toroidal buoyancy modes, and non-axisymmetric magnetorotational modes.  Winds having a sufficiently steep field gradient ($d\\ln{B}/d\\ln{R} < -0.75$ for a purely toroidal-field case) are globally unstable to the long wavelength fundamental mode concentrated at small radii; these promote the establishment of narrow dense jets in the centers of wider winds. Long-wavelength outer-wind modes are all stable for power-law wind equilibria. The toroidal buoyancy instabilities promote small-scale radial mixing provided the equilibrium has nonzero magnetic forces. For low-temperature toroidal-${\\bf B}$ winds, both axisymmetric and non-axisymmetric magnetorotational instabilities have very low growth rates. The stabilization of buoyancy instabilities by shear and of magnetorotational instabilities by compressibility may be important in allowing cold MHD winds to propagate over vast distances in space.  When magnetic fields are predominantly {\\it poloidal}, as may occur in protostellar winds close to their source or in astrophysical disks, we find the system is susceptible to four additional growing modes: axisymmetric magnetorotational (Balbus-Hawley), axisymmetric poloidal buoyancy, non-axisymmetric geometric buoyancy, and poloidal resonance modes.  The well-known axisymmetric Balbus-Hawley mode has the fastest growth rate.  When the magnetic field is nonuniform, the axisymmetric poloidal buoyancy mode promotes radial mixing on small scales.  The geometric poloidal buoyancy mode requires high $m$, thus is readily stabilized by shear.  Previous work on magnetorotational instabilities has concentrated on near-incompressible systems (accretion disks or stellar interiors). We extend this analysis to allow for compressibility (important in winds). We introduce a ``coherent wavelet'' technique (a WKB temporal approximation), and derive closed-form analytic expressions for instantaneous instability criteria, growth rates, and net amplification factors for generalized non-axisymmetric magnetorotational instabilities in compressible flows with both poloidal and toroidal fields. We confirm that these are in excellent agreement with the results of shearing-sheet temporal integrations, and that ``locally-axisymmetric'' perturbations have the largest amplifications only provided $(\\bold{k}\\cdot\\bold{\\Va})/\\Omega \\lesssim 1$. ",
        "introduction": "The ubiquity of energetic molecular outflows and atomic jets from young stellar objects (YSOs) ranging from deeply embedded infrared sources to classical T Tauri stars suggests that they are an inescapable by-product of star formation \\cite[e.g.,][and references therein]{ric00}. Winds from YSO disks play an important role in shedding angular momentum carried by inflowing material, thereby permitting further accretion in order for the central objects to attain stellar dimensions \\citep{har82,pud86,shu88}. These winds may also have strong effects on the dynamical evolution of the parent cloud by providing a source of turbulent energy \\citep{nor80}, and may help to determine the final masses of stars by reversing the infall of surrounding gas \\citep{shu87}. Therefore, understanding the physics of protostellar winds is of essential importance to the theory of star formation.  Among the various scenarios regarding the origin and nature of the protostellar winds, the most promising is magnetohydrodynamic (MHD) models in which winds are driven by the interaction of the centrifugal force with open magnetic fields threading rapidly rotating disks. These magnetocentrifugally driven wind models can account for observed high mass and momentum losses \\citep{lad85, bac96} and the kinematic and structural characteristics of bipolar molecular outflows in general \\citep{li96}, as well as in specific cases \\cite[e.g., HH111, cf.][]{nag97}. It is, however, still controversial whether the wind originates only from a small magnetosphere-disk interaction region near the central star \\citep{shu94,shu00}, or whether it emanates from an extended region of the disk (following the seminal model of Blandford \\& Payne 1982; see, e.g., K$\\ddot{\\rm o}$nigl \\& Pudritz 2000).  Although the role of magnetic fields in driving protostellar winds is by now well established (at least theoretically), their complementary role in governing the stability properties of winds is less well explored. In addition, azimuthal and vertical shear within winds may also affect their stability properties. Questions of wind stability are potentially important for both large scale and small-scale phenomena.  These include understanding the role of magnetic fields and velocity shear in (a) helping winds to propagate over enormous distances (up to a factor $\\sim 10^6$ in dynamic range) through the ISM in parsec-scale giant HH flows \\citep{rei97}; (b) creating bright HH ``knots'' spaced throughout optical jets \\citep{har00}; (c) governing the angular extent of the emergent wind and establishing the momentum distribution for driving molecular outflows; and (d) converting large-scale ordered flow energy to jet heating through small-scale instabilities.  In addition, it is important to study the dynamical properties of large classes of theoretical wind solutions to test whether they are stable equilibria which can represent real astronomical systems, or whether they are unstable equilibria which should rarely be observed in nature because they evolve rapidly into other configurations.  Previous studies of time-dependent behavior in steadily-input MHD winds have focussed primarily on the Kelvin-Helmholtz instabilities driven by the interaction of the wind surface layer with the ambient medium \\cite[see, e.g., the studies of][and references therein]{app92, har92,ros99}. Generally, heavy jets containing strong toroidal fields are relatively resistant to these instabilities. In addition to ``driven'' instabilities resulting from boundary conditions, winds may also be subject to ``free'' instabilities in their interiors.  Whether a given wind solution is internally unstable must depend on the velocity shear, magnetic geometry, and internal stratification.  One route to studying internal wind instabilities is via time-dependent numerical simulations of winds.  Although such studies \\cite[e.g.,] [and references therein]{ouy97} have yielded intriguing results on the development of episodic knots in MHD winds, the computational demands in carrying out simulations precludes extensive exploration of parameter space, large spatial dynamic range, or very long-term integration.  In addition, some of the time-dependent internal features found in simulations may be introduced by particular choices of inflow boundary conditions that are inconsistent with a steady-state flow, rather than occurring as a result of intrinsic instability of the wind.  Due to the importance of gaseous accretion disks in a wide variety of astrophysical systems, major attention has focussed on disk dynamics, and, in particular, the role of instability-driven turbulence in angular momentum transport \\cite[see, e.g.,][for reviews] {bal98,sto00}. Saturated magnetorotational instabilities \\cite[hereafter MRIs;][]{bal91,bal92,haw91,haw95} represent perhaps the most important local dynamical process affecting disk evolution. Magnetized disk winds share many generic properties with disks, so it is interesting to investigate the potential importance of MRIs in winds.  In this work, we investigate the internal stability of rotating, magnetized protostellar winds to (primarily) local shear and buoyancy modes. We also extend previous studies of local MRIs in accretion disks.  The fundamental difference between ``wind'' and ``disk'' systems in our idealized models is in the absence or presence of gravity as a confining force.  These systems may also be distinguished by the geometry of the magnetic field, with toroidally-dominant fields expected in the wind case, but either poloidal or toroidal fields possible for the disk case. Our most general analysis and results apply to cold flows, but we also perform separate calculations (see \\S 7) including thermal effects which specialize to local analysis of MRIs.  Whether a wind originates from a narrow boundary layer or an extended radial region, the radial expansion of the flow will lead to shear in the azimuthal velocity field of the asymptotic state. The total specific angular momentum of the flow, $J = R(v_\\phi - B_{\\rm pol}B_\\phi/(4\\pi\\rho v_{\\rm pol}))$, is conserved along streamlines (where $B_{\\rm pol}$, $B_\\phi$ and $v_{\\rm pol}$, $v_\\phi$ are the poloidal, toroidal components of the magnetic field and the flow velocity). If the \\Alf\\ mach number $M_{\\rm A}\\equiv v_{\\rm pol}/v_{\\rm A,pol} \\gg 1$, the kinetic part dominates the specific angular momentum and $v_\\phi \\approx J/R$. Thus, the angular velocity $\\Omega = v_\\phi/R \\approx J/R^2$ in the asymptotic wind will have a gradient $d\\ln \\Omega/d\\ln R= d \\ln J /d\\ln R - 2$. If the wind comes from a boundary layer, then $d \\ln J /d\\ln R$ may be small; if the wind originates from a large region with self-similar scalings, then $d \\ln J /d\\ln R = 1/2$. In either case, $d\\ln \\Omega/d\\ln R$ is expected to be a negative, order-unity quantity for wind systems. For thin disk systems (i.e., negligible pressure support), a Keplerian radial profile $d\\ln \\Omega/d\\ln R = -3/2$ is expected if the central mass is dominant. For the analysis of this paper (except where noted otherwise), we adopt $d\\ln \\Omega/d\\ln R=-3/2$ for both ``wind'' and ``disk'' systems, but our qualitative results are insensitive to this assumption. As discussed below, order-unity radial logarithmic gradients may also be expected in the magnetic field strengths; we allow for a range of magnetic gradients.  Significant shear may also exist in the poloidal velocities of jets if they originate from an extended radial region. The asymptotic outflow speed $v_{\\rm 0z}$ generally scales linearly with the rotational speed at the footpoint $R_{\\rm foot}$ of the streamline, so that \\begin{displaymath} \\frac{\\partial \\ln v_{\\rm 0z}}{\\partial \\ln R}= \\frac{\\partial \\ln v_{\\rm 0z}}{\\partial \\ln v_{\\rm foot}} \\frac{\\partial \\ln v_{\\rm foot}}{\\partial \\ln R_{\\rm foot}} \\frac{\\partial \\ln R_{\\rm foot}}{\\partial \\ln R} \\sim -\\frac{1}{2} \\frac{\\partial\\ln R_{\\rm foot}}{\\partial \\ln R}. \\end{displaymath} If the range of footpoint radii is small compared to the range of asymptotic radii (as for a wind from a boundary layer), then $|\\partial\\ln v_{\\rm 0z}/\\partial\\ln R| \\ll 1$; if the radial ranges are comparable, then $\\partial\\ln v_{\\rm 0z}/\\partial\\ln R$ is negative and order unity. For disks, the vertical velocity shear is negligible. We allow for a range of vertical shear rates in the present analysis.  Our analysis consists of developing and solving sets of linearized MHD equations for several general classes of background flows.  Both axisymmetric and non-axisymmetric disturbances are explored.  Even in linearized form, MHD problems present considerable technical challenges. Thus, instead of attacking sophisticated sets of steady-state wind solutions, for many specific examples we will take as an unperturbed configuration one of the power-law cylindrical equilibrium solutions recently identified by \\citet{ost97} as asymptotic states of self-similar disk winds. These have density $\\rho \\propto R^{-q}$, $\\bold{B} \\propto R^{-(1+q)/2}$, $\\bold{v} \\propto R^{-1/2}$, and sound speed $\\cs=0$. Although these adopted initial configurations are relatively simple, they retain general asymptotic characteristics of MHD disk winds in the sense that they have both azimuthal and vertical magnetic field and velocity components with arbitrary ratios, and all the physical variables have radial gradients. These gradients may also be thought of as representing local scalings within a more complex overall stratification.  We also include models without vertical motion but with significant equilibrium gravity to study stability in magnetized astrophysical disks. In our analysis of MRIs, we use equilibria with uniform $\\bold{B}, \\rho$, and $\\cs\\neq 0$, and take $\\Omega\\propto R^{-a}$ with arbitrary $a$.  Because the systems we are studying contain significant azimuthal shear, an arbitrary initial spatial planform is not maintained indefinitely. When $\\Omega'\\neq 0$ and the azimuthal wavenumber $m\\neq 0$ (and/or when $v_{\\rm 0z}'\\neq0$ and the vertical wavenumber $\\kz\\neq0$; the prime represents a differentiation with respect to $R$), spatial wavefunctions may become increasingly radially corrugated in time due to the kinematic shearing of the wavefronts imposed in the initial conditions. If we describe the radial spatial wavefunction in terms of the amplitudes of Fourier coefficients with radial wavenumbers $\\kr$, this corresponds at late times to a secular increase in the amplitudes of large-$\\kr$ terms and decrease in the amplitudes of small-$\\kr$ terms. As we shall show, all the disturbances we identify are stabilized at sufficiently large $\\kr$. Thus, if $m\\Omega'$ and/or $\\kz v_{\\rm 0z}' \\neq 0$, the net amplification factor for any arbitrary initial disturbance is limited by the rate of kinematic growth of radial corrugation compared to the growth rate of any dynamically-driven instabilities.  In previous work, two complementary analytical methods have been used to study small-amplitude disturbances in shearing astrophysical systems. One approach adopts the ``shearing-sheet'' formalism, and integrates the local, time-dependent, linearized wave equations directly to obtain the evolutionary behavior of shearing wavelets treated as an initial-value problem \\citep{gol65, jul66, bal92}. An alternative analytical approach uses WKB techniques to derive dispersion relations for spatial Fourier modes, superpositions of which represent local wavefunctions \\cite[e.g.,][]{shu74,shu92, ryu92,ter96}. This paper includes analyses using both approaches, and also introduces a hybrid technique which we term a ``coherent wavelet'' analysis. We adopt the ``modal'' strategy in order to identify characteristic instantaneous growth rates and physical mechanisms for a wide variety of spatial disturbances. By considering the time over which a spatial pattern is altered by shear, we can estimate the net amplification factor of a given initial modal disturbance. We use the shearing-sheet formalism for studying magnetorotational instabilities, which are cut off at relatively small values of $R\\kr/m$ (whereas the modal analysis applies to large $R\\kr/m$), and also for studying buoyancy instabilities in the high-$m$ regime where modes are short-lived. We show that the results obtained from the shearing-sheet integrations in both cases are in excellent agreement with the predictions of a coherent wavelet analysis, in which time-dependent growth rates $\\gamma(t)$ are computed by time-localizing the shearing-sheet equations and solving an analytic dispersion relation.  The organization of this paper is as follows: We begin by studying instabilities in cold, magnetized winds. In \\S 2, the basic MHD equations and the specific adopted wind equilibria are described. In \\S 3, we analyze the stability of winds to the simplest perturbation with $\\kz=m=0$, where $\\kz$ and $m$ are respectively the vertical and azimuthal wavenumbers of the perturbation. We term these the ``fundamental modes''; we present solutions for stable and unstable global modes under the assumption of free Lagrangian boundary conditions. The modal analysis and general local dispersion relation for cold flows with arbitrary $\\kz$ and $m$ are presented in \\S 4. We present numerical solutions of the dispersion relation for both axisymmetric and non-axisymmetric perturbations in \\S 5. In \\S 6, we classify the unstable or overstable modes and provide the physical interpretation for each mode. Next, we include (variable) thermal pressure terms to compare the susceptibility of cold winds vs. warm disks to shear instabilities. In \\S7 we analyze the axisymmetric Balbus-Hawley instability of poloidal fields and the non-axisymmetric MRI of toroidal fields, discuss the respective instability mechanisms, and provide the corresponding instability criteria. Here, we use the coherent wavelet technique to compute growth rates, and compare with direct shearing-sheet integrations. The generalized instability criteria and net amplification factors for the magnetorotational disturbances with both toroidal and poloidal background fields are also derived. Finally in \\S 8, we summarize and discuss conclusions of the present work. ",
        "conclusions": "\\subsection{General Conclusions Based on Linearized Analysis}  Through linear analyses of the ideal MHD equations, we have explored the stability of shearing, rotating flows to a wide range of (primarily local) disturbances. The chief motivation for this study is to characterize the internal instabilities that could develop in disk winds that emanate from an extended region of a differentially rotating protostellar disk around a young star. The dynamics of such winds has inspired intensive theoretical effort because they may be responsible for observed YSO jets and outflows. In our analysis, we include both results based on generic density, magnetic field, and flow profiles, and results which adopt as initial equilibrium configurations the power-law asymptotic solutions of self-confined cylindrically symmetric winds presented by Ostriker (1997): $\\rho \\propto R^{-q}, B_{\\phi} \\propto B_{\\rm z} \\propto R^{-(1+q)/2}$, and $v_{\\phi} \\propto v_{\\rm z} \\propto R^{-1/2}$. For most of our analysis (\\S\\S 2-6), the flows were assumed to be cold enough that thermal effects can be ignored compared with magnetic forces. To make contact with other studies of shear-induced MHD instabilities in rotating disks, we also consider stability of specific models which include non-zero thermal pressure (\\S7). For the lowest-order ``fundamental'' modes, we employ a normal-mode analysis with free Lagrangian boundary conditions to find eigenvalues and eigenfunctions of both stable and unstable modes (\\S 3). For higher-order modes, we employ three different local techniques to study growth of unstable disturbances: In \\S\\S 4-6, we present numerical and analytic solutions of dispersion relations obtained from normal mode analyses in the $R\\kr \\gg 1$ WKB limit. These are exact for $m\\Omega'+\\kz v_{\\rm 0z}'=0$ disturbances and are valid for a limited time for weakly-shearing circumstances where $\\kr\\gg m|\\Omega'/\\Omega|, \\kz|Rv_{\\rm 0z}'/v_{\\rm 0z}|$ (see also \\S 8.2 below). In \\S 7, we employ temporal integrations of the shearing-sheet equations to study MRI modes (which are cut off for $R\\kr \\gg 1$). We also introduce, in \\S 7.2, a ``coherent wavelet'' formalism which adapts modal analyses for situations where shear is considerable (i.e., small $R\\kr/m$); the coherent wavelet analysis is equivalent to a WKB approach in the temporal domain. We include a comparison of results from the shearing-sheet and coherent-wavelet techniques applied to MRIs, in \\S 7.2.3.  Applying these techniques we have identified a total of nine different unstable or overstable families of disturbances that occur for a wide range of flow parameters: five (FM, BH, ATB, PB, and TR) of them are axisymmetric and the other four (NTB, GPB, PR, and NMRI) are non-axisymmetric. Table 1 summarizes the properties of these modes. The main general conclusions drawn from the analysis in this work can be summarized as follows:  (1) Systems having a primarily azimuthal magnetic field, for example, disk winds far from their source, are susceptible to the fundamental (FM), axisymmetric (ATB) and non-axisymmetric (NTB) toroidal buoyancy, non-axisymmetric magnetorotational instability (NMRI) and toroidal resonance (TR) modes. Unstable fundamental modes (see \\S 3.2) are concentrated in the central parts of jets, and occur in $B_\\phi$-dominated flows when the logarithmic gradient of the magnetic field is steeper than $\\approx -0.75$ (cf.\\ eqs.\\ [24] and [25]). Growth rates of unstable FM are comparable to inner-wind \\Alf\\ frequencies. Long wavelength modes with large amplitudes at large radii are all stable, for power-law wind profiles. The TR mode (see \\S 6.1.4) is an overstability, with growth suppressed when $\\kr$ increases through shear of the vertical velocity, simply becoming oscillatory MHD waves. The axisymmetric toroidal buoyancy mode (ATB; \\S 6.1.2) is activated initially by the buoyancy force and subsequently by bending poloidal magnetic fields. In geometrical form, it is locally similar to the sausage mode of a plasma column confined by toroidal fields, and leads to radial mixing. Because growth rates are larger on smaller scales, ATB can contribute to the generation of local turbulence in disk winds. The non-axisymmetric toroidal buoyancy mode (NTB; see \\S 6.2.1) is much like the Parker instability, but with the centrifugal force replacing the role of external gravity. Although the normal-mode analysis for NTB has the largest temporal validity at small $m/R\\kr$, the instantaneous growth rate increases with increasing $m/R\\kr$ (cf.\\ eq.\\ [52] and Fig.\\ 8c). We thus return to the NTB in \\S 8.2, below, applying time-dependent techniques to study the $R\\kr/m \\ll 1$ limit. Because the NTB is present whenever radial magnetic forces are non-zero, it may be important in promoting radial mixing. Both of the toroidal buoyancy instabilities require non-zero magnetic forces in the equilibrium state. The rarefied and cold conditions of disk winds do not favor the development of the NMRI (see \\S 7.2.2 and \\S 7.3). Like the original (poloidal field) magnetorotational (BH) instability, NMRI requires $d\\ln{\\Omega^2}/d\\ln{R} <0$, but also requires a relatively incompressible medium, as is provided by the relatively dense and warm ($\\cs \\gtrsim \\Va$) conditions in an accretion disk. We show the NMRI vanishes in the limit of $\\cs/\\Va \\rightarrow 0$. We further discuss perturbations in cold, $B_\\phi$-dominated flows in \\S\\S8.2 and 8.3, below.  (2) Systems having primarily axial magnetic fields, for example, accretion disks or winds very near their origin, are susceptible to the Balbus-Hawley (BH), poloidal buoyancy (PB and GPB), and poloidal resonance (PR) modes. The well-known axisymmetric Balbus-Hawley instability (BH; see \\S 6.1.3 and \\S 7.1) is the most efficient member of the family of magnetorotational instabilities (MRIs; see \\S 7.3). It will work to produce channel flows, eventually generating fully-developed MHD turbulence through coupling to non-axisymmetric disturbances in the non-linear stage. Driven by background magnetic pressure and the centrifugal force, the axisymmetric poloidal buoyancy mode (PB; see \\S 6.1.1) requires a gradient in the magnetic field strength to be unstable. If the field distribution is steep enough, the poloidal buoyancy modes would also work effectively to generate radial mixing and turbulence over much smaller scales than the BH instability. Because of their overstable characteristics, the impact on the system of poloidal resonance modes (PR; see \\S 6.2.3) would be best evaluated with a global rather than local formalism. Configurations with shallower background magnetic gradients $(q<1)$ are also subject to a non-axisymmetric poloidal buoyancy instability (GPB; see \\S 6.2.2) which arises in part from geometric effects.  (3) In distinction to the original, incompressible, axisymmetric BH instability, we found that the compressible axisymmetric BH mode is strongly stabilized by the presence of an azimuthal magnetic field if the medium has substantially sub-Alfv$\\acute{\\rm e}$nic sound speeds. For example, in a cold rotating flow with $\\Omega \\propto R^{-3/2}$, the axisymmetric BH instability would be completely suppressed if the local pitch angle $i \\equiv \\tan^{-1}(B_{\\rm z}/B_\\phi)$ is less than 30$^{\\rm o}$ (cf.\\ eq.\\ [47]). In an incompressible medium (as provided by a disk with $\\cs \\gtrsim \\Va$), faster sound waves preserve perturbed toroidal fields from being dispersed by MHD wave motions, thereby providing a favorable condition for the BH instability. When the field configuration is purely poloidal, the compressible BH instability is identical with its incompressible counterpart, independent of temperature (cf.\\ eqs.\\ [58] and [59]).  (4) Even though they share the same instability criterion (cf.\\ eqs.\\ [59b] and [73]), the operating mechanisms for the NMRI of purely toroidal $\\bold{B}$-fields is entirely different from the axisymmetric BH instability of primarily poloidal $\\bold{B}$-fields. In the NMRI (see \\S 7.2), vertical MHD wave motions driven by magnetic pressure play an essential role in the feedback loop for induced radial disturbances, while the axisymmetric BH instability tends to be stabilized by vertical wave motions. Faster sound speeds produce higher growth rates in both instabilities, but for different reasons: in the NMRI by activating azimuthal fluid motions preceded by the vertical MHD wave motions; in the BH instability by maintaining the perturbed azimuthal fields generated by shear (when $\\bp\\neq 0$). Because of their non-axisymmetric nature, the NMRI has a transient growth, stabilized by the growth of $\\kr$ from kinematic azimuthal shear. For the NMRI mode, we show explicitly by comparison to direct temporal integrations of the shearing sheet equations that the growth rate at $\\kr=0$ can be used to provide a good estimate of the net amplification magnitude (see \\S 7.2.3).  (5) The coherent wavelet formalism we develop (\\S7.2) may be used to compute instability criteria and net amplification factors for generalized MRI disturbances with arbitrary magnetic field and wavevector orientations (\\S7.3). Eq.\\ (80) gives the instantaneous dispersion relation for generalized MRIs. For strongly compressible flows ($\\cs/\\Va\\rightarrow 0$), instability does not occur in $B_\\phi$-dominated configurations (cf.\\ eq.\\ [81a]); in this case, flows with an $\\Omega\\propto R^{-3/2}$ rotation law can be unstable only when the magnetic pitch angle $i>24^{\\rm o}$. Because MHD disk winds generally have very small pitch angles, this result has the important implication that such {\\it winds will not be subject to the development of strong internal turbulence that occurs as a consequence of nonlinear MRIs in disks}. The absence of MRIs in cold, $B_\\phi$-dominated winds may be crucial in enabling them to propagate over large distances from their sources. High-$\\kz$ and/or high-$m$ modes of the generalized MRI are stabilized by MHD waves, which is a sharp contrast with the NMRI of {\\it purely toroidal} fields in which maximum growth rates are attained at $\\kz \\rightarrow \\infty$. For incompressible flows, the amplification factor for all MRIs can be written analytically in terms of $\\qA = ({\\bf k \\cdot \\Va})/\\Omega$ and $\\theta = \\tan^{-1}(m/R\\kz)$ (eqs.\\ [83]$\\sim$[85]); within the unstable regime ($\\qA < |d\\ln\\Omega/d\\ln R|^{1/2}$, from eq.\\ [81b]), the amplification is $\\sim \\exp[2\\Omega/(\\kappa\\tan\\theta)]$, favoring ``locally-axisymmetric'' disturbances.  \\subsection{Effect of Shear on Dynamical Growth of Buoyancy Instabilities}  Apart from the results of \\S 7 where we adopted the shearing sheet formulation of the dynamical equations to study MRIs, the results in this work have been elicited on the basis of the local normal mode analyses. As described in \\S 4, these modes may have a limited range of temporal validity, due to the effects of background shear. For axisymmetric disturbances with negligible vertical shear (i.e., $m\\Omega'+\\kz v_{\\rm 0z}'\\rightarrow 0$), the results presented in \\S 5 and \\S 6 are acceptable for all time; the modes with pure imaginary $\\omega$ will show an exponential growth without interruption over arbitrarily long time until nonlinearity sets in. However, for non-axisymmetric disturbances, or for flows with non-negligible vertical shear, unstable modes identified in \\S5 and \\S6 are not purely growing. As time evolves, the differential velocities build up the radial wavenumber through the kinematic shear (cf.\\ eq.\\ [40]), which in turn tends to stop the further growth of disturbances. This can be seen directly through the suppression of instabilities in the local analysis when $\\kr$ is large (cf.\\ Fig.\\ 8). The characteristic time for the wave pattern to change by a fraction $\\epsilon$ is $t=\\epsilon\\kr/|m\\Omega'+\\kz v_{\\rm 0z}'|$; over this interval, the disturbance will be amplified by a factor $\\exp(\\epsilon\\kr{\\rm Im}(\\tom)/|m\\Omega'+\\kz v_{\\rm 0z}'|)$. When $\\kr \\gg m/R, \\kz$, Fig.\\ 8 shows that Im($\\tom) \\propto \\kr^{-1}$, so that the net amplification factor is nearly independent of $\\kr$. Since, however, Im($\\tom$) is not larger than $\\sim|m\\Omega'+\\kz v_{\\rm 0z}'|\\kr^{-1}$, only order-unity amplification can be expected for disturbances which are consistent with the requirements for quasi-steady normal mode analysis.  Because the normal-mode dispersion relations indicate larger values of growth rates when $R\\kr/m$ is small (which is however not self-consistent with the WKB treatment), it is desirable to extend investigation to allow for $R\\kr/m$ small. The coherent wavelet formalism used for the MRI in \\S 7 suggests that when $m\\gg 1$ (or $R\\kz \\gg 1$ for $v_{\\rm 0z}'\\neq 0$ cases), this can be done by regarding $\\kr$ as a time-dependent variable according to eq.\\ (40) and using the asymptotic dispersion relations of \\S 6 (i.e., eq.\\ [43] for PB, eq.\\ [45] for ATB, eq.\\ [52] for NTB, and eq.\\ [53] for GPB). To verify this argument, we specifically consider the NTB modes (which are one of the chief instabilities in $B_\\phi$-dominated winds) and compare the results with the shearing sheet temporal integrations. For the latter, we set $B_{\\rm z}=0$ and integrate eqs.\\ (26)$\\sim$(31) in time, setting all of the coefficients to constant values. The resulting instantaneous growth rates and time evolutions of variables are plotted in Fig.\\ 17 as functions of the normalized time $\\tau = t\\Omega$. We omitted the $\\kr$-dependent term in eq.\\ (26) in order to remove rapid oscillations arising from a phase mismatch between the density and radial velocity; the amplitude evolution is independent of this term. We also neglected the vertical velocity shear and and selected $q=\\kr(0)=\\kz=0$, $R\\Omega = 0.1 \\Vap$, and $m=100$. As initial conditions, we chose 0.1 for every variable except $b_{\\rm R}=0.01$ and integrated the system of the linearized equations. Various curves are computed from direct numerical integrations of shearing sheet equations, while the heavy solid lines are drawn from the normal mode solution, eqs.\\ (52) and (86) (see below), after taking allowance for the time dependence of $\\kr$. The rapid fluctuations of the perturbed variables for $\\tau<0$ are due to MHD waves with high $|\\kr|$, disappearing after variables grow substantially. Again, most of growth occurs over a relatively short period of time near $\\kr \\sim 0$. Note an excellent agreement between the results from two different approaches; we have also obtained similar results with integration from other initial conditions. This confirms that our normal mode results can also be applied to high-$m$ disturbances if $\\kr$ is allowed to vary with time.  Using eq.\\ (52) with $\\kr(t)$ from eq.\\ (40), we integrate $\\tom$ over time to estimate the net amplification for the NTB modes \\begin{equation} \\frac{\\chi_1(t)}{\\chi_1(0)} = \\exp{\\int_0^t {\\rm Im}(\\tom)\\, dt} = \\left[ \\frac{3}{2}\\Omega t + \\sqrt{1+\\left(\\frac{3}{2}\\Omega t\\right)^2} \\right]^{(1-q)\\Vap/3R\\Omega}, \\end{equation} where we put $v_{\\rm 0z}'=\\kr(0)=0$ and assume a Keplerian rotation. Thus instead of an exponential growth, at later time of evolution we have the power-law growth due to the kinematic shear. This behavior is distinct from the MRI modes, which are strictly stable for large enough $\\kr$. However, the continued local growth of buoyancy perturbations is offset by the role of kinematic shear in mixing phases of disturbances. Considering waves on $z$=constant plane in a square with sides $L$, the maximum averaged contrast in any variable relative to the mean value is $(\\lambda/L)\\chi_1(t)$, where $\\lambda$ is the local wavenumber of the waves. With $\\lambda \\sim \\kr(t)^{-1}$ and using eqs.\\ (6) and (86), the average contrast for the NTB modes evolve with time as $\\sim (\\Omega t)^{\\sqrt{2(1-q)}/3-1}$, vanishing as $t \\rightarrow \\infty$ for $0<q<1$. From eq.\\ (86), amplification factors are essentially scale-free. Although there may be a significant growth of the NTB modes on large scales, their dynamical effect on small scales is limited by the phase mixing due to shear.  In the presence of vertical shear, the evolution of ATB modes is also affected by the kinematic growth of $\\kr$ although they are axisymmetric modes. The net amplification of the ATB modes follows a power-law growth as that of the NTB mods does. In fact, eq.\\ (86) with $\\Omega$ replaced by $2v_{\\rm 0z}'/3$ gives the temporal behavior of the net amplification for the ATB modes.  \\subsection{Discussion of Applications to Protostellar Winds}  In order for a disk wind to overcome the gravitational barrier due to a central object and to be centrifugally launched from the surface of a Keplerian disk, the poloidal components of field lines should thread the disk at an angle of 30$^{\\rm o}$ or more from the axis \\citep{bla82}. Once material starts to flow outward along such field lines, it is accelerated primarily in the radial direction by the centrifugal force or by the pressure gradient in the toroidal field. Beyond the \\Alf\\ surface where the local, poloidal component of the flow velocity is equal to that of the \\Alf\\ wave velocity, the magnetic field is not strong enough to play a role of ``a rigid wire\", and the inertia of gas becomes important, winding up the field lines to be progressively more toroidal. In this process, the azimuthal flow velocity decreases below the corotation value. With the increase in the azimuthal component of the magnetic field, the associated hoop stress provides the collimation of the outflow and causes the streamlines to bend upward. The radial flow velocity of the outflow is still positive, although it decreases gradually, eventually becoming zero at the cylindrical asymptotic limit. The power-law solutions (with $\\Vr =\\br =0$) we adopted for many specific cases represent the asymptotic limit of each streamline.  \\citet{ost97} presented self-similar steady solutions for disk winds with cylindrical asymptotics and gave the asymptotic fluid and \\Alf\\ speeds and the location of the asymptotic streamlines, characterized by $q$ together with $\\RA/R_1$ or $R_0/R_1$, where $R_0, \\RA$, and $R_1$ denote the radii of the footpoint, the \\Alf\\ surface, and the asymptote of each given streamline, respectively. Typical numerical values for those solutions are $\\Omega = 0.2\\Omega_0$, $\\Vap/R = 0.42\\Omega_0$, and $\\Vap/v_{\\rm z}=6$ for $q=0.5$, and $\\Omega = 0.1\\Omega_0$ and $\\Vap/R = 0.45\\Omega_0$ for $q=0.9$, where $\\Omega_0$ is the Keplerian rotation rate at the streamline footpoint.  Using these values we can estimate the growth times of the global fundamental mode and the fastest growing (with $\\kr$ near 0) toroidal buoyancy modes. The foregoing analysis suggests that these disturbances will play the most significant dynamical role, given the ineffectiveness of MRIs in cold, $B_\\phi$-dominated flows. We define the time to grow by $\\Gamma$ orders of magnitude as $t_\\Gamma$. For FM, $t_{\\Gamma,\\rm FM} = \\Gamma/(|\\omega|\\log e)$, so from Fig.\\ 2 and eq.\\ (25) we have for $q=0.9$ \\begin{displaymath} t_{\\Gamma,\\rm FM} \\sim \\frac{14\\Gamma}{\\Omega_{\\rm 0,i}} = 25\\; \\Gamma \\;{\\rm days} \\; \\left(\\frac{M}{M_\\odot}\\right)^{-1/2} \\left(\\frac{R_{\\rm 0,i}}{0.1\\;\\rm AU}\\right)^{3/2}. \\end{displaymath} For ATB and NTB, from eqs.\\ (45) and (52), we have for $i\\rightarrow 0$ and $q=0.5$ \\begin{displaymath} t_{\\Gamma,\\rm ATB} \\sim \\frac{10^\\Gamma}{0.21\\Omega_0} = 24\\times 10^\\Gamma\\; {\\rm yrs} \\; \\left(\\frac{M}{M_\\odot}\\right)^{-1/2} \\left(\\frac{R_0}{10\\;\\rm AU}\\right)^{3/2}\\;\\;\\;{\\rm and} \\end{displaymath} \\begin{displaymath} t_{\\Gamma,\\rm NTB} \\sim \\frac{10^{1.4\\Gamma}}{3\\Omega_0} = 1.7\\times 10^{1.4\\Gamma} \\;{\\rm yrs} \\; \\left(\\frac{M}{M_\\odot}\\right)^{-1/2} \\left(\\frac{R_0}{10\\;\\rm AU}\\right)^{3/2}, \\end{displaymath} respectively. Here, $M$ is the mass of the central star and $\\Omega_{\\rm 0,i}$ is the angular speed of a disk at the footpoint $R_{\\rm 0,i}$ of the innermost streamline of winds. The fact that the growth of the FM by a factor $10^\\Gamma$ occurs within $\\sim \\Gamma$ times the rotation period of the disk at the inner radius, far shorter than the lifetime of winds ($\\sim 10^4-10^5$ yrs), suggests that the FM mode is dynamically important in the evolution of the disk winds. When $q$ is small, the radial turbulent mixing of the wind, caused by both axisymmetric and non-axisymmetric toroidal buoyancy modes over a relatively short time, is likely to cascade down into arbitrarily smaller scales to dissipate when the microscopic processes such as magnetic reconnection are included. The released energy in the dissipation processes may heat up the flow, potentially making a significant contribution to the heating of protostellar winds and jets. Because the growth rates of buoyancy modes are proportional to the equilibrium magnetic forces (cf.\\ eqs.\\ [45] and [52]), winds that have approached a force-free magnetic configuration will not be subject to the ATB and NTB instabilities.  The global fundamental mode affects only the inner region of the disk winds (e.g., the central tenth for the model shown in Fig.\\ 3). The logarithmic density gradient $\\partial\\ln\\rho/\\partial\\ln R$ changes relative to the equilibrium value by $\\partial(\\rho_1/\\rho_0)/ \\partial\\ln R$. Fig.\\ 3 shows that as a consequence of the fundamental mode, the very central region becomes more steeply stratified, a surrounding concentric region less steeply stratified, and the balance (most of the wind) remains nearly unchanged. Thus, the FM tends to enhance jetlike structure in the central parts of winds. In addition, because of their tendency to compress interior gas via the FM mechanism, disk winds may help to collimate any interior flows into narrow, fast jets, even when the disk winds themselves have relatively slow motion \\cite[cf.][]{ost97}.  What do the present results imply about the likely radial extent of protostellar winds? First, we note that observed optical jets are unlikely to be isolated structures, because if so they would be significantly overpressured relative to the ambient medium: Since magnetocentrifugal jet models typically predict internal \\Alf\\ speeds comparable to their flow speeds in the range $150\\sim400$ km s$^{-1}$, they have strong internal magnetic pressure $P_{\\rm wind} = B_\\phi^2/8\\pi \\sim \\rho_{\\rm w}\\Va^2/2 \\sim 2.8\\times 10^{-7}$ ergs cm$^{-3}$, which is about 6 orders of magnitude greater than the gas pressure of ambient medium, $P_{\\rm ext} = \\cs^2\\rho_{\\rm ext} \\sim 1.3\\times 10^{-13}$ ergs cm$^{-3}$. Here as reference values we adopted $\\rho_{\\rm w} = 350\\; m_{\\rm H}$ cm$^{-3}$, $\\rho_{\\rm ext} = 200\\; m_{\\rm H}$ cm$^{-3}$, $\\Va = 310$ km s$^{-1}$, and $\\cs = 0.2$ km s$^{-1}$ with $m_{\\rm H}$ being the mass of a hydrogen atom \\citep{har99}. Thus the pressure imbalance at the outer boundary of the jet would cause either the wind as a whole or only its surface layer to expand until a new balance is attained. Based on the results of this paper, if the magnetic field at the base of the wind is stratified less steeply than $R^{-1}$, then perturbations of the outer parts of the wind are stable. As a consequence, only the surface layers of such winds would expand in order to achieve a pressure-balanced condition with the ambient medium. If, on the other hand, the wind's magnetic field is stratified more steeply than $R^{-1}$ at its base, no equilibrium is even possible; the wind would expand as a whole to fill the entire $4\\pi$ steradians, with the inner parts having higher density observable as a narrow optical jet \\cite[cf.][]{shu94,sha98}. Numerical simulations presently underway \\citep{lee00} support previous work indicating that protostellar winds with a wide-angle component are better able to produce observed molecular outflow structures than purely jetlike winds \\cite[see also][]{li96, ost97,ost98,mat99} but further studies are required to determine just how distributed in angle the wind momentum should be - i.e., to discriminate between ``fully-expanded'' and ``surface-expanded'' models. Recent observations \\cite[see, e.g.,][]{ric00} showing a correlation in molecular outflow kinematics with age - with extremely high velocity, highly-collimated flows seen only in the youngest sources - may indicate an underlying temporal evolution from more-collimated to more-expanded protostellar winds."
    },
    "0004/astro-ph0004339_arXiv.txt": {
        "abstract": "We present high resolution spectroscopy of a sample of 24 solar-type stars in the young (15-40\\,Myr), open cluster, NGC 2547. We use our spectra to confirm cluster membership in 23 of these stars, determine projected equatorial velocities and chromospheric activity, and to search for the presence of accretion discs. We have found examples of both fast (\\vsini$>50$\\kms) and slow (\\vsini$<10$\\kms) rotators, but find no evidence for active accretion in any of the sample. The distribution of projected rotation velocities is indistinguishable from the slightly older IC 2391 and IC 2602 clusters, implying similar initial angular momentum distributions and circumstellar disc lifetimes. The presence of very slow rotators indicates that either long (10-40\\,Myr) disc lifetimes or internal differential rotation are needed, or that NGC 2547 (and IC 2391/2602) were born with more slowly rotating stars than are presently seen in even younger clusters and associations. The solar-type stars in NGC 2547 follow a similar rotation-activity relationship to that seen in older clusters. X-ray activity increases until a saturation level is reached for $\\vsini>15-20$\\kms. We are unable to explain why this saturation level, of $\\log (L_{\\rm x}/L_{\\rm bol})\\simeq -3.3$, is a factor of two lower than in other clusters, but rule out anomalously slow rotation rates or uncertainties in X-ray flux calculations. ",
        "introduction": "Open clusters are obvious laboratories in which to study the evolution of stellar X-ray activity. They contain stars with a variety of masses but a similar age, distance and composition. A large amount of {\\em Einstein}, and {\\em ROSAT} X-ray observatory time was spent looking at open clusters (see for instance the reviews of Randich 1997 and Jeffries 1999) and a major achievement of these missions was to show that solar analogues, and stars of even lower mass, in young open clusters, could have X-ray activity orders of magnitude greater than the Sun. This activity correlates with fast rotation and is hypothesized to be due to an internal convective dynamo which generates the magnetic fields that both heat and confine hot coronae. The evolution of X-ray activity is thought to obey an age-rotation-activity paradigm (ARAP). Young single stars are rapidly rotating and active, but lose angular momentum and spin-down as they age, resulting in a decrease of their X-ray activity.  The spin-down of cool stars as they age is not as simple as the $\\Omega \\propto t^{-1/2}$ law once proposed by Skumanich (1972). It appears that G and K stars stars arrive on the main sequence (for instance in the Pleiades -- age $\\simeq100$\\,Myr) with a spread in rotation rates from a few times to a hundred times the solar rotation rate. Subsequently, this spread almost converges by the age (600\\,Myr) of the Hyades. In recent times these phenomena have been understood in terms of braking caused by a magnetically coupled, ionized wind, discussed in detail by (for instance) Barnes \\& Sofia (1996) and Krishnamurthi et al. (1997).  The spread in rotation rates observed at the ZAMS combined with the magnitude and rather narrow distribution of rotation rates among pre main-sequence (PMS) stars implies that, in order to produce the fastest ZAMS rotators, substantial angular momentum must be lost during the PMS phase, and the wind braking mechanism must saturate at high rotation rates (Bouvier, Forestini \\& Allain 1997a).  This could be due to a saturation of the dynamo mechanism itself, or perhaps changes in the magnetic field geometry (Barnes \\& Sofia 1996).  It has long been supposed that during the initial stages of PMS rotation evolution, wind braking was negligible compared with the moment of inertia decrease as a star contracts towards the ZAMS. An alternative angular momentum loss mechanism is required and might be provided by magnetic torques transferring angular momentum to a circumstellar disc (K\\\"{o}nigl 1991; Shu et al. 1994). It can be shown that disc-regulated angular momentum loss leads to almost constant angular velocity as a PMS star shrinks.  A range of disc lifetimes, perhaps connected to the initial disc mass (Armitage \\& Clarke 1996), then results in stars becoming decoupled from their discs at varying ages. Spin up from this point to the ZAMS, combined with wind angular momentum loss that saturates above some threshold rotation rate, can provide a very wide spread of ZAMS rotation rates. This scenario has been modelled extensively by Keppens, MacGregor \\& Charbonneau (1995), Collier-Cameron, Campbell \\& Quaintrell (1995), Bouvier et al. (1997a) and Krishnamurthi et al. (1997).  \\begin{figure*} \\vspace{11.0cm} \\special{psfile=fig1.eps hscale=58 vscale=58 hoffset=30 voffset=320 angle=270} \\caption{X-ray activity, expressed as $L_{\\rm x}/L_{\\rm bol}$, in NGC 2547 (dots) as a function of intrinsic colour, compared with stars in IC 2391/2602 (triangles and crosses). The dotted line indicates the X-ray sensitivity limit for NGC 2547 and the dashed line indicates the ``saturation level'' for X-ray emission reached in all young clusters. The circled NGC 2547 points were selected for high resolution optical spectroscopy in this paper. Triangles represent members of IC 2391/2602 confirmed by spectroscopy in S97. Crosses are X-ray selected members of IC 2391/2602 so far unconfirmed by spectroscopy.} \\end{figure*}   Support for disc-regulated PMS angular momentum loss comes from observations that show a connection between rotation rates in PMS stars and the presence of circumstellar accretion discs. Choi \\& Herbst (1996 and references therein) claimed that rotation rates in the Orion nebula cluster were bimodal, with periods either shorter or longer than 4-5 days. They proposed that the slower group of rotators might be those that still maintained a circumstellar accretion disc. Bouvier et al. (1993) showed that in their sample of Taurus-Auriga PMS stars, the classical T-Tauri stars (CTTS - those stars showing optical signatures of active accretion and circumstellar discs) rotated more slowly on average than the weak T-Tauri stars (WTTS) that showed no sign of circumstellar discs. Furthermore, Edwards et al. (1993) were able to show that a similar situation held when the near-IR signatures of discs were considered for Orion PMS stars. More recently the whole disc-regulated angular momentum loss idea has been challenged by Stassun et al. (1999), who find little correlation between disc signatures and rotation rates in a larger sample of low mass (0.2-0.6\\msun) Orion PMS stars, as well as a population of extremely rapid rotators (periods $<2$ day) both with and without discs. They claim that an alternative to disc-locking is required, if the Orion rotation rate distribution is to evolve into that seen among the ZAMS stars of the Pleiades cluster.  These new ideas on the evolution of stellar rotation rate have been successfully applied to X-ray observations of young clusters of stars ({\\em e.g.} IC 2391 and IC 2602, age $\\simeq 30$\\,Myr -- Stauffer et al. 1997 [S97], Alpha Per, age $\\simeq 60$\\,Myr -- Randich et al. 1996, Pleiades, age $\\simeq100$\\,Myr -- Stauffer et al. 1994).  In every cluster studied so far the ARAP appears to hold, with rotation providing the parameter that determines the level of X-ray emission from a star. This was most clearly demonstrated in the Pleiades by Stauffer et al. (1994). They showed that X-ray activity (measured by $L_{\\rm x}/L_{\\rm bol}$) in solar-type stars, increases with $\\vsini$ until a saturation is reached at $L_{\\rm x}/L_{\\rm bol}\\simeq 10^{-3}$ for $\\vsini>15-20$\\kms, with perhaps a decrease in the \\vsini\\ threshold for lower mass stars with thicker convection zones (Krishnamurthi et al. 1998).  This relationship has since been confirmed in the younger Alpha Per and IC 2391/2602 clusters, with some hint that the X-ray activity may even decline at ultra-fast rotation rates ($\\vsini>100$\\kms -- Randich et al. 1998).    NGC 2547 is an excellent cluster with which to explore some of these issues. Jeffries \\& Tolley (1998 [JT98]) reported X-ray observations of the cluster, detecting a rich population of PMS stars with an age, deduced from fits to low mass isochrones, of $(14\\pm4)$\\,Myr -- somewhat younger than IC 2391/2602. In order to produce the large numbers ($\\sim$85\\% with $\\vsini<20$\\kms) of slow rotators at the age of the Pleiades, the solid-body rotation models of Bouvier et al. (1997a) predict that at $\\sim15$\\,Myr about 15-20\\% of solar-type stars should still possess circumstellar discs and that approximately 50\\% of stars should have a projected equatorial velocity (\\vsini) less than 20\\kms.  Shorter disc lifetimes are allowed if differential rotation between the core and convective envelope is possible (Keppens et al. 1995; Krishnamurthi et al. 1997), and would be more in accord with the (few) available measurements in T-Tauri stars, which suggest median and maximum disc lifetimes of about 1-3\\,Myr and $\\sim10$\\,Myr respectively (Strom et al. 1989; Skrutskie et al. 1990; Edwards et al. 1993; Kenyon \\& Hartmann 1995).  We would therefore expect (according to the ARAP) that approximately half of the solar-type stars in NGC 2547 should show saturated X-ray emission. In fact the distribution of X-ray emission in NGC 2547 does not meet these expectations. Figure 1 compares $L_{\\rm x}/L_{\\rm bol}$ as a function of intrinsic colour for NGC 2547 with that of IC 2391/2602 members (taken from Randich et al. 1995; Patten \\& Simon 1996; S97 and Simon \\& Patten 1998). The X-ray activity appears to peak or saturate in NGC 2547 at a level about a factor of two lower than in IC 2391/2602 for the late F to early K stars. The same is true when NGC 2547 is compared with the Alpha Per and Pleiades clusters (see also Fig.6 in Randich et al. 1995 and Fig. 13 in JT98).  JT98 argued that this was not a simple scaling error in the conversion from X-ray count rates to X-ray fluxes in NGC 2547, because cooler stars ($(B-V)_{0}>1.3$, $(V-I)_{0}>1.8$) appear to have a saturation level, $L_{\\rm x}/L_{\\rm bol}\\simeq 10^{-3}$, that is consistent with other clusters. They suggested that perhaps {\\em all} the solar-type stars in NGC 2547 were rotating slower than a speed of 20\\kms, above which, the X-ray emission of solar-type stars in other clusters appears to saturate.  This might be the case if the stars in NGC 2547 had either started life with much less angular momentum than stars in other clusters or had somehow retained circumstellar accretion discs for longer than usual.  The possibility that disc lifetimes could exceed 10\\,Myr in some circumstances would have important implications for stellar angular momentum loss and the possible formation of planetary systems.  In this paper we present the results of high resolution spectroscopy of a sample of solar-type stars in NGC 2547. Our aim is to test for the presence of active accretion discs, measure rotation rates and to see whether the ARAP can successfully explain the X-ray emission that is seen in these stars. In section 2 we outline the sample of stars we have observed, the observations that were made and their analysis. Section 3 presents the results of these analyses and compares the rotation and activity in NGC 2547 with that in other young clusters. These results are discussed in section 4. ",
        "conclusions": "\\subsection{The age of NGC 2547}  How we interpret the rotation data for NGC 2547 depends a great deal on what we assume its age is. JT98 obtained $14\\pm4$\\,Myr from low mass isochrone fits. The same isochrones would yield ages of 25\\,Myr for IC2391/2602, 50\\,Myr for Alpha Per and 90\\,Myr for the Pleiades, in reasonable agreement with the traditional nuclear turn-off ages determined from high mass stars in the Hertzsprung-Russell diagram (Mermilliod 1981). In the last couple of years these ages have been challenged by measurements of the lithium depletion boundary (LDB) in very low mass cluster stars. The luminosity at which lithium remains unburned in a fully convective star that is contracting towards the ZAMS can be mapped onto an age with reasonable precision (see Ushomirsky et al. 1998). This method has been used to obtain ages of $53\\pm5$\\,Myr for IC 2391 (Barrado y Navascues, Stauffer \\& Patten 1999), $85\\pm10$\\,Myr for Alpha Per (Stauffer et al. 1999) and $125\\pm8$\\,Myr for the Pleiades (Stauffer, Schultz \\& Kirkpatrick 1998). This older age scale implies a modest amount of convective core overshoot to bring the nuclear turn-off ages into agreement (Mazzei \\& Pigatto 1988).   Jeffries et al. (2000) have attempted to find the LDB in NGC 2547, but could only establish a {\\em lower} limit to the age of about 23\\,Myr. If the relative positions of the low mass stars in IC 2391 and NGC 2547 are accepted as an indication of an age difference between the two clusters and we assume that the LDB age for IC 2391 is correct, then NGC 2547 could be as old as 35-40\\,Myr. This would make a substantial difference in the interpretation of the rotation data because at 40\\,Myr, a solar-type star would have completed the vast majority of its contraction (and consequent change in moment of inertia) towards the ZAMS. At 15\\,Myr, the surface rotation rate could still increase by about a factor 1.8 (neglecting angular momentum loss) due to changes in radius and moment of inertia (from the models of D'Antona \\& Mazzitelli 1997). Also of course, any deductions about the lifetimes of circumstellar discs are crucially dependent upon the assumed ages of the younger clusters.  Without  making a judgement on the relative merits of the isochronal and LDB ages, we will need to consider the case of both the younger and older ages scales. i.e. where the ages of NGC 2547 and IC 2391 are roughly 15\\,Myr and 25\\,Myr and where they are roughly 40\\,Myr and 55\\,Myr respectively.   \\subsection{Rotational evolution of solar-type stars}  NGC 2547 (along with the IC 2391/2602 clusters) occupies an important age position between the older, well studied Alpha Per and Pleiades clusters and PMS stars in star forming regions. Previous attempts to study rotation in this age range (10-40\\,Myr) have concentrated on dispersed populations of X-ray selected objects in and around OB and T associations ({\\em e.g.} Bouvier et al. 1997b). These studies appear to show a lack of the slow rotators that are needed to explain the older Pleiades rotation distribution where 50\\% of solar-type stars have $\\vsini<10$\\kms\\ (Queloz et al. 1998). The problems with these investigations of scattered PMS populations are that the stellar ages rely on rather uncertain distances and that X-ray selection might be quite severe, so biasing against the presence of slow rotators. In our NGC 2547 study and in the IC 2391/2602 study of S97, it seems likely that for solar-type stars at least, this selection effect is absent or weak.   The rotation rates we have measured in NGC 2547 largely confirm the results found in IC 2391/2602 by S97. If the clusters had ages of 40\\,Myr and 55\\,Myr respectively and assuming the initial angular momentum distributions and circumstellar disc lifetimes were similar, then we would expect to see little difference in their rotation rate distributions. This is because solar type stars in both clusters would have reached the ZAMS, there would be little moment of inertia change and angular momentum losses over the course of 15\\,Myr might be too small to be measured, except perhaps in the few most rapid rotators. However, if the clusters were aged 15\\,Myr and 25\\,Myr, then assuming solid-body rotation and ignoring angular momentum losses, we might expect to see a rotational spin-up of 50\\% between NGC 2547 and IC 2391/2602. The median \\vsini\\ among the solar-type stars ($0.6<(V-I_{\\rm c})_{0}<1.1$) in NGC 2547 is about 20\\kms and very similar in IC 2391/2602. These figures are based on relatively small numbers but one could view this similarity as a (very) weak argument for the older age scale.  The upper quartile of rotation in the Pleiades solar type stars occurs for $\\vsini>15$\\kms. In IC 2391/2602 it is 50\\kms\\ and about 40\\kms\\ in NGC 2547. The numbers here really are too small to analyse any difference between NGC 2547 and IC 2391/2602. If the solar type stars in all three clusters have completed their PMS contraction, then taken together, the results indicate that if they rotate as solid bodies, the fastest rotating stars must lose 60-70\\% of their angular momentum between 40-50\\,Myr and $\\sim125$\\,Myr. If the clusters are younger then the amount of angular momentum loss must be even greater (80-90\\%) to allow for some contraction and spin up onto the ZAMS. The similarity in $\\vsini$ of the rapid rotators of NGC 2547 and IC 2391/2602 argues for similar initial conditions and circumstellar disc lifetimes.   Overall what we have measured is in very good agreement with the models put forward by Bouvier et al. (1997a). These models start with an observed rotation rate in T-Tauri stars and evolve this using wind angular momentum loss (which saturates at fast rotation rates) and early coupling to a circumstellar disc which is responsible for the predominance of slow rotators at later times. These solid-body rotation models predict maximum rotational velocities of order 150-200\\kms\\ for ages 15-40\\,Myr and that roughly half the solar-type stars in NGC 2547 should have $\\vsini<20$\\kms, which is what we have measured. This high proportion of slow rotators is achieved by assuming circumstellar disc lifetimes as old as 40\\,Myr in some stars and that about 15\\% of stars (the slowly rotating ones) are locked to their discs at 15\\,Myr.  We have no evidence that any of the solar-type stars in NGC 2547 still have discs. This could argue that the cluster has an age of $\\sim40$\\,Myr, but other interpretations are possible. Cameron \\& Campbell (1993) and Armitage \\& Clarke (1996) show that discs with mass accretion rates of only $10^{-10}$\\msun/yr can still enforce rotational equilibrium, whereas mass accretion rates at least an order of magnitude higher are needed to provoke the optical accretion signatures we have searched for in this paper. Some of our slowest rotators may still have remnant discs and it would be worth searching for these in more detail at infrared wavelengths. Another possibility is that the radiative core and convective envelope are not perfectly coupled and that interior differential rotation is possible. This would allow less angular momentum loss to produce slow rotators and thus requires shorter disc lifetimes. However, the core-envelope coupling timescale must be substantially greater than the 10\\,Myr proposed by Keppens et al. (1995), who found that short disc lifetimes ($\\sim6$\\,Myr) could not produce enough slow rotators on the ZAMS in these circumstances. The differential rotation models put forward by Krishnamurthi et al. (1997) show that sufficient slow rotators {\\em can} be produced with disc lifetimes of 3-10\\,Myr if the core-envelope coupling timescale is of order 100\\,Myr.  Alternatively, NGC 2547 and IC 2391/2602 may have been born with a population of slower rotators than are typically seen in even younger clusters and star forming regions. Barnes et al. (1999) show that to explain the slow rotators in IC 2602 requires that 20-30\\% of solar-type stars in IC 2602 needed to have initial periods as slow as 16 days if disc lifetimes are to be limited to $\\leq3$\\,Myr. However, Choi \\& Herbst (1996) and Kearns \\& Herbst (1998) find that 90\\% of stars in the 3\\,Myr old NGC 2264 cluster and 1\\,Myr Orion Nebula and Trapezium clusters have rotation periods faster than 10 days. Thus either long disc lifetimes ($>10$\\,Myr), internal differential rotation with long core-envelope coupling timescales ($\\sim100$\\,Myr), or anomalously small initial angular momenta are required to explain the slow rotators in NGC 2547, IC 2391 and IC 2602.     In any of these scenarios it is at least clear that the NGC 2547 (and IC 2391/2602) data have partially solved the problem of the lack of slow rotators at ages between the PMS stars in Taurus and Orion and the older Alpha Per and Pleiades clusters. The slow rotators are there, it just requires observations of optically selected samples, or at least complete X-ray selected samples which are not biased against slow rotators. We can also say that the combination of initial angular momentum distribution and disc lifetimes must be reasonably similar in NGC 2547 and IC 2391/2602 in order to produce similar \\vsini\\ distributions at their current ages. This is an important result because in the disc-regulated angular momentum loss model, even a small variation in disc locking timescales would produce big changes in the \\vsini\\ distribution as solar-type stars reached the ZAMS.     \\subsection{Anomalous X-ray emission in NGC 2547}  Our main reason for performing high resolution spectroscopy in NGC 2547 was to explain the anomalously low level of X-ray emission in the most active solar type stars in the cluster. Our working hypothesis was that all of the NGC 2547 late F and G stars were rotating slower than 20\\kms\\ and hence none of them showed the saturated level of X-ray emission seen in other young clusters. Our data conclusively reject this hypothesis. There are several examples of very rapid rotators in NGC 2547 and the rotation rate distribution is indistinguishable from the IC 2391 and IC 2602 clusters.  We had also suggested that slow rotation in NGC 2547 might be caused by long lived circumstellar discs or an unusually small amount of initial angular momentum. Neither of these is now required and we have evidence that none of the solar type stars in NGC 2547 possess active accretion discs. This latter discovery in itself places an upper limit on the lifetime of such discs (that could be detected with H$\\alpha$, O\\,{\\sc i} emission or optical veiling) of 40\\,Myr and possibly as low as 15\\,Myr depending on what is finally concluded about the age of NGC 2547.  NGC 2547 appears to follow the ARAP discussed in the introduction up to a point. The levels of activity are appropriate for its age and measured rotation rates, except for the case of coronal activity in the fastest rotators. There, the ``saturation level'' for X-ray emission is a factor of two lower than seen in all other young clusters, but appears to occur at similar rotation rates of 15-20\\kms.  The fact that the X-ray saturation occurs at a similar rotation rate and that the chromospheric activity in NGC 2547 concurs with that in IC 2391 and IC 2602, leads us to suspect that perhaps the X-ray fluxes in NGC 2547 have been underestimated by a factor of two. A further piece of evidence in support of this view is that RX12a, which is almost certainly a background field star, has $\\vsini=35.5$\\kms and $L_{\\rm x}/L_{\\rm bol}=10^{-3.42}$. The spectral type of this object (determined from the cross-correlations) is about K0V, so it does not appear to suffer from much more reddening or absorption than the NGC 2547 stars. If the low X-ray saturation levels in NGC 2547 are peculiar to that cluster for some reason, then it is difficult to explain why a similar phenomenon occurs in an unconnected fast rotating field star in the same direction.  A factor of two underestimate of intrinsic X-ray fluxes might be possible if the interstellar absorption towards the cluster is higher than assumed by JT98, or if the assumed stellar X-ray spectrum is incorrect. Another factor to consider is that the IC 2391, IC 2602, Alpha Per and Pleiades {\\em ROSAT} data we have discussed were obtained with the Position Sensitive Proportional Counter (PSPC), rather than with the High Resolution Imager (HRI) as in the case of NGC 2547. However, we do not believe there could be as much as a factor of two error produced here. The assumed interstellar column density for NGC 2547 of $3\\times10^{20}$\\,cm$^{-2}$ was estimated from the cluster reddening (Bohlin, Savage \\& Drake 1978) and from Lyman $\\alpha$ measurements of early type stars at similar galactic coordinates (Fruscione et al. 1994). A likely error in these estimates is 50\\%, but the column density would have to be as high as (2-3)$\\times10^{21}$\\,cm$^{-2}$ to double the unabsorbed X-ray fluxes in NGC 2547. Such a large column density probably exceeds the column density out of the Galaxy in this direction, as deduced from 21\\,cm maps (Marshall \\& Clark 1984). JT98 also assumed a 1\\,keV optically thin plasma, but again, the consequences of this coronal temperature being wrong by as much as 50\\% only changes X-ray fluxes by 15\\% (David et al. 1996). We also cannot appeal to a mismatch in the calibrations of the PSPC and HRI. Simon \\& Patten (1998) have measured X-ray fluxes for stars in IC 2391 with the HRI and compared them with fluxes measured for the same stars with the PSPC by Patten \\& Simon (1996). They find excellent agreement with essentially no systematic difference and little variability in the X-ray fluxes. The HRI count rate to flux conversion factor used by Simon \\& Patten (1998) is 10\\% smaller than that used for NGC 2547 by JT98, but this is consistent with the smaller assumed value of interstellar absorption for IC 2391. Lastly, we can also say that the bolometric corrections used by various authors to calculate $L_{\\rm x}/L_{\\rm bol}$ are the same to within a few hundredths of a magnitude at all colours, so that none of the discrepancy can be attributed to differences in these.  Our other line of argument for claiming that JT98 calculated the fluxes correctly is more indirect. It seems that the peak levels of X-ray activity in the M stars of NGC 2547 are very similar to those in IC 2391/2602 and other clusters. Fig.1 shows that this is the case. Note that we have not tried to compare mean X-ray emission in these cool stars because the samples are incomplete and heavily biased towards the most active stars. A global factor of two increase in all the X-ray fluxes in NGC 2547 would shift its M dwarfs to activity levels higher than seen in other clusters. The only way of escaping this problem would be if the solar type stars were more absorbed or had radically different coronal spectra to the M dwarfs. Neither of these seem likely but we will be able to rule them out when we have X-ray spectra from the Chandra or XMM satellites.   We have also considered compositional differences between the clusters as a possible solution. Different metallicities could affect convection zone depths and dynamo activity or may simply alter the coronal abundances and emission measure distributions. The metallicities of these three young open clusters are not expected to strongly depart from solar values but they remain undetermined at the present time. However, we have interpreted the strong turn-on of X-ray activity at $0.5<(V-I)_{0}<0.7$ as due to the onset and deepening of convection zones in F stars. The fact that this occurs at a very similar intrinsic colours in NGC 2547, IC 2391 and IC 2602 argues that any compositional differences between the clusters are small and do not greatly influence the magnetic dynamo efficiency.  We are left with a puzzle, perhaps akin to the very different X-ray luminosity functions in the older (600\\,Myr) Hyades and Praesepe clusters (Randich \\& Schmitt 1995; Barrado y Navascues, Stauffer \\& Randich 1998), which also remains unexplained and does not seem likely to result from very different rotational properties. That clusters with similar ages can have different X-ray properties means we must be careful about using X-ray data and the ARAP to draw conclusions about the age distributions of arbitrary samples of stars, whether in clusters or the field."
    },
    "0004/astro-ph0004080_arXiv.txt": {
        "abstract": "Newton's law gets modified in the presence of a cosmological constant by a small repulsive term (antigarvity) that is proportional to the distance. Assuming a value of the cosmological constant consistent with the recent SnIa data ($\\Lambda \\simeq 10^{-52} m^{-2}$) we investigate the significance of this term on various astrophysical scales. We find that on galactic scales or smaller (less than a few tens of kpc) the dynamical effects of the vacuum energy are negligible by several orders of magnitude. On scales of 1Mpc or larger however we find that vacuum energy can significantly affect the dynamics. For example we show that the velocity data in the Local Group of galaxies correspond to galactic masses increased by 35\\% in the presence of vacuum energy. The effect is even more important on larger low density systems like clusters of galaxies or superclusters. ",
        "introduction": "Almost two years ago two groups (the Supernova Cosmology Project \\cite{scpc99} and the High-Z Supernova Team \\cite{hz98,Riess:1998cb} presented evidence that the expansion of the universe is accelerating rather than slowing down. These supernova teams have measured the distances to cosmological supernovae by using the fact that the intrinsic luminosity of Type Ia supernovae, while not always the same, is closely correlated with their decline rate from maximum brightness, which can be independently measured.  These measurements combined with redshift data for these supernovae has led to the prediction of an accelerating universe. A non-zero and positive cosmological constant $\\Lambda$ with \\be \\Lambda \\simeq 10^{-52} m^{-2} \\label{valcc} \\ee  could produce the required repulsive force to explain the accelerating universe phenomenon. A diverse set of other cosmological observations also compellingly suggest that the universe posesses a nonzero cosmological constant corresponding to vacuum energy density of the same order as the matter energy density\\cite{Weinberg:1996xe,Krauss:1995yb,Carroll:1992mt}.  In addition to causing an acceleration to the expansion of the universe the existence of a non-zero cosmological constant would have interesting gravitational effects on various astrophysical scales\\cite{Ostriker:1995su}. For example it would affect gravitational lensing statistics of extragalactic surveys\\cite{Quast:1999fh}, large scale velocity flows\\cite{Zehavi:1999fm} and there have been some claims that even smaller systems (galactic\\cite{Whitehouse:1999rs} and planetary\\cite{ct98}) could be affected in an observable way by the presence of a cosmological constant consistent with cosmological expectations. Even though some of these claims were falsified\\cite{Wright:1998bc,Neupane:1999hr,Roberts:1987ch} the scale dependence of the dynamical effects of vacuum energy remains an interesting open issue.  The effects of the cosmological constant on cosmological scales and on local dynamics can be obtained from the Einstein equations which in the presence of a non-zero cosmological constant are written as \\be R_{\\mu \\nu}-{1\\over 2} g_{\\mu \\nu} R = 8 \\pi G T_{\\mu \\nu} - \\Lambda g_{\\mu \\nu} \\label{einst}\\ee These equations imply the Friedman equation \\be {\\dot R}^2={{8 \\pi G}\\over 3} R^2 \\rho_M - k c^2 + {{\\Lambda R^2}\\over 3}\\equiv {{8 \\pi G}\\over 3} R^2 (\\rho_M + \\rho_\\Lambda) - k c^2 \\label{fried} \\ee  where the vacuum energy density $\\rho_\\Lambda$ is defined as $\\rho_\\Lambda \\equiv \\Lambda /8\\pi G$. Since the vacuum energy does not scale with redshift it is easily seen from eq. (\\ref{fried}) that it can cause acceleration (${\\ddot R}>0$) of the universe expansion.  The observational evidence for accelerating expansion along with constraints on the matter density as derived from dynamical measurements of galaxies and clusters, and additional constraints from the anisotropies of the cosmic microwave background\\cite{White:1998fn}, lead to consistent picture with ${{\\rho_\\Lambda} \\over {\\rho_M}}\\simeq 2$, with the total energy density  approximately equal to the critical density necessary to solve (\\ref{fried}) with k=0 ($\\rho_\\Lambda + \\rho_M\\simeq \\rho_c$).  However, the vacuum energy implied from eq. (\\ref{valcc}) ($10^{-10} erg/cm^3$) is less by many orders of magnitude than any sensible estimate based on particle physics. In addition, the matter density $\\rho_M$  and and the vacuum energy $\\rho_\\Lambda$ evolve at different rates, with $\\rho_M / \\rho_\\Lambda \\simeq R^{-3}$ and it would seem quite unlikely that they would differ today by a factor of order unity. Interesting attempts have been made during the past few years to justify this apparent fine tuning by incorporating evolving scalar fields (quintessence\\cite{Zlatev:1999tr}) or probabilistic arguments based on the anthropic principle\\cite{Garriga:1999bf}).  In addition to the prediction for accelerating universe, the presence of a non-zero cosmological constant  also affects the form of gravitational interactions. The generalized spherically symmetric vacuum solution of eq. (\\ref{einst}) may be written as \\be ds^2=A(r)c^2 dt^2 -dr^2 /A(r) - r^2(d\\theta^2 + \\sin^2\\theta d\\phi^2) \\label{sds}\\ee where $A(r)=1-2GM/c^2 r - \\Lambda r^2/3$. This metric is known as the Schwarzschild-de-Sitter metric\\cite{Gibbons:1977mu} and describes a space that is not asymptotically flat but has an asymptotic curvature induced by the vacuum energy corresponding to $\\Lambda$. In the weak field limit we may use the Schwarzschild-de-Sitter metric to find the corresponding Newtonian potential $\\phi$ \\be g_{00}=A(r)=1+2 \\phi/c^2  \\label{g00}\\ee which leads to \\be \\phi =  {{G M}\\over r} + {{\\Lambda c^2 r^2}\\over 6} \\label{ptl}\\ee This generalized Newtonian potential leads to a gravitational interaction acceleration  \\be g=-{{GM}\\over r^2} + {{\\Lambda c^2}\\over 3} r \\label{acc}\\ee This generalized force includes a repulsive term \\be g_r = {{\\Lambda c^2}\\over 3} r \\label{rep}\\ee which is expected to dominate at distances larger than \\be r_c = ({{3 G M}\\over {\\Lambda c^2}})^{1\\over 3}\\simeq 10^2 ({{{\\bar M}_1}\\over {{\\bar \\Lambda_{52}}}})^{1\\over 3} pc \\simeq 2\\times 10^7 ({{\\bar M}_1\\over {{\\bar \\Lambda_{52}}}})^{1\\over 3} AU \\label{crdist}\\ee where ${\\bar M}_1$ is the mass within a sphere of radius $r_c$  in units of solar masses $M_\\odot = 2 \\times 10^{30} kg$ and  ${\\bar \\Lambda_{52}}$ is the cosmological constant in units of $10^{-52} m^{-2}$.  The question we address in this paper is the following: `What are the effects of the additional repulsive force $g_r$ on the various astrophysical scales?' This issue has been addressed in the literature for particular scales. For example it was shown\\cite{Wright:1998bc} that the effects of this term in the solar system could only become measurable (by modifying the perihelia precession) if the cosmological constant were fourteen orders of magnitude larger than the value implied by the SnIa observations.  A recent study\\cite{Whitehouse:1999rs} has also addressed this issue for galactic scales. In that study an attempt was made to explain the flat rotation curves of galaxies without the existence of dark matter. It was found that a cosmological constant $\\Lambda \\simeq 10^{-48}m^{-2}$ can account for the flat rotation curves of M33 and other galaxies  and is consistent with a theoretical value obtained from the Extended Large Number Hypothesis. This value of $\\Lambda$ however is four orders of magnitude larger than that of eq. (\\ref{valcc}) indicated by supernova and other cosmological observations.  In the next section it will be shown that the vacuum energy required to close the universe (eq. (\\ref{valcc})) has negligible effects on the dynamics of galactic scales (few tens of kpc). The dynamically derived mass to light ratios of galaxies obtained from velocity measurements on galactic scales are modified by less than $0.1\\%$ due to the vacuum energy term of eq. (\\ref{valcc}). This is not true however on cluster scales or larger. Even on the scales of the Local Group of galaxies (about 1Mpc) the gravitational effects of the vacuum energy are significant. We show that the dynamically obtained masses of M31 and the Milky Way must be increased by about 35\\% to compensate the repulsion of the vacuum energy of eq. (\\ref{rep}) and produce the observed relative velocity of the members of the Local Group. The effects of vacuum energy are even more important on larger scales (rich cluster and supercluster). ",
        "conclusions": "We conclude that the Local Group of galaxies is a system that is large enough and with low enough matter density to be a sensitive probe of the gravitational effects of a cosmological constant with value consistent with cosmological expectations and recent SnIa obsrevations. Smaller and denser systems do not have this property. On the other hand the gravitational effects of $\\Lambda$ should be even more pronounced on larger cosmological systems.  The specific cluster dynamic effect which is discussed here offers an independent determination of the cosmological constant contribution to a galactic `dark matter' halo. This is consistent with the fact that galactic masses of galaxies like our Milky Way can be deduced from cluster dynamics such as our Local Group as well as by using non-dynamical methods (e.g. gravitational lensing).  It is legitimate to wonder whether the distinctive action of the cosmological constant over small versus larger scales persists if a different form of vacuum energy such as `quintessence'\\cite{quint} is dominant. This type of effective `scalar matter' possesses an equation of state $p=w\\rho$ ($-1<w<0$) and is much softer than the one associated with the cosmological constant ($w=-1$). As such it interpolates between the latter a normal dark matter component ($w\\geq 0$). Moreover we should expect that it causes a much smaller repulsion effect smoothly passing over to a purely dark matter for increasing $w$. As a consequence our analysis for such a `reduced' type of antigravity would most probably imply a successively smaller discrepancy for the relative total galactic mass contribution of quintessence as derived from the Local Group versus the galactic velocity data for increasing $w$ ($-1 < w < 0$). A recent analysis of the dynamical effects of quintessence on flat galactic rotation curves\\cite{matos} in fact corroborates to this point of view.  The gravitational effects discussed here can only be used as an independent detection method of the cosmological constant, if the galactic masses of systems like the Local Group are measured independently using non-dynamical methods (eg gravitational lensing).  In that case $\\Lambda$ could be obtained using plots like the one shown in Fig. 2. This type of investigation is currently in progress."
    },
    "0004/astro-ph0004343_arXiv.txt": {
        "abstract": "We analyse the radiation pressure instability of the standard, optically thisc disc and propose it as an explanation of the obseved variability of the superluminal microquasar GRS 1915+105. The quantitave agreement with the observed  outburst amplitudes is obtained under the assumption that substantial part of the energy is carried away by the jet. ",
        "introduction": "The X-ray source GRS 1915+105 discovered by GRANAT observatory exhibits both the superluminal jet (Mirabel \\& Rodriguez 1994) and very complex variability pattern (Belloni et al. 1997). The outburst X-ray luminosity is about $10^{39}$ erg/s, being of the order of Eddington luminosity in case of stellar mass black hole system. The spectral analysis reveals the domination of disk-like component during the outbursts, and therefore the variability of the source may be connected with disk evolution. ",
        "conclusions": "We used a physically viable instability due to radiation pressure to model the time dependent behaviour of the optically thick disk itself. The presence of the hot optically thin medium was accounted for by assuming assuming that a certain part of gravitational energy  ($f_{cor}=0.5$) is released in hot corona outside the disc. The mechanism operates for accretion rates above $\\dot M_{0}/ \\dot M_{Edd} \\sim 0.16$ which is in agreement with the observations of GRS 1915+105 - the source does not exhibit the outburst if the mean luminosity temporarily drops below $2.1 \\times 10^{38}$ erg/s. The viscosity coefficient of order of 0.01 is appropriate to model the typical outburst duration. We conclude that the radiation pressure instability is a promising model of the basic instability mechanism underlying the observed variability of this microquasar.  \\vspace{1cm}"
    },
    "0004/astro-ph0004175_arXiv.txt": {
        "abstract": "From analysis of the DIRBE weekly averaged sky maps, we have detected substantial flux in the $60\\micron$ and $100\\micron$ channels in excess of expected zodiacal and Galactic emission.  Two methods are used to separate zodiacal light from more distant emission.  Method I makes use of the time-dependence of the North-South annual variation observed at the ecliptic poles.  This method is robust against errors in the inter-planetary dust (IPD) model, but does not demonstrate isotropy of the background.  Method II measures the ecliptic latitude dependence of the dust over a range of ecliptic latitudes ($|\\beta| > 35\\degree$) at solar elongation $e=90\\degree$.  This allows the excess to be determined in each week of the DIRBE mission for high redundancy, but the results depend weakly on the IPD model.  Both methods give consistent results at $60\\micron$ and $100\\micron$.  The observed signal is consistent with an isotropic background at the level $\\nu I_\\nu = 28.1 \\pm 1.8 \\pm 7(\\rm{syst})\\nWpMMSr$ at $60\\micron$ and $ 24.6 \\pm 2.5 \\pm 8\\nWpMMSr$ at $100\\micron$.  The IR excess detected at 140 and $240\\micron$ by these methods agrees with previous measurements, which are thought to be the cosmic infra-red background (CIB).  The detections at 60 and $100\\micron$ are new.  The integrated IR excess in the window $45-125\\micron$ is $23 \\pm 8\\nWpMMSr$, to be added to the $18 \\pm 4\\nWpMMSr$ previously measured with the DIRBE and FIRAS instruments in the window $125 - 2\\mm$.  While this new excess is not necessarily the CIB, we have ruled out all known sources of emission in the solar system and Galaxy.  We therefore tentatively interpret this signal as the CIB and consider the implications of such energy production from the viewpoint of star formation efficiency and black hole accretion efficiency. However, the IR excess exceeds limits on the CIB derived from the inferred opacity of the IGM to observed TeV photons, thus casting doubt on this interpretation.  There is currently no satisfactory explanation for the $60-100\\micron$ excess. ",
        "introduction": "\\label{sec_intro}  The extra-galactic background light (EBL), from optical to sub-millimeter wavelengths, records the energetics history of galaxy formation.  This background is the cosmic relic of star formation, AGN, and black hole formation. The existence of such a background was discussed first in the optical and near IR (\\cite{partridge67}) and then at other wavelengths (\\cite{low68}; \\cite{peebles69}; \\cite{harwit70}; \\cite{kaufman76}; \\cite{dube79}).  However, measurement of an unresolved background is very difficult at most wavelengths because of numerous foregrounds which may be many times brighter.  Direct measurement of the individual background sources by deep, high resolution imaging has only recently become possible, and only at selected wavelengths.  \\subsection{Current Knowledge of Extragalactic Background}  The Infra-Red Astronomical Satellite (\\IRAS) in 1983 obtained the first full-sky census of far infrared (FIR) point sources.  Some $300,000$ point sources, including $\\sim 20,000$ galaxies, were detected in four bands from $12-100\\micron$ (see \\cite{iras88}). Optical follow-up indicated that typically 30\\% of the bolometric luminosity of these galaxies is radiated in the FIR (\\cite{soifer91}), presumably as thermal radiation from dust heated by optical/UV radiation.  In the case of ultra-luminous IR galaxies (ULIRGs) up to 95\\% of the bolometric luminosity is radiated in the FIR (e.g., Sanders \\& Mirabel 1996), suggesting that optically obscured galaxies might produce a substantial fraction of the extragalactic background light.  Mid IR ($12-25\\micron$) emission from most ULIRGs is centrally concentrated, consistent with AGN activity (\\cite{soifer99}).  More recently, Puget \\etal\\ (1996) used the \\COBE\\ Far InfraRed Absolute Spectrophotometer (FIRAS) data to constrain the FIR/SMM background at longer wavelengths.  They found that the integrated energy of the EBL in the $200\\micron - 2$mm window is comparable to that emmited at optical/near IR wavelengths.  This picture was confirmed by Guiderdoni \\etal\\ (1997) who concluded that the majority of high-$z$ star formation may be hidden by dust.  The FIRAS measurement was greatly refined by Fixsen \\etal\\ (1998) who obtained a fit to the CIB of the form \\BE I_\\nu = (1.3\\pm 0.4) \\times 10^{-5} (\\nu/\\nu_0)^{0.64\\pm0.12} B_\\nu(18.5\\pm1.2\\K) \\EE in the interval $150 < \\nu < 2400\\GHz$ ($2000-125\\micron$), where $\\nu_0=3000\\GHz$ and $B_\\nu(T)$ is the Planck function.  The integrated intensity observed in this frequency interval is $14 \\pm 4\\nWpMMSr$.  This curve is the dotted line shown in Figure \\ref{fig_cibfig}.  Recent efforts to resolve this background into discrete sources with the SCUBA detector on the James Clerk Maxwell Telescope have been very successful (\\cite{smail97}; \\cite{hughes98}; \\cite{barger98}; \\cite{eales99}).  By extending the counts to a limit of 0.5 mJy in cluster-lensed fields, enough counts are found to account for most of the expected $850\\micron$ background of $\\nu I_\\nu = 5\\pm 2 \\times 10^{-1} \\nWpMMSr$ (\\cite{blain99b}).  At shorter wavelengths, the background is not resolvable with current instruments, and zodiacal emission from interplanetary dust (IPD) hampers detection of the unresolved background.  The zodiacal emission peaks at $\\sim 25\\micron$ and dominates any expected signal from the CIB in most DIRBE channels.  Nevertheless, the background has been measured at $140\\micron$ and $240\\micron$ by Schlegel, Finkbeiner, \\& Davis (1998).  A more thorough analysis by the DIRBE team (\\cite{hauser98}) improved those results, and also provided upper limits in the other 8 DIRBE channels (see Figure \\ref{fig_cibfig}). The FIRBACK survey at $175\\micron$ (\\cite{puget99}) using the \\emph{Infrared Space Observatory} (\\ISO) detected point sources to a flux limit of $200\\mJy$, yielding an integrated flux significantly higher than a simple extrapolation of IRAS counts would have predicted, but still much lower than the observed background.  Puget \\etal\\ (1999) claim that a plausible extrapolation of the counts down to $10\\mJy$ would account for the entire background at $175\\micron$.  At wavelengths $\\lambda \\le 3.5\\micron$, emission from Galactic stars dominates that from zodiacal dust.  Using ground-based $2.2\\micron$ counts to remove the stellar foreground, and adopting a value of $7.4\\nWpMMSr$ at $2.2\\micron$, Dwek \\& Arendt (1998) measured the background at $3.5\\micron$ to be $\\nu I_\\nu \\sim 10\\pm 3 \\nWpMMSr$. Gorjian, Wright, \\& Chary (1999) make different assumptions about the stellar foregrounds, yet arrive at a similar result: $16.2\\pm6.4 \\nWpMMSr$ at $2.2 \\micron$ and $9.3\\pm 3\\nWpMMSr$ at $3.5 \\micron$.  At optical wavelengths the Hubble Space Telescope can resolve a large fraction of the extragalactic background.  By integrating the light of the resolved galaxies in the HDF, Pozzetti \\etal\\ (1998) find that the brightness of the extragalactic sky is $2.1^{+0.4}_{-0.3}\\times 10^{-20}\\ergcmHzsr$ in I-band ($\\sim 0.8\\micron$), which is equivalent to $\\nu I_\\nu = 8 \\nWpMMSr$.  In the other Hubble broadband filters, this background obeys $\\nu I_\\nu \\sim \\nu^{-1}$ for 2000 - 8000\\AA. The HDF counts in each filter appear to flatten at the faint end, indicating a possible convergence, although other authors find hints of an optical/UV background as much as twice this large (\\cite{bernstein97}).  Comparing the latest FIR background with the optical/UV background has supported the view that $\\sim 2/3$ of the starlight in high-$z$ galaxies is reprocessed by dust into FIR radiation, a much higher percentage than in local galaxies.  Where, then, are the sources of the FIR background?  It is well established that the ULIRGs are typically interacting galaxies (see Sanders \\& Mirabel 1996), and Williams \\etal\\ (1996) find that most high-$z$ objects in the HDF are interacting.  What is not very well established is the dominant energy source in these objects, especially at the high-luminosity end.  Certainly, many of these objects are powered by starbursts; only the very brightest objects so far observed are primarily powered by AGN (\\cite{genzel98,lutz98}). But this scenario may change at high redshift.  Measurements of the X-ray background provide some constraint on AGN activity, unless the very brightest AGN are so obscured that even $10\\keV$ photons cannot escape.  \\subsection{Interpretation} The above observational evidence allows for a coherent interpretation. Interacting ULIRGs at high-$z$ undergo violent episodes of dusty star formation, with AGN as a minor contributor to the energy.  Roughly 2/3 of this energy is reprocessed to the FIR.  Accounting for present day metallicity of galaxies and the IGM enrichment, the energy from stars allows for an integrated extragalactic background light (EBL) of $I_{\\rm bol} = 50 / (1+z_f) \\nWpMMSr$ (see \\S \\ref{sec_crisis}), which is roughly consistent with current measurements if the formation redshift $z_f \\la 2$.   The IR excess measurements presented in this paper, if they are interpreted as an extragalactic background, would push this paradigm to the limit.  They would indicate an integrated EBL flux in the far IR of $\\sim 40 \\nWpMMSr$ with a hot spectrum, possibly suggesting that AGN dominate the energy input in early galaxies, at least at short wavelengths.  This interpretation might also imply the presence of up to $\\sim 0.15\\%$ of all baryons in black holes, and possibly violate constraints on the X-ray background.  Furthermore, recent measurements of high energy gamma rays place limits on the opacity of the IGM, which is primarily due to pair production on CIB photons. The interpretation of the $60-100\\micron$ excess as CIB raises serious problems; however, we have been unable to identify any alternative source in the Galaxy, local bubble, or Solar system that can account for the emission.  \\subsection{Organization of Paper} We review previous foreground models in \\S\\ref{sec_foregrounds}, and present our two procedures for extracting the IR excess signal in \\S \\ref{sec_methods}.  Our assessment of various systematic errors is given in \\S 4.  In \\S 5 we discuss the energy crisis resulting if this emission is interpreted as an extragalactic background.  In addition, we discuss the constraint on the density of CIB photons implied by the observation of TeV photons from nearby AGN.  These difficulties, as well as general conclusions and prospects for the future appear in \\S 6. ",
        "conclusions": "\\label{sec_discussion} Because we are unable to find any emission component within the Galaxy to explain the observed IR excess, we tentatively interpret the emission as an isotropic extragalactic background (CIB).  In this section we discuss the consequences of such an interpretation.  \\subsection{The EBL Energy Crisis} \\label{sec_crisis}  The integrated energy in the CIB for $60\\micron < \\lambda < 1$mm is $40\\pm12\\nWpMMSr$ according to our measurements and those of Fixsen \\etal\\ (1998) - roughly twice the FIR energy content inferred from Hauser \\etal\\ (1998), and much higher than that predicted by Malkan \\& Stecker (1998).  Where is this energy coming from? A comprehensive analysis by Bond, Carr, \\& Hogan (1986) reviews several possible sources for the FIR background, including primeval galaxies, pregalactic stars, black hole accretion, and decaying particles.  Nuclear fusion within stars at epoch $z$ contributes a radiation energy density $\\Omega_R$ today (in units of the critical density) of \\BE \\Omega_{R}\\approx 0.007 {\\Delta Z \\over 1+z}\\Omega_B {\\cal F} \\EE where $\\Omega_B$ is the  baryon density of the Universe today (e.g. $\\Omega_B h^2 = 0.019 \\pm 0.001$; \\cite{burles98}),  $\\Delta Z$ is the mean metallicity of {\\it all} baryons in the Universe, and ${\\cal F}$ is the fraction of the emitted radiation which is reprocessed into the FIR.  Black hole accretion is another source of energy for the CIB and yields a radiation density \\BE \\Omega_{R}= {\\left(\\epsilon\\over 1-\\epsilon\\right)} (1+z)^{-1} \\Omega_{BH,acc} {\\cal F} \\EE where $\\epsilon\\approx 0.1$ is the efficiency of rest mass to energy conversion and $\\Omega_{BH,acc}$ is the mass density of accreting black holes.  With the definition of the critical density, $\\rho_{crit}=1.9\\times10^{-29}h^2 {\\rm g}\\cm^{-3}$ and $h=H_0/$(100 km s$^{-1}$ Mpc$^{-1})$, the detected background of $40\\nWpMMSr$ becomes $\\Omega_{CIB}=9.8\\times10^{-7}h^{-2}$. If this radiation was generated within stars and mostly reprocessed into the FIR at $z\\approx 2$, then $\\Delta Z \\approx .02$, (regardless of $h$) averaged over the entire Universe.  This seems rather large, and would indicate that most star formation must be well hidden from view.  Another possibility  is generation of the CIB by accretion onto black holes, again at $z\\approx 2$, with most of the radiation emerging in the FIR. If the black holes were formed by the accretion of baryonic matter ($\\Omega\\approx 0.019h^{-2}$), with an efficiency $\\epsilon=0.1$, then the cosmological density of black holes today must be $\\Omega_{BH} \\approx 2.9\\times10^{-5}h^{-2}$ and the fraction of all baryons that must have fallen into massive black holes (assuming no accretion of dark matter) is approximately 0.15\\%.  While this number is large, it is smaller than the black hole mass fraction of 0.5\\% in galactic bulges (\\cite{magorrian98}).  It is also possible that energy extraction from black hole accretion is more efficient than assumed above.  Recent work by Gammie (1999) indicates that for rotating black holes with a magnetic field, accretion efficiencies of $\\epsilon \\approx 0.5$ are possible.  This reduces the requirement to 0.015\\% of baryon mass in black holes, consistent  with the Magorrian \\etal\\ measurement and a disk/bulge mass ratio of $\\sim 30$.  Recent work (\\cite{almaini99}) shows that AGN observed by Beppo-Sax can explain the 30 keV X-ray background, and concludes that most of the energy generation takes place in obscured AGN.  However, those AGN have only been shown to contribute 10-20\\% of the extragalactic background at 240 and $850\\micron$, and are unlikely to produce the measured flux at 60-$100\\micron$.  A recent model by Fabian (1999) suggests that the majority of black holes undergo a highly obscured growth phase, and predicts a population of objects at $z>1$ which emit predominantly hard ($E>30\\keV$) X-rays and FIR/submillimeter photons. The Chandra telescope has recently resolved the majority of the 2-10 keV background, finding two new classes of objects: 1) optically ``faint'' galaxies ($I >> 23$) with very high X-ray to optical ratios, and 2) point-like, luminous hard X-ray sources in the nuclei of normal bright galaxies showing no other sign of activity (\\cite{mushotzky2000}).  The former group is consistent with either early quasars, or extremely dust enshrouded AGN at $z>2$.  Conceivably some combination of accretion onto black holes and powerful starbursts within the core regions of merging galaxies is capable of explaining the IR excess reported here.  The spectral shape of the IR excess also has interesting implications.  Because the $60/100\\micron$ ratio is $\\sim 1$ (in $\\nu I_\\nu$ units) the source must have a high dust temperature of $T \\approx (1+z) 28\\K$.  For $z\\approx 2-3$ this temperature indicates a violent process is at work, and argues in favor of a few hot, bright sources.  SIRTF will provide valuable information about the luminosity function of such sources at 24, 70, and $160\\micron$.    \\subsection{TeV Gamma Crisis} \\label{sec_tev}  Another observational conflict is inevitable if the observed IR excess is of extragalactic origin: the observation of gamma rays by HEGRA (\\cite{konopelko99}, \\cite{aharonian99}) at energies up to $E \\approx 20$TeV from Mkn 501, and by Whipple (\\cite{samuelson98}) up to $E \\approx 10$TeV. The $\\gamma$-ray opacity on CIB photons may be approximated by \\BE \\tau_{\\gamma\\gamma}(E_{\\gamma}) \\approx 0.24\\left(\\frac{E_\\gamma}{1\\TeV}\\right) \\left(\\frac{u(\\epsilon_*)}{10^{-3}\\eV\\cm^{-3}}\\right) (z_S/0.1)h^{-1}_{60} \\EE where $u(\\epsilon_*)=\\epsilon_*^2n(\\epsilon_*)$ is the typical energy density in an energy band centered on $\\epsilon$, $h_{60}$ is the Hubble constant, and $z_s$ is the source redshift (see \\cite{coppi99}).  Integrating the exact cross section over the DEBRA (Diffuse Extragalactic Background Radiation, including CMBR, CIB and optical/UV EBL), Stecker \\& de Jager (1998) find an optical depth of 2.5 at $20\\TeV$ photons using a CIB prediction (Malkan \\& Stecker 1998) based on an extrapolation of \\IRAS\\ counts and other data.  The HEGRA data are consistent with this value at 20 TeV (\\cite{konopelko99}).  However, using the currently accepted measurements of the CIB, Coppi \\& Aharonian (1999) obtain an optical depth $\\tau\\approx5$ for $20\\TeV$ photons.  If we modify the CIB spectrum to take account of the 60-$100\\micron$ CIB measurements given in this paper, $\\tau$ doubles to $\\sim10$.  Even without this doubling, the TeV observations imply that the intrinsic Mkn 501 spectrum is concave upward in the 10-30 TeV range, contradicting the synchrotron self-Compton emission model which requires it to be concave downward. Including the current CIB measurements would imply the intrinsic Mkn 501 spectrum increases a factor of 1000 in $\\nu F_\\nu$ from 10 to $20\\TeV$.  Although very little is yet known about the true intrinsic spectrum of these blazars, this seems an unlikely explanation (however, see Mannheim 1999 for an alternative emission model).  There has also been speculation (\\cite{harwit99}) that multiple TeV photons may be emitted coherently by blazars such as Mkn 501, and might arrive in Earth's atmosphere so close in time and space that they are confused with a single-photon event.  Such coherent emission seems implausible, given the very large phase space available to TeV photons, but some similar mechanism might yet resolve the apparent conflict between TeV gamma observations and the expected opacity of the CIB.  Still more radical explanations have been proposed.  Coleman \\& Glashow (1997, 1999) have proposed that quantum gravity effects cause a small violation of the invariance principle for very high energy particles.  This effect may be large at the Planck scale ($10^{19}\\GeV$) but even at 20 TeV could have measurable consequences. Kifune (1999) has shown that one possible effect of such a violation is a sudden drop in the effective $\\gamma$-ray / IR cross section for $E > 10\\TeV$.  The resulting dispersion relation may also have observable consequences for $\\gamma$-ray bursts at cosmological distances (\\cite{amcam98}).  Future TeV $\\gamma$-ray data may rule out these quantum gravity theories, or perhaps demand a further exploration.    \\subsection{Future Data}  Data to be gathered within the next few years should be adequate to resolve the problems discussed in this paper.  The Space Infra-Red Telescope Facility (SIRTF) will obtain deep number counts at 25, 70, and $160\\micron$ with the Multiband Imaging Photometer (MIPS) (\\cite{rieke96}).  If the integrated flux from measured sources sums to the observed IR excess, then the TeV $\\gamma$-ray measurements must be reinterpreted.  However, if the source density of objects generating the CIB is sufficiently high, then the deep SIRTF images will be confusion limited, and constraints on this diffuse emission will have to come from a fluctuation analysis, in which separation of the extragalactic component from the diffuse cirrus emission of our galaxy will be a limiting factor.  Non-detection by SIRTF would confine the emission to our Galaxy, indicating either Galactic halo dust emission, a serious systematic calibration problem in the COBE/DIRBE instrument, a new emission component, or a serious flaw in the present analysis.  Another test is to observe more blazars at still higher energies.  By combining data from blazars at different distances and different energies, one can place reliable limits on the DEBRA intensity, as long as the intrinsic blazar spectra are at all similar, or can be predicted by their relation to the X-ray emission from these objects.  Work by the HEGRA and Whipple collaborations is underway.    Previous attempts by the DIRBE team to measure the cosmic IR background made use of a sophisticated model of the Inter-Planetary Dust (IPD) (\\cite{kelsall98}).  Such detailed modeling was necessary to establish the isotropy of the CIB signal detected at 140 and $240\\micron$.  The DIRBE team found excess emission at $100\\micron$ also, but doubted the isotropy of the emission and declined to call it a measurement.  In order to recover the valuable information about galaxy formation and evolution contained in the CIB, we have measured the FIR excess at 60 and $100\\micron$.  We analyze the excess DIRBE emission using two different methods. Each of these methods uses a dimensionless parameter derived from the DIRBE data in each week of the mission, parameters that are robust with respect to dust temperature and density variation and detector gain drift.  These statistics are nearly insensitive to details of the IPD model.  It is not necessary to know the IPD emission for every line of sight at every time, although we use the Kelsall \\etal\\ (1998) model as a reference.  Method I uses time variation observed in the flux at the ecliptic poles to measure the background at the poles. Method II uses the spatial ``shape'' of the $e=90\\degree$ data for each week to remove it and yields an independent measurement of the background in each week.  Results derived from these two methods are consistent with each other, giving a background of $\\nu I_\\nu = 28.1 \\pm 1.8 \\pm 7\\nWpMMSr$ at $60\\micron$ and $ 24.6 \\pm 2.5 \\pm 8\\nWpMMSr$ at $100\\micron$.  A variety of arguments rule out alternative sources of emission. Analyses of pulsar dispersion measures and WHAM \\Halpha\\ data demonstrate that signal from the WIM-correlated dust is already accounted for in the cirrus zero point determined from \\HI.  After removing the cirrus, there is no additional component of this emission correlated with $\\csc|b|$, arguing against any additional dust slab aligned with the Galaxy.  The absence of extended far-IR emission halos around nearby galaxies (e.g. M31) rules out dust emission from extended halos.  Unable to find an alternative emission mechanism, we cautiously consider the implications if this excess is an extragalactic background.  The energy required to produce a $44 \\pm 9 \\nWpMMSr$ integrated FIR background is large compared to the energy expected from stellar fusion.  If the observed flux is indeed of extragalactic origin, then stellar fusion is probably not the dominant source of energy in the universe.  Alternative sources such as highly obscured AGN at moderate redshift are a possibility, but would predict that 0.15\\% of all baryons are in black holes at the present time (although this figure is uncertain by a factor of $\\sim 10$ because of uncertainty in accretion efficiency).  A large X-ray background would also be predicted, unless the obscuration is of sufficient optical depth ($N(H) > 10^{24}$) to block it.  The most serious problem with an extragalactic origin of the IR excess is the observation of TeV gamma rays.  The opacity of the measured CIB to 20~TeV photons coming from Mkn 501 is 10 optical depths, much greater than the apparent absorption measured by HEGRA (see \\S \\ref{sec_tev}).  Because of these inconsistencies, there is currently no satisfactory explanation for the observed excess, especially at $60\\micron$.  We continue to search for possible sources of emission in the solar system or Galaxy that could account for the observed emission, and urge caution in the use of these results.  We eagerly await source counts from SIRTF and the X-ray observatories that might help to solve this mystery."
    },
    "0004/astro-ph0004260_arXiv.txt": {
        "abstract": "We present first results from an X-ray study of the Hubble Deep Field North (HDF-N) and its environs obtained using 166~ks of data collected by the Advanced CCD Imaging Spectrometer (ACIS) on board the \\chandra\\ X-ray Observatory. This is the deepest X-ray observation ever reported, and in the HDF-N itself we detect six X-ray sources down to a 0.5--8~keV flux limit of $\\approx 4\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$. Comparing these sources with objects seen in multiwavelength HDF-N studies shows positional coincidences with the extremely red object NICMOS~J123651.74 +621221.4, an active galactic nucleus (AGN), three elliptical galaxies, and one nearby spiral galaxy. The X-ray emission from the ellipticals is consistent with that expected from a hot interstellar medium, and the spiral galaxy emission may arise from a `super-Eddington' X-ray binary or ultraluminous supernova remnant. Four of the X-ray sources have been detected at radio wavelengths. We also place X-ray upper limits on AGN candidates found in the HDF-N, and we present the tightest constraints yet on X-ray emission from the SCUBA submillimeter source population. None of the 10 high-significance submillimeter sources reported in the HDF-N and its vicinity is detected with \\chandra\\ ACIS. These sources appear to be dominated by star formation or have AGN with Compton-thick tori and little circumnuclear X-ray scattering. ",
        "introduction": "The Hubble Deep Field North (HDF-N; Williams et~al. 1996, hereafter W96) provides an unprecedented view of the distant Universe, and the incredible amount of work focused on this patch of sky has led to fundamental insights about the formation and evolution of galaxies. Intense follow-up studies have been performed at a variety of wavelengths including the radio, submillimeter, far-infrared, near-infrared, and optical bands; these investigations have discovered a fascinating array of cosmic objects. Notably lacking until now, however, is a deep X-ray survey of this area.\\footnote{The only previous pointed X-ray observation of the HDF-N was a 21~ks exposure made with the \\rosat\\ High Resolution Imager (HRI). No X-ray sources in the HDF-N were detected due to the limited sensitivity of this observation, and to our knowledge the \\rosat\\ results have never been published.} X-ray surveys provide the most direct and unbiased probe of massive black hole accretion activity throughout the Universe. This is particularly true in the hard X-ray band above 2~keV where most of the energy density of the extragalactic X-ray background (XRB) resides. In addition, X-ray surveys allow studies of starburst galaxies, elliptical galaxies, clusters and groups of galaxies, and other objects. Early observations with the \\chandra\\ X-ray Observatory (hereafter \\chandra; Weisskopf, O'Dell \\& van~Speybroeck 1996) have resolved much ($\\simgt$~60--75\\%) of the 2--8~keV XRB into point sources (Brandt et~al. 2000; Mushotzky et~al. 2000); an even higher fraction ($\\simgt$~70--80\\%) of the 0.5--2~keV X-ray background had already been resolved by \\rosat\\ (e.g., Hasinger et~al. 1998). The key goals now remaining are to resolve the rest of the XRB and to understand in detail the nature and evolution of the sources creating it.  In this letter, we present first results from a deep X-ray survey of the HDF-N area using the \\chandra\\ Advanced CCD Imaging Spectrometer (ACIS; G.P. Garmire et~al. 2000, in preparation). We focus on X-ray sources in the HDF-N itself and in its immediate vicinity where substantial follow-up work has already been performed. The Galactic column density along this line of sight is $(1.6\\pm 0.4)\\times 10^{20}$~cm$^{-2}$ (Stark et~al. 1992), and $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ are adopted. ",
        "conclusions": ""
    },
    "0004/astro-ph0004110_arXiv.txt": {
        "abstract": "Low mass companions to high mass stars are difficult to detect, which is partly why the binary fraction for high mass stars is still poorly constrained.  Low mass companions can be detected more easily however, once high mass stars turn into white dwarfs. These systems are also interesting as the progenitors of a variety of intensely-studied interacting binary systems, like novae, CVs, symbiotics, Ba and CH giants, Feige~24-type systems, and dwarf carbon stars.  We describe a near-IR photometric search for cool red dwarf companions to hot white dwarfs (WDs).  IR photometry offers a sensitive test for low mass main sequence (MS) companions. Our sample of EUV-detected WDs offers several advantages over previous (largely proper motion-selected) WD samples: (1) the high WD temperatures ($24000<T_{\\rm eff}<70000K$) insure excellent IR flux contrast with cool dwarfs (2) the range of evolutionary parameter space occupied by the WDs is considerably narrowed and (3) the random effects of the intervening ISM provide a complete but reasonably-sized sample. While some composite systems have been found {\\it optically} among WDs detected in recent EUV All-Sky Surveys, we develop an IR technique that probes farther down the main sequence, detecting yet more companions.  We use detailed DA model atmosphere fits to optical spectra to predict $K$ magnitudes and distances, against which we contrast our near-IR observations.  Our photometric survey reveals 10 DAs with a significant excess in both $J$ and $K$.  Half are newly discovered, and are most likely previously unrecognized binary systems. Neither the frequency of infrared excess nor the mass estimate of the red dwarf companion correlate with white dwarf mass, as might be expected if either the EUV detectability or mass of the white dwarfs were significantly affected by a companion.  Infrared spectra of these systems should help to determine the mass and spectral type of the cool companions presumably causing the IR excess, leading to better estimates of the mass ratio distribution in binaries.  Counting previously known binaries, and resolved pairs, we find the total binary fraction of the sample is at least a third. Since most WD progenitors had initial masses $\\geq 2\\msun$, we thus provide a photometric measure of the binary fraction of high mass stars that would be difficult to perform in high mass main sequence stars.  We estimate that 90\\% of the companions are of type K or later. ",
        "introduction": "\\label{intro}  EUV-detected stars have revealed in their UV spectra the presence of about 15 hot white dwarf (WD) companions to bright stars in non-interacting binary systems  (e.g., Burleigh, Barstow, \\& Fleming 1997). These WDs are hidden at optical wavelengths, because of their close proximity to much more luminous companions - main sequence stars of spectral type K or earlier, or evolved stars.  Many interacting binary systems where the WD is the primary (i.e., optically brightest) star have also been found among EUV-detected systems (e.g., 6 close, interacting white dwarf/red dwarf binaries by Vennes \\& Thorstensen 1994).  Optical or ultraviolet spectral observations are most commonly used to detect companions to WD primaries, by searching for (1) the presence of narrow Balmer line emission overlying the broad smooth Balmer absorption of the WD, (2) a composite WD + main sequence spectrum, or (3) radial velocity (RV) variations.  However, only WDs with very close, or intrinsically active companions will be found by method (1).  For hot WD systems, composite spectra (2) are only expected to be visible if the companion's spectral type is early. RV variations (3) require multiple observations at high spectral resolution, and detection strongly favors close and/or massive companions.  Such discoveries have been strongly dominated by these selection effects, with companions biased to earlier types than predicted by the simulations of deKool \\& Ritter (1993) and others. Scaling from deKool \\& Ritter's results, Vennes \\& Thorstensen (1994) estimate that ``at least twice as many close binary systems remain to be identified from EUV surveys, most of them with a low mass secondary.'' The resulting sample of binaries known to date therefore must diverge strongly from the intrinsic distribution, in overall normalization, as well as in mass and spectral type of the main sequence companions.  Large proper motion-selected samples (including those observed in the IR e.g., Probst 1983a,b and Zuckerman \\& Becklin 1992), are kinematically biased {\\em against} luminous (statistically more distant) young, hot WDs, and span a bewildering range of temperatures, ages, and metallicities.  Our EUV-selected sample is instead dominated by freshly-minted disk DAs, providing a narrower slice of parameter space for stellar evolution (in the wide systems) or for dynamical evolution (in any new close systems).  Objects known or proposed to contain white dwarf (WD) stars in currently or previously interacting binary systems constitute diverse laboratories for the study of stellar evolution.  A partial list includes novae, cataclysmic variables, symbiotic stars, Ba and CH giants, Feige~24-type systems and dwarf carbon stars (Green \\& Margon 1994).  We describe here a near-IR photometric survey for low mass companions to hot white dwarfs (WDs).  By investigating only EUV-detected WDs, we obtain a very reasonably-sized but complete sample of young WDs, next to which very late-type dwarf companions can be detected in the near-infrared by searching for a $K$ excess. Many hot white dwarfs ($T_{\\rm eff}>24000$K; Finley et al. 1993) have been detected in the recent EUV all-sky surveys.  EUV detection of these hot WDs depends primarily on their temperature, distance, and the intervening Galactic ISM.  Our sample of EUV WDs (whose selection we define below) offers excellent flux contrast in the IR relative to optical; cool companions will almost always be brighter in the $K$ band than the hot WDs.  IR photometry of hot white dwarfs is nearly non-existent in the literature.  To know what $K$ mag to expect for the WDs, we benefit from constraints on log\\,$g$, radius, and $T_{\\rm eff}$ derivable from optical spectra for the WDs in our sample, using NLTE model atmosphere fits (Napiwotzki, Green, \\& Saffer 1999; hereafter NGS99).  The resulting predictions for $K$ magnitudes allow a direct search for any IR excess from a cool companion.  In some cases, IR colors also provide a preliminary spectral type for the companion. ",
        "conclusions": ""
    },
    "0004/astro-ph0004395_arXiv.txt": {
        "abstract": "A continuous stochastic background of gravitational waves (GWs) for burst sources is produced if the mean time interval between the occurrence of bursts is smaller than the average time duration of a single burst at the emission, i.e., the so called duty cycle must be greater than one. To evaluate the background of GWs produced by an ensemble of sources, during their formation, for example, one needs to know the average energy flux emitted during the formation of a single object and the formation rate of such objects as well. In many cases the energy flux emitted during an event of production of GWs is not known in detail, only characteristic values for the dimensionless amplitude and frequencies are known. Here we present a shortcut to calculate stochastic backgrounds of GWs produced from cosmological sources. For this approach it is not necessary to know in detail the energy flux emitted at each frequency. Knowing the characteristic values for the ``lumped'' dimensionless amplitude and frequency we show that it is possible to calculate the stochastic background of GWs produced by an ensemble of sources. ",
        "introduction": "The detection of gravitational radiation will probably mark a new revolution in the history of astronomy. It is worth mentioning that the detection of gravitational waves (GWs) will directly verify the predictions of the general relativity theory concerning the existence or not of such waves, as well as other theories of gravity \\cite{thor87}.  The realm of astrophysics is the place where one finds sources of GWs detectable by the GW observatories. There is a host of possible astrophysical sources of GWs: namely, supernovas, the collapse of a star or star cluster to form a black hole, inspiral and coalescence of compact binaries, the fall of stars and black holes into supermassive black holes, rotating neutron stars, ordinary binary stars, relics of the big bang, vibrating or colliding of monopoles, cosmic strings, etc., among others \\cite{thor87}. Nowadays there is a great effort to study, from the theoretical point of view, which are the most promising sources of GWs to be detected, in particular, their wave forms, characteristic frequencies, and the number of sources a year that one expects to observe.  In a few years, instead of building models trying to understand how the sources of GWs work, it will be possible, starting from the observations (wave forms, amplitudes, polarizations, etc.) to really understand how the GW emission takes place.  Because of the fact the GWs are produced by a large variety of astrophysical sources and cosmological phenomena it is quite probable that the Universe is pervaded by a background of such waves. Binary stars of a variety of stars (ordinary, compact, or combinations of them), population III stars, phase transitions in the early Universe, and cosmic strings are examples of sources that could generate such putative background of GWs.  As the GWs possess a very weak interaction with matter passing through it with impunity, relic radiation (spectral properties, for example) once detected can provide information on the physical conditions from the era in which the GWs were produced. In principle it will be possible, for example, to get information from the epoch when the galaxies and stars started to form and evolve.  Here we present, in particular, a shortcut to the calculation of stochastic background of GWs. For this approach it is not necessary to know in detail the energy flux of the GWs produced in a given burst event. If the characteristic values for the dimensionless amplitude and frequency are known and the event rate is given it is possible to calculate the stochastic background of GWs produced by an ensemble of sources of the same kind.  This paper is organized as follows. In Sec. II we show how to calculate the stochastic background of GWs starting from characteristic values for the dimensionless amplitude and frequency as well as the burst event rate. In Sec. III we apply the idea presented in Sec. II to the calculation of a stochastic background of GWs from a cosmological population of black holes, and finally in Sec. IV we present the conclusions. ",
        "conclusions": "Here we present a shortcut to the calculation of stochastic background of GWs. For this approach it is not necessary to know in detail the energy flux at each frequency of the GWs produced in a given burst event, if the characteristic values for the ``lumped'' dimensionless amplitude and frequency are known, and the event rate is given, it is possible to calculate the stochastic background of  GWs produced by an ensemble of sources of the same kind.  Since one knows the dominant processes of GWs emission one can calculate the stochastic background of an ensemble of black holes. We argue that the same holds for other processes of GWs production, particularly those involving cosmological sources, since the number of sources could be big enough to produce stochastic backgrounds.  We apply this formulation to the study of a stochastic background of GWs produced during the formation of a cosmological population of stellar black holes. We compare the results obtained here with a study by Ferrari {\\it et al.} \\cite{ferr99a}, who take into account in their calculations an average energy flux for the GWs emitted during the formation of black holes obtained from simulations by Stark \\& Piran \\cite{star86}. Our results are in good agreement.  For most sources of GWs only characteristics values for the dimensionless amplitude and frequency are known, if these sources are numerous, a stochastic background of GWs could be produced. We argue that the formulation  presented here could be applied to other calculations of stochastic backgrounds as well."
    },
    "0004/hep-ph0004205_arXiv.txt": {
        "abstract": "\\noindent We calculate the dominant loop corrections from both standard and supersymmetric QCD to the effective coupling of neutralinos to nucleons. The potentially largest corrections come from gluino--squark loop contributions to the Higgs boson couplings to quarks; these corrections also affect the leading spin--independent squark exchange contribution. Ordinary QCD corrections to the effective coupling of CP--even Higgs bosons to two gluons are also sizable. For large $\\tanb$ values, i.e. in the region of parameter space probed by current and near--future direct Dark--Matter search experiments, the total corrections can exceed a factor of three. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004381_arXiv.txt": {
        "abstract": "We explore observed dynamical trends in a wide range of dark-matter-dominated systems (about seven orders of magnitude in mass) to constrain hypothetical dark matter candidates and scenarios of structure formation.  First, we argue that neither generic warm dark matter (collisionless or collisional) nor self-interacting dark matter can be responsible for the observed cores on all scales. Both scenarios predict smaller cores for higher mass systems, in conflict with observations; some cores must instead have a dynamical origin.  Second, we show that the core phase space densities of dwarf spheroidals, rotating dwarf and low surface brightness galaxies, and clusters of galaxies decrease with increasing velocity dispersion like $Q\\propto \\sigma^{-3}\\propto M^{-1}$, as predicted by a simple scaling argument based on merging equilibrium systems, over a range of about eight orders of magnitude in $Q$.  We discuss the processes which set the overall normalization of the observed phase density hierarchy.  As an aside, we note that the observed phase-space scaling behavior and density profiles of dark matter halos both resemble stellar components in elliptical galaxies, likely reflecting a similar collisionless, hierarchical origin. Thus, dark matter halos may suffer from the same systematic departures from homology as seen in ellipticals, possibly explaining the shallower density profiles observed in low mass halos.  Finally, we use the maximum observed phase space density in dwarf spheroidal galaxies to fix a minimum mass for relativistically-decoupled warm dark matter candidates of roughly 700$\\eV$ for thermal fermions, and 300$\\eV$ for degenerate fermions. ",
        "introduction": "\\label{sec:intro}  Recent work has drawn attention to the apparent conflict between predictions of collisionless Cold Dark Matter (CDM) on small scales and observations of rotation curves of dark matter dominated galaxies.  Numerical simulations suggest that in a CDM cosmology, dark matter halos should have steeply rising central cusps ($\\rho\\propto r^{-1.5}$) and high densities.  While the observational conclusions are somewhat ambiguous on the innermost profile shapes (e.g.\\ Swaters et al.\\ 2000, van den Bosch et al.\\ 2000, van den Bosch \\& Swaters 2000, Borriello \\& Salucci 2000, Dalcanton \\& Bernstein 2000, Burkert 1997), rotation curves consistently imply low characteristic halo densities in the central regions.  Because $\\rho(<r)\\propto (V(r)/r)^2$, the characteristic slope of the rotation curve is proportional to the square root of the mean enclosed density. Observations of dark matter dominated galaxies consistently find rotation curves which rise with $V/R \\sim 10-20\\kms$, suggesting much lower characteristic densities than implied by simulations of halos in viable CDM cosmologies, where $V/R \\sim 30-50\\kms$ (Moore et al.\\ 1998, 1999, Navarro et al.\\ 1996, 1997).  The growing belief that there truly is a conflict between theory and observations has led to a renaissance in exploring alternative models for dark matter.  By violating either the ``collisionless'' or ``cold'' properties of traditional CDM, or by considering additional exotic properties, many authors have sought to preserve the successes of CDM on large scales, while modifying the manifestations of dark matter on small scales (Spergel \\& Steinhart 1999, Hogan \\& Dalcanton 2000, Sommer-Larson \\& Dolgov 1999, Mohapatra \\& Teplitz 2000, Peebles 2000, Goodman 2000, Riotto \\& Tkachev 2000, Hu et al.\\ 2000, Shi \\& Fuller 1999, Colin et al. 2000.)  In this paper, we   place broad constraints upon these alternatives to CDM, by revisiting observations of the structure and phase space density of halos over a wide range of scales.  Many of the above alternative dark matter scenarios make specific predictions for the sizes of dark matter cores as a function of mass scale (see Spergel \\& Steinhart 1999, Hogan \\& Dalcanton 2000, Hannestad 1999, Burkert 2000, Kochanek \\& White 2000, Yoshida et al.\\ 2000, Moore et al.\\ 2000, although some of these are in conflict with each other, particularly regarding the long term stability of self-interacting cores).  We confront these predictions with existing limits on the scale of inner halo cores in \\S\\ref{sec:cores}, and argue that neither packing of phase space nor highly collisional dark matter can be primarily responsible for the observed behavior of dark matter cores at all scales.  In addition to discussing the scaling behavior of dark matter cores, we focus upon trends of phase space density $Q$ as a probe of structure formation history.  In \\S\\ref{sec:evolution} we summarize the statistical and dynamical behavior of $Q$ in hierarchical clustering, including a simple argument predicting the decrease of $Q$ with mass and velocity dispersion.  Current observational limits on the variation of phase space density with mass are reviewed and summarized in \\S\\ref{sec:observations}.  We show in \\S\\ref{sec:interp} that the phase space density of the dark matter halos is a very strongly declining function of mass (consistent with the nearly constant density seen across a similar mass scale in Firmani et al.\\ 2000).  This behavior is exactly as predicted by models where halos formed hierarchically, with successive mergers leading to phase mixing and dilution of the coarse-grained distribution function.  We also note that the observed decline in phase space density for dark matter halos and their predicted density profiles resembles that observed in the baryon-dominated central regions of giant elliptical galaxies, and that by analogy, the structure of dark matter halos may also suffer from the systematic departures from homology similar seen in ellipticals.  Finally, we use the maximum observed phase space densities to   derive a limit on particle mass of $> 700\\eV$ for relativistically decoupled thermal relics. ",
        "conclusions": "\\label{sec:conclusion}  The above examinations of the existing data on the structure and phase space density of dark matter halos yields a number of conclusions which may be relevant to constraining the nature of the dark matter.  First, the behavior of halo core size with increasing mass suggests that it is unlikely that either phase-space packing or highly collisional dark matter is sufficient for simultaneously explaining the dark matter cores of dwarf spheroidals, rotating dwarf galaxies, and clusters of galaxies.  The generic $r_{core}\\propto1/\\sqrt{\\sigma}$ behavior of these scenarios for core formation would predict far larger cores for the dwarf spheroidals than for larger systems, in contrast to observational evidence.  Second, if there were still any doubt, the dramatic decrease in the characteristic phase space density with increasing mass is extremely strong evidence for a ``bottom-up'' hierarchical buildup of bound structures.  Given that the phase space density can never increase with successive mergers (in the absence of dissipation), smaller structures cannot generically have fragmented from larger ones. The trend points to dwarf spheroidals as the lowest entropy, and therefore dynamically most primitive observable systems.  Third, the observed dependence of $Q\\propto \\sigma^{-3}$ approximately agrees (over eight orders of magnitude in $Q$!) with a simple scaling relation that assumes the minimal decrease compatible with a gentle merging hierarchy with virial equilibrium and homologous tidal stripping at each stage.  The same scaling is also compatible with a simple synchronous collapse of different mass systems to constant virial density. The two descriptions are both appropriate, at different stages, for the collisionless hierarchy predicted in CDM models. We also discuss the physics which sets the normalization of the $Q$-$\\sigma$ scaling.  Fourth, examination of the phase space density for dark matter halos suggests some parallels to elliptical galaxies.  In most scenarios, both dark matter halos and cluster elliptical galaxies are thought to be formed through collisionless merging and accretion.  We show that observationally, both systems show similar decreases in the coarse-grained phase space density with increasing velocity dispersion.  They also show surprisingly similar density profiles.  If elliptical galaxies can be used as a rough analog to dark matter profiles, then they suggest that the structure of dark matter halos may undergo subtle, systematic deviations from homology, leading to somewhat flatter inner cores (and steeper fall-off at large radii) for low mass halos.  The amplitude of these deviations are sufficiently small that they are unlikely to be well resolved in current numerical simulations, though they may have already been detected observationally. This non-homology could help to alleviate some of the discrepancies between observations of rotating dwarfs and predictions of dark matter simulations on the smallest, most poorly resolved scales.  This parallel also provides suggestive, although not conclusive, evidence that dark matter halos are indeed collisionless on the scale of halo cores.  Fifth, the very high phase space densities of dwarf spheroidals can be used to place constraints on the masses of potential dark matter candidates.  For dark matter particles with 2 degrees of freedom, decoupled while still relativistic, masses of $m_X > 700\\eV$ (thermal fermions) or $m_X > 300\\eV$ (degenerate fermions) are preferred. Because systems with smaller masses and higher phase space densities than dwarf spheroidals may exist, the actual particle masses may be substantially higher than these limits. This seems likely since smaller $\\sigma$ halos would be invisible even if they exist; in other words we see halos populated with stars right up to the cooling limit, which would be a coincidence if they also correspond to the phase density limit. This interpretation is again consistent with the view that dwarf spheroidals are the most primitive bound systems so far observed.  Finally, we conjecture that the addition of primordial velocity dispersion can help to reconcile the discrepancies between numerical predictions of dense central cores in hierarchical clustering (e.g.\\ Navarro et al.\\ 1996,1997, Fukushige \\& Makino 1997, Moore et al.\\ 1998\\&1999, Ghinga et al.\\ 1999, Jing \\& Suto 2000) and observations of much lower central densities (e.g.\\ Flores \\& Primack 1994, Moore 1994, Burkert 1995, Navarro et al.\\ 1996, Stil 1999).  Regardless of the origin of the ``universal'' density profile, simulations routinely predict a denser, more concentrated halo than is actually observed; in other words, the predicted central phase space density for CDM is too high, which in turn suggests that the CDM initial conditions themselves have too high a phase density.  Instead, if the primordial phase space density is lower than the CDM case, then even if the final density structure is set by merging and relaxation, the final phase space density should be lowered as well, provided the primordial $Q_0$ is not much higher than that which occurs as a result of virialization at the earliest nonlinear collapse in CDM.  Numerical simulations by Huss et al.\\ (1999) approximately explore this conjecture through studying monolithic collapse of spherical overdensities with varying velocity dispersion.  However, numerical relaxation is clearly a problem for these simulations, and we draw no conclusion from them at this time.  Future numerical work will certainly shed light on this hypothesis."
    },
    "0004/astro-ph0004348_arXiv.txt": {
        "abstract": "In this article we investigate the morphology and stellar populations of high-redshift galaxies through multi-waveband HST imaging and ground-based spatially-resolved spectroscopy.  We study the redshift evolution of galaxy morphology in the Hubble Deep Field, using the deep IDT-NICMOS near-infrared HST imaging coupled with spectroscopic and photometric redshifts. Using the multi-waveband data to compare the appearance of galaxies at the same rest-frame wavelengths reveals that morphological $k$-corrections (the change in appearance when viewing high-$z$ objects at shorter rest-frame wavelengths) are only important in a minority of cases, and that galaxies were intrinsically more peculiar at high redshift. One example of significant morphological $k$-corrections is spiral galaxies, which often show more pronounced barred structure in the near-infrared than in the optical. Therefore, the apparent decline in the fraction of barred spirals at faint magnitudes in the optical HDF may be due to band-shifting effects at the higher redshifts, rather than intrinsic evolution.  Using such features as the age-sensitive Balmer$+$4000\\,\\AA\\ break, the spatially-resolved colours of distant galaxies in optical/near-infrared imaging can also be used to study their component stellar populations. We supplement this with deep Keck/LRIS spectroscopy of two extended sources: a chain galaxy at $z=2.8$ (HDF\\,4-555.1, the ``Hot Dog'' -- the brightest $U$-drop Lyman-break galaxy in the HDF) and a pair of $z=4.04$ gravitationally lensed arcs behind the cluster Abell~2390. The absence of measurable rotation across the $z=2.8$ chain galaxy implies that it is unikely to be a disk viewed edge on. With the resolution enhancement from lensing, we detect stellar populations of different ages in the $z=4$ arcs. The Ly-$\\alpha$ emission powered by the H{\\scriptsize II} regions is spatially offset from the star-forming knots in these arcs, possibly as a result of resonant scattering by neutral hydrogen. ",
        "introduction": "Until recently, the study of the most distant galaxies was very restricted: unresolved ground-based imaging gave measurements of the global colours; noisy spectra revealed AGN- or starburst-driven emission lines; and WFPC\\,2 imaging with the Hubble Space Telescope (HST) showed the shape of high-redshift galaxies in their unfamiliar rest-frame ultraviolet light. However, our knowledge has radically increased through the light-gathering power of the Keck telescopes, and the availability of the near-infrared NICMOS camera on HST. Deep Keck spectra showing the rest-ultraviolet continuum and stellar- and ISM-absorption features in $z\\sim 3$ galaxies (\\eg, Steidel \\etal\\ 1996b; Pettini \\etal\\ 2000) has opened the door to spectroscopic study of stellar populations in the distant Universe. A complementary approach to exploring the composition of high-redshift objects comes from multi-waveband HST imaging from the ultraviolet to the near-infrared, which provides the resolution necessary to study their {\\em spatially-resolved} stellar populations.  Morphology offers a window on the evolutionary status of galaxies. However, because of band-shifting effects, the interpretation of the shape of galaxies hinges on knowledge of their redshifts and spectral energy distributions (SEDs). In this article we explore this through high resolution optical/near-infrared imaging and deep spectroscopy of distant galaxies.  To resolve the issue of whether the peculiar galaxies which dominate the number counts at faint magnitudes are the counterparts of local irregulars or whether their morphological peculiarity is due mainly to band-shifting effects at high redshift, an unbiased study of the rest-frame optical morphological properties is demanded, where the effects of dust and recent star formation are less dominant than in the rest ultraviolet. In Section~\\ref{sec:HDFmorphology} we address this by analysing the IDT-NICMOS images of the northern Hubble Deep Field (HDF). In Section~\\ref{sec:hotdog} we present a case study of the brightest $z\\sim 3$ galaxy in the HDF. This example, HDF\\,4-555.1 (known as the ``Hot Dog''), is a highly elongated system and is sufficiently extended to allow resolved ground-based spectroscopy, which we have obtained with Keck/LRIS (Oke \\etal\\ 1995). We use a similar approach to explore the stellar populations in a pair of $z\\approx 4$ gravitationally-lensed arcs behind the cluster Abell~2390, which is described in Section~\\ref{sec:abell2390}. Throughout, we assume a cosmology where $h_{50}=H_{0}~/~50~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$, $q_{0}=0.5$ and $\\Lambda=0$, unless otherwise stated. {All magnitudes in this paper are with respect to the AB system (Oke \\& Gunn 1983) where $m_{AB} = -48.57 - 2.5\\log_{10}~f_{\\nu}/{\\rm (erg~cm^{-2}~s^{-1}~Hz^{-1})}$.} ",
        "conclusions": "In this article, we have explored spatially-resolved stellar populations at high-redshift, and addressed the impact on studies of galaxy morphology.  The deep, high-resolution IDT-NICMOS near-infrared imaging of a portion of the northern Hubble Deep Field has been combined with the WFPC\\,2 data and photometric redshift estimates to study the redshift evolution of morphology, comparing galaxy appearance at the same rest-wavelengths.  Some Hubble tuning-fork galaxies only reveal their true morphology in the near-infrared images. This is particularly so for galaxies with a large dispersion in stellar ages and spatially-distinct stellar populations, such as spiral galaxies which sometimes exhibit galactic bars in the NICMOS images which are invisible at shorter wavelengths.  However, galaxies which do undergo a morphological metamorphosis from the WFPC\\,2 to NIC\\,3 images are in the minority; most galaxies retain the same appearance in all wavebands, or are too compact for the structural parameters to be determined. Once the morphological $k$-corrections have been accounted for, it appears that the fraction of galaxies falling outside the Hubble sequence does increase at faint magnitudes/high-$z$.  Many of these ``true peculiars'' show evidence of being dynamically disturbed (possibly through mergers) with recent star formation activity. From the HST imaging and resolved spectroscopy with Keck/LRIS, we have shown that a $z=2.8$ chain galaxy in the HDF has a predominantly young stellar population and no significant rotation, and is thus unlikely to be an edge-on disk galaxy. Using gravitational amplification to increase our resolution, we have also resolved the stellar populations on sub-kpc scales in a system of $z=4.04$ lensed arcs.  The analysis of galaxy morphologies and colours in multi-waveband imaging, coupled with resolved spectroscopy, provides a valuable probe into the stellar populations and evolution of galaxies.  A natural progression is to use the integral field unit spectrographs currently being developed. Spectroscopy of spatially-resolved stellar populations in the high-redshift Universe will be a major scientific goal for NGST and the next generation of large ground-based telescopes with adaptive optics."
    },
    "0004/astro-ph0004104_arXiv.txt": {
        "abstract": "We report on the X-ray observation of the radio selected supernova SN1979C carried out with ASCA in December 1997. The supernova of type II$_{L}$ was first observed in the optical and occurred in the weakly barred, almost face on spiral galaxy NGC 4321 (M100) which is at a distance of 17.1 Mpc, and contains at least three other supernovae discovered in this century. No point source was detected at the radio position of SN1979C in a 3' diameter half power response circle in a 27.3 ks SIS exposure. The background and galaxy subtracted SN signal had a 3$\\sigma$ upper limit to the count rate of 1.2$\\times 10^{-3}$ cps in the full ASCA SIS band (0.4-10.0 keV).  These measurements give the first ever x-ray flux limit of a Type II$_{L}$ SN in the higher energy band ($\\geq$ 2 keV) which is an important diagnostic of the {\\it outgoing} shock wave ploughing through the circumstellar medium. ",
        "introduction": "The observation of X-rays provides the most direct view of the region of interaction between the ejecta from a core-collapse supernova (SN) and the circumstellar medium of the progenitor star (see e.g. Chevalier \\& Fransson 1994).  This circumstellar interaction gives rise to copious X-ray emission within days to months of the supernova explosion, in contrast to X-ray emission that occurs in the later supernova remnant (SNR) during the free expansion, adiabatic or radiative phases.  The circumstellar matter is heated by the outgoing SN shock wave while the SN ejecta are heated by a reverse shock wave. At the same time, this interaction is expected to produce nonthermal radio radiation; in recent years several young SNe have been found to be emitting in the radio bands (see Weiler et al. 1998). Relativistic electrons and magnetic fields in the region, which give rise to the radio emission, may be built up by acceleration in the shock wave and field amplification through Rayleigh-Taylor instabilities. The detection of strong radio emission from a SN provides good empirical evidence of circumstellar interaction of ejecta, and such a SN is a probable emitter at X-ray wavelengths as well. The temperature behind the outgoing shock is typically 10 keV and therefore the bulk of this radiation detected by ASCA would appear in its higher energy bands.  In the circumstellar interaction model, the shocked supernova ejecta is the dominant source of X-rays because of its higher density and lower temperature. ",
        "conclusions": "Until now almost all SNe detected in the X-ray were measured in the lower energy band (0.1 - 2.0 keV) of ROSAT.  Our reported measurement of SN1979C here gives the flux limit in the higher energy band ($>$ 2 keV) for the first time for a Type II$_{L}$ SN (the only other X-ray detected Type II$_{L}$ SN is SN1980K).  Because the outgoing shock wave expanding through circumstellar material generally creates a higher temperature emission spectrum compared to the reverse shock wave, the measurement of SN X-ray fluxes in the higher energy bands constitutes as important goal.    Both the ROSAT detection of the x-ray emission as well as our ASCA upper limit referred to above are well below its expected value if we scale its X-ray emission with respect to another type II-L SN active in the radio and X-ray, e.g. SN1980K. This lower than expected $L_x / L_R$ may mean that the X-ray and radio emission properties of SNe could be quite variable from one environment to another, i.e. the emission in the two bands are not directly scalable in a universal way or that the X-ray emission could be more strongly time variable than the radio emission. The shock may be interacting with local cloudlets to give enhanced thermal X-ray emission at certain epochs whereas the radio radiation may be coming from a more global arrangement of amplified magnetic fields and synchrotron radiating accelerated electrons with long lifetimes.  One of us (A.R.) was a Senior Research Associate of the National Research Council at NASA/Goddard Space Flight Center during the course of this work and acknowledges support from the NRC and ASCA Guest Observer Program. He also thanks the Aspen Center for Physics for hospitality where this manuscript was completed."
    },
    "0004/astro-ph0004332_arXiv.txt": {
        "abstract": "Continuing with previous works, we present a cosmological model in which dark matter and dark energy are modeled by scalar fields $\\Phi $ and $\\Psi $% , respectively, endowed with the scalar potentials $V(\\Phi )=V_{o}\\left[ \\cosh {(\\lambda \\,\\sqrt{\\kappa _{o}}\\Phi )}-1\\right] $ and $\\tilde{V}(\\Psi )=% \\tilde{V_{o}}\\left[ \\sinh {(\\alpha \\,\\sqrt{\\kappa _{o}}\\Psi )}\\right] ^{\\beta }$. This model contains 95\\% of scalar field. We obtain that the scalar dark matter mass is $m_{\\Phi }\\sim 10^{-26}eV.$\\ The solution obtained allows us to recover the success of the standard CDM.  The implications on the formation of structure are reviewed. We obtain that the minimal cutoff radio for this model is $r_{c}\\sim 1.2 \\, kpc.$ ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004042_arXiv.txt": {
        "abstract": "I present a somewhat selective review of microlensing theory, covering five major areas: 1) the derivation of the basic formulae, 2) the relation between the observables and the fundamental physical parameters, 3) binaries, 4) astrometric microlensing, and 5) femtolensing. ",
        "introduction": "All of gravitational microlensing is reducible to a single equation, the Einstein formula for $\\alpha$, the deflection of light from a distant source passing by a lens of mass $M$ at an impact parameter $b$, \\begin{equation} \\alpha = {4 G M\\over b c^2}. \\end{equation} This equation has been verified experimentally by Hipparcos to an accuracy of 0.3\\% (Froeschle, Mignard, \\& Arenou 1997). Despite its apparent simplicity, equation (1) generates an incredibly rich phenomenology.  The aim of this review is to present the reader with a concise introduction to microlensing phenomena from the standpoint of theory.  It is impossible to cover all aspects of microlensing in the space permitted.  Several good reviews of microlensing can be consulted to obtain a deeper appreciation for various aspects of the subject (Paczy\\'nski 1996; Gould 1996; Roulet \\& Mollerach 1997; Mao 1999b).  By definition, microlensing is gravitational lensing where the images are too close to be separately resolved.  The main microlensing effect that has been discussed in the literature (and the only one that has actually been observed) is {\\it photometric microlensing}: the magnification of the source due to the convex nature of the lens (Einstein 1936; Refsdal 1964; Paczy\\'nski 1986; Griest 1991; Nemiroff 1991). However, there are two other effects that deserve attention from the standpoint of theory: {\\it astrometric microlensing}, the motion of the centroid of the images relative to the source, and {\\it femtolensing}, which refers to interference effects in microlensing. ",
        "conclusions": "When microlensing experiments began in the early 1990s, few participants expected that there was much room for theoretical development at all.  The physical effect (eq.\\ 1) was completely understood, and the basic equations of microlensing had all been worked out.  A decade later, microlensing theory has shown itself to be a very dynamic field. The original problems turned out to be much richer than expected, while new observations and new developments in instrumention have raised new problems.  Thus, microlensing theory promises to remain vibrant.  \\noindent{\\bf Acknowledgements}: This work was supported by grant AST 97-27520 from the NSF."
    },
    "0004/astro-ph0004274_arXiv.txt": {
        "abstract": "We consider possible geometries of magnetic fields in GRB outflows, and their evolution with distance from the source. For magnetically driven outflows, with an assumed ratio of magnetic to kinetic energy density of order unity, the field strengths are sufficient for efficient production of $\\gamma$-rays by synchrotron emission in the standard internal shock scenario, without the need for local generation of small scale fields. In these conditions, the MHD approximation is valid to large distances ($\\ga 10^{19}$\\,cm). In outflows driven by nonaxisymmetric magnetic fields, changes of direction of the field cause dissipation of magnetic energy by reconnection. This dissipation takes place outside the photosphere of the outflow, and can convert a significant fraction of the magnetic energy flux into radiation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/hep-ph0004230_arXiv.txt": {
        "abstract": "Systems that exhibit pattern formation are typically driven and dissipative. In the early universe, parametric resonance can drive explosive particle production called preheating. The fields that are populated then decay quantum mechanically if their particles are unstable. Thus, during preheating, a driven-dissipative system exists. We have shown previously that pattern formation can occur in two dimensions in a self-coupled inflaton system undergoing parametric resonance. In this paper, we provide evidence of pattern formation for more realistic initial conditions in both two and three dimensions. In the one-field case, we have the novel interpretation that these patterns can be thought of as a network of domain walls. We also show that the patterns are spatio-temporal, leading to a distinctive, but probably low-amplitude peak in the gravitational wave spectrum. In the context of a two-field model, we discuss putting power from resonance into patterns on cosmological scales, in particular to explain the observed excess power at $100 h^{-1}$Mpc, but why this seems unlikely in the absence of a period of post-preheating inflation. Finally we note our model is similar to that of the decay of DCCs and therefore pattern formation may also occur at RHIC and LHC.\\\\  Fermilab Preprint: Pub-98/231-A \\qquad \\qquad \\, BROWN-HET-1201 \\qquad \\qquad \\, Imperial/TP/99-0/006\\\\ ",
        "introduction": "\\addtocounter{footnote}{1} \\footnotetext{\\tt{ats@camelot.mssm.edu}} \\addtocounter{footnote}{1} \\footnotetext{\\tt{mparry@ic.ac.uk}}  Much recent work has been devoted to the topic of preheating in inflationary cosmology. Preheating is a stage of explosive particle production which results from the resonant driving of particle modes by an inflaton oscillating in its potential at the end of inflation \\cite{trasbran,kls94,stb95}.  In regions of parameter space where parametric resonance is effective, much of the energy of the inflaton is transferred to bands of resonant wave modes. This energy transfer is non-thermal and can lead to interesting non-equilibrium behavior. Two examples of the non-equilibrium effects that can be produced are non-thermal phase transitions \\cite{kofmlind,tkac,rt96,khlekofm1} and baryogenesis \\cite{andelind,kolblind,kolbriot,ggks99}. The non-thermal phase transitions induced during preheating can sometimes lead to topological defect formation \\cite{kasukawa,tkackhle,parrsorn,drss99}, even at energies above the eventual final thermal temperature. Furthermore, non-linear evolution of the field when quantum decay of the resonantly produced particles is negligible leads to a chaotic power-law spectrum of density fluctuations \\cite{khletkac,pr97}.  In a previous letter \\cite{sornparr}, we presented evidence for a new phenomenon that can arise from preheating: pattern formation. It has long been known that many condensed matter systems exhibit pattern formation\\footnote{For an extensive review of pattern formation in condensed matter systems, see \\cite{croshohe,newepass}.}. Examples of pattern forming systems which have been studied are ripples on sand dunes, cloud streets and a variety of other convective systems, chemical reaction-diffusion systems, stellar atmospheres and vibrated granular materials. All of these physical systems have two features in common. They are all driven in some manner, i.e. energy is input to the system, and they are all dissipative, usually being governed by diffusive equations of motion. Typically, patterns are formed in these systems in the weakly non-linear regime before the energy introduced into the system overwhelms the dissipative mechanism. Sometimes, patterns persist beyond the weakly non-linear regime as well.  At the end of inflation, the inflaton $\\phi$ is homogeneous and, in most commonly studied models, oscillating about the minimum of its potential. This oscillation gives an effectively time-dependent mass to fields which are coupled to the inflaton, including fluctuations of the inflaton itself. The time-dependent mass drives exponential growth in particle number in certain bands of wave modes. However the fields into which the inflaton can decay resonantly are also unstable to quantum decay. For these reasons, at the end of inflation, we are considering fields which are driven, due to resonant particle creation, and also dissipative, due to quantum decay. In \\cite{sornparr}, we were able to show that pattern formation occurs in a chaotic inflationary model with a self-coupled inflaton in the weakly non-linear regime.  In this paper, we primarily consider a $\\lambda \\phi^4$ theory with the addition of a phenomenological decay term to mimic the inflaton's quantum decay. This model without the decay term has been studied extensively in the literature\\cite{kls94,stb95,khletkac,bvhs96,pr97,k97,greekofm}, and a similar model including the decay term has also been studied \\cite{kolbriot}.  In \\cite{sornparr}, we used restricted initial conditions. We only seeded the resonant band with small fluctuations of order $\\sim 10^{-3}$ of $\\langle \\phi \\rangle$, then simulated the field's evolution. We found that the resonant modes interacted and formed patterns. Here we use initial conditions appropriate to the vacuum at the end of chaotic inflation. We also extend our study to a 3-dimensional volume. In \\cite{sornparr}, we did not point out the spatio-temporal nature of the patterns, which was not evident at the time due to an unfortunate coincidence in the form of the pattern at the timesteps at which we viewed the data. Here we note the spatio-temporal behavior. We also discuss resonance giving rise to patterns on cosmological scales.  The paper is ordered as follows: In section \\ref{equations}, we present the model we are investigating. In section \\ref{ic}, we discuss the initial conditions appropriate for the end of inflation. In sections \\ref{results}, we present the results of our simulations in two and three spatial dimensions. In section \\ref{diss}, we extend our analysis to the two-field case and discuss the possible implications of our results. This is followed by the conclusions. ",
        "conclusions": "In this paper, we have presented new evidence that there is a pattern forming regime in a $\\lambda\\phi^4$ theory during preheating at the end of inflation. We have used vacuum initial conditions appropriate to sub-Hubble modes at the end of inflation. We have shown that patterns arise in both two- and three-dimensions, and that, in the one-field case, they may be thought of as a network of domain walls. Furthermore since the patterns vary spatio-temporally, gravitational waves will be produced. However, relative to the gravitational waves produced in an undamped model, their amplitude will be small, and therefore, unless the directionality of the pattern can aid in detection, extremely difficult to detect.  We have speculated on the possibility, in the context of a two-field model, of putting a resonant band at cosmological scales, in particular at $100 h^{-1}$Mpc. We conclude that this is not feasible because it is impossible to amplify only the required modes. However this conclusion may change in a model which allows a subsequent period of inflation.  Finally, we point out that heavy ion collisions at RHIC and LHC may be another physical system which exhibits pattern forming behavior.  \\vskip0.5cm  {}\\noindent{\\bf Acknowledgements}  It is a pleasure to thank Ewan Stewart, Rocky Kolb, Robert Brandenberger, Bruce Bassett and Dan Boyanovsky for a number of useful discussions. We would also like to thank Bob Rosner and the Astrophysical Thermonuclear Flash Group at the University of Chicago for the use of their computational facilities. Much of this work was completed at Fermilab, supported by the DOE and the NASA grant NAG 5-7092, and at Brown under DOE contract DE-FG0291ER40688, Task A. A portion of the computational work in support of this research was performed at the Theoretical Physics Computing Facility at Brown University."
    },
    "0004/astro-ph0004019_arXiv.txt": {
        "abstract": "An excellent candidate for a young elliptical, or `protoelliptical' galaxy is NGC~1700. Here we present new $B$, $V$ and $I$ band imaging using the Keck telescope and reanalyse existing $V$ and $I$ band images from the \\emph{Hubble Space Telescope}. After subtracting a model of the galaxy from the Keck images NGC~1700 reveals two symmetric tidal tail-like structures. If this interpretation is correct, it suggests a past merger event involving two spiral galaxies.  These tails are largely responsible for the `boxiness' of the galaxy isophotes observed at a radius of $\\sim13$ kpc.  We also show that the $B-I$ colour distribution of the globular cluster system is bimodal.  The mean colour of the blue population is consistent with those of old Galactic globular clusters.  Relative to this old, metal poor population, we find that the red population is younger and more metal rich.  This young population has a similar age and metallicity as that inferred for the central stars, suggesting that they are both associated with an episode of star formation triggered by the merger that may have formed the galaxy.  Although possessing large errors, we find that the majority of the age estimates of NGC~1700 are reasonably consistent and we adopt a `best estimate' for the age of $3.0\\pm1.0$ Gyr. This relatively young age places NGC~1700 within the age range where there is a notable lack of obvious candidates for protoellipticals.  The total globular cluster specific frequency is rather low for a typical elliptical, even after taking into account fading of the galaxy over the next 10 Gyr.  We speculate that NGC~1700 will eventually form a relatively `globular cluster poor' elliptical galaxy. ",
        "introduction": "Merging is now thought to be a key process in the evolution of galaxies. The hypothesis that two colliding spiral galaxies will eventually form an elliptical galaxy \\cite{toomre72} has gained much observational and theoretical support over the years. The `smoking gun' of this type of merger is the presence of two tidal tails, formed from the progenitor's discs. A number of well known examples are found in the local Universe, such as {\\it The Mice}, {\\it The Antennae}, Arp~220, NGC~3256 and NGC~7252. For such galaxies a variety of methods are available (e.g. spectroscopy of the stellar populations, dynamical measurements, comparison with models etc.) to estimate the time since the merger occurred. The derived ages, for these classic tidal tail systems, is up to 1--2 Gyr since nuclear coalescence. It was Ivan King in 1977 who first pointed out the general lack of obvious candidates for older merger remnants (i.e. 2--5 Gyr old). These galaxies have been referred to as `King gap objects' or protoellipticals. Identifying, and age dating, these protoellipticals could provide the missing link between late stage spiral mergers and elliptical galaxies.  Indeed, a crucial step in testing the merger hypothesis would be to show an `evolutionary consistency', i.e. that spiral mergers and protoellipticals evolve to have the same energetic, structural, dynamical and chemical properties as normal, old ellipticals. This has been problematic due to the difficulty of estimating the age of old stellar populations, without telltale morphological signatures such as tidal tails.  Recently, several different methods have become available. These include breaking the age--metallicity degeneracy with spectral synthesis (e.g. \\pcite{gonzalez92}) and new models (e.g. \\pcite{wortheymod}), quantifying optical fine structure \\cite{ss92} and using the colours of globular clusters (e.g. \\pcite{whitmore97}).  An ideal candidate for a nearby protoelliptical is NGC~1700. It possesses a kinematically distinct core (Franx, Illingworth \\& Heckman 1998a), reveals evidence for extensive morphological disturbance \\cite{ft92} and it has a high rotational velocity to velocity dispersion ratio, relative to other ellipticals (Bender, Burstein \\& Faber 1992).  As well as being a possible protoelliptical, it offers the opportunity to compare various age estimates. For example, age estimates can be made from its two faint tidal tails, optical fine structure, globular cluster colours, kinematic structure, and its stellar component.  We have adopted the same distance to NGC~1700 as \\scite{whitmore97}, i.e.  51.4 Mpc (which includes a correction for Virgocentric infall and assumes $H_{\\circ}=75$ km~s$^{-1}$~Mpc$^{-1}$). This corresponds to 249 pc per arcsec. The total $B$ band magnitude for NGC~1700 is $M_{\\mathrm{B}}=-21.56$ (RC3:\\pcite{rc3}). It is classified as an E4 elliptical in the RC3 and E3 in the RSA. Photometric studies have been carried out by Franx, Illingworth \\& Heckman (1998b) and \\scite{g94}. A recent kinematic study is that of Bender, Saglia \\& Gerhard (1994).  In this paper we reanalyse {\\it Hubble Space Telescope} (\\emph{HST}) WFPC2 $V$, $I$ images and present new Keck $B$, $V$ and $I$ images. We focus on the morphology of the galaxy and photometry of the globular cluster system, both of which provide new age estimates. We compare our age estimates with a variety of alternative methods and discuss the implications of our results for the formation and evolution of elliptical galaxies. ",
        "conclusions": "We have presented results from new $B$, $V$ and $I$ band imaging of NGC~1700, taken using the Keck telescope and reanalysed previous \\emph{HST} imaging. From the morphology of the galaxy and photometry of its globular cluster system, we have derived new estimates of the galaxy age.  Subtraction of an elliptical model from the Keck images revealed the presence of two symmetric tidal tail-like structures extending $\\sim40$ kpc to the North-West and South-East of the galaxy.  The presence of tidal tails is thought to be a classic signature that a galaxy has undergone a merger event, involving the collision of two spiral galaxies during the last few Gyr. If the observed tidal features are indeed genuine tidal tails, then this would suggest that NGC~1700 has undergone a major merger event during its recent history. If they are merely plumes or other fine structure, then the situation of a disc galaxy merging with an elliptical is also possible.  Based on the fraction of galaxy light contained within these tails we have estimated their age at $3.2\\pm1.5$ Gyr.  We have also shown that NGC~1700 possesses boxy isophotes for radii of $7.5<r<20.0$ kpc, and that this boxiness is largely caused by the presence of the tail structures.  From the \\emph{HST} imaging, we detected 146 globular clusters.  Of these, 34 were in common with our GC candidates detected in the Keck data. We then used these common GCs to refine our Keck GC selection. After further restricting our sample in magnitude and colour we obtained a final Keck object list of 312 GCs. These show a bimodal colour distribution with peaks at $B-I=1.54$ and 1.98.  The colour of the blue population is consistent with that of Galactic GCs.  Assuming that the blue population is indeed an old metal poor system and measuring the offset in magnitude and colour between the two populations, we find that the red GCs are younger and more metal rich. The \\scite{bc96} stellar population models suggest that the red population is 2.5--5.0 Gyr old and has super-solar metallicity. The equivalent models of \\scite{wortheymod} predict a significantly older age (5.0--8.0 Gyr) and a lower metallicity ($[{\\mathit{Fe}}/H]\\sim-0.2$). The bimodality in $B-I$ is supported by a hint of bimodality seen in the $V-I$ colours from \\emph{HST} data.  We discuss other age estimates from the literature based on fine structure, globular cluster colours, Fundamental Plane residuals, stellar dynamics, and spectroscopic line strengths. We find that, although possessing significant errors, the various age estimates generally indicate a young age for NGC~1700 of about $3.0\\pm1.0$ Gyr. The fact that the `dynamical/structural' and `star formation' ages are similar suggests that NGC~1700 has undergone an episode of enhanced star formation triggered by a merger event, which may have created the galaxy itself.  The surface density profile of GCs reveals a flatter (i.e.  more extended) profile than the underlying galaxy starlight. The total number of GCs and present total specific frequency are estimated to be $N_{\\mathrm{T}}=1320\\pm270$ and $S_N({\\mathrm{total}})=1.4^{+1.0}_{-0.6}$ respectively. The value for $S_N({\\mathrm{total}})$, even considering the large errors, is quite low compared to the majority of typical ellipticals and is more consistent with spiral galaxies. We predict that after 10 Gyr the total $S_N$ will have increased to a maximum of $\\sim2.8$ if NGC~1700 has undergone a starburst that constitutes 50 per cent of its stellar mass. This is consistent with low $S_N$ ellipticals. If the starburst constitutes only 10 per cent of the total stellar mass, the predicted $S_N$ is somewhat lower at $S_N\\sim1.6$. This suggests that NGC~1700 will form a relatively `globular cluster poor' elliptical galaxy once it reaches a comparable age to typical `old' ellipticals."
    },
    "0004/astro-ph0004369_arXiv.txt": {
        "abstract": "We present an observational and theoretical study of the spiral structure in galaxy NGC 1566. A digitized image of NGC 1566 in $I$-band, obtained with the Australian National University 30in telescope, was used for measurements of the radial dependence of amplitude variations in the spiral arms. The azimuthal variations of the surface brightness in the $I$-band are about $\\pm 7$ \\% at $50^{\\prime \\prime}$ radius and increase up to $57$ \\% at $100^{\\prime \\prime}$.  We use the known velocity dispersion in the disk of NGC 1566, together with its rotation curve, to construct linear and 2D nonlinear simulations which are then compared with observations. The linear stability analysis of the disk of NGC 1566, made under the assumption that the ratio of the vertical and the radial velocity dispersions is equal to 0.6, shows that the disk is stable with respect to spiral perturbations. We confirm this conclusion with 2D hydrodynamical simulations. The disk of NGC 1566 is unstable if the ratio of the vertical and the radial velocity dispersions $c_z/c_r$ is of order $0.8 - 1.0$. Under this assumption, a two-armed spiral constitutes the most unstable global mode in the disk of NGC 1566.  The 2D hydrodynamic simulations of the unstable disk seeded with random perturbations show the exponential growth of two unstable modes, $m=2$, and $m=3$. The growth rates of the global modes seen in the nonlinear simulations are in good agreement with the results of the linear modal analysis. The $m=3$ global mode is less important, however, in comparison with the main $m=2$ global mode, and the overall evolution of the perturbations is determined by the two-armed spiral mode, which saturates at a level $\\log_{10} A_2 \\approx -0.5$. The theoretical surface amplitude and the velocity residual variations across the spiral arms calculated during the nonlinear phase of instability are in qualitative agreement with the observations.  The spiral arms found in the linear and nonlinear simulations are considerably shorter than those observed in the disk of NGC 1566. The nonlinear phase of instability is characterized by the transport of angular momentum towards the disk center, making the surface density distribution quite steep in comparison to the observations. We argue therefore, that the surface density distribution in the disk of the galaxy NGC 1566 was different in the past, when spiral structure in NGC 1566 was linearly growing. ",
        "introduction": "There is a general consensus that the regular spiral arms observed in disk galaxies are a manifestation of the global density waves propagating in galactic disks (Binney \\& Tremaine 1987). In the approach advocated for many years by C. C. Lin and his collaborators (e.g. Bertin et al. 1989), such global waves arise from the gravitational instability of exponentially growing internal modes which are stabilized at some ``quasi-steady'' level due to nonlinear effects. Another explanation assumes that the global spiral arms arise from swing-amplified chaotic fluctuations originating from N-particle noise (James \\& Sellwood 1978), or from the gravitational drag of the co-orbiting mass clumps within the disk (Toomre 1990). A long-standing controversy over the physical mechanisms responsible for generation of the global density waves has recently been resolved by Shu et al (1999). Those authors showed that global structures arising from the internal normal modes of a gravitating galactic disk, and the spiral patterns driven by external tidal influences, are both manifestations of the fundamental underlying process of modal selection and growth.  The discovery of the protoplanetary disks around T Tauri stars initiated a new wave of interest in global instabilities in self-gravitating disks. A number of authors have developed various techniques to study linear density perturbations in gaseous and the collisionless self-gravitating disks (Adams, Ruden, \\& Shu 1989; Noh, Vishniac, \\& Cochran 1991; Laughlin \\& R\\'o\\.zyczka 1996; Vauterin \\& Dejonghe 1996; Kikuchi, Korchagin, \\& Miyama 1997). Linear stability analysis methods for self-gravitating disks are well complemented by direct two-dimensional simulations (e.g., Tomley et al 1994; Miyama et al. 1994; Nelson et al. 1998). Zhang (1996), and independently Laughlin, Korchagin, \\& Adams (1997, 1998) have examined nonlinear self-interaction of a single strong global mode. They found that nonlinear self-interaction is responsible for the nonlinear saturation of the mode, and concluded that the saturation mechanism provides a key for understanding the long-term behavior of spiral density waves.  An ultimate aim of any theory of spiral structure is the construction of models of particular galaxies on the basis of the available observations. A number of authors have undertaken such comparisons (e.g., Lin, Yuan, \\& Shu 1969; Roberts, Roberts, \\& Shu 1975; Mishurov et al 1979; Elmegreen \\& Elmegreen 1990). The procedure in these studies has relied on empirical estimates of the positions of the corotation and Lindblad resonances using ``optical tracers'' in the galactic images. Local dispersion relations were then used to calculate the spiral response under additional assumptions regarding the radial behavior of the stability parameter $Q$ derived by Toomre (1964). These two foregoing assumptions, and especially the stipulated choice of Toomre's $Q$-parameter, considerably reduce the predictive power of the spiral density wave theory.  Long-slit spectroscopic observations of spiral galaxies can provide a stronger basis for the comparison of density wave theory with observations. Knowledge of the $z$-velocity dispersion in a galactic disk provides a quantitative estimate of the surface density distribution under the realistic assumption that disks are self-gravitating in the $z$-direction (Bottema 1992, hereafter RB). The velocity dispersion, together with the known rotation curve and the surface density distribution, uniquely determine the axisymmetric background state of the galactic disk, and can be used for linear and nonlinear analyses aimed at the modeling of individual spiral galaxies.  The aim of our paper is to model the spiral structure in the nearby grand design spiral galaxy NGC 1566. The galaxy is apparently isolated, and it has been the subject of many photometric and spectroscopic studies. In particular, the rotation curve and the radial profile of the stellar velocity dispersion have been accurately measured (RB).  NGC 1566 represents one of the best candidates for the comparison of spiral density wave theory with observations.  In contrast to previous comparisons between theory and observations, our study models the spiral arms in the disk of NGC 1566 from ``scratch'', without making any additional assumptions with regards to the positions of the principal resonances in its disk or the behavior of Q-parameter. Basing our work on both our own, as well as previously published observations, we undertake the linear global modal analysis of this system, and extend the analysis of the spiral structure into the nonlinear regime using 2D numerical simulations. We then compare the results of our linear, and nonlinear simulations with the observed properties of the spiral arms in NGC 1566. The purpose of this comparison is to determine the extent to which the internal global modes emerging from random perturbations imposed on the gravitating disk of NGC 1566 can be responsible for the observed properties of its spiral arms.  In Section 2 we present the optical maps in $B$- and $I$-bands for NGC 1566, along with the results of our measurements of the amplitude variations in the spiral arms. In Section 3, we use the available observational data to determine the background equilibrium properties in the disk of NGC 1566. In sections 4 and 5 we describe the basic equations and our computational methods. Section 6 gives the results of the linear global modal analysis. Section 7 describes the results of the nonlinear 2D simulations and discusses our comparison of the theoretical results with observations. ",
        "conclusions": "This paper has focused on the observational study of the spiral arms in galaxy NGC 1566, and on the theoretical study of the stability properties of this galaxy. We have applied a global modal approach and have used two-dimensional one-component and multi-component simulations to study the dynamics of the self-gravitating disk in the galaxy NGC 1566 and we have extended our analysis to the nonlinear stage using 2D numerical simulations. In our theoretical analysis of the spiral structure in NGC 1566 we have not made any additional assumptions used in previous comparisons, and we have followed the development of the spiral arms emerging from random perturbations all the way through to the nonlinear saturation stage. The general conclusions which can be made from our work are as follows: As C.C. Lin and his collaborators envisioned many years ago, the most unstable linear global mode emerging from stochastic noise in a galactic disk determines the appearance of the spiral structure on linear, and on nonlinear stages of the evolution of the spiral pattern. We have found a good agreement between the linear global modal analysis of the stability of the disk of galaxy NGC 1566, and direct 2D numerical simulations. The theoretical spiral pattern obtained from the linear global modal analysis, and the pattern emerging from the random perturbations in the nonlinear simulations are in agreement, at least qualitatively, with the observed two-armed spiral structure in the disk of NGC 1566.  We did not aim to discriminate between the two principal physical mechanisms explaining the physics of spiral instabilities in gravitating disks. We concur however with the conclusion of Shu et al (1999) that ``...there is a coexistence of the two principal mechanisms that produce the beautiful structures that astronomers observe in the universe of spiral galaxies''.  The specific results of our analysis are specified below.  1. The CCD images of NGC 1566 in $B$ and $I$-bands obtained with the Australian National University 30in telescope were used for measurements of the radial dependence of the amplitude variations in the spiral arms of NGC 1566. The azimuthal variations of the surface brightness in the $I$-band increase with radius up to $57$ \\% at $100^{\\prime \\prime}$. Our results concur with the previous measurements of the amplitude variations within the spiral arms of NGC 1566.  2. The linear stability analysis of the disk of NGC 1566 made under the standard assumption used in the galactic dynamics, namely that the ratio of the vertical and the radial velocity dispersions is equal to 0.6, shows that disk is stable towards spiral perturbations within observational error bars. We confirm this conclusion with help of 2D simulations of the disk evolution seeded with random perturbations.  3. By increasing the $c_z/c_r$ ratio up to $0.8 - 1.0$, the disk, when seeded with random perturbations, becomes unstable with respect to $m=2$, and $m=3$ spiral modes. The growth rates and the shapes of the global modes found in the linear analysis, and those seen in the nonlinear simulations are in good agreement.  The two-armed spiral mode prevails over its competitors, and thus determines the behavior of perturbations during the linear, and during the nonlinear phases (see Figures 16 and 17). At the nonlinear stage, the two-armed spiral saturates at an amplitude $\\log_{10} A_2 \\approx -0.5$. The surface density, and the velocity residual variations in the arms are in a qualitative agreement with observations.  4. The nonlinear phase of instability is characterized by the transport of angular momentum towards the disk center making the surface density distribution more steep than the observed surface brightness profile in the disk of NGC 1566. The spiral arms found in the linear modal analysis, and seen in the nonlinear simulations, are considerably shorter than the observed spiral arms in the disk of NGC 1566. We argue therefore, that the surface density distribution in the disk of the galaxy NGC 1566 was different in the past when spiral structure in NGC 1566 was growing linearly."
    },
    "0004/astro-ph0004196_arXiv.txt": {
        "abstract": "We study observational constraints on cosmological models with a quintessence arising from moduli fields. The scalar field potential is given by a double exponential potential $V=V_0\\exp(-Ae^{\\sqrt{2}\\kp\\ph})$. After reviewing the properties of the solutions, from a dynamical systems phase space analysis, we consider the constraints on parameter values imposed by luminosity distances from the 60 Type IA supernovae published by Perlmutter \\etal, and also from gravitational lensing statistics of distant quasars. We also update the constraints on models with a single exponential potential $V=V_0e^{-\\la\\kp\\ph}$. ",
        "introduction": "A cosmological constant $\\LA$ has been considered as the missing energy of the universe for a long time. Recently, a varying vacuum energy or ``quintessence'' \\cite{CDS} has become popular as an alternative candidate. Since more parameters are involved in such models, they are more flexible in explaining some of the problems left by the cosmological constant model. One such problems is the ``cosmic coincidence problem'' \\cite{ccp}: the missing energy and the matter energy densities decrease at different rates as the universe expands, so it seems purely coincident that they are comparable today with the concordant values from several observational tests of $\\Omp\\goesas{2\\over3}$ and $\\Om_{m0}\\goesas{1\\over3}$ \\cite{WCOS}. Several candidates for quintessence fields have been proposed. Typically quintessence models possess attractor solutions with common evolutionary properties for a wide range of initial conditions. For example, in some cases the scalar field energy maintains constant ratio to the matter density at late time \\cite{PNGB}, where as in other models it tracks the dominant matter component in some other sense \\cite{plp,ss,CLW,FJ,LS}.  In this paper, I will study quintessences arising from string moduli. In string or Kaluza-Klein type models the moduli fields associated with the geometry of the extra dimensions may have effective potentials which depend exponentially on the moduli fields, due to the curvature of the internal spaces, or alternatively through the interaction of the moduli with the form fields on the internal spaces (see, e.g., \\cite{CLW} and references therein). Single exponential potentials of the form \\be \\label{sep} V=V_0e^{-\\la\\kp\\ph} \\ee which give rise to a scaling solution have been well-studied in the literatures (see, e.g., \\cite{ss,CLW,FJ} and references therein). In this paper I will concentrate on double expontial potentials of the form \\be \\label{dep} V=V_0\\exp(-Ae^{\\sqrt{2}\\kp\\ph})\\,  \\ee which typically arise from supersymmetry breaking via gaugino condensation \\cite{gc,BCC,Bin}. Furthermore, it has also recently been argued that if the superpotential of the moduli field in string theory is T-duality invariant, then it cannot be approximated by a single exponential function, but must depend on double exponentials \\cite{delaM}.  Binetruy \\cite{Bin}, and later de la Macorra \\cite{delaM}, have argued that scalar fields with double exponential potentials cannot act as quintessences for two reasons. Firstly, such models have a global attractor solution leading to a matter-dominated universe at late times. Secondly, near the attractor the equation of state is positive, $w_\\ph>0$ and approaching zero as $t\\rarr\\infty$. This would appear to contradict the latest observational results from type IA supernovae that a quintessence is dominating over matter and the universe is entering a phase of accelerated expansion \\cite{Per,Rie}, similar to inflation. The purpose of this paper is to show that for parameter values away from the attractor, there can exist models which are consistent with the observational tests, although some tuning is required for the universe not to have reached the attractor at present. This paper also updates the observational constraints on the single exponential potential model, which have been given by Frieman and Waga \\cite{FW}.  The organization of this paper is as followed: in section II, I will present a phase-space analysis of the double exponential potential model. In section III, I will discuss the numerical integration of the evolution equations for both the single and double exponential potential models, and obtain values for $\\Omp$ and $H_0t_0$. In section IV, I will constrain both models using the light-curve calibration luminosity distances of type IA supernovae (see \\cite{Per,Dre} and references therein) and the gravitational lensing statistics of high luminosity quasars by intervening galaxies (see \\cite{Koc,WM} and references therein). Similar constraints on other quintessence models have been presented in \\cite{FW,WM,WF,NW}. ",
        "conclusions": "We have performed a phase-space analysis on the double exponential potential model, its properties can be understood in terms of a single exponential potential model with a varying coefficient $\\la$. The analysis shows that unlike the single exponential potential model which has two late times attractor depending on $\\la$, the double exponential potential model has only one global attractor which leads to a matter dominated universe at late times. However, it is always possible for the model to be dominated by the scalar field at present and dominated by matter in the future. The problem with this model is we need to fine-tune the parameters in order not to have reached the attractor at the present epoch, thereby circumventing the objections of Binetruy \\cite{Bin} and de. la. Macorra \\cite{delaM}. However, all other quintessence models appear to also require a degree of fine-tuning \\cite{PNGB,plp,ss,CLW,FJ,LS}.  We studied the observation constraints for both the single and double exponential potential models using a more updated type IA supernovae data and the gravitational lensing statistics. The results show that there are regions in the parameter spaces for which the models are consistent with the observations at the same time giving appropriate values for the scalar field energy density and the age of the universe at present."
    },
    "0004/astro-ph0004313_arXiv.txt": {
        "abstract": "We point out that the baryon-loading problem in Gamma-Ray Burst (GRB) models can be amelioriated if a significant fraction of the baryons which inertially confine the fireball are converted to neutrons.  A high neutron fraction in some circumstances can result in a reduced transfer of energy from relativistic light particles in the fireball to baryons.  The energy needed to produce the required relativistic flow in the GRB is consequently reduced, in some cases by orders of magnitude. This could be relevant to GRB models because a high neutron-to-proton ratio has been calculated in neutron star-merger fireball environments. Significant neutron excess also could occur near compact objects with high neutrino fluxes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004255_arXiv.txt": {
        "abstract": "Two  studies  have  recently  reported  the  discovery  of  pronounced  Halo substructure  in the  Sloan Digital  Sky Survey  (SDSS)  commissioning data. Here we  show that this Halo substructure  is almost in its  entirety due to the expected tidal  stream torn off the Sagittarius  dwarf galaxy during the course of its many close encounters with the Milky Way.  This interpretation makes strong  predictions on  the kinematics and  distances of  these stream stars.  Comparison of the structure in old horizontal branch stars, detected by  the SDSS  team, with  the carbon  star structure  discovered in  our own survey, indicates  that this halo stream  is of comparable age  to the Milky Way. It  would appear that  the Milky Way  and the Sagittarius  dwarf galaxy have been a  strongly interacting system for most  of their existence.  Once complete, the SDSS will provide a unique dataset with which to constrain the dynamical evolution of  the Sagittarius dwarf galaxy, it  will also strongly constrain the mass distribution of the outer Milky Way. ",
        "introduction": "In the standard hierarchical picture  of galaxy formation, galaxies like the Milky Way  grew by repeated mergers  with other galaxies,  and in particular with many  smaller galaxy ``fragments''.   In the outer regions  of galactic halos, disruption timescales can be long,  and we can hope to observe galaxy remnants as  well as  to catch  galaxies in the  process of  being accreted. That this accretion  is still an on-going process is  clear, as evidenced by the  many studies  of tidal  features in  the dwarf  satellite  galaxies and globular clusters.  However, what is of prime interest  is to compare the observed structures to the results of galaxy formation  theories. These now make strong predictions about the  structure and substructure  of galaxy halos. Indeed  the standard cold  dark  matter  models  appear  to be  inconsistent  with  observational constraints,   as   they    over-predict   halo   clumpiness   (see,   e.g., \\citealt{moore}).  In an  attempt to  constrain such theories,  several groups  have undertaken large  scale surveys  to uncover  substructure in  the halo.  These include, among others,  the ``spaghetti''  survey \\citep{morrison00}, the  APM carbon star  survey  \\citep{totten98,  totten00,  ibata00},  and  the  SDSS  survey \\citep{yanny, ivezic}.  The APM carbon star survey, an almost all high-latitude sky survey, revealed the existence  of a gigantic band  of these intermediate  age ($\\sim 6\\Gyr$) stars around  the sky (\\citealt{ibata00};  hereafter paper~1). \\footnote{Extant   optical   and  NIR   photometry   and  detailed   optical spectroscopy of  the APM  colour-selected cool Halo  carbon stars  show that these stars  are identical in  nature to those  found in AGB  populations in other Local  Group galaxies, particuarly  the Magellanic Clouds and  many of the  dwarf   spheroidal  Galactic  satellites   \\citep{totten98,  totten00}. Furthermore the central  part of Sgr has a known  AGB carbon star population \\citep{whitelock96}   that    also   has   NIR    and   optical   properties indistinguishable from  the APM  Halo sample.  Normal  AGB carbon  stars are carbon-enriched during  dredge-ups driven  by helium burning  thermal pulses during the post  main sequence phase of 1-2 solar mass  stars.  As such they are thought to be  $> 1-2$~Gyrs old but less than 10~Gyrs old (see e.g. the reviews by Da Costa 1997 and  also Azzopardi \\& Lequeux 1992).} Approximately half of the ``Halo''  carbon stars detected in this survey are located in  this ten  degree wide  band, which is  centered on  the expected orbit  of  the  Sagittarius  dwarf  galaxy in  a  spherical  Halo  potential (paper~1). Sagittarius, the closest dwarf  satellite galaxy to the Milky Way \\citep{ibata94}, is clearly in the  process of being tidally torn apart.  As shown in paper~1, those carbon stars trace a gigantic tidal stream of debris disrupted  off the  Sagittarius dwarf  galaxy which  has now  wrapped itself around the Halo.   Interestingly, this stream is observed  to follow a great circle on the sky, which means that it is probably a coplanar rosette-shaped structure  that we  observe from  within (since  it should  trace relatively closely  the  orbit of  the  center  of mass  of  the  dwarf galaxy).   Most importantly,  however, the  fact that  the stream  is confined  to  a plane, implies  that it  does  not  suffer significant  precession,  which is  only possible if the Galactic halo potential is not significantly oblate.  Two  SDSS teams  have recently  presented  results from  their first  year's commissioning  data   \\citep{yanny,  ivezic}.   Due   to  initial  technical constraints,  the drift-scanning  camera  took data  only  on the  celestial equator during the commissioning period.  Their data consists of two wedges: a NGC  sample $87^\\circ$ long  and a SGC  sample $60^\\circ$ long,  with each covering  a declination  range  of only  $-1.25^\\circ  \\simlt \\delta  \\simlt 1.25^\\circ$.  Despite  the fact  that their data  provide essentially  a two dimensional slice through the  Galaxy, they find astonishing sub-structure in the distribution of A-stars.   Sharply-defined concentric rings of stars are detected.   The  number  of  RR-lyrae  variables also  cuts  off  sharply  at $50\\kpc$.   The nature of  these structures  is important,  as they  are the dominant structure seen in old  Halo stars, comprising approximately half of the total number of Halo stars seen in the survey.  In this contribution, we show that these A-star  Halo sub-structures are morphologically very similar to  the expected  structure of  the tidal  debris of  the  Sagittarius dwarf galaxy, are  at similar distances  and are highly  likely to be part  of the Sagittarius stream.  At first sight it may seem  surprising that such a limited survey can detect Halo  sub-structure.   However,  since  most long-lived  Galactic  satellite systems  appear to be  on close  to polar  orbits the  chances of  a stellar debris stream  intersecting the current SDSS  zero-declination survey strips is  $\\approx$50\\%.  The  Sagittarius stellar  stream is  currently  the only known large  relatively intact debris system  and is within  $13^\\circ$ of a polar orbit. ",
        "conclusions": "We have shown  that numerical simulations of the  dynamical evolution of the Sagittarius  dwarf  galaxy,  previously  reported in  paper~1,  predict  the presence of sharp shell-like features in the Galactic halo.  This shell-like structure is due  to the effect of orbital precession  acting on stars which were  disrupted  from  the  Sagittarius  dwarf  galaxy  many  billion  years ago. Making a  cut through our simulated dataset  along the equatorial plane of the sky, we find a distribution  very similar to that found in A-stars in the SDSS data \\citep{yanny, ivezic}.  The  A-stars furthest  away from  the dominant  stream detected  in paper~1, belong  to  those populations  that  were  first  disrupted from  the  dwarf galaxy. By comparing  the relative densities of young to  old stars in these tidal structures, it may be  possible to reconstruct the dynamical evolution of the Sagittarius dwarf galaxy stream in great detail \\citep{johnston}.  The SDSS data that will soon  become available along other scan-lines on the sky offer the  very useful opportunity to distinguish  between the different dynamical models of the Sagittarius  dwarf galaxy. Such data will also allow us to  distinguish competing  models for the  mass distribution of  the dark halo. For instance, a comparison  of the particle distributions of Figures~2 and 3,  shows that the enhanced  stream precession in the  flattened halo of Figure~3 gives rise  to a substantial density between  $270^\\circ < \\alpha < 0^\\circ$, which is not seen in  Figure~2. Data in that right ascension range will provide a means to further constrain the shape of the Galactic halo.  With  the almost  all  high-latitude sky  coverage  of the  APM carbon  star survey, only  two Halo streams were  found (the other probably  being due to the Large Magellanic Cloud). Given  that SDSS has also intersected the tidal stream from the  Sagittarius dwarf galaxy, this stream appears  to be by far the dominant  contributor to  Halo substructure at  Galactocentric distances less  than  $\\sim 50\\kpc$.  It  seems  unlikely,  therefore, that  any  more accreted systems of  such magnitude will be found by  SDSS or other surveys. However,  there  is  hope  to  discover the  much  smaller  accreted  galaxy fragments, if they exist.  The simulation result of Figure~1, has a cautionary tale to tell about using starcounts  to determine  the shape  of  the ``Halo''.   That a  substantial fraction  of the  outer  Galactic halo  stars  appear to  be  clumped on  an inclined band  on the  sky clearly  does not mean  that the  Halo is  a thin inclined flattened band! Such structures just tell a complicated story about the process of stellar accretion into the Halo.  Since stars only contribute a miniscule amount to  the total halo mass, the only way  to reveal the dark matter distribution  is through dynamical  analyses.  For this  purpose, the dynamically cold streams discussed in this paper are ideal."
    },
    "0004/astro-ph0004063_arXiv.txt": {
        "abstract": "A sample of 35 SNe\\,Ia with good to excellent photometry in B and V, minimum internal absorption, and 1200 $<$ $v \\stackrel{_<}{_\\sim}$ 30\\,000\\,km\\,s$^{-1}$ is compiled from the literature. As far as their spectra are known they are all Branch-normal. For 29 of the SNe\\,Ia also peak magnitudes in I are known. The SNe\\,Ia have uniform colors at maximum, i.e. $<$$B$-$V$$>$=$-$0\\magn012 ($\\sigma$=0.051) and $<$$V$-$I$$>$=$-$0\\magn276 ($\\sigma$=0.078). In the Hubble diagram they define a Hubble line with a scatter of $\\sigma_M$=0\\magn21-0\\magn16, decreasing with wavelength. The scatter is further reduced if the SNe\\,Ia are corrected for differences in decline rate $\\Delta m_{15}$ or color ($B$$-$$V$). A combined correction reduces the scatter to $\\sigma$\\,$\\stackrel{_<}{_\\sim}$\\,0\\magn13. After the correction no significant dependence remains on Hubble type or galactocentric distance. The Hubble line suggests some curvature which can be differently interpreted. A consistent solution is obtained for a cosmological model with $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7, which is indicated also by much more distant SNe\\,Ia. Absolute magnitudes are available for eight equally blue (Branch-normal) SNe\\,Ia in spirals, whose Cepheid distances are known. If their well defined mean values of $M_B$, $M_V$, and $M_I$ are used to fit the Hubble line to the above sample of SNe\\,Ia one obtains $H_0$=58.3 km\\,s$^{-1}$\\,Mpc$^{-1}$, or, after adjusting all SNe\\,Ia to the average values of $\\Delta m_{15}$ and ($B$$-$$V$), $H_0$=60.9 km\\,s$^{-1}$\\,Mpc$^{-1}$. Various systematic errors are discussed whose elimination tends to decrease $H_0$. The finally adopted value at the 90-percent level, including random and systematic errors, is $H_0$=58.5$\\pm$6.3 km\\,s$^{-1}$\\,Mpc$^{-1}$. Several higher values of $H_0$ from SNe\\,Ia, as suggested in the literature, are found to depend on large corrections for variations of the light curve parameter and/or on an unwarranted reduction of the Cepheid distances of the calibrating SNe\\,Ia. ",
        "introduction": "Supernovae of type Ia (SNe Ia) are the prime distance indicators for the determination of the Hubble constant $H_0$ since they can be followed out to large distances, and since it is possible to determine Cepheid distances with HST for the nearest of their host galaxies and hence to calibrate their luminosity at maximum \\citep{sand92}. This immediately yields the distances of the more distant SNe Ia because they are -- if restricted to blue objects -- the best standard candles known with a luminosity scatter of less than $\\sigma_M$\\,=\\,0\\magn25. As standard candles they are now so heavily relied upon that they are even used for the much more sensitive test for the cosmological constant $\\Lambda$ \\citep{ries98b,perl99}. A wealth of excellent photometric data for SNe Ia has been accumulated since 1985 by the Cal\\'an/Tololo team \\citep{hamu96b} and others. They reach out to \\( \\sim30\\,000\\) km\\,s$^{-1}$, i.e. far enough to avoid the effects of peculiar and streaming motions. In parallel, the HST Supernova Project has by now provided Cepheid-calibrated luminosities for six nearby SNe Ia \\citep{saha99}, not counting SN 1895 B for which only the B magnitude is known. These are augmented by two additional calibrators from \\citet{tanv95} and \\citet{turn98}. With a total of eight calibrators and three dozen more distant SNe Ia, the statistics rests on solid ground.  The one remaining question is whether the calibrators and the distant SNe Ia are genuine twins, or whether the different selection criteria cause systematic differences between the samples. For instance, the Cepheid distances of the calibrators imply that they lie in galaxies containing young-population stars, whereas the distant SNe Ia have also been observed in S0 and E galaxies. Various second parameters have been proposed to correlate with SN luminosity; they can serve as a control of luminosity differences between the nearby calibrators and the distant SNe Ia. From the point of view of the physicist the most interesting second parameters are several spectral features which are known to correlate with luminosity \\citep{nuge95,fish95,fish99,ries98a,mazz98}. But the available data are too sparse to be useful in the present context. One is therefore reduced to empirical second parameters such as light curve shape, SN color, Hubble type and position in the parent galaxy. Second parameters have been discussed by, e.g., \\citet{phil93,tamm95,vaug95,hamu96a,hamu96c,trip98,trip99,saha99,ries99,jha99,phil99,sunt99,gibs00a}.  The difficulty to find a correlation between SN Ia luminosities and the second parameters in face of an intrinsic scatter of $\\sigma_M$$\\stackrel{_<}{_\\sim}$0\\magn25 is that very accurate {\\em relative} distances are required. Cepheid distances and, e.g., Tully-Fisher distances are not sufficiently accurate for the purpose. SNe Ia in the Virgo cluster cannot be assumed to lie at the same distance because of the important depth effect of the cluster. Regress must therefore be taken to the {\\em relative} distances that are indicated by recession velocities. All velocities $v$\\,$>$\\,1200 km\\,s$^{-1}$ are taken as indicative of the relative distances. This is permissible because the errors assigned to the relative distances and to the resulting relative absolute magnitudes make allowance for reasonable values of the peculiar velocities (Section 2.1).  The purpose of the present paper is to discuss the correlation of second parameters with the peak luminosity of SNe\\,Ia --- using an enlarged and well-defined sample of distant SNe Ia --- and to determine the value of the Hubble constant. The SNe Ia magnitudes, after correction for decline rate $\\Delta m_{15}$ and color ($B$-$V$), have a scatter of only $\\sigma_M$$\\stackrel{_<}{_\\sim}$0\\magn13. In fact the magnitude-corrected SNe Ia define the Hubble diagram so well that a flat Universe model with $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7 gives a marginally better fit than an $\\Omega_M$=1 model. The corrected magnitudes, if combined with the corrected absolute magnitudes of eight Cepheid-calibrated SNe\\,Ia, determine $H_0$ with a very small {\\em statistical} error.  The organization of the present paper is as follows. In Section 2 the available data for blue SNe Ia are compiled, and their luminosity calibration by means of eight SNe Ia with known Cepheid distances is discussed.   The SNe Ia colors and extinctions are discussed in Section 3. The Hubble diagram is shown in Section 4. In Section 5 the SN Ia luminosities are discussed in function of decline rate $\\Delta m_{15}$, SN color ($B$-$V$), Hubble type $T$, and position in the parent galaxy. The effective Hubble diagram with decline-rate and color corrected magnitudes is shown in Section 6, and the resulting values of $H_0$ are derived. Alternative solutions are explored in Section 7, and the conclusions are given in Section 8. ",
        "conclusions": "A fiducial sample of 35 well observed, blue (Branch-normal) SNe\\,Ia with minimum absorption in their parent galaxies and with 1200\\,$<$\\,$v$\\,$\\stackrel{_<}{_\\sim}$\\,30\\,000\\,km\\,s$^{-1}$ has been compiled from the literature. For 29 SNe\\,Ia also I-magnitudes are known. The intrinsic color of SNe\\,Ia is found to be ($B$-$V$)=-0\\magn012$\\pm$0.008, ($V$-$I$)=-0.276$\\pm$0.016 from objects in E/S0 galaxies and outlyers in spirals; the scatter in color of $\\sigma$=0\\magn05 for individual objects is surprisingly small.  The fit of the SNe\\,Ia sample, although quite local ($z$$\\stackrel{_<}{_\\sim}$0.1), to a Hubble line is somewhat dependent on the cosmological model adopted. Among various models, the marginally best fit is obtained for a flat Universe with $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7, which is presently also favored by external data. In this case the scatter about the Hubble line is $\\sigma$=0\\magn21-0\\magn16, depending on the pass band. But the magnitude residuals correlate with decline rate $\\Delta m_{15}$ and color ($B$-$V$). If the dependency on either of these parameters is removed, the scatter is reduced to $\\sigma$=0\\magn19-0\\magn14. A simultaneous correction for $\\Delta m_{15}$ and color ($B$-$V$) decreases the scatter to $\\sigma$$\\stackrel{_<}{_\\sim}$0\\magn13 in all three colors. This illustrates the unique potential of SNe\\,Ia as distance indicators.  If the SNe\\,Ia of the fiducial sample are assigned the absolute magnitudes $M_B$, $M_V$, and $M_I$ of eight local SNe\\,Ia whose Cepheid distances are known, one obtains closely the same value of $H_0$=58.3$\\pm$2.0 for the three passbands. This solution is remarkably robust against various assumptions and corrections. Bootstrapping the calibrators by excluding two or three objects, does not change their mean absolute magnitude. The eight calibrators lie essentially all in spirals and they have correspondingly a somewhat smaller average $\\Delta m_{15}$ than the fiducial sample. If they are adjusted to the latter's mean $\\Delta m_{15}$, $H_0$ is increased by 4 percent to $H_0$=60.9$\\pm$2.0. A correction for ($B$-$V$) does not change this value because the calibrators and the fiducial sample have closely the same mean color. Any remaining dependencies on Hubble type and galactocentric distance are insignificant and have no effect on $H_0$. Alternative solutions considering either SNe\\,Ia only in spirals or SNe\\,Ia with small $\\Delta m_{15}$ --- making the sample more similar to the calibrators --- yields $H_0$$\\approx$59. The effect of systematic errors is mainly one-sided (photometric zeropoint, skewness of photometric errors, the LMC zeropoint of the PL relation, selection effects of Cepheids, and possibly metallicity effects), leading to too high values of $H_0$. If they are taken into account one obtains at the 90 percent confidence level $H_0$=58.5$\\pm$6.3.  External results on $H_0$ from SNe\\,Ia should be compared with $H_0$=60.9$\\pm$2.0 from equation (12), because the subsequent corrections for systematic errors have not been applied yet in other papers. The results of \\citet{saha99} and \\citet{trip98} are indistinguishable from the above result despite of the re-definition of the ``fiducial sample'' here, of variations of other input data, and of independent corrections for $\\Delta m_{15}$ and  ($B$-$V$). The value of $H_0$=62($\\pm$2) derived by \\citet{trip99} from a sample of six calibrators and 13 field SNe\\,Ia in spirals and fully corrected for $\\Delta m_{15}$ and ($B$-$V$)  is in satisfactory statistical agreement with equation (12), too.  The somewhat higher value of $H_0$=63.9$\\pm$2.2 by \\citet{sunt99}, based on only five of the calibrators in Table 3, is fully explained by their adopted steeper luminosity dependence on $\\Delta m_{15}$, which is closely the same as was derived by \\citet{phil99}, who included spectroscopically peculiar SNe\\,Ia like SN\\,1986\\,G, 1991\\,T, 1991\\,bg, and 1995\\,ac. Had Suntzeff's et al. (1999) $\\Delta m_{15}$-correction been applied to our fiducial sample, we would have obtained $H_0$=59.7 for the SNe\\,Ia with $\\Delta m_{15}$$<$1.2 ($n$=17), and $H_0$=64.6 for the ones with $\\Delta m_{15}$$\\ge$1.2 ($n$=18). Since seven of the eight calibrators fall into the first category the lower value of $H_0$ is more nearly correct. The 3-$\\sigma$ discrepancy demonstrates that their $\\Delta m_{15}$-correction introduces an over-correction, making the calibrators with a small mean $\\Delta m_{15}$ too faint and the field SNe\\,Ia with large $\\Delta m_{15}$ too bright. This is the price to be paid if one attempts to derive a global correction formula which encompasses also peculiar objects that are extraneous of the SNe\\,Ia sample to which it should be applied. --- G00 have followed closely the line of \\citet{sunt99}, but they have reduced in addition the mean absolute magnitude of the calibrators by $\\sim$0\\magn2 on the basis of quite objectionable additional Cepheids (cf. Section 2.2); their high value of $H_0$=68 (based on six calibrators) then follows by necessity. The repetition of one or two calibrating galaxies with the forthcoming Advanced Camera System on HST will bring a definitive decision on the merits of the additional Cepheids of G00, which were rejected here because their photometry relies so far only on ALLFRAME.  \\citet{jha99} have derived $H_0$=64.4$\\pm$5.4 from a sample of 42 distant SNe\\,Ia and only four Cepheid-calibrated SNe\\,Ia. Seven of these objects have not been used here as either being too red or having peculiar spectra (SN 1992\\,K, 1995\\,ac). The SNe\\,Ia are standardized by a ``multi-color light curve shape'' (MCLS) method giving simultaneously the absorption and a light-curve parameter substituting $\\Delta m_{15}$ \\citep{ries96,ries98b}. Unfortunately the individual data are not given, and the discussion is restricted to $V$ magnitudes denying a control on the consistency of colors. In spite of this, there is some question as to the grip of the MLCS light-curve parameter. The calibrators as well as the distant SNe\\,Ia have still a magnitude scatter of $\\sigma_V$=0\\magn16 after the correction, which is larger than $\\sigma$=0\\magn14 obtained by \\citet{phil99} and significantly larger than $\\sigma_V$$\\le$0\\magn13 found here by means of the corrections due to equation (10).  This discussion confirms previous investigations suggesting that if one attempts to derive $H_0$ from SNe\\,Ia to within $\\stackrel{_<}{_\\sim}$10\\,percent, a major stumbling block are the second-parameter corrections. If the relatively mild $\\Delta m_{15}$-corrections found here can be agreed upon it is clear that $H_0$ will settle close to 60 with a tendency for a small downward correction due to various (small) systematic errors.  Future Branch-normal SNe\\,Ia in the field, preferably with $v$\\,$\\stackrel{_>}{_\\sim}$\\,3000\\,km\\,s$^{-1}$, will help to settle the problem of light-curve shape corrections. The importance of these corrections will decrease as it will become possible to define sufficiently large sub-samples of field SNe\\,Ia which have the same mean second parameters as the calibrating SNe\\,Ia. Additional field SNe\\,Ia are forthcoming \\citep{nuge99}. Finally, the advent of the new Advanced Camera System on HST will allow to check the present Cepheid distances and to increase the number of calibrating SNe\\,Ia.  At present the evidence from SNe\\,Ia is best reflected at the 90-percent confidence level by a value of the Hubble constant, including random and systematic errors, of $H_0$=58.5$\\pm$6.3."
    },
    "0004/astro-ph0004241_arXiv.txt": {
        "abstract": "Measured time delays between the images of a gravitationally lensed source can lead to a determination of the Hubble constant ($H_o$), but only if the lensing mass distribution is well understood. The inability to sufficiently constrain galaxy mass models results in large uncertainties on the derived $H_o$, and severely hampers the cosmological application of this otherwise elegant method. At the same time, lensing must compete with new techniques that have the potential to measure the Hubble constant to within a few percent by the middle of the decade. In this letter we re-evaluate the role of strong gravitational lensing in the age of precision cosmology, and present a series of Monte Carlo simulations that demonstrate the effect of the galaxy mass distribution on Hubble constant determination.  Though most gravitational lens systems are unlikely to contribute significantly to precision studies of the cosmological distance scale, the strong dependence of predicted time delays on the galaxy mass profile suggests that the most useful and interesting results may be obtained by running the traditional lensing problem in reverse -- namely, combining measured time delays with a well-constrained value of $H_o$ to study galaxy mass distributions. ",
        "introduction": "Gravitational lensing offers a potent tool for determining the cosmological distance scale (Refsdal 1964), but only if the mass distribution of the deflector is well constrained. Since the mass model predicts time delays between the lensed images up to a factor of $h^{-1}$, where $H_o = 100h\\kmsm$, a simple comparison with measured time delays can yield the Hubble constant. This method is particularly tempting as it would determine $H_o$ directly out to cosmological distances, forgoing the calibration errors associated with traditional distance-ladder techniques (e.\\ g.\\ Mould et al.\\ 2000).  Much of the recent interest in finding arcsecond-scale gravitational lens systems through large, dedicated programs such as the Cosmic Lens All-Sky Survey (CLASS; Myers et al.\\ 2000) has been motivated by this goal.  Measurements of $H_o$ from gravitational lenses are only as good as our understanding of the galaxies doing the lensing, as the model-predicted time delays contribute an equal fractional uncertainty to the derived Hubble parameter. Unfortunately, experience has shown that our knowledge of lensing mass distributions is not very good at all. This problem has been greatly exacerbated by the observation that a significant fraction of lenses, including several of those with measured time delays, are quite complex in structure (e.\\ g.\\ CLASS B1359+154, Rusin et al.\\ 2000; QSO 0957+561, Barkana et al.\\ 1999; PG 1115+080, Schechter et al.\\ 1997; CLASS B1608+656, Koopmans and Fassnacht 1999). Mass models are difficult to construct for such systems due to limited constraints and parameter degeneracies (e.\\ g.\\ the mass-sheet degeneracy, Gorenstein et al.\\ 1988).  Significant uncertainties in the mass model can remain even in systems lensed by a single isolated galaxy. This is because the positions and flux density ratios of the lensed images constrain derivatives of the potential, rather than the potential itself. Often a range of models can provide sufficient fits to the lens system while placing different potentials at the image positions. The predicted time delays will be greatly affected by model degeneracies, as they depend on the potential directly (Schneider, Ehlers and Falco 1992). This is particularly discouraging for many two-image lens systems, which do not offer a sufficient number of constraints to even attempt fits to the profile shape. For such systems the profile must be assumed, and is commonly taken to be isothermal. The uncertainty in the slope of the mass profile and its effect on Hubble constant determination are rarely included in most preliminary modeling analyses.  Attempts to constrain galaxy mass profiles through strong lensing generally focus on image positions, flux density ratios and, if available, the ratios of measured time delays. Studies of this type have thus far been performed in only a small handful of lens systems, and have had limited success in pinning down the mass profile. Assuming a power-law surface mass density ($\\Sigma \\sim r^{-\\gamma}$), it has been found that the merging lens galaxies of CLASS B1608+656 are best fit by $0.8 \\leq \\gamma \\leq 1.2$ (Koopmans and Fassnacht 1999), while the large elliptical galaxy in QSO 0957+561 is best fit by $0.82 \\leq \\gamma \\leq 0.93$ (Grogin and Narayan 1996). However, each of these systems is affected by complicated deflectors. A profile analysis has also been done for MG 1654+134 (Kochanek 1995), which is a radio ring system lensed by an isolated galaxy. Because the lensed extended emission probes a wide range of radii (as opposed to the images of a compact source), ring systems can produce tighter constraints on the mass profile ($0.9 \\leq \\gamma \\leq 1.1$ for MG 1654+134).  While isothermal ($\\gamma=1$) mass profiles are largely compatible with lensing, a range of near-isothermal profiles ($0.8 \\leq \\gamma \\leq 1.2$) are equally feasible. In this Letter I outline the systematic effects of this uncertainty on the predicted time delays, and evaluate the possibility of measuring $H_o$ with the set of lens systems currently known. I then consider the largely-ignored reverse problem in lensing -- namely, using measured time delays along with a well-determined value of $H_o$ to study galaxy mass profiles. The results offered by this technique have the potential to provide greatly improved constraints on the inner mass distribution for galaxies at intermediate redshift. ",
        "conclusions": "Gravitational lensing has already contributed to our knowledge of both galaxy mass distributions and the Hubble constant. The profiles of several lens galaxies have been constrained at the $10-20\\%$ level, while the ensemble of lenses with measured time delays may have determined the Hubble constant to $\\sim 10\\%$ (Koopmans and Fassnacht 1999). In this age of precision cosmology and astrophysics, however, new techniques and observations promise very accurate measurements of a wide range of quantities. It is therefore necessary to reassess the goals and merits of more familiar methods.  The traditional application of gravitational lensing seeks to measure the Hubble constant, and requires that the mass distribution of the lensing galaxy be modeled independently of the absolute time delays. Uncertainties in the mass profile dominate the error budget in $H_o$ for every system in which time delays are currently measured. At the same time, lensing must compete with new techniques that have the potential to determine the Hubble constant to within a few percent.  The time delay method, however, can be elegantly inverted to provide powerful constraints on galaxy mass profiles, once the Hubble constant has been accurately measured by some other means. Whereas highly constrained lens systems, such as quads or doubles with well-separated substructure, would be needed to derive $H_o$ independent of any assumptions regarding the mass profile for a given galaxy, time delay determinations in almost any simple two-image system can make interesting statements about the mass distribution present. In time this approach may provide the best tool for probing the inner few $h^{-1}$ kpc in elliptical galaxies, and will nicely complement weak lensing techniques that are sensitive to the mass profile at large radius (Fischer et al.\\ 2000). Because photometric studies by the CASTLES collaboration (Kochanek et al.\\ 2000) have found that lensing galaxies appear to be consistent with an ordinary population of passively evolving ellipticals, the results derived from lensing can be applied to typical galaxies at intermediate redshift. In conclusion, the precision studies of galaxy mass profiles offered by lensing suggest that the most significant results from Refsdal's method may be astrophysical in nature, rather than cosmological."
    },
    "0004/astro-ph0004288_arXiv.txt": {
        "abstract": "We propose a physical model for nonlinear stochastic biasing of one-point statistics resulting from the formation epoch distribution of dark halos. In contrast to previous works on the basis of extensive numerical simulations, our model provides for the first time an analytic expression for the joint probability function.  Specifically we derive the joint probability function of halo and mass density contrasts from the extended Press-Schechter theory. Since this function is derived in the framework of the standard gravitational instability theory assuming the random-Gaussianity of the primordial density field alone, we expect that the basic features of the nonlinear and stochastic biasing predicted from our model are fairly generic. As representative examples, we compute the various biasing parameters in cold dark matter models as a function of a redshift and a smoothing length. Our major findings are (1) the biasing of the variance evolves strongly as redshift while its scale-dependence is generally weak and a simple linear biasing model provides a reasonable approximation roughly at $R\\simgt 2(1+z)\\himpc$, and (2) the stochasticity exhibits moderate scale-dependence especially on $R\\simlt 20\\himpc$, but is almost independent of $z$.  Comparison with the previous numerical simulations shows good agreement with the above behavior, indicating that the nonlinear and stochastic nature of the halo biasing is essentially understood by taking account of the distribution of the halo mass and the formation epoch. ",
        "introduction": "\\label{sec:intro}  The universe after the last scattering can be probed only through observing the distribution of luminous objects, either directly or indirectly via the weak lensing effect. This is why several wide-field and/or deep surveys of galaxies, clusters and quasars are planned and operating at various wavelengths. The purpose of such cosmological surveys is two-fold; to understand the nature of the astronomical objects themselves, and to extract the cosmological information. From the latter point of view, which we will pursue throughout this paper, the objects serve merely as tracers of dark matter in the universe. Since such {\\it luminous} objects should have been formed as a consequence of complicated astrophysical processes in addition to the purely gravitational interaction, it is quite unlikely that they faithfully trace either the spatial distribution of dark matter or its redshift evolution.  Rather it is natural to assume that they sample the dark matter distribution in a biased manner.  To describe the {\\it biasing} more specifically, define the density contrasts of galaxies and mass at a position ${\\bf x}$ and redshift $z$ smoothed over the scale $R$ as \\begin{equation} \\label{eq:dgdm} \\deltag({\\bf x},z|R) ={n_{\\rm gal}({\\bf x},z|R) \\over \\bar{n}_{\\rm gal}} - 1, \\qquad \\deltam({\\bf x},z|R) ={\\rho_{\\rm mass}({\\bf x},z|R) \\over \\bar\\rho_{\\rm mass}} - 1. \\end{equation} Here and in what follows we use the words, ``mass'' and ``dark matter'', interchangeably, and ``galaxies'' to imply luminous objects (galaxies, clusters, quasars, etc.) in a collective sense.   The simplest, albeit most {\\it unlikely}, possibility is that they are proportional to each other: \\begin{equation} \\label{eq:linbias} \\deltag({\\bf x},z|R) = b~ \\deltam({\\bf x},z|R) . \\end{equation} While the proportional coefficient was assumed to be constant when the concept of the biasing was first introduced in the cosmology community (Kaiser 1984; Davis et al. 1985; Bardeen et al.1986), it has been subsequently recognized that $b$ should depend on $z$ and $R$ (e.g., Fry 1996; Mo \\& While 1996). As a matter of fact, it is more realistic to abandon the simple linear biasing ansatz (\\ref{eq:linbias}) completely, and formulate the biasing in terms of the conditional probability function $P(\\deltag|\\deltam)$ of $\\deltag$ at a given $\\deltam$. Then equation (\\ref{eq:linbias}) is rephrased as \\begin{equation} \\label{eq:linbias2} \\overline\\deltag(\\deltam) = \\int_{-1}^{\\infty} \\deltag ~P(\\deltag|\\deltam) ~d\\deltag . \\end{equation} Obviously the relation between $\\deltag$ and $\\deltam$ is neither linear nor deterministic.  This general concept -- nonlinear and stochastic biasing -- was introduced and developed in a seminal paper by Dekel \\& Lahav (1999), which inspired numerous recent activities in this field (e.g., Pen 1998; Taruya, Koyama \\& Soda 1999; Blanton et al. 1999; Matsubara 1999; Taruya 2000; Somerville et al. 2000; Taruya et al. 2000).  Then the crucial question is the physical interpretation of $P(\\deltag|\\deltam)$.  There are (at least) two different interpretations for its physical origin. The first is based on the fact that $\\deltag$ is meaningless unless one specifies many {\\it hidden} variables characterizing the galaxies in the catalogue, for instance, their luminosity, mass of dark matter and gas, temperature, physical size, formation epoch and merging history, among others. This list is already long enough to convince one for the inevitably broad distribution of $P(\\deltag|\\deltam)$.  In this spirit, Blanton et al.  (1999) proposed that the gas temperature of a local patch is the important variable to control $P(\\deltag|\\deltam)$ on the basis of cosmological hydrodynamical simulations.  The other adopts the view that our universe is fully specified by the primordial density field of dark matter. According to this interpretation, all the properties of any galaxy should be in principle computable given an initial distribution of dark matter field in the entire universe, at least in a statistical sense.  Actually this is exactly what the cosmological hydrodynamical simulations attempt to do. The gas temperature of a local patch, for instance, should be determined by a non-local attribute of the dark matter fluctuations. Clearly the above interpretations are not conflicting, but rather stress the two different aspects of the physics of galaxy formation which is poorly understood at best.  In this paper, we present an analytical model for nonlinear stochastic biasing by combining the above two interpretations in a sense. Specifically we derive the joint probability function of $\\deltah$ and $\\deltam$, $P(\\deltah, \\deltam)$, from the distributions of the formation epoch and mass of halos on the basis of the extended Press-Schechter theory. In contrast to previous work which were based on the results of numerical simulations (Kravtsov \\& Klypin 1999; Blanton et al. 1999; Somerville et al. 2000), our model provides, for the first time, an analytic expression for $P(\\deltah, \\deltam)$. Thus one can compute the various biasing parameters in an arbitrary cosmological model at a given redshift and a smoothing length in a straightforward fashion.  We derive the joint probability function assuming the primordial random-Gaussianity of the dark matter density field, and thus the results are sufficiently general.  We note here, however, that our primary purpose is to present a general formulation to predict biasing properties of halos, and not to make detailed predictions at this point. In fact we adopt a few approximations with limited validity in presenting specific examples, but this is not essential in our paper and the resulting predictions can be improved in a straightforward manner if other analytical/numerical approximations become available.  Nevertheless we would like to emphasize that our simple analytical prescription largely explains the basic features of the biasing parameters reported in the previous numerical simulations (Kravtsov \\& Klypin 1999; Somerville et al. 2000). Thus our prescription is supposed to capture the most important processes in the halo biasing.  We organize the paper as follows. In \\S \\ref{sec:formalism}, we describe a general formalism for the one-point statistics of the galaxy and the mass distributions from a point of view of the hidden variable interpretation of the nonlinear and stochastic biasing theory.  Applying this general formalism, we develop a model for dark matter halo biasing treating the halo mass and formation epoch as the hidden variables in \\S \\ref{sec:model}.  The resulting expression for the conditional probability function, $P(\\deltah|\\deltam)$, can be numerically evaluated using the extended Press-Schechter theory, and we show various predictions for the halo biasing in \\S \\ref{sec:results}.  In particular, we pay attention to their scale-dependence and redshift evolution, and compare our model predictions to previous simulation results. Finally section \\ref{sec:conclusions} is devoted to the conclusions and discussion. ",
        "conclusions": "\\label{sec:conclusions}   In this paper, we presented a general formalism to describe the nonlinear stochastic biasing from the hidden variable interpretation. According to this formulation we proposed a physical model for the halo biasing assuming that mass and formation epoch of dark halos act as the major hidden variables. In particular, this model for the first time provides an analytic expression for the joint probability distribution function of the halos and the dark matter density fields. The specific functional forms for the PDFs can be derived, or are at least consistent with the phenomenological results, from the standard gravitational instability theory and the assumption of the random-Gaussianity of the primordial density field.  Therefore we expect that the basic features of the nonlinear and stochastic biasing predicted from our model are fairly generic. In fact, detailed comparison with the previous numerical simulations by Kravtsov \\& Klypin (1999) and Somerville et al. (2000) shows good agreement with our predictions, indicating that the nonlinear and stochastic nature of the halo biasing is essentially understood by taking account of the distribution of the halo mass and the formation epoch.  Then we introduced a set of biasing parameters from second-order statistical quantities following Dekel \\& Lahav (1999), which properly quantify the one-point statistics of nonlinear and stochastic biasing. As specific examples, we explicit compute those parameters in LCDM model, and examined their scale- and time-dependence. Our major findings are summarized as follows: \\begin{enumerate} \\item The biasing of the variance, $\\bvar \\equiv \\sqrt{\\sigmahh/\\sigmamm}$, evolves strongly as redshift. While its scale-dependence becomes noticeable on small scales and/or at high redshifts, a simple linear biasing model provides a reasonable approximation roughly at $R\\simgt 2(1+z)\\himpc$. \\item The stochasticity, $\\rcorr \\equiv \\sigmahm^2/\\sqrt{\\sigmahh\\sigmamm}$ exhibits moderate scale-dependence especially on $R\\simlt 20\\himpc$, but is almost independent of $z$. Since $\\scatt \\gg \\nonl$ in general, the stochasticity of halo biasing is mainly generated by the scatter which is dominated by the formation epoch distribution at lower redshifts and by the halo mass distribution at higher redshifts. Importantly, the stochasticity still remains constant even on large scales, which yields the small deviation of $\\rcorr$ from unity. \\end{enumerate} The fact that biasing approaches linear and scale-independent on large scales is just expected from previous theoretical work (e.g., Coles 1993; Scherrer \\& Weinberg 1998) and was recently suggested observationally (Seaborne et al. 1999). Our model further implies that the scale-independence still holds even in the quasi-linear regime, and more importantly, that biasing is still stochastic to some extent on those scales.   Before closing we would like to comment on several limitations of our model and on future prospects in order. First, our analysis is entirely dependent on the formation redshift distribution derived by Lacey \\& Cole (1993). It is known, however, that the distribution becomes negative in some models, suggesting that this function is not fully correct. Another proposal for the distribution function by Kitayama \\& Suto (1996a) also suffers from  a different conceptual problem. This might be ascribed to the difficulty of distinguishing the merger and accretion unambiguously within the framework of the extended PS theory.  Despite this limitation, the two proposals by Lacey \\& Cole (1993) and by Kitayama \\& Suto (1996a) lead to very similar predictions and thus we believe it unlikely that this significantly affects the final results.  Second, several authors have claimed that a better agreement with numerical simulations is attained by incorporating the non-spherical effect into the Press-Schechter theory (Lee \\& Shandarin 1998; Sheth, Mo, \\& Tormen 1999). Although we do not think that the non-spherical effect drastically changes the statistical predictions of our model, this is definitely a well-defined and interesting possibility to improve our model provided that the corresponding distribution function of the formation epoch can be derived. Third, for the proper comparison with observation, our predictions should be corrected for the redshift-space distortion. Since the distortion induced by peculiar velocity field also leads to the stochasticity, the resulting biasing becomes object-dependent (Taruya et al. 2000).  Finally and more importantly, the astrophysical effects other than gravity must be taken into account. Since a halo identified in our scheme carries its formation epoch, it is not so difficult to combine with the phenomenological galaxy formation model.  Such an improved model will enable more generic and testable predictions for the {\\it galaxy} biasing.  \\bigskip \\bigskip   A.T. gratefully acknowledges support from a JSPS (Japan Society for the Promotion of Science) fellowship.  Numerical computations were carried out at RESCEU (Research Center for the Early Universe, University of Tokyo) and ADAC (the Astronomical Data Analysis Center) of the National Astronomical Observatory, Japan.  This research was supported in part by the Grant-in-Aid by the Ministry of Education, Science, Sports and Culture of Japan (07CE2002) to RESCEU.  \\bigskip \\bigskip"
    },
    "0004/astro-ph0004131_arXiv.txt": {
        "abstract": "We have investigated the effects of resonant scattering of emission lines on the image morphology and intensity from coronal loop structures. It has previously been shown that line of sight effects in optically thin line emission can yield loop images that appear uniformly bright at one viewing angle, but show ``looptop sources'' at other viewing angles.  For optically thick loops where multiple resonant scattering is important, we use a 3D Monte Carlo radiation transfer code.  Our simulations show that the intensity variation across the image is more uniform than the optically thin simulation and, depending on viewing angle, the intensity may be lower or higher than that predicted from optically thin simulations due to scattering out of or into the line of sight. ",
        "introduction": "Two effects of resonant line scattering on emission lines are to change the morphology of images and to change line ratios from what would be expected from optically thin analyses.  Qualitatively, the effects can be split into three broad optical depth categories, low, moderate, and large.  At low optical depths ($\\tau < 5$) resonant scattering primarily changes the photon direction and photons emerge into the directions of lowest optical depth.  At moderate optical depths ($\\tau \\sim 10 - 100 $) multiple scattering is important and there is a large probability of photon energy conversion, e.g., the conversion of Ly$\\beta$ photons to H$\\alpha$ photons.  Large optical depths ($\\tau >> 100$) result in photon trapping and collisional de-excitation, eventually leading to thermalization.  The goal of this paper is to investigate the effects of resonant scattering on images and intensity ratios in the low optical depth regime. Situations where our work is applicable include UV, EUV, and soft X-ray emission line imaging and spectroscopy of the Sun and stellar coronae.  While we have chosen a geometry based on coronal magnetic loops, similar considerations would apply, for instance, to some of the strong X-ray emission lines in older supernova remnants. We have simulated an X-ray line of Fe~{\\sc xvii}, but the predicted effects on image morphology should be visible in the narrow-band EUV images of EIT and TRACE.  Ultraviolet and soft X-ray emission line ratios are frequently used as diagnostics of temperature, density, elemental abundances, and the ionization state of astrophysical plasmas.  Several assumptions are often implicitly made when emission line ratios are analyzed.  In particular, the optical depth is generally assumed to be negligible. However, solar coronal observations of the Fe~{\\sc xvii} line at 15.01{\\AA} show line ratios that are lower than theoretical predictions (Phillips et al. 1997; Schmelz et al. 1997; Saba et al. 1999; Waljeski et al. 1994) or laboratory measurements (Brown et al. 1998). Rugge \\& MacKenzie (1985) suggested that optically thick resonant scattering of Fe~{\\sc xvii} 15.01{\\AA} photons out of the line of sight may be responsible for the small line ratios. Saba et al. (1999) employed the escape probability method of Kastner \\& Kastner (1990) in their analysis of various line ratios and found that up to 50\\% of the Fe~{\\sc xvii} flux is unaccounted for. Bhatia \\& Kastner (1999) found similar results for planar and cylindrical geometries. This explanation for the intensity decrements is likely, given the large oscillator strength of the Fe~{\\sc xvii} 15.01{\\AA} line (Table~1). For typical densities ($n_e \\sim 10^{10}$) and sizescales ($\\sim10^9$cm) in the solar corona, the optical depth of the Fe~{\\sc xvii} 15.01{\\AA} line is around 1. This is in the low optical depth regime in which the major effect is scattering of photons into optically thin directions.  In this paper we present three dimensional Monte Carlo radiation transfer simulations of the transfer of Fe~{\\sc xvii} 15.01{\\AA} photons. Our analysis does not treat transfer within the line itself, instead considering all photons to be emitted and resonantly scattered at a single wavelength.  This approximation is valid so long as the optical depth is not large enough that photons scatter out of the Doppler core. The advantage of our Monte Carlo technique is that it naturally accounts for arbitrary illumination and multiple scattering in complex geometries.  The results of our Monte Carlo simulations are images that may be compared directly with optically thin simulations to determine differences in image morphology and optically thick to optically thin intensity ratios.  Previous simulations by Alexander \\& Katsev (1996) produced optically thin images of loops in which the density and temperature follow the scaling law prescription of Rosner, Tucker, \\& Vaiana (1978, hereafter RTV). The simulations of Alexander \\& Katsev showed that it is crucial to consider the geometry and temperature dependence of the emissivity along with the thermal response of the imaging telescope when analyzing solar emission line images.  Such images may be in the light of a single line or an ensemble of lines as is the case with the broad wavelength response of Yohkoh SXT. In particular, a loop viewed face-on may appear uniformly bright, but line-of-sight effects can yield an apparent ``looptop'' source when the same loop is viewed close to edge-on (see Fig.~5 of Alexander \\& Katsev).  Our work is a natural extension of the Alexander \\& Katsev analysis to optically thick, multiple scattering situations, and for continuity purposes we also adopt RTV loops for our simulations.  Section~2 presents the model ingredients (loop structure, Fe~{\\sc xvii} emissivity and opacity, and radiation transfer technique).  Results of our simulations (images and intensity ratios of optically thick/optically thin models) are presented in \\S~3, and in \\S~4 we discuss our results with respect to observational data. ",
        "conclusions": "Our simulations have shown that optical depth effects, both within loops and in a background plasma, can change the amount of intensity emerging into a given direction and also the image morphology.  Intensity ratios may vary in the range 50\\% to more than 120\\% of their optically thin values.  If one observes an ensemble of loops, for instance spread over the surface of another star, the effects tend to average out.  Equal numbers of photons are scattered into and out of the line of sight.  However, a net effect may still arise if the emissivity and opacity vary in different ways.  Given the electron density factor in the emissivity and its absence in the opacity, such differences are likely. To account for the general trend of reduced Fe~{\\sc xvii} 15.01 \\AA\\/ intensity (e.g. Saba et al. 1999) by scattering rather than errors in atomic rates, the typical structure must preferentially scatter photons down to the solar surface.  Such an example may be  emission from dense, low-lying loops that is scattered in a more diffuse, larger scale region.  This is expected from RTV models and is observed in solar active regions. Models similar to those in the righthand column of Fig.~5 can easily be constructed to reduce the intensities of lines having large oscillator strengths.  They may be invoked when line ratio anomalies appear in Chandra or XMM data.  Proving that optical depth is indeed the explanation for a given anomaly will undoubtably be a challenge."
    },
    "0004/astro-ph0004307_arXiv.txt": {
        "abstract": "A search for radio recombination lines near 20 cm at z=0.193 and z=0.886 towards the gravitational lens system PKS 1830-211 has yielded upper limits of $|\\tau_L| \\leq$~5 $\\times$ 10$^{-5}$~and $\\leq$~5 $\\times$ 10$^{-4}$~ at the two redshifts respectively. Based on the non-detections, we derive upper limits to the emission measure of the ionized gas in the absorbing systems. We also present continuum flux density measurements over the frequency range 0.3$-$45 GHz made at a single epoch. ",
        "introduction": "\\vspace{-0.1in} There is evidence that the gravitational lens system PKS 1830-211 has two absorbers at two different redshifts of 0.886 and 0.193. The main lens, a normal galaxy at z=0.886, has been studied in HI, OH and a host of molecular lines (Chengalur, de Bruyn, \\& Narasimha 1999 and references therein). The only evidence for an additional absorber at z=0.193 is from HI absorption (Lovell et al. 1996; Verheijen et al. 1999) studies. Searches for molecular lines at this redshift have been unsuccessful.  \\vspace{-0.6cm} ",
        "conclusions": ""
    },
    "0004/astro-ph0004182_arXiv.txt": {
        "abstract": "The starburst galaxy M\\,82 has been mapped at $\\lambda$1.2~mm using the IRAM 30--m telescope and the MPIfR 19--channel bolometer. The influence of the sidelobe pattern and the $^{12}$CO(2$-$1) line has been carefully analyzed. Based on this analysis, we conclude that the cold dust does not extend far into the halo as has been previously claimed. \\changed{The total mass of dust in the inner 3~kpc of the galaxy is $\\changedm{7.5 \\cdot 10^6}$ M$_\\odot$. Assuming a solar gas-to-dust ratio, the inferred total mass of gas is $\\changedm{7.5 \\cdot 10^8}$ \\Msun.}  Unlike in any other galaxy we find the CO(2$-$1) line emission to be clearly more extended than the $\\lambda$1.2\\,mm dust emission. A likely explanation for this exceptional behaviour, seen for the first time in galaxies, is heating of molecular gas by low--energy cosmic rays and soft X--rays in the halo of M\\,82. ",
        "introduction": "\\label{sec:intro}  The prototypical starburst galaxy M\\,82 is the prime example of a galaxy in which the violent star formation activity gives rise to the formation of a bipolar outflow and an associated extended halo, which is visible in various regimes of the electro--magnetic spectrum. Outflowing material was first evident in the H$\\alpha$ light (Lynds \\changed{ \\& } Sandage 1963, McCarthy et al. 1987, Bland \\changed{ \\& } Tully 1988). The kinematics of the outflowing material and the geometry of the cone have been thoroughly worked out by McKeith et al. (1995). The recent detection of H$\\alpha$ emission $\\sim$11~kpc away from the plane of M\\,82 (Devine \\changed{ \\& } Bally 1999) determines the cap of this outflow. The whole scenario is further corroborated by the existence of vertical magnetic field lines (Reuter et\\,al. 1994), along which relativistic particles partake in the outflow, thus forming a radio halo (Seaquist \\& Odegard 1991). This halo exhibits a filamentary structure away from the plane (Reuter et\\,al. 1992). Hot gas has been also found in some kind of halo which extends for several arcminutes along the minor axis (Fabbiano 1988, Schaaf et\\,al. 1989, Bregman et\\,al. 1995). This X--ray emission correlates well with the H$\\alpha$ if observed with sufficiently high angular resolution (Watson et\\,al. 1984).  In this context, the search for neutral gas and dust away from the plane of M\\,82 has seen a number of attempts. Owing to the close interaction between M\\,81 and M\\,82, neutral hydrogen is seen enveloping M\\,82, but whether this gas stems from the M\\,81 or has been expelled from M\\,82 is still a matter of debate (Yun et\\,al. 1994). Molecular gas associated with the dusty filaments outside the plane of M\\,82 has been reported by Stark \\changed{ \\& } Carlson (1984) and has been observed by Sofue et\\,al. (1992) out to a projected distance of $\\pm$2~kpc. Clear evidence for the existence of scattering dust away from the plane of M\\,82 comes from measurements of optical polarization (Bingham et\\,al. 1976, Notni et\\,al. 1981, Neininger et\\,al. 1990).  Owing to its high luminosity in all spectral bands, M\\,82 has been a prime target for first--light experiments in the mm-- and submm--regime (Elias et\\,al. 1978, Jura et\\,al. 1978). With the improvement of bolometric measurements detailed studies of the distribution of the cold dust in this galaxy have become feasible; thus the recent past has seen an increasing number of such investigations ($\\lambda$2~mm: Kuno \\changed{ \\& } Matsuo 1997; $\\lambda$1~mm: Thronson et\\,al. 1989, Hughes et\\,al. 1990, Kr\\\"ugel et\\,al. 1990a, Alton et\\,al. 1999; $\\lambda$0.5~mm: Jaffe et\\,al. 1984, Smith et\\,al. 1990, Alton et\\,al. 1999). Carlstrom \\changed{ \\& } Kronberg (1991) have shown that at $\\lambda$3~mm thermal free--free emission still dominates the overall spectrum.  Apart from these dedicated bolometric measurements, continuum maps at millimeter wavelengths have been produced as by-products in various observations of spectral lines (e.g. Neininger et\\,al. 1998). As a result of the discoveries of a galactic wind and halo in M\\,82, the more recent  bolometric measurements aimed at detecting a dust halo at mm wavelengths. CO emission away from the plane was reported by Stark \\changed{ \\& } Carlson (1984), Nakai et\\,al. (1986) and by Sofue et\\,al. (1992). An outflow of molecular gas was claimed by Nakai et\\,al. (1987). A dust continuum halo was first mentioned by Hughes et\\,al. (1990). Observations with higher resolution (Kuno \\changed{ \\& } Matsuo 1997) suggest emission at $\\lambda$2~mm out to 400~pc from the plane. Recently Alton et\\,al. (1999) reported a dust outflow from the central region of M\\,82, based on their sub-mm images.  Here we report observations of M\\,82 in the $\\lambda 1.2$ mm continuum performed with the MPIfR 19-channel bolometer at the IRAM 30-m telescope. In Sect.~\\ref{sec:obse} we describe the observations and data analysis, with consideration of the possible influence of the error beam. In Sect.~\\ref{sec:result} the distribution of the cold dust in M\\,82 will be presented, with a discussion conducted in Sect.~\\ref{sec:discuss}, along with a comparison with other published measurements. In Sect.~\\ref{sec:sum} we give a short summary of our results. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{A dust halo?} \\label{sec:discuss:conti}  The $\\lambda$2\\,mm map (Kuno \\& Matsuo 1997) shows continuum emission with a box--like shape, indicating a dust halo around M\\,82. Although the emission at $\\lambda$1.2\\,mm and $\\lambda$2\\,mm should be similar, we do not see any /*-evidence for a significant halo contribution to the cold dust. This difference could be due to the error beam of the Nobeyama 45-m telescope.  Effects of the error beam and the antenna pattern have been investigated and largely removed from our map (see Sect. \\ref{sec:obse:errorbeam}), while this is not mentioned in case of the Nobeyama measurements. The total flux density in the Nobeyama map is 1.7\\,Jy, which is much more than expected at this wavelength (see Fig.~\\ref{fig:spektrum}\\changed{)}. This extra flux might be due to emission entering through sidelobes and the error beam.  The 850\\,$\\mu$m map made by Alton et\\,al. (1999) with the JCMT shows continuum emission  at the 50\\,mJy level up to 40\\arcsec\\ above the galactic plane. This translates to  15\\,mJy at $\\lambda$1.2\\,mm, corresponding to the second contour (3$\\sigma$) in our  map. Considering only the overall extent of the emission, the JCMT map is consistent with ours. The different shapes might be due to problems with the baseline subtraction, because the area mapped by Alton et\\,al. is only slightly larger than the extent of the continuum emission.   \\subsection{Comparison of dust and CO} \\label{sec:discuss:co}  In non--active galaxies like NGC\\,4565 (Neininger et\\,al. 1996) and NGC\\,5907 (Dumke et\\,al. 1997) the molecular line emission is more concentrated towards the central region than the thermal dust emission. As can be seen from Fig.~\\ref{fig:slice} and~\\ref{fig:vergleich} the galaxy M\\,82 exhibits the opposite behaviour. The CO emission appears to be much more extended than the continuum. Since the maps in Fig. \\ref{fig:vergleich} are made with the same telescope at the same frequency, the difference can not be due to different beam patterns and error beams.  \\changed{The dust is heated by the UV radiation field, which is concentrated towards the dense central region, because the UV photons can hardly escape without being absorbed. Only a small fraction of the UV photons makes its way out to kpc distances from the galactic plane, where it might contribute to the observed polarized radiation. }  The molecular gas can be excited by the radiation field, but also by low--energy cosmic rays and by soft X--ray photons (see e.g. Glassgold \\changed{ \\& } Langer 1973)\\changed{, which have a much larger scale height than the UV photons (Shopbell \\& Bland-Hawthorn 1998)}.  Fig.~\\ref{fig:slice} shows that the synchrotron intensity at $\\lambda$20~cm drops faster with distance from the galaxy plane than the CO and even the dust emission. However, at low radio frequencies there may be a significant radio halo (low--energy cosmic rays have longer lifetimes), which is already indicated by the 327--MHz map of Reuter et\\,al. (1992). Hence, there may be abundant cosmic rays for heating. Such cosmic rays (E $\\la$100~MeV) remain, however, invisible in the radio window, since even in the strong ($\\rm{B}\\approx 10 \\ldots 50~\\mu$G) magnetic field of M\\,82 they produce synchrotron radiation at frequencies below about 10~MHz. Soft X--rays are also seen far out of the plane of M\\,82 (e.g. Bregman et\\,al. 1995). These circumstances \\changed{altogether} provide the likely heating sources for any molecular material transported out of the disk of M\\,82 into the halo, rendering the CO ''overluminous''. This is also indicated by the low conversion ratio of CO line intensity to molecular gas column density, X$_{\\rm CO} = \\rm N_{\\rm H_2}/\\rm I_{\\rm CO}$ derived by Smith et\\,al. (1991). The existence of a soft X--ray and a low--frequency radio halo around M\\,82 readily explains the larger extent of the CO emission as compared to the cold dust.  \\begin{figure} \\begin{center} \\psfig{file=h1816f1b.eps,width=8.5cm,angle=270}  \\vspace{1cm}  \\psfig{file=h1816f5b.eps,width=8.5cm,angle=270} \\caption{Comparison of the $\\lambda$1.2~mm continuum and the CO(2$-$1) line emission. The upper image shows the CO--corrected $\\lambda$1.2~mm continuum map of M\\,82. Contours start at 2$\\sigma$ (10~mJy/beam). The lower image shows the integrated CO(2$-$1) line intensity. Contours start at 3$\\sigma$ (20~K\\,km\\,s$^{-1}$/beam). The beam size of the IRAM 30-m telescope is given in the lower left corner of each image.} \\label{fig:vergleich} \\end{center} \\end{figure}"
    },
    "0004/hep-ph0004102_arXiv.txt": {
        "abstract": "Ultra-high-energy, $>10^{19}$~eV, cosmic-ray and high energy, $\\sim10^{14}$~eV, neutrino production in GRBs is discussed in the light of recent GRB and cosmic-ray observations. Emphasis is put on model predictions that can be tested with operating and planned cosmic-ray and neutrino detectors, and on the prospects of testing for neutrino properties. ",
        "introduction": "The origin of GRBs, bursts of 0.1 MeV---1 MeV photons lasting for a few seconds, remained unknown for over 20 years, primarily because GRBs were not detected prior to 1997 at wave-bands other than $\\gamma$-rays \\cite{Fishman}. The isotropic distribution of bursts over the sky suggested that GRB sources lie at cosmological distances, and general phenomenological considerations were used to argue that the bursts are produced by the dissipation of the kinetic energy of a relativistic expanding fireball (see [2] for review).  Adopting the cosmological fireball hypothesis, it was shown that the physical conditions in the fireball dissipation region allow Fermi acceleration of protons to energy $>10^{20}{\\rm eV}$ \\cite{W95a,VMU}, and that the average rate at which energy is emitted as $\\gamma$-rays by GRBs is comparable to the energy generation rate of UHECRs in a model where UHECRs are produced by a cosmological distribution of sources \\cite{W95a}. Based on these two facts, it was suggested that GRBs and UHECRs have a common origin (see [5] for a recent review).  In the last two years, afterglows of GRBs have been discovered in X-ray, optical, and radio wave bands \\cite{AG_obs_review}. Afterglow observations confirmed the cosmological origin of the bursts, through the redshift determination of several GRB host-galaxies, and confirmed standard model predictions of afterglows that result from the collision of an expanding fireball with its surrounding medium \\cite{AG_th_review}. These observations therefore provide strong support for the GRB model of UHECR production.  In this review, UHECR and neutrino production in GRBs is discussed in the light of recent GRB and UHECR observations. The fireball model is briefly described in \\S2.1, and proton acceleration in GRB fireballs is discussed in \\S2.2. Implications of recent afterglow observations to high energy particle production are discussed in \\S3. Model predictions are shown to be consistent with the observed UHECR spectrum in \\S4. Predictions of the GRB model for UHECR production, that can be tested with future UHECR experiments, are discussed in \\S5. High energy neutrino production in fireballs and its implications for future high energy neutrino detectors are discussed in \\S6. ",
        "conclusions": ""
    },
    "0004/astro-ph0004092_arXiv.txt": {
        "abstract": "Numerous methods for finding clusters  at moderate to high redshifts have   been proposed in recent   years,  at wavelengths ranging from radio  to X-rays.   In  this  paper we  describe  a new  method  for detecting clusters in two-band  optical/near-IR  imaging data.   The method relies upon the observation  that all  rich clusters, at  all redshifts observed     so far, appear   to have   a red  sequence of early-type galaxies.  The emerging picture is that all rich clusters contain a core population of  passively evolving elliptical galaxies which  are coeval and formed  at high redshifts. The proposed search method exploits this strong empirical fact by using the red sequence as a direct indicator of  overdensity.  The fundamental advantage of this approach is that  with appropriate filters, cluster  elliptical galaxies at a given redshift are  redder than all normal galaxies at lower  redshifts.  A simple color cut  thus virtually eliminates all foreground contamination,  even  at significant redshifts.   In this paper,   one of   a  series of    two, we   describe the  underlying assumptions and   basic  techniques of  the method   in  detail, and contrast the method with those used by other  authors.  We provide a brief demonstration of  the effectiveness of  the  technique using a real photometric sample with redshift  data, and from this  conclude that the method offers a powerful yet simple  way of identify galaxy clusters. We find that the method  can reliably detect structures to masses as   small as   groups  with velocity   dispersions  of  only $\\sim300{\\rm\\thinspace    km}{\\rm\\thinspace   s}^{-1}$, with redshifts  for all detected  structures estimated  to an accuray  of $\\sim$10\\%. ",
        "introduction": "The detection  and characterization  of rich  clusters  of galaxies from low ($z\\sim0$)  to  high ($z\\gtrsim1.0$)  redshifts provides  a crucial     test     of     both    cosmological    models (e.g., \\markcite{carlberg1}Carlberg         et       al.                1996; \\markcite{fan}Fan, Bahcall, \\& Cen 1997; \\markcite{oukbir}Oukbir, Bartlett \\&  Blanchard  1997; \\markcite{eke}Eke et al. 1998; \\markcite{borgani}Borgani et al. 1999; \\markcite{ledlow}Ledlow et al.   1999; \\markcite{reichart}Reichart et al. 1999; \\markcite{viana}Viana \\& Liddle 1999; \\markcite{henry4}Henry 2000, to list a recent subset of an extensive lterature) and galaxy evolution (see \\S3 for an extensive discussion). Clusters trace  structure   in    the   Universe   to  large    scales (e.g., \\markcite{bahcall2}Bahcall \\& Soneira 1983; \\markcite{gramann}Gramann  et  al.      1995; \\markcite{tadros}Tadros, Efstathiou, \\&  Dalton 1998; \\markcite{borgani2}Borgani, Plionis \\& Kolokotronis 1999 ), and provide many  examples of both strong and  weak   lensing (see \\markcite{mellier}Mellier 1999 for a recent review). The obvious importance of galaxy clusters  has led to numerous galaxy  cluster surveys in recent years, both optically (e.g., \\markcite{gho}Gunn, Hoessel \\& Oke 1986; \\markcite{abell89}Abell, Corwin \\& Olowin 1989; \\markcite{couch}Couch et al. 1991 \\markcite{lumsden}Lumsden et al. 1992; \\markcite{lidman}Lidman \\& Peterson 1996; \\markcite{post}Postman et al. 1996; \\markcite{dalton}Dalton et al. 1997; \\markcite{zaritsky}Zaritsky et al. 1997; \\markcite{ostrander}Ostrander et al. 1998; \\markcite{scodeggio}Scodeggio et al. 1999; \\markcite{gal}Gal et al. 2000 ) and with X-rays (\\markcite{gioia}Gioia \\& Luppino 1994; \\markcite{scharf1}Scharf et al. 1997; \\markcite{rosati}Rosati  et al. 1998; \\markcite{vikh}Vikhlinin et al. 1998; \\markcite{degrandi}de Grandi et al. 1999; \\markcite{romer}Romer et al. 2000).  As pointed out by \\markcite{kepner}  Kepner  et al.  (1999),  each  of these surveys has defined its  own  detection algorithm,  considered  by  the individual authors to  be appropriate to  their particular data. In essence, each of these detection algorithms makes assumptions about ``what a cluster looks like''  and then searches the  relevant data  to identify likely cluster locations.  \\paragraph{} All cluster finding methods suffer from selection effects. At minimum, each will be biased against any clusters which do not fit the particular cluster definition used. Even if cluster samples appear essentially congruent at different wavelengths, scatter in the individual properties of individual clusters will still engender somewhat different samples if surveys are done at different wavelengths. For example, results from the CNOC1 survey show that the optical richness (important for optically selected catalogs) and X-ray luminosity (important for X-ray selected catalogs) of a sample of 16 X-ray selected clusters at $0.1<z<0.55$ are only moderately correlated (\\markcite{cnoc1global}Yee et al. 2000a).  X-ray surveys, which ultimately define clusters as diffuse X-ray sources, are potentially biased against gas-poor clusters, and clusters in which the gas distribution is compact and hence unresolved.  There is some limited evidence for optically rich and massive, but X-ray weak clusters.  For example, \\markcite{bower1}Bower et al. (1994) found that a sample of 14 optically selected rich clusters at $z\\sim0.4$ showed relatively weak X-ray emission, though the significance of this result is dependent on the interpretation of masses from minimal spectroscopic data. \\markcite{castander}Castander et al.  (1994) found similar results for a sample of optically selected clusters at $z\\sim0.8$. Optical surveys are potentially biased against optically dark clusters -- deep cluster-sized potential wells which, for whatever reason, have failed to form a significant galaxy population.  Notably, there is little evidence for many such clusters.  The X-ray selected sample of \\markcite{vikh} Vikhlinin et al. (1998) showed no evidence for any optically dark clusters, with a sample of $\\sim$ 200 X-ray selected clusters from $z=0.015$ to $z>0.5$. A similar result is found in the Bright SHARC sample (\\markcite{romer}Romer et al. 2000). At higher redshift, the cluster AXJ2019+112 at $z=1.01$, claimed by \\markcite{hattori}Hattori et al.  (1997) to be a dark cluster, has been shown by \\markcite{benitez} Ben\\'{\\i}tez et al.  (1999) to not be dark at all. Clusters which are under-luminous in galaxian light may exist, but appear to be the exception rather than the norm. Another low-mass cluster type likely missed by optical techniques are the so-called fossil groups (\\markcite{ponman2}Ponman et al. 1994; \\markcite{vihk2}Vikhlinin et al. 1999; \\markcite{romer}Romer et al. 2000) which appear to be relatively common though not massive (\\markcite{vihk2}Vikhlinin et al. 1999).  \\paragraph{} More complicated redshift dependent biases and uncertainties also exist in most cluster catalogs, quite apart from the dominant cosmological effects of redshift.  For example, optical-IR selected samples are subject to the effects of cluster and field galaxy evolution, which are, in general, not well determined at $z>1$. Note, however, that the evolution of the {\\it elliptical} galaxy population in clusters has been shown by numerous authors (e.g., \\markcite{stanford98} Stanford, Eisenhardt, \\& Dickinson 1998 and references therein, hereafter SED98) to be remarkably simple and homogeneous.  The apparent stability and smooth evolution of this dominant cluster population implies the selection functions for clusters defined by these galaxies can be computed with some confidence, even at $z>1$.  For a given cluster finding technique to be most useful, the reliable computation of selection functions is as critical as the sensitivity of the algorithm (\\markcite{post}Postman et al. 1996).  \\paragraph{} The general selection functions for clusters in most surveys are typically quite complicated functions of numerous parameters, including, but not limited to, the cluster redshift and the survey flux limits.  However, in most cases, the detection rate for clusters of a given parameter set can generally be computed directly from the survey data, by inserting and attempting to recover fiducial test clusters. So long as the test clusters are a realistic representation of the cluster populations, this procedure will produce reliable results, and such techniques have been used successfully in several cluster surveys (e.g., \\markcite{post}Postman et al. 1996; \\markcite{ostrander}Ostrander et al. 1998; \\markcite{bramel}Bramel, Nichol \\& Pope 2000). False positive detections are however more difficult to quantify, as a thorough understanding of such false cluster detections requires some $a$ $priori$ knowledge of the contaminants -- such as projection in optical catalogs (e.g., \\markcite{post}Postman et al. 1996) or contamination by active galactic nuclei in X-ray selected catalogs (e.g., \\markcite{vikh} Vikhlinin et al. 1998).  Current attempts to quantify the false positive contaminations rates in samples in the absence of detailed follow-up data are typically based on some form of resampling strategy (e.g., \\markcite{bramel}Bramel, Nichol \\& Pope 2000).  \\paragraph{} Indeed, the dominant cause of false-positive cluster candidates in optical cluster catalogs is projection.  It has long been suggested that there are significant projection effects in the Abell catalog (e.g., \\markcite{lucey}Lucey 1983), though the exact degree of contamination is a matter of some debate (Collins et al. 1995). Even more modern cluster catalogs suffer significant false cluster detections due to random projections (\\markcite{post}Postman et al. 1996; \\markcite{oke1}Oke, Postman, \\& Lubin 1998; \\markcite{holden}Holden et al.  1999).  The projection-induced contamination in the Abell catalog has been studied with exhaustive n-body simulations, which show this to be primarily an artifact of the large aperture used to initially define the Abell catalog (\\markcite{vanhaarlem}van Haarlem, Frenk, \\& White 1997). These n-body results demonstrate that the false-positive rates due to projection for simulated cluster catalogs detected by either X-rays or 2-D galaxy overdensities is in fact quite similar, provided a similar detection aperture is used (\\markcite{vanhaarlem}van Haarlem, Frenk, \\& White 1997; also see \\markcite{omaryee}L\\'{o}pez-Cruz \\& Yee 2000a, hereafter LCY00).  \\paragraph{} Given the extreme importance of galaxy clusters for a variety of studies, it is clear that large, well-understood samples of clusters are needed. In such a setting, it desirable to have samples which are selected in as many ways as possible, at as many wavelengths as possible, as each sample will probe a somewhat different population with somewhat different selection effects.  Furthermore, the cluster finding algorithms used should be as efficient as possible in terms of the required data, since clusters are rare objects and thus only found in large numbers in observationally expensive wide-field surveys. The method proposed below addresses these requirements. The method, which we term the Cluster-Red-Sequence (CRS) method, exploits the observational fact that the bulk of the early-type galaxies in all rich clusters lie along a linear color-magnitude relation. This relation, referred to hereafter as the red sequence, has been shown to have an extremely small scatter (e.g., \\markcite{bowera} \\markcite{bowerb} Bower, Lucey, \\& Ellis 1992b) and appears to be extremely homogeneous from cluster to cluster  (LCY00).  In the  CRS method  clusters  are  detected  as over-densities  in  projected  angular position,  color  and magnitude simultaneously. The  color  of  the  red sequence provides   a precise redshift estimator for  the detected  clusters.  The color  constraint makes the CRS method  extremely insensitive to projection effects,  as random projections do not exhibit the necessary red sequence signature in the color-magnitude plane.  This makes the CRS method fundamentally different from previous optical cluster finding algorithms. The concept of using the red sequence as a cluster marker, or to delineate cluster superstructures, has some history in the literature (e.g., \\markcite{omarthesis}Lopez-Cruz 1997; \\markcite{cfht1}Gladders \\& Yee 1998; \\markcite{nick3}Kaiser et al. 1998; \\markcite{pasedena}Yee, Gladders, \\& L\\'{o}pez-Cruz 1999; \\markcite{lubin3}Lubin et al. 2000; \\markcite{omaryee}LCY00), and other cluster-finding techniques have been suggested which exploit the dominance of early-type galaxies in clusters in some way (\\markcite{ostrander}Ostrander et al. 1998). However, the critical point of the technique suggested here is the recognition that the cluster red sequence can be effectively isolated from survey data with only two filters, so long as the filter pair samples the 4000\\AA\\thinspace break. This technique precludes the need for many-band imaging, such as is required for full photometric redshift analyses (\\markcite{kodama2}Kodama, Bell \\& Bower 1999) and so is efficient while still exploiting the available color information. This paper represents the first effort to cast the CRS method as a well-defined, well-motivated algorithm. The purpose is to describe an algorithm which is amenable to stringent testing, and useful for finding clusters in a wide range of optical-IR imaging data.  \\paragraph{} We have begun a large observational program, the Red-Sequence Cluster Survey (RCS), to identify and characterize clusters using the CRS method.  This paper, the first of a series of two, describes the basic observational motivation for the CRS method, and provides a detailed algorithm and a test of that algorithm using existing data with spectroscopic redshifts.  The second paper in the series discusses the CRS method in the specific context of the RCS dataset, focussing on a detailed derivation of the RCS selection functions for clusters using extensive simulations.   \\paragraph{} This paper is arranged as follows.  We describe and motivate the CRS method in the context of previous work and current observational constraints in \\S2 and \\S3.  A detailed implementation of the method is given in \\S4.  In \\S5 we provide a simple but powerful test of the method, using real redshift data from the CNOC2 Redshift Survey (\\markcite{yeecnoc2}Yee et al.  2000b).  Section 6 provides a brief discussion of possible refinements to the basic method outlined here. We summarize our findings in  \\S7.  A cosmology of $H_0$=70 \\hbox{${\\rm\\thinspace km}{\\rm\\thinspace s}^{-1}{\\rm\\thinspace Mpc }^{-1}$}, $\\Omega_M$=0.2 and $\\Omega_\\lambda$=0.0 is used throughout. ",
        "conclusions": "Numerous large studies to $z\\sim1$ have now demonstrated that the red sequence of early-type galaxies is a universal and homogeneous feature of galaxy clusters. Based on this strong observational fact, we have developed a new technique, the CRS method, for locating clusters in 2-filter optical or infrared imaging data. Using appropriate filters, the CRS method circumvents problems with foreground, and to some extent background, projection effects which have been so problematic in previous optical cluster surveys. A specific algorithm has been constructed to implement this general method. This algorithm constructs 2-D smoothed density maps of color slices corresponding to redshifts bin, which are then normalized and reassembled into a 3-D density data-cube. Peaks in this data-cube are then defined as cluster candidates. Using data from the CNOC2 survey we have tested the CRS method over a significant redshift range. From this test, and the general data on the evolution of cluster early-type galaxies available in the literature, we conclude the following: \\begin{quote}  1) The apparent universality of the red sequence in available data, even to redshifts greater than one, indicates that the CRS method should be applicable to the entire redshift range accessible to CCDs ($0<z<1.4$).  2) The CRS method is a powerful new approach to finding galaxy clusters in imaging data at optical and infrared wavelengths. The performance of the method shows that it is able to reliably detect rather poor bound structures ($\\sigma_v \\leq 400$ km~s$^{-1}$) with a very low contamination rate of less than 5\\%, using data of only moderate depth.  3) The redshift accuracy of the method is $\\Delta z\\sim 0.025$ over the redshift range $0.1<z<0.5$, or equivalently 10\\% in redshift. This accuracy is comparable to that achieved for single galaxies using four (or more) filter photometric redshift methods.  \\end{quote}  We have begun a large survey, the RCS, designed to exploit the power of the CRS method (e.g., \\markcite{mex1}Gladders \\& Yee 2000a). The primary survey data are 100  square degrees of imaging in two filters ($R_c$ and $z'$), to the depth necessary to find clusters to $z\\sim1.4$. Paper II of this series (Gladders \\& Yee 2000b) discusses the application of the CRS method to the RCS data specifically, and further explores the applicability and completeness of the method to redshifts as high as $z\\sim1.4$. Unlike this paper, the main results of paper II are based on a broad and comprehensive set of tests performed on simulated data, created using detailed and realistic galaxy and cluster models."
    },
    "0004/astro-ph0004167_arXiv.txt": {
        "abstract": "The good energy resolution (3--4 keV FWHM) of the Transient Gamma Ray Spectrometer (TGRS) on board the {\\em WIND\\/} spacecraft makes it sensitive to Doppler-shifted outbursts of 511 keV electron-positron annihilation radiation, the reason being that the Doppler shift causes the cosmic line to be slightly offset from a strong instrumental background 511 keV line at rest, which is ubiquitous in space environments.  Such a cosmic line (blueshifted) is predicted to arise in classical novae due to the annihilation of positrons from $\\beta$-decay on a timescale of a few hours in an expanding envelope.  A further advantage of TGRS --- its broad field of view, containing the entire southern ecliptic hemisphere --- has enabled us to make a virtually complete and unbiased 3--year search for classical novae at distances up to $\\simeq 1$ kpc.  We present negative results of this search, and estimate its implications for the highly-uncertain Galactic classical nova rate and for future space missions. ",
        "introduction": "Classical novae are rather frequently observed in our Galaxy (Liller \\& Mayer 1987, Shafter 1997), and have also been studied in external galaxies ; typically $N \\sim$ 3--4 per year are detected in our Galaxy (Duerbeck 1995, Warner 1995).  Most of the discoveries and observations of Galactic novae have been made by amateur astronomers with little access to spectroscopic and photometric equipment.  Sky coverage has been episodic and extremely hard to calculate.  Classification attempts have also been hindered.  As a result, many of the most basic properties involving their global rate and distribution are surprisingly uncertain. For example, a number of arguments suggest that the Galactic rate of novae must be much higher than $N$:  (a) The typical limiting apparent magnitude obtainable with amateur apparatus and methods has been increasing steadily in recent years, but for the period covered by this paper may be taken to be $m_{V} \\sim 8$, within a very wide range, and with extremely uneven coverage.  Application of the expanding-photosphere method to a subset of relatively nearby and bright novae has yielded the empirical relation \\begin{eqnarray} M_{V} & = & 2.41 \\log t_{2} -10.7 ~~~\\mbox{for 5 d $< t_{2} < 50$ d} \\\\ & = & -9 ~~~\\mbox{for $t_{2} \\le 5$ d} \\\\ & = & -6.6 ~~~ \\mbox {for $t_{2} > 50$ d} \\end{eqnarray} (Warner 1995) for the absolute magnitude, where $t_{2}$ (the {\\em speed class\\/}) is the time taken for $m_{V}$ to increase by 2 from discovery.  It follows that the distance out to which amateur astronomers are detecting typical novae is $\\sim 10$ kpc, or only about one-half the volume of the Galaxy.  Furthermore, the rate of discoveries at the faintest magnitudes ($m_{V} > 6$) is greater than what would be extrapolated from brighter novae.  This indicates that a new population --- presumably associated with the Galactic bulge rather than the disk --- is present and poorly sampled (Duerbeck 1990; see below).  (b) Even within that part of the Galaxy which is effectively searched for novae, the discovery rate is blatantly incomplete.  Not only does the discovery rate for novae with $3 < m_{V} \\leq 6$ fall below the extrapolated rate for brighter events (thus, in contrast to the preceding argument, suggesting that many events in this range are missed: Duerbeck 1990), but there is a marked deficiency of discoveries in the southern celestial hemisphere (Warner 1995).  This is relevant to our work, since the TGRS detector is permanently pointed at the southern sky (\\S 2.1).  During its period of operation (1995--1997) five novae were discovered in the southern hemisphere (Harris et al. 1999, hereafter Paper I), but there is no way of knowing how many were missed.\\footnote{ The discovery process is particularly liable to miss \"fast\" novae ($t_{2} \\simeq$ a few days) which rise and fall in between successive visits to a given location.}  The possibility of detecting undiscovered novae as bright as $m_{V} = 3$ (marginally within TGRS's capabilities) is one of the justifications for the present work.  (c) In Galactic latitude, the distribution of classical novae is somewhat concentrated toward the equatorial plane (scale heights for disk and bulge populations 125 and 500 pc respectively: Duerbeck 1984, 1990).  They must therefore be affected to some degree by interstellar extinction, and a deficiency of discoveries close to the plane is indeed observed (Warner 1995).  In terms of the composition of their ejecta, novae are classified into CO-rich and ONe-rich; it is thought that the distinction reflects the composition of the underlying white dwarf material, with the ONe class coming from more massive progenitors whose cores burned beyond the early He-burning stage which yields C and O.  Different levels of positron annihilation line flux are expected from each class (\\S 4).  If the progenitors of the ONe subclass are really more massive, they will tend to lie closer to the Galactic plane, and the resulting novae will be more strongly affected by extinction and relatively under-represented in the discovered sample (of which they compose $\\sim 1/3$: Gehrz et al. 1998).  Evidence of this has been detected by Della Valle et al. (1992).  (d) The three preceding factors would all tend to enhance the true Galactic nova rate above that observed.  However, a second, quite distinct approach to the problem tends to produce systematically lower rates. In this approach, several external galaxies (particularly the Magellanic Clouds, M31 and M33) have been monitored for novae, and their observed rates extrapolated in some fashion to the Milky Way (Ciardullo et al. 1987, Della Valle \\& Livio 1994). The usual basis for extrapolation is absolute blue luminosity (Della Valle \\& Claudi 1990).  As can be seen in Table 1, the results from this approach are systematically smaller than attempts to correct for the missing Galactic novae directly.  The original explanation for this effect was provided by Duerbeck (1990), who postulated two different classes of event by spatial distribution --- disk and bulge novae.  It was claimed that the bulge population has a systematically slower speed class, and is therefore generally less luminous by Equations (1--3), which might account for the discrepancy, given a larger bulge in the main external source of novae, M31.  As will be seen (\\S 4.1), our search method is probably relevant only to a disk population.  A third approach to the problem is theoretically possible, by which classical nova outbursts are assumed to be part of a life-cycle of which other cataclysmic variables are manifestations.  The Galactic nova rate is then derived from the assumed space densities of these related objects, together with some model for the outburst recurrence time (Warner 1995).  This approach is more reliable at predicting the Galactic space density rather than the global rate, which is more directly related to the measurements we shall present.  It is important to correct for and combine these various factors into an overall global Galactic nova rate, which would govern the input of novae into Galactic chemical evolution, dust grains and interstellar radioactivity (Gehrz et al. 1998).  However attempts to do so have yielded wildly discordant results, ranging from 11--260 novae yr$^{-1}$ (see Table 1).  We have therefore adopted in this work yet a fourth (and simplest) approach which is to make an {\\em unbiased\\/} search for novae in our Galaxy. The detection of $\\gamma$-ray lines from radioactive decays of the nucleosynthesis products produced in novae is such an approach; these decays in general emit positrons, whose annihilation with electrons produces a line at 511 keV.  An obvious advantage of this approach is the very small absorption of $\\gamma$-rays in the Galaxy.  We will also see that problems of uneven coverage and sensitivity are minimal.  These advantages are realised when the $\\gamma$-ray detector TGRS, on board the {\\em WIND\\/} mission, is used (\\S 2.1).  In Paper I we determined that TGRS does indeed have the capability to perform a sky survey for classical novae.  The target of Paper I was to detect the positron annihilation line in five known novae; although none was detected, the viability of such a method was established.  The key to the method (see \\S 2 below) is that the line arises in nova material expanding towards the observer, and is therefore broadened and blueshifted (Leising \\& Clayton 1987).  Its peak is therefore shifted away from a strong background line at exactly 511 keV, which arises in the instrument itself from decays of unstable nuclei produced by cosmic ray spallation.  In the next section we give a brief description of the detector and data, and of our analysis.  None of these is substantially different from that of Paper I, where the reader may find a more detailed description. ",
        "conclusions": "\\subsection{Limit on the Galactic nova rate}  Recent developments in the theory of nucleosynthesis in classical novae (Hernanz et al. 1999) have been discouraging for our purpose of a positron annihilation $\\gamma$-ray search, since new measurements of nuclear reaction rates have led to much lower predictions of the flux in this line after 6 and 12 hr.  The discussion in Paper I of the capability of constraining the global Galactic nova rate using our present results was therefore over-optimistic.  Nevertheless, we will discuss the application of our method in general terms, so that even though important constraints cannot now be derived, it may be useful for more sensitive future experiments (e.g. {\\em INTEGRAL\\/}) or for more optimistic theoretical predictions.  A formal expression for the number of novae detectable by TGRS is \\begin{eqnarray} N_{obs} & = & R_{gal} ~T_{tot} ~\\int_{0}^{1.4 M_{\\sun}} ~\\Phi(M) \\int ~f(\\phi > \\phi_{min}) ~w(\\phi > \\phi_{min}) ~d\\phi_{min} ~dM \\end{eqnarray} where $R_{gal}$ is the Galactic nova rate; $\\phi_{min}$ is a given (time varying) threshold flux for detection by TGRS; $f(\\phi > \\phi_{min})$ is the fraction of the mass of the Galaxy within TGRS's detection radius $r_{d}$ and $r_{d}  = \\sqrt{\\phi_{pred}(M)/\\phi_{min}}$; $T_{tot}$ is the total TGRS live time; $w(\\phi > \\phi_{min})$ is the fraction of TGRS live time for which $\\phi > \\phi_{min}$; $\\Phi$ is the distribution of white dwarf masses in classical novae; and $\\phi_{pred}(M)$ is the predicted 511 keV line flux at 1 kpc for mass $M$.  The white dwarf mass distribution in novae, $\\Phi(M)$, is very poorly known.  Whereas field white dwarf masses appear to peak at $\\simeq 0.6 M_{\\sun}$ and to decline in number for higher masses up to the Chandrasekhar limit $1.4 M_{\\sun}$ (Warner 1990), the mass distribution in nova systems must be weighted towards higher masses.  This is because the thermonuclear runaway occurs when the basal pressure of the material accreted onto the white dwarf exceeds some critical value.  The critical pressure is proportional to the -4 power of the white dwarf radius, to the white dwarf mass, and to the accreted mass.  Since white dwarf radii decrease with increasing white dwarf mass, the accreted mass necessary to reach critical pressure is a strongly decreasing function of white dwarf mass.  If the accretion rate from the secondary star is roughly independent of white dwarf mass, it follows that explosions on more massive white dwarfs will recur after much shorter intervals (Gehrz et al. 1998). There have not been reliable measurements of this effect, although theory indicates that the ratio of ONe:CO novae of 1:2 is compatible with a distribution peaking at about $1.2 M_{\\sun}$ (Truran \\& Livio 1986).  Further, the mass ranges corresponding to the CO and ONe compositions are poorly known and may well overlap (Livio \\& Truran 1994).  Theoretical predictions of 511 keV line emission are only available for a few values of $M$.  In Table 2 we show the parameters of the most recent models suitable for use in Eq. (4) (Hernanz et al. 1999).  Earlier models suggest that emission from lower-mass CO white dwarf events is considerably less (G\\'{o}mez-Gomar et al. 1998). In view of the remarks above about the ONe:CO ratio, we will make the crude assumption that the ratio of \"low mass\" CO objects to \"high mass\" CO objects to ONe objects is 1:1:1, where \"high mass\" CO objects have the properties given in Table 2 and \"low mass\" CO objects are assumed to produce no 511 keV line emission at all.  This eliminates the integral over {\\em M\\/} in Eq. (4).  The remaining integral in Eq. (4) can be approximated by the value of the integrand when $\\phi_{min}$ has its mean value --- this follows from our result in \\S\\S 2.1, 3 that the variation of live times in our sample (and therefore of the errors in Fig. 2) is very small. For a given model in Table 2, therefore, taking $\\phi_{min} = 4.6 \\times 8.2 \\times 10^{-4} = 3.8 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$, the problem is reduced to the computation of the fraction $f$ of Galactic mass which lies within the radius $r_{d} = \\sqrt{\\phi_{pred}/3.8 \\times 10^{-3}}$ kpc.  As an example, let us consider the Hernanz et al. (1999) model of a $1.15 M_{\\sun}$ CO nova from Table 2.  Here $\\phi_{pred} = 3.1 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$, so that $r_{d} = 0.9$ kpc.  Within this value of $r_{d}$ we determined that 0.61\\% of the Galaxy's mass resides, according to the widely used Bahcall-Soneira model of the Galaxy (Bahcall \\& Soneira 1984): one-half of this value (i.e. the southern hemisphere) gives $f = 0.00305$.  For the $1.15 M_{\\sun}$ CO model, Eq. (1) then reduces to $R_{gal} =\\frac{N_{obs}}{T_{tot} ~f}$.  Our upper limit for $N_{obs}$ is $<1$, with 63\\% probability (for Poisson-distributed events: Gehrels 1986), and the live time is $7.7 \\times 10^{7}$ s (\\S 2.1), giving us a 63\\% upper limit on the rate of \"high-mass\" CO novae of $R_{gal} <134$ yr$^{-1}$. This is quite close to the exact value 123 yr$^{-1}$ obtained by explicitly integrating Eq. (4) over $\\phi_{min}$ (Table 2).  In the same way we obtain an upper limit of 238 yr$^{-1}$ on the rate of novae occurring on ONe white dwarfs from Hernanz et al.'s (1999) prediction. Our best result comes from the CO model in Table 2, which we have assumed to be one-third of the total, from which we derive a global Galactic nova rate of $< 369$ yr$^{-1}$.  Uncertainties in this value clearly arise from uncertainties in the nova models, in the fraction of white dwarfs in nova systems of each type, in the Bahcall-Soneira model, and in the possibility of distinct spheroid and disk nova populations having differing rates, since our typical detection radius $< 1$ kpc includes almost none of the Bahcall-Soneira spheroid.  Since our results do not significantly constrain previous measurements of the nova rate (Table 1), we do not make estimates of these errors, which will require attention from other, more sensitive experiments (see next section).  \\subsection{Implications for nova detection by other instruments}  An attempt has been made, using the BATSE instrument on the {\\em Compton\\/} Observatory, to detect 511 keV line emission from a recent nearby nova (V382 Vel) by a similar method to that used here and in Paper I (Hernanz et al. 2000).  The advantage of observing with BATSE over TGRS is its much larger effective area.  Its disadvantages are much poorer energy resolution with a NaI spectrometer, and a background varying on very short ($<< 90$ min) timescales.  The sensitivities achieved are comparable to those obtained here.  Degradation of the Ge detector (\\S 2.1) prevented TGRS from achieving comparable sensitivity on V382 Vel (Harris et al. 2000), so future efforts in this field will rely on BATSE and on the {\\em INTEGRAL\\/} mission, which is scheduled for launch in 2001 September carrying a Ge spectrometer (SPI) with resolution comparable to TGRS but a much larger effective area.  Hernanz et al. (1999) estimated that SPI could detect the model novae of Table 2 out to $\\sim 3$ kpc. However they also pointed out that the short duration of the 511 keV line emission would make it difficult for {\\em INTEGRAL\\/} to slew to a candidate event.  Thus the detection rate would be limited to novae within the SPI field of view, which is $\\sim 25 \\deg$ FWHM.  The search method which we have used, i.e. an {\\em ex post facto\\/} search in background spectra, ought to be perfectly feasible with {\\em INTEGRAL\\/}. The chief requirements for this method are very high energy resolution and a sufficiently low and stable background.  While the SPI detector has excellent resolution, the background level in it has not yet been rigorously computed.  Nevertheless, qualitative arguments suggest that the background will be no worse than that in TGRS.  Like {\\em WIND\\/}, {\\em INTEGRAL\\/} will be in a high-altitude elliptical orbit which avoids extensive exposure to Earth's trapped radiation belts and to albedo $\\gamma$-rays from Earth's atmosphere.  The main disadvantages of {\\em INTEGRAL\\/} for nova detection are the small SPI field of view and the planned observing strategy which cuts down the amount of time spent pointing towards the main concentration of novae near the Galactic center.  We can make use of the planned program of {INTEGRAL\\/} observations of the central Galactic radian in the first year of operation (Winkler et al. 1999) to estimate the rate at which novae might be detected in the SPI data.  As previously, we assume that novae follow the Bahcall \\& Soneira (1984) Galactic distribution. The planned first-year {\\em INTEGRAL\\/} observations may be approximated by a $31 \\deg \\times 11 \\deg$ grid with $2 \\deg$ spacing between $-30 \\deg \\ge l \\ge 30 \\deg$ and $-10 \\deg \\ge b \\ge 10 \\deg$, the exposure to each point being 1180 s per pass, with 12 passes per year covering the whole grid.  Thus the live time for the whole grid is 0.153 yr.  From the Bahcall-Soneira model, the pointing geometry, and the SPI aperture $\\sim 25 \\deg$ we calculated that the {\\em INTEGRAL\\/} detection radius $\\sim 3$ kpc intercepts $\\sim$0.75\\% of the Galactic nova distribution.  The live time 0.153 yr is then multiplied by a typical Galactic nova rate $\\sim 50$ yr$^{-1}$, (of which 2/3 are practically detectable, as assumed in \\S 4.1), and by the intercepted fraction, to imply that {\\em INTEGRAL\\/} ought to detect 0.04 novae yr$^{-1}$. Unless theoretical estimates of the 511 keV line flux turn out to be considerably larger, the prospects for such a detection appear to be small.  The same conclusion probably applies to a different method of detecting 511 keV line emission indirectly, by observing the 170--470 keV continuum produced by Compton scattering in the nova envelope using SPI's large-area CsI shield (Jean et al. 1999)."
    },
    "0004/astro-ph0004203_arXiv.txt": {
        "abstract": "We present multiband photometric observations of three spiral galaxies selected in a sample suited for the study  of young stellar groupings and their relationship with the parent galaxy and the galactic environment.  Star forming regions have been identified using an objective technique based on a multivariate statistical analysis. Maps of young star groupings are given for each galaxy. The luminosity functions of the young star group populations show a remarkable similarity with a slope in the range $-1.52$ to $-1.33$ range. The size distributions peak around the {\\it classical} 100pc value of the Local Group associations for two out of the three galaxies.  NGC 1058 shows smaller associations (peak at $\\sim 50$pc). The total number of young groups per unit B absolute luminosity of the galaxy is significantly greater in UGC 12732. The activity of star formation is in all three galaxies clearly stronger in the central regions. ",
        "introduction": "A few recent papers have pointed out the importance of the determination of correlations among star forming region characteristics (size, luminosity, etc.)  with those of the parent galaxy to gain insights on star formation processes and modes.  For instance, Elmegreen et al.  (1994, 1996) suggest the existence of a linear scaling between the largest complex size in a galaxy and the size of the galaxy itself. Moreover, Elmegreen and Salzer (1999) find that the largest complex size in each galaxy (in their sample of 11 galaxies) scales with the square root of the total galaxy B--luminosity.  These scaling laws are  consistent with the largest complex size being equal to 1/2 of the characteristic Jeans length in the ambient gas, if the gas velocity dispersion is always the same in a marginally stable ISM.  In this picture, the time--scale of star formation is longer (about 50 Myr) for larger galaxies than for smaller ones when the size difference is a factor of 10.  This means that in large galaxies star forming regions should contain both evolved stars and young OB stars, while in smaller galaxies only OB stars are expected in such regions.  As a consequence, a correlation between the average age (or colour) of the star forming regions and galaxy luminosity is expected.  A related prediction is that star forming regions in small galaxies should be much brighter, but not much larger relative to the host galaxy than those in giant spirals. It has been clearly shown by Hodge (1986) that any astrophysical considerations deduced from the collection of heterogeneous data should be carefully handled.  Indeed, many biases are introduced by the difference in the quality of observational data as well as by some degree of subjectivity unavoidably present in any {\\it by--eye} identification of star groupings.  In order to minimize the uncertainties just described, it would be worth dealing with a homogeneous set of observational data (taken with the same telescope and equipment) and using an objective method for identifying blue star groupings and determining their main properties. Although OB associations in Local Group galaxies are important to study, it is not possible to derive significant correlations among star-forming region properties and their parent galaxy characteristics using only Local Group galaxies due to the relative scarcity of galaxies in the Local Group. For the reasons and purposes discussed above we decided to start a project which consists of: {\\it i)} the collection of multiband photometric data for a large set of galaxies beyond the Local Group; {\\it ii)} their treatment with the cluster analysis algorithm described in Adanti et al. (1994). Some preliminary results are presented in Wyder et al. (1998).  This paper is the first of a series in which we will present and discuss  our determination of the young stellar groupings (hereafter YSGs) in various spirals, chosen among almost face--on galaxies with a radial velocity limit that corresponds to a moderate distance ($\\lappreq 16$ Mpc) and small enough in angular size to be included in a single CCD frame.  A relevant by--product of this work is an accurate list of YSG targets to be investigated with higher resolution observations (e.g. HST--ACS).  We can note that the philosophy of the present project is the same followed by Bresolin et al. (1998), who analysed HST data of several nearby galaxies with an automatic technique suited just for resolved stellar fields. ",
        "conclusions": "This paper presents the first results of an ongoing project whose main aim is to collect homogeneous data suited to the study of regions of recent star formation in a variety of galaxies.  We apply a cluster analysis algorithm to obtain, in a semi--automatic  way, an objective identification of young star groupings (YSGs) in UBVRI  frames of NGC 7217, NGC 1058 and UGC 12732.  These are the first spirals we observed from a sample (selected from the RC3) of nearly face--on galaxies, with radial velocities less than 1200 km/sec and small enough to fit in just one CCD frame of the Apache Point Observatory Spicam. This sample will allow us to check in a statistically significant way whether correlations among properties of YSGs and their parent galaxies exist. Such correlations are useful to understand better how star formation in galaxies proceeds (see e.g. Elmegreen 1994, 1996).  We have identified 149, 71 and 108 YSGs in NGC 7217, NGC 1058 and UGC 12732, respectively. These numbers correspond to specific frequencies $S_N\\equiv N_t \\cdot 10^{0.4(M_B^0+15)}$ equal to $1.0$,  $3.9$, $19.7$.  If we consider just YSG brighter than the likely completeness limit ($L_B=10^5L_{B\\odot}$), specific frequencies reduce to $1.0,~3.1$ and $12.9$ for NGC 7217, NGC 1058 and UGC 12732, respectively.  In any case, the high $S_N$ value for UGC 12732 indicates an overabundance of YSGs in this galaxy; however, since distances are from radial velocities, peculiar motions may significantly affect the total luminosity and consequently the $S_N$ value. For instance, a variation of $50\\%$ in the UGC 12732 distance makes its $S_N$ value comparable to those of the other two galaxies.   As explained in Section 3, determining an accurate background for the YSGs is an inherently difficult problem since they usually lie in areas with a lot of structure. This makes the colours of the YSGs unreliable but the the luminosities can still be used to determine a statistically reliable luminosity function. We found that the slope of the bright tails of the YSG  luminosity functions in these three galaxies are similar, within the errors,  and lie in the range $-1.36 to -1.52$. In Section 4 we also present the size distribution and surface density of the YSGs versus galactocentric radius: no significant variation is found. From the radial distribution of our identified YSGs and their integrated luminosities we could extract information about differences in star formation rate within each galaxy and between galaxies. We found a much steeper radial slope for the SFR in NGC 1058  while in both NGC 7217 and UGC 12732 the variation of the SFR between the inner and outer galactic regions is up to a factor of about 10.  The YSG size distributions (which are of course affected  by the poorly known galactic distances)    are similar in NGC 7217 and UGC 12732 (peaks around 100 pc), while NGC 1058 shows smaller YSGs (the peak is at $50$ pc). Since the latter galaxy is much  closer than NGC 7217 and UGC 12732, one could believe the smaller sizes of YSGs as an artifact of the intrinsic better spatial resolution (pc/pixels) of the images. A simple test we performed clearly proves this is not the case: the identification of YSGs on images of NGC 1058 artificially re-binned to simulate larger distances shows that the peak of the size distribution is almost insensitive on a variation up to $40\\%$ of the distance. More specifically, increasing the distance implies a loss of just the smallest objects (those smeared out by the image re-binning)."
    },
    "0004/astro-ph0004279_arXiv.txt": {
        "abstract": "We present simulations of dwarf nova outbursts taking into account realistic variations of the mass loss rate from the secondary. The mass transfer variation has been derived from 20 years of visual monitoring and from X-ray observations covering various accretion states of the discless cataclysmic variable AM\\,Herculis. We find that the outburst behaviour of a fictitious dwarf nova with the same system parameters as AM\\,Her is strongly influenced by these variations of the mass loss rate. Depending on the mass loss rate, the disc produces either long outbursts, a cycle of one long outburst followed by two short outbursts, or only short outbursts. The course of the transfer rate dominates the shape of the outbursts because the mass accreted during an outburst cycle roughly equals the mass transferred from the secondary over the outburst interval. Only for less than 10\\% of the simulated time, when the mass transfer rate is nearly constant, the disc is in a quasi-stationary state during which it periodically repeats the same cycle of outbursts. Consequently, assuming that the secondary stars in non-magnetic CV's do not differ from those in magnetic ones, our simulation indicates that probably all dwarf novae are rarely in a stationary state and are constantly adjusting to the prevailing value of the mass transfer rate from the secondary. ",
        "introduction": "Cataclysmic variable systems (CV's), in which a white dwarf accretes material from a Roche lobe filling secondary star (see Warner 1995 for an enyclopaedic review) can be broadly divided into three subclasses: in non-magnetic systems, the white dwarf accretes via an accretion disc; in the magnetic polars, the infalling matter couples directly onto the strong magnetic field of the white dwarf before it can build a disc and is funnelled to accretion region(s) near the magnetic pole(s) of the white dwarf; finally, in CV's containing a weekly magnetic white dwarf (intermediate polars), a partial disc may exist, with the mass flowing from the inner edge of the disrupted disc through magnetically funnelled accretion curtains onto the white dwarf.  Dwarf novae are non-magnetic cataclysmic variable stars which show characteristic eruptions with photometric amplitudes in the range of 2--8 magnitudes {\\footnote{Also intermediate polars are known to show outbursts in their partial discs; e.g. EX Hya, GK Per}}. These eruptions typically last a few days to several weeks and recur quasi-periodically on timescales of weeks up to many years. They are generally thought to result from thermal instabilities associated with hydrogen ionization in the accretion disc (see Cannizzo 1993a for a review and Ludwig et al. 1994 for a detailed parameter study).  An important boundary condition of the disc-limit-cycle model is the mass transfer rate from the secondary to the accretion disc. In most studies of dwarf nova outbursts this mass transfer rate is kept constant (Cannizzo 1993b; Ludwig \\& Meyer 1997; Hameury et al. 1998, 1999). Duschl and Livio (1989) were the first to examine combined mass transfer and disc outbursts, though within the framework of the mass transfer instability model where individual and short-lived mass transfer events are capable of producing single outbursts. Smak (1991, 1999) discussed mass transfer variations in the context of the superoutburst phenomenon in dwarf novae of the SU\\,UMa type and dwarf nova outbursts with enhanced mass transfer during outburst.  King \\& Cannizzo (1998) and Leach et al. (1999) tested how the accretion disc in a dwarf nova system behaves if the mass transfer from the secondary varies abruptly between different levels.  They found that these mass transfer variations produce only subtle effects on normal dwarf novae, including variations in the outburst shape, and that the dwarf novae keep on having outbursts even if the transfer rate is reduced close to zero.  In spite of these efforts we are left with the question of what the real variations of the mass tranfer rates from the secondaries in dwarf nova systems are and how they can influence the outburst behaviours of the accretion discs. Fortunately, nature provides an answer to the first question in the form of discless cataclysmic variables, the polars or AM\\,Her systems. In these systems, the mass transfer rate can be estimated directly from the observations because there is no accretion disc acting as a mass buffer. Thus, we can use the long-term light curve of an AM\\,Her system as a measure for the mass transfer variations in a fictitious but realistic dwarf nova system with the same system parameters. In the present paper, we present the results of such a numerical experiment.  In the following section, we briefly review the equations for the viscous and thermal evolution of the accretion disc, discuss the numerical methods used in our code, and compare the results of two standard models with other fine mesh computations.  Thereafter we derive the mass transfer rate in AM\\,Her as a function of time, $\\dot M_\\mathrm{tr}(t)$. Finally, we apply this mass transfer rate to a fictitious dwarf nova and discuss the effects that can be observed in the outburst behaviour. ",
        "conclusions": "There is no reason to assume that the mass transfer variations of the secondary observed in AM\\,Her are not present in non-magnetic systems.  Our numerical experiment including realistic variations of the mass transfer rate in a dwarf nova system is, therefore,  a significant step towards a better understanding of dwarf nova light curves, and, thereby, of the underlying disc limit cycle.  The light curve produced by our fictitious system switches between three states depending on the actual mass transfer rate. High transfer rates lead to only long outbursts where the entire disc is transformed into the hot state. If the transfer rate is near the average value the disc goes through a cycle of three outbursts, one long outburst followed by two short ones. Even long periods of low transfer rates do not force the disc to stop its outburst activity: long outbursts are suppressed and the duration of quiescence increases but the disc always produces short outbursts. From this follows that the low-states of VY Sculptoris stars (a subgroup of novalike variables) could not be caused by low transfer rates alone (see also Leach et al. 1999).  We find that in our fictitious system the mass accreted during an outburst cycle is dominated by the course of the mass transfer rate if the mass transfer rate varies significantly. The disc always relaxes to equilibrium with the mass input from the secondary. Thus, our experiment strongly supports King \\& Cannizzo's (1998) claim that dwarf nova accretion discs are probably never in a stationary state but are constantly adjusting to the prevailing value of $\\dot{M}_{\\mathrm{tr}}$. Only during periods where the mass transfer rate is nearly exactly constant the disc periodically repeat the quasy stationary outburst cycle. Such periods are rare but occur in AM\\,Her.  The strong influence of the mass transfer rate on the outburst behaviour of the fictitious system clearly indicates that probably most (if not all) the deviations from periodic outburst cycles seen in the light curves of dwarf novae are caused by variations of the mass transfer rate.  \\begin{figure} \\includegraphics[angle=270,width=9cm]{fig11.ps} \\caption[]{\\label{exam-4} The same as in Fig.\\,\\ref{exam-3} but assuming a more massive primary $M_{\\mathrm{wd}}=1.0\\,M_{\\mathrm{\\sun}}$ } \\end{figure}   \\begin{acknowledgement} We thank Daisaku Nagami and the referee for their helpful comments and suggestions. MRS would like to thank the Deutsche Forschungsgemeinschaft for financial support (Ma\\,1545\\,2-1). BTG thanks for support from the DLR under grant 50\\,OR\\,99\\,03\\,6. \\end{acknowledgement}"
    },
    "0004/astro-ph0004086_arXiv.txt": {
        "abstract": "The bispectrum is the lowest-order statistic sensitive to the shape of structures generated by gravitational instability and is a potentially powerful probe of galaxy biasing and the Gaussianity of primordial fluctuations.  Although the evolution of the bispectrum is well understood theoretically from non-linear perturbation theory and numerical simulations, applications to galaxy surveys require a number of issues to be addressed.  In this paper we consider the effect on the bispectrum of stochastic non-linear biasing, radial redshift distortions, non-Gaussian initial conditions, survey geometry and sampling.  We find that: 1) bias stochasticity does not affect the use of the bispectrum to recover the mean biasing relation between galaxies and mass, at least for models in which the scatter is uncorrelated at large scales. 2) radial redshift distortions do not change significantly the monopole power spectrum and bispectrum compared to their plane-parallel values.  3) survey geometry leads to finite-volume effects which must be taken into account in current surveys before comparison with theoretical predictions can be made. 4) sparse sampling and survey geometry correlate different triangles leading to a breakdown of the Gaussian likelihood approximation.  We develop a likelihood analysis using bispectrum eigenmodes, calculated by Monte Carlo realizations of mock surveys generated with second-order Lagrangian perturbation theory and checked against N-body simulations.  In a companion paper we apply these results to the analysis of the bispectrum of \\IRAS galaxies. ",
        "introduction": "It is widely accepted that the large-scale structure of the universe routinely observed in galaxy surveys is the result of gravitational instability that operates on (initially small) fluctuations in the density distribution. Essentially, initially dense (underdense) regions contract (expand) as a consequence of the attractive nature of gravity. As a result, most of the volume of the universe is in underdense regions whereas the rest contains dense structures such as clusters, galaxies, etc. This asymmetry between underdense and overdense regions implies non-Gaussianity in the galaxy distribution. Understanding the dynamics of gravitational instability thus gives us a powerful tool that can be used to extract useful information from the non-Gaussian clustering pattern of galaxies (Peebles 1980; Fry 1984; Goroff et al. 1986; Bernardeau 1992; Bouchet et al. 1995; Scoccimarro \\& Frieman 1996).  A quantitative description of galaxy clustering is obtained by applying statistical methods, in particular, multi-point correlation functions (Peebles 1980). In a Gaussian field, the statistical properties are fully characterized by its two-point correlation function or power spectrum, all higher-order (connected) correlation functions being zero. Information on non-Gaussianity is thus accessible by calculating higher-order correlations functions. An alternative to multi-point correlation functions is to consider higher-order moments (i.e. smoothed multi-point correlation functions collapsed at a single point); these characterize the one-point distribution of the galaxy distribution and can be measured with larger signal to noise than the corresponding multi-point functions. However, the latter contain more information, in particular, being evaluated at three or more points they are sensitive to the shapes of large-scale structures, which is the key to disentangle degeneracies present otherwise.  The bispectrum, the three-point correlation function in Fourier space, is the lowest order statistic sensitive to the shape of structures generated by gravitational instability. It can be used to probe the nature of primordial fluctuations (Gaussian versus non-Gaussian), galaxy biasing (i.e.  the relation between the galaxy distribution and the underlying dark matter distribution), and help to break the degeneracy between linear bias and the matter density parameter $\\Omega_m$ present in power spectrum measurements in galaxy redshift surveys (Fry 1994; Hivon et al. 1995; Matarrese, Verde \\& Heavens 1997; Scoccimarro et al. 1998).   Pioneering work on three-point statistics from the Zwicky (Peebles \\& Groth 1975) and Lick (Groth \\& Peebles 1977; Fry \\& Seldner 1982) angular catalogs showed that the bispectrum, $B(k_1,k_2,k_3)$, scaled in terms of the power spectrum, $P(k)$, as $B(k) \\propto P(k)^2$. This motivatied the so-called hierarchical model and was also in agreement with the scaling at small scales inferred from the BBGKY equations of motion assuming stable clustering and self-similarity (Davis \\& Peebles 1977; Peebles 1980). At these small scales deep into the non-linear regime, however, limited understanding of the non-linear physics including the possibility of complicated galaxy biasing prevents using higher-order statistics to extract quantitative knowledge.   It was not until somewhat later, with the systematic development of non-linear perturbation theory (hereafter PT; Fry 1984), that it was possible to derive quantitatively how non-Gaussianity is generated by gravity. In particular, the scaling $B(k) \\propto P(k)^{2}$ (and in general, the $n$-point function $\\xi_n$ scales as $\\xi_n \\propto \\xi_2^{n-1}$) is recovered at large scales, being simply a result of quadratic non-linearities in the equations of motion of gravitational instability. An important signature predicted by PT is that the bispectrum at large scales should be a rather strong function of triangle configuration, with the reduced bispectrum $Q_{123}$  \\beq Q_{123} = \\frac{B_{123}}{P_{1}P_{2}+ P_{2}P_{3}+ P_{3}P_{1}} \\label{q}, \\eeq \\noindent where $B_{123}\\equiv B(k_1,k_2,k_3)$ and $P_i\\equiv P(k_i)$, showing higher amplitude for colinear configurations than equilateral configurations, corresponding to the fact that gravity generates anisotropic large-scale structures.  These results have been confirmed since then by N-body simulations (Fry, Melott \\& Shandarin 1993; Scoccimarro et al. 1998).   These predictions hold for the density field, in order to compare with observations one must necessarily deal with the possibility of galaxy biasing. At scales larger than those relevant to galaxy formation, one can make the simplifying assumption that biasing is a local function of the underlying {\\em smoothed} density field, and thus expandable perturbatively at large scales as (Fry \\& Gazta\\~naga 1993; see also Coles 1993)  \\beq \\dt_g(\\x) = \\sum_{n=0} b_n\\ \\dt^n(\\x), \\label{locbias} \\eeq \\noindent where $b_n$ are the bias parameters, and a tilde denotes the smoothing operation.  Thus the galaxy bispectrum is calculable in terms of the density bispectrum at leading order in PT  \\beq Q_g = {Q \\over b_1} + {b_2 \\over b^2_1} \\label{Qg} \\eeq \\noindent in terms of linear ($b_1$) and non-linear ($b_2$) biasing parameters. Basically, a large linear bias decreases the dependence of the bispectrum on triangle, whereas antibias ($b_1<1$) enhances the configuration dependence of the galaxy bispectrum relative to that of the density.  Since the density bispectrum depends to a good approximation only on spectral index, comparison of the density and galaxy bispectrum can give constraints on the biasing parameters independent of $\\Omega_m$ (Fry 1994). Note that $Q_g$ and $Q$ are {\\em functions} that depend on triangle configuration; this is what allows to break the degeneracy between $b_1$ and $b_2$ present in analogous measurements of one-point moments using the skewness. Note that Eq.~(\\ref{Qg}) assumes Gaussian initial conditions, otherwise there are additional contributions (Scoccimarro 2000).   A first attempt to use this procedure in the Lick catalog data, concluded that the lack of configuration dependence on the observed bispectrum required a large bias, $b_{1} \\approx 3$, assuming the absence of systematic errors (Fry 1994). However, other effects such as projection (Fry \\& Thomas 1999; Frieman \\& Gazta\\~naga 1999; Buchalter, Kamionkowski \\& Jaffe 2000) and most importantly non-linear evolution (Scoccimarro et al.  1998) also wash out the configuration dependence of the bispectrum, so application of this idea to observational data had to await until the availability of larger galaxy surveys. Recent determination of the angular three-point correlation function from the APM survey (Frieman \\& Gazta\\~naga 1999) shows a configuration dependence consistent with APM galaxies being unbiased ($b_1\\approx 1$, $b_2\\approx 0$), but with large error bars. This result is however consistent with the determination of one-point higher-order moments in the APM (Gazta\\~naga 1994), which show remarkable agreement with the predictions of PT from Gaussian initial conditions and no bias (Juszkiewicz, Bouchet \\& Colombi 1993; Bernardeau 1994a,1994b).  Redshift surveys map the three-dimensional distribution of galaxies in a large volume, and are thus ideally suited to use the bispectrum as a probe of galaxy biasing and non-Gaussian initial conditions. In this case, the observed distribution is actually distorted by peculiar velocities along the line of sight which contribute to redshifts. However, these redshift-space distortions are now well understood, at least in the plane-parallel approximation (Hivon et al. 1995, Verde et al. 1998; Scoccimarro, Couchman \\& Frieman 1999). Determination of three-point statistics from redshift surveys has been carried out for the CfA survey and a sample of redshifts in the Pisces-Perseus supercluster (Baumgart \\& Fry 1991) and the LCRS survey (Jing \\& B\\\"orner 1998). In both cases, however, there is a rather limited range of scales in the weakly non-linear regime to obtain quantitative constraints on biasing and non-Gaussian initial conditions. On the other hand, the next generation of large redshift surveys such as 2dF and SDSS will provide excellent conditions for the determination of the bispectrum at large scales.   There are however a number of issues which are potentially important for current and future surveys and must be addressed to reliably use the bispectrum to extract useful cosmological information.  In this paper we consider effects that can potentially influence theoretical predictions of the bispectrum in redshift surveys: stochastic biasing, radial nature of redshift distortions, finite survey volume and sparse sampling.  The main result of this study is a likelihood analysis that allows extracting accurate constraints on galaxy biasing and non-Gaussian initial conditions.  In this work, we consider \\IRAS mock surveys, which map a large enough volume in the weakly non-linear regime. Application of our results to the actual data is considered in a companion paper (Scoccimarro et al. 2000).  This paper is organized as follows. In Section~2 we briefly review the behavior of the bispectrum at large scales from PT. In Section~3 we discuss the effects of galaxy biasing, including stochasticity. In Section~4 we review second-order Lagrangian PT, and discuss its regime of validity by comparing to N-body simulations. Section~5 presents results on the effects of redshift distortions in the plane-parallel and radial case. Section~6 discusses the generation of mock catalogs and the optimal weighting for higher-order statistics. In Section~7 we discuss the effects of finite volume of the survey and sampling on the bispectrum. Section~8 develops the likelihood analysis using bispectrum eigenmodes.  Finally, Section~9 offers a summary of the results and our conclusions. ",
        "conclusions": "We studied the use of the bispectrum to recover bias parameters and constrain non-Gaussian initial conditions in realistic redshift surveys. We considered the effects of stochastic non-linear biasing, radial redshift-distortions, survey geometry and sampling on the shape dependence of the bispectrum.  We found that bias stochasticity does not seem to affect the use of the bispectrum to recover the mean biasing relation between galaxies and mass, at least for models in which the scatter is uncorrelated at large scales. Clearly more work is necessary along these lines, we have just considered a very simple extension to the local deterministic bias model. We also found that the radial nature of redshift distortions changes the bispectrum monopole compared to the plane-parallel values by only a few percent, well below current error bars.  On the other hand, survey geometry leads to finite volume effects which must be taken into account in current surveys before comparison with theoretical predictions can be made.  Similarly, sparse sampling and survey geometry correlate different triangles leading to a breakdown of the Gaussian likelihood approximation.  We developed a likelihood analysis using bispectrum eigenmodes, calculated by Monte Carlo realizations of mock surveys generated with second-order Lagrangian perturbation theory (2LPT) and checked against N-body simulations.  This can be used to derive quantitative constraints from current and future redshift surveys. In a companion paper (Scoccimarro et al. 2000) we apply the results obtained here to the analysis of the bispectrum in \\IRAS surveys.  Our treatment can be improved in several ways for applications to surveys that demand more accuracy. In particular, we relied on the simple relation in Eq.~(\\ref{Qg}) to relate the galaxy bispectrum to the mass bispectrum; in practice, however, redshift-distortions and biasing do not conmute, and this relation is more complicated (Scoccimarro, Couchman \\& Frieman 1999).  Related to this issue is a formulation, so far lacking, of a joint likelihood for the power spectrum and bispectrum, which involves dependence on $\\Omega_m$ and biasing parameters. In addition, we have only considered closed triangles in Fourier space. In practice, due to breakdown of translation invariance by the survey geometry, selection function, and radial redshift distortions, open triangles also contain useful information. In order to do this, one possibility is to construct a general cubic estimator for the bispectrum, analogous to the quadratic estimator for the power spectrum (e.g. Tegmark et al. 1998), which seems a rather difficult task for non-Gaussian fields. However, future large redshift surveys such as 2dF and SDSS will be the ideal datasets to take full advantage of the cosmological information available from higher-order statistics, and thus the effort in developing more accurate statistical methods will likely be rewarded."
    },
    "0004/astro-ph0004109_arXiv.txt": {
        "abstract": "As part of a global analysis of deep star counts to constrain scenarii of galaxy formation and evolution, we investigate possible links between the galactic spheroid and the dark matter halo. A wide set of deep star counts at high and intermediate galactic latitudes is used to determine the large scale density law of the spheroid. Assuming a power density law, the exponent, flattening, local density and IMF slope of this population are estimated. The estimation is checked for robustness against contamination of star counts by the thick disc population. Contamination effects are derived from a model of population synthesis under a broad variety of thick disc parameters. The parameter fit is based on a maximum likelihood criterion. The best fit spheroid density law has a flattening of 0.76, a power index of 2.44. There is a significant degeneracy between these two parameters. The data are also compatible with a slightly less flattened spheroid (c/a = 0.85), in combination with a larger power index (2.75). A flatter spheroid (c/a = 0.6) with a power index of 2 is not excluded either. We also constrain the spheroid IMF slope $\\alpha$ to be 1.9 $\\pm$ 0.2, leading to a local density of 1.64  10$^{-4}$ stars pc$^{-3}$ and a mass density of 4.15 10$^{-5}$ \\Msun pc$^{-3}$. With this slope the expected mass density of brown dwarfs in the halo makes a negligible part of the dark matter halo, as already estimated from microlensing surveys.  So, as star count data progresses in depth and extent, the picture of the spheroid star population that comes out points to a shape quite compatible with what we know about the distribution of baryonic dark matter if it is made of stellar remnants, suggesting a common dynamical origin. ",
        "introduction": "This paper is part of a global analysis of star counts developed to constrain consistently scenarii of galaxy formation and evolution.The central tool of this approach is the \"Besan\\c{c}on\" model of population synthesis. This model is gradually tuned to fit an increasing number of observational constraints while keeping compatibility with previous fits and theoretical prescriptions. In the present paper we address the problem of the halo (dark or visible) by trying to compare the properties of the spheroid population (the visible halo) with the dark matter halo, as traced by microlensing at high galactic latitudes and by the rotation curve. If the dark matter is made at least partly of stellar remnants, as shown by recent statistics of microlensing at high galactic latitudes \\cite{Aubourg1993Natur.365..623A,Alcock1997ApJ...486..697A,1998ApJ...499L...9A}, the density trend of this matter should be close to a power law with index of 2 (as expected from a flat rotation curve). It is natural to think of a similar shape for the stellar spheroid.  Constraints on the overall shape of the dark halo are poor. Cosmological simulations of halo formation generally predict that halos are flattened by about c/a $\\sim$ 0.7 \\cite{1996IAUS..169...23R}. But the axis ratio depends on how much the halo matter is dissipative, the more dissipative, the flatter the halo. Direct determinations of the dark matter distribution in polar ring galaxies show flattened halos with c/a $\\sim$ 0.5 \\cite{sackett94,1996IAUS..169...23R}.  Concerning the spheroid population, most previous analyses suggest rather steep density slopes with power indices between 3.0 and 3.5. However, these analyses are based on rather small samples of well identified tracers. The estimated flattening also cover a wide range between 0.6 to 1.0:  The distribution of galactic globular clusters appears to be well fitted by a power law density with index $n \\sim 3.5$ and flattening of 1. \\cite{harris76,zinn85}. Hawkins RR Lyrae observations \\cite*{hawkins84}  showed $n = 3.1 \\pm 0.2$ with a flattening of 0.9. Saha \\cite*{saha85}, using a spherically symmetric model, found $n \\sim 3$ out to 25 kpc but then the RR Lyrae density falls off more rapidly beyond 25 kpc. Another study of RR Lyrae by Wetterer \\cite*{wetterer96} showed that a spherically symmetric model yields $n \\sim 3$ whereas an ellipsoidal distribution yields $n \\sim 3.5$. Sluis \\cite*{sluis98} counted blue horizontal branch (BHB) stars and RR Lyrae and found $c/a \\sim 0.5$ and $n = 3.2 \\pm 0.3$. Still from BHB star counts, Sommer-Larsen \\cite*{sommer87} derived $c/a \\sim 0.8$ and $n \\sim 3$ up to 40 kpc, Preston \\cite*{preston91} found that $c/a$ increases from 0.5 to 1 up to 20 kpc with $n = 3.5$. Soubiran \\cite*{soubiran93} showed that $n = 3.5 \\pm 0.5$ is compatible with the kinematical behavior of a star sample near the north galactic pole. K dwarf counts with HST yield $c/a = 0.8 \\pm 0.1$ and $n = 3.06 \\pm 0.22$ \\cite{gould98}. All of these studies were based on a few hundred objects at most.   In order to find new constraints on the spheroid density law, we undertook a photometric and astrometric sample survey in various galactic directions. We complemented these data with existing deep photometric star counts in several high and intermediate latitude fields. Most such counts contain large numbers of halo dwarfs, but they cannot be distinguished from thick disc dwarfs by their colours but at faint magnitudes. Since no large optical surveys were available at magnitude fainter than 20, we used heterogeneous data coming from various studies (often of extragalactic aim) in various photometric systems.  The population synthesis model used here permits to perform a global analysis of these heterogeneous data, since observational data can be simulated in each field with the true observational conditions (photometric system, errors and selection effects). The synthetic approach allows also to estimate the biases and expected contaminations by other populations.  In section~2 we describe the model of population synthesis and external constraints on the spheroid population. In section~3 we describe the data sets and the comparison method. In section~4 we discuss the results and their implications for the dark matter halo. ",
        "conclusions": "We have shown that, with the data available up to now, the spheroid star distribution follows a power law with an index smaller than previously thought. It is moderately flattened, as already found by several investigations. The best fit power law index is found 2.44 for a flattening of 0.76. We cannot exclude a spheroid population having a power law as low as 2 and an axis ratio of 0.5 at the 2 sigma level. Assuming a high exponent of 3.5, as suggested by some globular cluster and RR Lyrae data, model predictions deviate significantly from observations in the external parts of the galaxy whatever reasonable flattening is adopted.  The IMF slope of the spheroid is found to be $\\alpha=1.9\\pm0.2$, value which gives a local density of 1.64  10$^{-4}$ stars pc$^{-3}$ and a mass density of 4.15 10$^{-5}$ \\Msun pc$^{-3}$ for the stellar halo, yet this value ignores possible old white dwarfs. With this slope the expected mass density of brown dwarfs in the halo makes a negligible part of the dark matter halo, as already estimated from microlensing surveys.  Recent searches for ancient halo white dwarfs have given a hope to identify the microlensing events with such objects \\cite{1999ApJ...524L..95I,Ibata2000,2000Natur.403...57H}. Ibata et al. \\cite*{Ibata2000} conclude that old white dwarfs may constitute a significant fraction of about 10\\% of the dark matter halo. In the mean time, microlensing experiments have narrowed the range of the estimated halo baryonic fraction to 20 to 50\\% \\cite{alcock2000}. These two results are well in agreement according to the uncertainties.  So, as star count data progresses in depth and extent, the picture of the spheroid star population that comes out points to characteristics quite compatible with what we know about the distribution of baryonic dark matter if it is made of stellar remnants, suggesting a common dynamical origin. The visible spheroid and its heavy counterpart of dark remnants can make a significant but not dominant part of the so-called dark matter halo."
    },
    "0004/astro-ph0004364_arXiv.txt": {
        "abstract": "We report on the strengths of three spectral indicators --- Mg$_2$, H$\\beta$, and Hn/Fe --- in the integrated light of a sample of 100 field and cluster E/S0 galaxies.  The measured indices are sensitive to age and/or and metallicity variations within the galaxy sample. Using linear regression analysis for data with non-uniform errors, we determine the intrinsic scatter present among the spectral indices of our galaxy sample as a function of internal velocity dispersion. Our analysis indicates that there is significantly more intrinsic scatter in the two Balmer line indices than in the Mg$_2$ index, indicating that the Balmer indices provide more dynamic range in determining the age of a stellar population than does the Mg$_2$ index.  Furthermore, the scatter is much larger for the low velocity dispersion galaxies, indicating that star formation has occurred more recently in the lower mass galaxies. ",
        "introduction": "An enormous effort has been made in recent years to disentangle the degenerate effects of age and metallicity on the integrated light of a galaxy.  This age-metallicity degeneracy (Worthey 1994, 1996) plagues both the photometric colors and the integrated spectrum of a galaxy, making it extremely difficult to separate the effects of the two. Because broadband colors cannot break the degeneracy (Bica et al.\\ 1990; Charlot \\& Silk 1994), the best way to determine a unique solution to the problem is to use a combination of population synthesis models and spectral line indices. Absolute ages and metallicities determined using this method are naturally subject to systematic errors, yet relative values are still of great interest and are generally more reliable (e.g., Trager et al.\\ 2000a). Although no single spectral index can completely separate the effects of age and metallicity, it is possible to determine the relative range empirically by plotting a single spectral index with a predominant age or metallicity sensitivity versus a mass indicative parameter such as the absolute magnitude or internal stellar velocity dispersion.  The spread in index values which is observed as a function of mass is then directly related to the range in age and/or metallicity in the galaxy sample.  The most widely used spectral index system is the set of indices originally defined in Burstein et al.\\ (1984) and further refined in Worthey et al.\\ (1994) and Trager et al.\\ (1998), now commonly known as the Lick system.  In this system, the H$\\beta$ index is used as the primary age-sensitive spectral indicator, whereas the Mg and Fe indices are used as the primary metallicity indicators. Using the Lick indices, many groups have been able to place constraints on the ages and metallicities of various galaxy samples. For example, Trager et al.\\ (2000a) find that for a sample of early-type galaxies in low density environments there is a large spread in H$\\beta$ values (i.e., age), but very little variation in metallicity.  For galaxies in the Fornax cluster, though, Kuntschner (2000) finds the opposite effect, in that a large spread in metallicity is present with little variation in age.  Whether such contrasting results are a product of the different environments in the two samples is an interesting and still unanswered question.  It is clear, however, that a comparison of the range in age and metallicity of galaxies in different environments will lead to a better understanding of the parameters which govern the evolution of galaxies.  To constrain galaxy evolution scenarios, it is imperative that we first derive reliable values for the ages and metallicities of galaxies. In this paper we present first results of a sample of 100 early-type galaxies covering a large baseline in mass and a large range in wavelength in order to use both the Lick/IDS age/metallicity indicators as well as additional indices centered around the higher order Balmer lines. Although the data will eventually be used to determine absolute ages and metallicities, and thereby place constraints on the theories governing galaxy formation and evolution, these goals are beyond the scope of this paper. In this letter, we demonstrate the importance of the lower mass galaxies and higher order Balmer lines for obtaining a complete view of the ages of early-type galaxies.  The paper is organized such that the data and reduction procedures are summarized in \\S 2, the spectral indices and the index diagrams are presented in \\S 3, and the analysis of these diagrams is given in \\S 4 and discussed in \\S 5. ",
        "conclusions": "The results of this analysis bring forth two key issues that must be addressed. First, since the two Balmer line indices exhibit more intrinsic scatter than the Mg$_2$ index, this qualifies them as more robust age indicators.  The relatively little intrinsic scatter observed in the Mg$_2$-log $\\sigma$ diagram indicates that this index is not well suited for constraining the age of a stellar population. The fact that significant age variations may be hidden in the Mg$_2$-log $\\sigma$ relation has been previously noted by both Worthey et al. (1996) and J$\\o$rgensen (1999) and has been studied in detail by Trager et al.\\ (2000b). This finding may have a significant impact on the previous results of authors who have used the Mg$_2$ index as their principal age diagnostic (e.g., Bernardi et al.\\ 1998), and their conclusions may need to be re-evaluated using a more age-sensitive index such as H$\\beta$ or Hn/Fe.  The second key result is the amount of scatter observed for the low-mass galaxies in the index diagrams.  The dramatic increase in scatter, particularly at the low-mass end of the Balmer line diagrams, indicates that the smaller galaxies have experienced a more varied star formation history and that these galaxies have a larger range in age than their brighter counterparts.  This finding places significant constraints on a hierarchical formation scenario such as the CDM model (Baugh, Cole \\& Frenk 1996; Kauffmann 1996), since in this case one would naively expect to see that low-mass galaxies are older than the larger galaxies (Kauffmann \\& Charlot 1998). However, because the majority of the low-mass galaxies in our sample are in the Virgo cluster, it is difficult to determine whether the observed variation in age is simply a product of the different environments.  Even if this were the case, though, in a hierarchical scenario we would expect that galaxies in clusters would be older than field galaxies, since merging will happen more rapidly in clusters.  Thus, regardless of the environmental influence, our results appear to place significant constraints on the hierarchical picture.  In a forthcoming paper, we will report on use of the higher order Balmer index and population synthesis models to determine the spread of age and metallicity of our sample, and thereby place quantitative constraints on formation and evolution scenarios."
    },
    "0004/astro-ph0004014_arXiv.txt": {
        "abstract": "Temporal and spectral analyses from BeppoSAX observations of AM Her performed during both an intermediate and a high state are presented and discussed. Optical observations taken a few days after the X--ray ones are also presented.  During the intermediate state observation, the source was in its ``normal\", one--pole accretion mode. In the high state it switched to an hitherto unobserved atypical ''two-pole'' accretion mode, with significant soft and hard X-ray emission from both poles. The emission from the second pole is much softer than that from the primary pole, while the soft X-ray excess of the primary pole is fairly modest in this accretion mode. These facts suggest that accretion onto the secondary is mainly due to blobs penetrating deeply in the photosphere, while that on the primary pole is mostly via a more homogeneous column, giving rise to the classical standing shock. A strong X-ray  flaring activity is also  observed in the soft X-ray band, but not the hard X-ray and optical emissions indicating that flares are due to inhomogeneous blobby-accretion events. ",
        "introduction": "Polars, a subgroup of magnetic Cataclysmic Variables (mCVs), contain a highly magnetized white dwarf with polar field strengths ranging from  $\\sim$ 10 MG $\\sim$ to 230 MG (see Beuermann 1998 and references therein), and which accretes from a late--type main sequence star. The magnetic field of the white dwarf is strong enough to phase--lock its rotation with the orbital period. These systems are strong X-ray emitters in both soft and hard X-ray bands (see review by Cropper 1990).  While hard X-rays and optical cyclotron radiation are emitted from a standing shock above the white dwarf surface, soft X-rays originate from hot photospheric regions heated either by irradiation from the post-shock plasma (Lamb \\& Masters 1979) or by dense plasma blobs carrying their kinetic energy deep into the atmosphere (Kuipers \\& Pringle 1982). Irradiation is important primarily for flow rates sufficiently low for the shock to stand high above the surface. However, a large fraction of the reprocessed radiation appears in the UV rather than at soft X-rays, as shown quantitatively in the prototypal system AM\\,Her (G\\\"ansicke et al. 1995).  While the soft X--ray component is in general adequately fitted by a black--body spectrum with a temperature of a few tens of eV (even if more sophisticated models are sometimes needed: see e.g. Van Teeseling et al. 1994), it has recently become clear that a simple thermal plasma model is no longer adequate in reproducing the hard X--ray emission of Polars. Reflection from the white dwarf surface, complex absorption and multi--temperature emission may contribute significantly to the spectrum above a few tenths of keV (e.g. Cropper et al. 1998 and references therein).  Polars are known to display long--term luminosity variations on several timescales ranging from tens of days to years. These variations are thought to be related to changes in the mass transfer rate from the secondary star. The best studied example is again the brightest polar AM\\,Her (Cropper 1990; G\\\"ansicke et al. 1995). Such changes come along with spectral variations in all energy bands, indicating changes not only in the spectral parameters of the primary radiation but also in the reprocessed emission. The relative proportion of hard and soft X-ray emissions is expected to be sensitive to the actual accretion luminosity. This long--term variability is scarcely monitored and important questions related to changes in accretion parameters and in accretion modes (such as switching from one--pole to two--pole accretion), as well as to the processes of energy release, are still unanswered.  To properly address some of these issues, a  long--term monitoring program of the prototype AM\\,Her was set  up with the BeppoSAX satellite in order to investigate its long--term behaviour on a wide (0.1--100 keV) X-ray energy band simultaneously. These observations have been complemented with optical photometry for a more comprehensive study. ",
        "conclusions": "The BeppoSAX data presented here have shed new light onto the temporal and spectral behaviour of AM\\,Her.   The observations carried out at different brightness levels indicate that the accretion mode can be substantially different. During intermediate states, accretion occurs mainly in the ``normal\" mode, with the primary pole dominating both hard and soft X-ray emissions. A 20\\,eV soft black body and a 14\\,keV  temperature for the thermal plasma are found, provided the inclusion of partial absorption, of a reflection component and of an iron $\\rm K{\\alpha}$ 6.4\\,keV line with EW= 150\\,eV.  High states instead can be very different from epoch to epoch. Differently from previous observations, our BeppoSAX high  state observation shows a markedly different variability in the soft and hard X-rays. A strong flaring activity characterized by an increase in flux by a factor of 5--10 with typical exponential rise and decay time-scales of the order of tens of minutes is observed below 0.4\\,keV. In hard X-rays the variability is essentially related to the orbital modulation. Differently from intermediate and previous high state observations, the presence of significant emission during the faint phase in both soft and hard X-rays provide evidence that the secondary pole is emitting. Furthermore, the lack of an orbital variability in the soft X-ray ``quiescence'' emission contrasts with the strong modulation observed in hard  X-rays. This suggests that soft  X-ray emission from the primary pole (whatever produced by blobs or by reprocessing of hard X-rays) is highly inefficient. An atypical low soft X-ray emission period was indeed observed in the past in AM\\,Her during a high state  (Priedhorsky et al. 1987) but accompanied by the lack of the X-ray minimum in both soft and hard bands. While this was interpreted by a shift in the accretion column with a corresponding change of the accretion pattern, the behaviour observed in August 1998 appears to be better explained by the onset of accretion onto the secondary pole whose emission is mainly soft together with a less active soft X-ray production from the  primary pole.  It is worth noting that the soft flaring activity appears to be related to both accreting poles. The observed flaring activity is not uncommon for AM\\,Her  but with much shorter time scales. Furthermore similar variations are not observed  in the UBV coordinated photometry indicating that the soft X-ray flares are unrelated to the optical cyclotron post-shock emission. Moreover, spectral analysis of flares indicates a higher temperature of the soft black body ($\\sim$ 30\\,eV) while the hard X-ray emission has not varied.    The August high state BeppoSAX data have also revealed for the first time a strong energy dependence of the orbital modulation which moves from a single-humped shape at high energies (4--30\\,keV) to a double-humped shape around 1\\,keV. This effect can be explained in terms of absorption in the accretion column when pointing towards the observer.     For this high state, it has been possible to perform a phase--resolved spectral analysis, and to measure the parameters of both black body and optically thin thermal plasma emission along the orbital phase. The temperatures of both components are consistent with being constant,  and compatible with those inferred for the May intermediate state. The plasma temperature (14 keV) also agrees with  the ASCA data,  but not so the black body temperature, which is 25$\\pm$1 eV (ASCA: 32$\\pm$1 eV). The latter is however in agreement with EUVE (24.5 eV) (Ishida et al. 1997). The phase dependence of the partial absorber is indeed compatible with an increase of the amount of cold matter when the accretion column is along the line of sight. The dip observed at $\\phi_{mag} \\sim 0.5$ is consistent with such an increase, as both covering fraction and column density are higher. The phase dependence of the EW of the neutral iron line is consistent with an origin in the white dwarf surface illuminated by the hard X-rays from the post-shock regions.   \\begin{figure} \\epsfig{file=H1952_f13.ps, height=9.5cm, width=9.cm, angle=-90} \\caption{Best fit model for the high state, bright phase (model 3 in Table~3). } \\label{model_high} \\end{figure}   Fixing the black body temperature of both intermediate and high states to 25\\,eV and that of the optically thin thermal plasma to 14\\,keV, the bolometric luminosities of these components at the maximum, and adopting $D$=91\\,pc, are: $\\rm L_{BB} = 0.81(^{+0.31}_{-0.17}) \\times 10^{32}\\, erg\\,s^{-1}$ and $\\rm L_{Th.} = (0.683^{+0.027}_{-0.027}) \\times  10^{32}\\, erg\\,s^{-1}$ for the intermediate state; whilst for the high state these are $\\rm L_{BB} = (5.32^{+0.31}_{-0.60}) \\times 10^{32}\\, erg\\,s^{-1}$ and $\\rm L_{Th.} = (3.333^{+0.047}_{-0.037}) \\times 10^{32}\\, erg\\,s^{-1}$. We note that while the hard X-ray luminosities are typical of intermediate and high states, that of the black body component during the high state is lower than previously measured. This is shown in Fig.~\\ref{model_high}, which can be compared with that reported by Ishida et al. (1997) for the intermediate state in 1993.  The inferred soft-to-hard X-ray ratio is 1.2 and 1.6 during the intermediate and high state, respectively. It has been shown that in AM\\,Her the bulk of reprocessed hard X-ray and cyclotron radiations emerges in the UV range rather than in the soft X-rays (G\\\"ansicke et al. 1995), the latter being mainly produced by the blobby accretion mechanism. During the high state observed in August, both the inferred black body luminosity and the lack of a clear modulation in the soft X-ray band suggest that discrete accretion onto the primary pole is not as efficient as in other epochs. Given the low level of soft X-ray flux during the faint phase it is not possible to infer the black body luminosity from the secondary pole. From the high state black body luminosity  during the bright phase, information on the mass accretion rate can be estimated. A value of $\\rm 8.5 \\times 10^{-11}\\, M_{\\odot}\\,yr^{-1}$ for a white dwarf mass of $\\rm 0.6\\,M_{\\odot}$  is derived. A fractional accreting area $ f \\sim  2.4 \\times 10^{-4}$ is obtained, which is a factor of $\\sim$ 3-10 larger than those derived by G\\\"ansicke et al. (1995) and Ishida et al. (1997) respectively during typical high states. The local mass accretion rate then results to be $\\sim$ 4\\,g\\,s$^{-1}$cm$^{-2}$ which is too low to make blobby accretion efficient.  As far as flares are concerned, the changes are essentially related to an increase in the temperature of the soft component, as also observed during other flaring epochs (Ramsay et al. 1996).  The increase in black body luminosity and then in the mass accretion rate is a factor of 2 with respect to ``quiescence'', and comparable to usual high state values. Our estimates of the  fractional area of the emitting black body during flares results to be similar to that during ``quiescence'' which gives a low local mass accretion rate of $\\sim$ 8\\,g\\,s$^{-1}$cm$^{-2}$. However, since  we cannot exclude that accretion occurs onto a more localized area, we consider this value as a lower limit.    Different from the flaring variability are the rapid fluctuations detected in the hard X-rays and optical ranges, the latter likely to be associated with inhomogeneities in the accretion flow itself.  Furthermore, the optical orbital modulation observed in August is typical of AM\\,Her during high states, where the V band is dominated by cyclotron beaming from the primary pole (G\\\"ansicke et al. in preparation). The secondary pole instead appears to dominate at optical blue wavelengths. This is very different from low and other high states where the main accreting pole usually dominates the UV and the blue portion of the optical spectrum.  \\medskip   In summary, the present observations reveal that AM\\,Her has switched from a ``normal\" accretion mode during the intermediate state of May 1998 into a ``two-pole\" accretion mode during the high state of August 1998, with a less efficient  ``blobby\" accretion onto the primary pole."
    },
    "0004/astro-ph0004222_arXiv.txt": {
        "abstract": "This paper is concerned with the Laplace boundary-value problem with the directional derivative, corresponding to the specific nature of measurements of the longitudinal component of the photospheric magnetic field. Boundary conditions are specified by a distribution on the sphere of projection of the magnetic field vector unto a given direction. It is shown that the solution of this problem exists in the form of a spherical harmonic expansion, and uniqueness of this solution is proved. A conceptual sketch of numerical determination of the harmonic series coefficients is given. The field of application of the method is analyzed having regard to the peculiarities of actual data. Finally, we present differences in results derived from extrapolating the magnetic field from a synoptic map and a full-disk magnetogram. ",
        "introduction": "Until the present time the practical potential-field extrapolation of solar magnetic fields has been based on several thoroughly studied Laplace boundary-value problems when applied to two cases of geometry: spherics, and planar geometry. The spherics, as a rule, is used in investigating large-scale long-lived global fields. The planar geometry is exploited to study magnetic active regions occupying small areas on the solar surface.  When employing the spherics, two approaches are distinguished to specifying boundary conditions on the solar surface. One is that it is specified the synoptic distribution of a so-called component $B_l$ (a projection of the magnetic field measured along the line-of-sight direction is assigned to each point of the surface at the time when a given point passes through the central meridian (\\cite{Altschuler} and \\cite{Hoeksema}). In the other approach, a normal component $B_r$ (\\cite{Wang and Sheeley}) is specified, which is constructed from $B_l$ assuming that the full vector of the magnetic field is radial. The latter case corresponds to the setting of a classical Neumann boundary-value problem, and the former case refers to a particular variant of the boundary-value problem with the directional derivative (the direction of the derivative with respect to a normal varies according to the heliographic latitude). Data of daily magnetograms do not satisfy these conditions. They are only approximately satisfied by points lying in the neighborhood of the central meridian. Synoptic maps when generated from series of daily magnetograms accumulated during one solar rotation provide input information corresponding to the above-mentioned boundary conditions.  Using the planar geometry, unlike the spherics, provides insight into the three-dimensional spatial structure at the time of measurement thus enabling one to track the temporal evolution of magnetic active regions. This approach has two important limitations. In the first place, the region under study must be sufficiently small; secondly, it must lie in the middle of the solar disk where the normality approximation of the measured component of field is valid. Boundary conditions in this case correspond to the Neumann problem. A variant of the problem with a constant slope of the derivative (for noncentral regions) can also be used but at a still greater sacrifice in the size of the region if variations in the line-of-sight slope to a normal inside the region are to be neglected.  This paper offers the solution of a Laplace boundary-value problem (in the spherics) with boundary conditions corresponding to measurements of the longitudinal magnetic components throughout the disk (data from daily magnetograms). Such a setting of the problem conforms to a Laplace problem with a variable directional derivative (by choosing a rigorously specified projection direction corresponding to the direction of the line of sight existing at a given time of measurement, the direction, with respect to which a derivative of the scalar magnetic potential is taken, changes relative to a normal on the boundary spherical surface). This let to use completely longitudinal magnetogram data which are the most sensitive and generally applied. We demonstrate the existence of a strictly definite harmonic expansion of the potential field in terms of spherical eigenfunctions, corresponding to the above boundary conditions. The construction of such a solution is based on the combined use of Galerkin's projection method and Trefftz' method (the solutions are sought in a class of solutions of the Laplace equation). The proof of the existence and uniqueness of the solution is constructed. We present and analyze results of numerical tests simulating the extrapolation of actual magnetograms. The limitations of the approach used in this study are discussed, which are associated with the resolution of the magnetograms employed and with the uncertainty of specifying the magnetogram of the averted side of the Sun (''back magnetogram''). A comparison is made of results derived from extrapolating from a synoptic magnetogram and from a full-disk magnetogram. ",
        "conclusions": "Main results and conclusions may be summarized as follows.  -- A new method has been developed for extrapolating the solar magnetic fields from daily full-disk magnetograms.  -- A mathematical rationale for the new method was given.  - An analysis was made of the applicability of the method, with due regard for the particular characteristics of existing data.  - The most suitable technique has been found for specifying the boundary conditions on the back side of the Sun.  - It has been shown that this method can be used successfully only when high-resolution daily magnetograms are employed (with the exception of Stanford magnetograms).  - Calculations confirmed that in the confidence region defined above, one can expect a good correspondence between calculated and real magnetic field structures.  - In contrast to other techniques, this method makes it possible to describe current conditions of a global magnetic structure (not averaged over time and not local).  - The analysis made in this paper has revealed hidden limitations of high-resolution magnetograms which manifest themselves when reconstructing the magnetic structure on the source surface. It may be concluded that with the proviso that they are eliminated or their influence is adequately reduced, it is possible to predict a change of sign of the near-terrestrial interplanetary magnetic field from the reconstructed picture of the radial magnetic field on the source surface."
    },
    "0004/hep-th0004109_arXiv.txt": {
        "abstract": "We show that in brane-world scenarios with warped extra dimensions, the Casimir force due to bulk matter fields may be sufficient to stabilize the radion field $\\phi$. In particular, we calculate one loop effective potential for $\\phi$ induced by bulk gravitons and other possible massless bulk fields in the Randall-Sundrum background. This potential has a local extremum, which can be a maximum or a minimum depending on the detailed bulk matter content. If the parameters of the background are chosen so that the hierarchy problem is solved geometrically, then the radion mass induced by Casimir corrections is hierarchycally smaller than the $TeV$. Hence, in this important case, we must invoke an alternative mechanism (classical or nonperturbative) which gives the radion a sizable mass, to make it compatible with observations. ",
        "introduction": "It has been suggested that theories with extra dimensions may provide a solution to the hierarchy problem \\cite{gia,RS1}. The idea is to introduce a $d$-dimensional internal space of large physical volume ${\\cal V}$, so that the the effective lower dimensional Planck mass $m_{pl}\\sim {\\cal V}^{1/2} M^{(d+2)/2}$ is much larger than $M \\sim TeV$- the true fundamental scale of the theory. In the original scenarios, only gravity was allowed to propagate in the higher dimensional bulk, whereas all other matter fields were confined to live on a lower dimensional brane. Randall and Sundrum \\cite{RS1} (RS) introduced a particularly attractive model where the gravitational field created by the branes is taken into account. Their background solution consists of two parallel flat branes, one with positive tension and another with negative tension embedded in a a five-dimensional Anti-de Sitter (AdS) bulk. In this model, the hierarchy problem is solved if the distance between branes is about $37$ times the AdS radius and we live on the negative tension brane. More recently, scenarios where additional fields propagate in the bulk have been considered \\cite{alex1,alex2,bagger}.  In principle, the distance between branes is a massless degree of freedom, the radion field $\\phi$. However, in order to make the theory  compatible with observations this radion must be stabilized \\cite{gw1,gw2,gt,cgr,tm}. Clearly, all fields which propagate in the bulk will give Casimir-type contributions to the vacuum energy, and it seems natural to investigate whether these could provide the stabilizing force which is needed. In this paper, we shall calculate the radion effective potential $V_{\\hbox{\\footnotesize\\it \\hspace{-6pt} eff\\,}}(\\phi)$ due to KK gravitons and other massless bulk fields. As we shall see, this effective potential has a rather non-trivial behaviour, which generically develops a local extremum. Depending on the detailed matter content, the extremum could be a maximum or a minimum, where the radion could sit. For the purposes  of illustration, here we shall concentrate on the background geometry discussed by Randall and Sundrum, although our methods are also applicable to other geometries, such as the one introduced by  Ovrut et al. in the context of eleven dimensional supergravity with one large extra dimension \\cite{ovrut}.  The paper is organized as follows. In Section II, we introduce the model and conventions. In Section III we calculate the effective potential due to conformally invariant bulk fields, and discuss the stabilization mechanism. Conformally invariant fields are convenient because of their simplicity, which allows a quick derivation of their contribution. Moreover, the backreaction of the Casimir energy on the geometry can be taken into consideration in this case. Section IV is devoted to the graviton contribution. For gravitons, conformal invariance is lost, and the evaluation of the effective potential is not straightforward. A method will be developed which relates the curved space effective potential to a suitable effective potential in flat space. The later takes the form of a sum over the contributions of an infinite tower of massive KK fields, whose mass  spectrum $m_n(\\phi)$ is a function of the brane separation (the discrete index $n$ labels the infinite KK tower). Our conclusions are summarized in Section V.  Related calculations of the Casimir interaction amongst branes have been presented in an interesting paper by Fabinger and Horava \\cite{FH}. In the concluding section we shall comment on the differences between their results and ours. ",
        "conclusions": "We have shown that in brane-world scenarios with a warped extra dimension, it is in principle possible to stabilize the radion $\\phi$ through the Casimir force induced by bulk fields. In particular, conformally invariant fields induce an effective potential of the form (\\ref{confveff}) as measured from the positive tension brane. From the point of view of the negative tension brane, this corresponds to an energy density per unit physical volume of the order $$ V_{\\hbox{\\footnotesize\\it \\hspace{-6pt} eff\\,}}^{-}\\sim m_{pl}^4 \\left[{A\\lambda^4 \\over (1-\\lambda)^4}+\\alpha+\\beta\\lambda^4\\right], $$ where $A$ is a calculable number (of order $10^{-3}$ per degree of freedom), and $\\lambda \\sim \\phi/(M^3 \\ell)^{1/2}$ is the dimensionless radion. Here $M$ is the higher-dimensional Planck mass, and $\\ell$ is the AdS radius, which are both assumed to be of the same order, whereas $m_{pl}$ is the lower-dimensional Planck mass. In the absence of any fine-tuning, the potential will have an extremum at $\\lambda \\sim 1$, where the radion may be stabilized (at a mass of order $m_{pl}$). However, this stabilization scenario without fine-tuning would not explain the hierarchy between $m_{pl}$ and the $TeV$.  A hierarchy can be generated by adjusting $\\beta$ according to (\\ref{consts}), with $\\lambda_{obs}\\sim (TeV/m_{pl}) \\sim 10^{-16}$ (of course one must also adjust $\\alpha$ in order to have vanishing four-dimensional cosmological constant). But with these adjustement, the mass of the radion would be very small, of order \\begin{equation} m^{2\\ (-)}_{\\phi} \\sim\\lambda_{obs}\\ M^{-3} \\ell^{-5} \\sim \\lambda_{obs} (TeV)^2. \\label{smallmass} \\end{equation} Therefore, in order to make the model compatible with observations, an alternative mechanism must be invoked in order to  stabilize the radion, giving it a mass of order $TeV$.  Goldberger and Wise \\cite{gw1,gw2}, for instance, introduced a field $v$ with suitable classical potential terms in the bulk and on the branes. In this model, the potential terms on the branes are chosen so that the v.e.v. of the field in the positive tension brane $v_+$ is different from the v.e.v. on the negative tension brane $v_-$. Thus, there is a competition between the potential energy of the scalar field in the bulk and the gradient which is necessary to go from $v_+$ to $v_-$. The radion sits at the value where the sum of gradient and potential energies is minimized. This mechanism is perhaps somewhat {\\em ad hoc}, but it has the virtue that a large hierarchy and an acceptable radion mass can be achieved without much fine tuning. It is reassuring that in this case the Casimir contributions, given by (\\ref{smallmass}), would be very small and would not spoil the model.  We have also calculated the graviton contribution to the radion effective potential. Since gravitons are not conformally invariant, the calculation is considerably more involved, and a suitable method has been developed for this purpose. The result is that gravitons contribute a negative term to the radion mass squared, but this term is even smaller than (\\ref{smallmass}), by an extra power of $\\lambda_{obs}$.  In an interesting recent paper \\cite{FH}, Fabinger and Horava have considered the Casimir force in a brane-world scenario similar to the one discussed in this paper, where the internal space is topologically $S^1/Z_2$. In their case, however, the gravitational field of the branes is ignored and the extra dimension is not warped. As a result, their effective potential is monotonic and stabilization does not occur (at least in the regime where the one loop calculation is reliable, just like in the original Kaluza-Klein compactification on a circle \\cite{ac}). The question of gravitational backreaction of the Casimir energy onto the background geometry is also discussed in \\cite{FH}. Again, since the gravitational field of the branes is not considered, they do not find static solutions. This is in contrast with our case, where static solutions can be found by suitable adjustment of the brane tensions.  Finally, it should be pointed out that the treatment of backreaction (here and in \\cite{FH}) applies to conformally invariant fields but not to gravitons. Gravitons are similar to minimally coupled scalar fields, for which it is well known that the Casimir energy density diverges near the boundaries \\cite{bida}. Therefore, a physical cut-off related to the brane effective width seems to be needed so that the energy density remains finite everywhere. Presumably, our conclusions will be unchanged provided that this cut-off length is small compared with the interbrane separation, but further investigation of this issue would be interesting."
    },
    "0004/astro-ph0004291_arXiv.txt": {
        "abstract": "We briefly review BeppoSAX observations of X-ray bright radio-galaxies. Their X-ray spectra are quite varied, and perhaps surprisingly, any similarity between radio-loud AGN and Seyfert galaxies is the exception rather than the rule. When detected, reprocessing features (iron line and reflection) are generally weak, suggesting two possible scenarios: either: (1) non-thermal (jet?) radiation dilutes the X-ray emission from the disk in radio-loud objects, or (2) the solid angle subtended by the X-ray reprocessing material is smaller in radio-loud than in radio-quiet AGN due to different characteristics of the accretion disk itself. ",
        "introduction": "\\label{intro}  X-ray observations of radio-galaxies can play an important role in understanding the AGN radio-loud and radio-quiet dichotomy. Several authors have speculated upon the possibility that the AGN distinction in two big families reflects different physical properties of the nuclear engine. Several years ago, Rees et al. (1982), suggested that in the nuclei of radio-galaxies a spinning black hole is surrounded by a hot thick accretion flow, in contrast with the radio-quiet picture, which assumes the accreting material in shape of cold thin disk. Later, other authors (Blandford 1990, Meier these proceedings) indicated the black hole spin as the physical parameter responsible of the AGN dichotomy, being more rapid in radio-louds and therefore more efficient in producing powerful jets.  X-ray photons are produced in the inner regions ($<1$ pc) of AGN and provide crucial information on physical processes occurring very near to the black hole. X-ray studies of radio-galaxies are then essential to answer a fundamental question: are the same accretion processes at work in radio-loud and radio-quiet AGNs?  In this paper, we present an on-going analysis of radio-galaxies observations performed with the BeppoSAX satellite (0.1-150 keV) and discuss the results of a comparative study between our sources and a sample of 12 Seyfert galaxies \\cite{mat99} observed with the same satellite. ",
        "conclusions": "BeppoSAX analysis of radio-galaxies has pointed out important differences between radio-loud  and radio-quiet AGN and has clearly shown, for the first time,  that in BLRGs the iron line equivalent widths (and probably the reflection strengths) are significantly smaller than in Seyferts 1.  In some radio-galaxies the cold material  responsible for the line production is located far from the X-ray primary source. It is possible that the innermost part of the accretion flow is not in shape of cold disk. The accretion flow might be then  hot and geometrically thick in the inner regions \\cite{rees82,sha76,na98} and cold optically thick \\cite{shak73} only at larger radii, subtending relatively smaller solid angles (compared to Seyfert 1s). Featureless spectra or  weak iron lines could be signatures of low efficiency accretion processes.  In sources with Seyfert-like spectra (i.e, with soft excess, strong UV bump and broad iron lines), the Fe line might be produced by a cold thin optically thick disk (Seyfert-like accretion flow). In this case, the iron line EW might be diluted by the Doppler-enhanced non-thermal jet continuum. This explanation predicts that the observed continuum (disk +jet) is more variable than the Fe line (disk)."
    },
    "0004/astro-ph0004258_arXiv.txt": {
        "abstract": "We have observed the Be/X-ray pulsar binary system X Per/4U~0352+30 on 61 occasions spanning an interval of 600 days with the PCA instrument on board the {\\it Rossi X-ray Timing Explorer} (\\RXTE\\,).  Pulse timing analyses of the 837-s pulsations yield strong evidence for the presence of orbital Doppler delays.  We confirm the Doppler delays by using measurements made with the All-Sky Monitor on \\RXTE\\,.  We infer that the orbit is characterized by a period $P_{orb} = 250$ d, a projected semimajor axis of the neutron star $a_x \\sin i = 454$ lt-s, a mass function $f(M) = 1.61$ M$_{\\sun}$, and a modest eccentricity $e = 0.11$.  The measured orbital parameters, together with the known properties of the classical Be star X Per, imply a semimajor axis $a = 2.2$ AU, and an orbital inclination $i \\sim 23\\arcdeg$--$30 \\arcdeg$.  We discuss the formation of the system in the context of the standard evolutionary scenario for Be/X-ray binaries.  The orbital eccentricity just after the supernova explosion was almost certainly virtually the same as at present, because the Be star is much smaller than the orbital separation.  We find that the system most likely formed from a pair of massive progenitor stars, and probably involved a quasi-stable and nearly conservative transfer of mass from the primary to the secondary.  We find that the He star remnant of the primary most likely had a mass $\\lesssim\\,6$ M$_{\\sun}$ after mass transfer.  If the supernova explosion was completely symmetric, then the present orbital eccentricity indicates that $\\lesssim$\\,4 $M_{\\sun}$ was ejected from the binary.  If, on the other hand, the birth of the neutron star was accompanied by a ``kick'' of the type often inferred from the velocity distribution of isolated radio pulsars, then the resultant orbital eccentricity would likely have been substantially larger than 0.11.  We have carried out a Monte Carlo study of the effects of such natal kicks, and find that there is less than a 1$\\%$ probability of a system like that of X Per forming with an orbital eccentricity $e \\lesssim 0.11$.  Finally, we speculate that there may be a substantial population of neutron stars formed with little or no kick. ",
        "introduction": "X Per is a bright and highly variable star with a visual magnitude that ranges from $\\sim$\\,6.1 to $\\sim$\\,6.8 (\\cite{moo74,roc97}).  When the star is particularly bright, its spectrum displays strong emission in H$\\alpha$ and other Balmer lines which marks this clearly as a Be-type system; when the star is faint, the emission lines disappear and it appears to be a normal early-type star (\\cite{fab92,roc97} and references therein).  The variability is commonly supposed to be caused by the formation and dissipation of a disk around the star, as in other Be stars, and appears to require free-free and free-bound emission within the disk in addition to electron scattering of the photospheric flux (e.g., \\cite{kun95,tel98}).  The spectral class of the underlying OB star has been estimated to be O9.5\\,III to B0\\,V (\\cite{sle82,fab92,lyu97}).  Based on spectroscopic parallax, distance estimates range from $700 \\pm 300$ pc up to $1.3 \\pm 0.4$ kpc (\\cite{fab92,lyu97,roc97,tel98}).  Lyubimkov et al. (1997) have used both spectroscopic and photometric observations of the low-luminosity diskless phase of X Per to infer the characteristics of the visible component; their results indicate that X Per is likely to have a mass of $\\sim$\\,13-20 M$_{\\sun}$ and a radius of 5-10 R$_{\\sun}$.  X Per is also the optical counterpart of the low-luminosity X-ray source 4U~0352+30 (\\cite{bm72,vdb72}).  X-ray observations have, furthermore, revealed pulsations with a period of $\\sim$\\,835 s (\\cite{whi76}).  This indicates that X Per must be a binary system containing both a Be star and either a slowly rotating neutron star or a white dwarf. For the first 5 years after the discovery of pulsations, the pulsar exhibited apparently erratic pulse frequency variations superposed on a long-term trend in which it was spinning up at the rate of $\\dot{P}_{pulse}/P_{pulse} = -1.5 \\times 10^{-4}$ yr$^{-1}$.  This was followed by $\\sim$\\,20 years of spindown at a similar absolute rate (\\cite{whi76,rob96,dis98} and references therein).  In all, there have been only about two dozen determinations of the pulse period of 4U~0352+30, so that its pulse period behavior on short timescales (e.g., weeks, months, and even a year) is not well documented.  The X-ray luminosity varies on long timescales (years) from as high as $\\sim 3 \\times 10^{35}$ ergs s$^{-1}$ to as low as its current value of $\\sim 5 \\times 10^{34}$ ergs s$^{-1}$ (for an assumed distance of 1.3 kpc; \\cite{roc93}; also see \\cite{dis98}).  The observed pulse period variations strongly suggest that the X-ray source is an accreting neutron star, with the period variations largely due to accretion torques (see, e.g., \\cite{bil97} and references therein).  White et al. (1982) have argued that considerations of the X-ray luminosity and spectrum in the context of accretion from a stellar wind strongly indicate that 4U~0352+30 is a neutron star.  Haberl et al. (1998) have noted that the 837 second rotation period is one of the longest known for any neutron star.  The only report of an orbital period for X Per ($\\sim$\\,580 days) was given by Hutchings et al. (1974) and was based on optical spectroscopy (also see \\cite{hut75,hut77}).  Penrod \\& Vogt (1985) explained the radial velocity results reported by these authors by emission variably filling in the Balmer absorption lines, so that the apparent radial velocity curve did not represent the velocity of X Per.  This cast doubt on the interpretation of $\\sim$\\,580 days as the orbital period. Moreover, this period has never been securely confirmed by any subsequent optical or X-ray observations (\\cite{wei84,rey92,roc93,smi99}).  There have been other suggestions of orbital periods based on weak evidence.  In retrospect, one of the most interesting of these was a possible $\\sim$\\,250 d periodicity in pulse period measurements noticed by White et al. (1982).  There is one line of evidence that qualitatively supports the idea of a long orbital period for X Per, viz. the ``Corbet diagram'' (\\cite{cor86,vdh87}) in which the pulse period is plotted versus orbital period for the known accretion powered X-ray pulsars.  For the $\\sim$\\,10 systems containing Be stars there is a good correlation between pulse period and orbital period which can be represented by the relation \\begin{equation} P_{orb}  \\sim  18 (P_{pulse}/1\\,\\rm{s})^{2/5}\\  \\rm{days} \\end{equation} with a residual scatter of about a factor of two.  For a pulse period of 837 s, eq. (1) predicts an orbital period of $\\sim$\\,266 d.  We have observed 4U~0352+30 with the Proportional Counter Array (PCA) on the {\\it Rossi X-ray Timing Explorer} (\\RXTE\\,) in order to track the pulse phase of the neutron star over a substantial interval of time and to thereby allow us to search for a Doppler signature of orbital motion.  We have also utilized observations of 4U~0352+30 made with the \\RXTE\\ All-Sky Monitor (ASM) in our search.  In Section 2 we describe the \\RXTE\\ observations, and provide a log of observation dates, durations, and pulse arrival times.  The pulse timing analysis of the PCA data and the results which indicate an orbit with a period of 250 days are described in Section 3.  In Section 4 we examine the possibility that the observed pulse arrival time delays arose from a random walk in the pulse phase due to accretion torques, and conclude that this is quite unlikely.  In this section, we also present the results of our successful search for orbital Doppler delays in observations made with the ASM.  The astrophysical implications of a wide yet only mildly eccentric orbit in a Be star X-ray binary are discussed in the context of binary evolution scenarios and natal neutron star kicks in Section 5. ",
        "conclusions": "Bildsten et al. (1997) list 5 Be/X-ray binaries with measured orbital parameters.  These are shown schematically, but to scale, in Figure 11.  For each system, the Be star is fixed, and the orbit of the neutron star is drawn with the {\\it full} semimajor axis $a$.  To estimate the semimajor axis we utilized the measured {\\em projected} semimajor axis and mass function for the neutron star, and then simply assumed masses for the Be and neutron stars of 18 M$_{\\sun}$ and 1.4 M$_{\\sun}$, respectively.  Also shown on the figure for comparison is the X Per system.  The average eccentricity for the 5 Be/X-ray binaries other than X Per is $e = 0.4$, in sharp contrast with the much smaller value of 0.11 for X Per.  This is especially noteworthy in that the X Per system has the smallest value of $R_{Be}/a\\ (\\lesssim 0.05)$, which indicates that there is virtually no chance that the system has circularized significantly since the birth of the neutron star (see, e.g., \\cite{zah77,zah89,ver95,sok98}).  Here we take 20 R$_{\\sun}$ as the upper limit on the radius of the Be star.  Thus, for any plausible circularization time for such a Be star filling its Roche lobe (e.g., a few thousand yr), the tidal circularization timescale, which is proportional to $(a/R)^8$, in the widely separated X Per system would exceed the lifetime of X Per by orders of magnitude.  If we assume that the orbit of X Per has an eccentricity of 0.11 which has remained unchanged since the birth of the neutron star, then we can set interesting and potentially important constraints on the dynamics of the supernova explosion and the mass that was ejected from the binary system.  First, however, we consider the origin of the X Per/4U0352+30 system.  In the conventional scenario for forming a Be X-ray binary (\\cite{rap82,vdh87,hab87}), the more massive star (e.g., $\\sim$\\,10-25 $\\rm{M}_{\\sun}$) in the progenitor binary evolves first and fills its Roche lobe.  Depending on the evolutionary state of the primary at the onset of mass transfer (see, e.g., \\cite{kip67,pod92} and references therein), the Roche lobe overflow of mass onto the secondary may proceed either on a thermal timescale or be dynamically unstable.  In the former case, relevant to the formation of Be stars, it is unclear how conservative the mass transfer may be.  In general, the orbital separation will first decrease, and then, if conditions are right (see below), increase until the mass transfer terminates.  In this scenario, detailed binary evolution calculations are required to determine how much envelope mass is retained by the He core of the original primary star (see, e.g., Podsiadlowski et al. 1992). However, simple orbital dynamics provide an analytic relation among the final orbital separation, the initial separation, and the final and initial masses of the system (see equation 5 of Podsiadlowski et al. 1992).  There are two ``free'' parameters in this relation, viz., $\\alpha$, the specific angular momentum carried away by matter ejected from the binary system in units of the binary angular momentum per reduced mass, and $\\beta$, the fraction of mass lost by the primary that is retained by the secondary during mass transfer.  We plot in Figure 12 some illustrative examples of the ratio of final to initial orbital separations as a function of fractional mass loss of the primary star for different assumed values of $\\beta$.  The initial stellar masses in this particular example are 15 and 10 $\\rm{M}_{\\sun}$, and we adopted a value $\\alpha = 1.5$ which might be typical for mass lost through the outer Lagrange points.  (We remind the reader that in the above paragraph as well as in the next paragraph the words ``initial'' and ``final'' refer only to the process of quasi-conservative mass transfer from the primary to the secondary {\\em prior} to the supernova.)  As one can see from Figure 12, for conservative mass transfer the final orbit can be substantially larger than the initial orbit. However, for $\\beta \\lesssim 0.6$ the orbit tends to shrink dramatically.  This is true for a range of plausible initial mass ratios.  If the initial orbital separation in the X Per system had been larger than $\\sim\\,5$ AU, the primary would not have filled its Roche lobe and transferred mass to the secondary (and the system would have become unbound after the supernova). Moreover, the orbit just prior to the supernova explosion was probably not much closer than 2 AU (see Fig. 14 and the associated discussion below).  Thus, since the orbit did not shrink by more than a factor of $\\sim\\,5/2$ during the mass-transfer phase from the primary, we can see from Fig. 12 that $\\beta$ cannot be substantially less than $\\sim\\,0.6$.  Finally, we conclude that the initial orbital separation of the primordial binary probably could not have been much smaller than about 0.5 AU. This latter conclusion is based upon the fact that in order for the final to initial orbital separation to increase by a factor of $\\sim$\\,4, the value of $\\beta$ must be close to unity; then one must also satisfy the condition that the computed current-epoch mass of X Per should not be too large.  In summary, we infer that the primordial binary which formed the current X Per system had an initial orbital separation of between 0.5 and 5 AU, experienced quasi-stable Roche-lobe overflow when the primary evolved, and that most of the mass lost by the primary was retained by the secondary.  We now consider the effects on the orbit of the collapse of the Fe core of the remnant primary and the ensuing supernova explosion. First, we examine the case where the supernova explosion was spherically symmetric in the rest frame of the progenitor star, i.e., no kick velocity was imparted to the newly born neutron star.  If the orbit of the progenitor of the neutron star and its companion (i.e., the present-day X Per) was circular, and there was no kick imparted to the neutron star at its birth, then there is a simple relationship among the current orbital eccentricity, the total mass of the progenitor system, $M_{tot}$ (just before the supernova explosion), and the mass ejected by the explosion, $\\Delta M$: \\begin{equation} e = \\Delta M/(M_{tot} - \\Delta M) \\end{equation} Some illustrative examples taken from this relationship are plotted in Figure 13.  Here we show the eccentricity induced by the supernova explosion as a function of the mass ejected, for a plausible range of values (10 - 30 M$_{\\sun}$) for the total mass of the binary just prior to the supernova explosion.  Note that for the supernova to have produced an orbital eccentricity $e < 0.15$ (i.e., the $2\\sigma$ observational upper limit), no more than $\\sim4\\,\\rm{M}_{\\sun}$ can have been ejected from the system.  Thus, we conclude that the progenitor of the neutron star was a He star with mass $< 5-6\\,\\rm{M}_{\\sun}$. Estimates of the minimum mass of a He star that can produce a neutron star are close to 2.3\\,M$_{\\sun}$ (see, e.g., \\cite{hab86a,hab86b}).  The results shown in Figure 13 are for the case where the neutron star was {\\em not} given a kick when the Fe core of its progenitor star collapsed.  However, it is conventional wisdom that most neutron stars are given substantial kicks at their birth (e.g., \\cite{ll94,cor97,cor98,fry98,han97}).  The measured 3-dimensional space velocity distribution for isolated neutron stars is uncertain, but is generally characterized by a function with a mean speed of $\\sim$\\,300 km $s^{-1}$ and which tends to vanish at low speeds. There seems to be a growing consensus that these large space velocities result from natal kicks to the neutron star, and that the ``slingshot'' effect from neutron stars born in binary systems is insufficient to produce the observed velocities (see, e.g., Cordes \\& Chernoff 1997, 1998; Fryer, Burrows, \\& Benz 1998; Hansen \\& Phinney 1997; cf. \\cite{ibe96}).  The small observed eccentricity in X Per seems to be qualitatively at odds with this ``universal'' kick distribution.  In order to study the problem more quantitatively, we have carried out a modest Monte Carlo study of supernovae in binary systems of the type that are directly relevant to the formation of 4U~0352+30/X Per.  We considered circular orbits for the progenitor He star and its massive companion (the current X Per).  We chose the mass of the He star to be 6 M$_{\\sun}$. For more massive He stars even a no-kick supernova explosion would leave the orbit more eccentric than is observed -- unless the supernova kick fortuitously ``corrects'' the eccentricity that would otherwise be induced.  We take the mass of the companion star to be 18 M$_{\\sun}$, which is approximately the current mass of X Per.  We also considered 4 different initial orbital separations for the pre-supernova binary: 0.5, 1, 2, and 4 AU.  For each Monte Carlo case we chose at random a kick speed from a Maxwellian distribution (Hansen \\& Phinney 1997): \\begin{equation} p(v) = \\sqrt{ \\frac{\\pi}{2} } \\frac{v^2}{v_0^3} e^{-v^2/2v_0^2}, \\end{equation} while the direction of the kick was chosen from an isotropic distribution (see, e.g., \\cite{bra95}). In this expression we adopted a value for $v_0$ of 190 km s$^{-1}$ which yields a mean kick velocity of 320 km s$^{-1}$ (\\cite{han97}).  These quantities then uniquely determine the final post-supernova orbital parameters; these were recorded for each system.  For each separation this procedure was repeated $10^6$ times, and the distributions of the results are shown in Figure 14.  We find that $\\sim$\\,3-24\\% of the binaries remain bound following the supernova explosion, with the exact percentage depending sensitively on the pre-supernova semimajor axis.  The best match to the current orbital period of 250 days is an initial orbital separation of 2 AU.  However, as can readily be seen from the distribution of eccentricities, there is only a {\\it very small} probability of finding the post-supernova orbit with an eccentricity of $\\lesssim$\\,0.11; typically $\\sim$\\,1\\% of the binaries which remain bound have an eccentricity this small.  From our analysis it seems reasonable to conclude that, at least for the X Per binary, the Be/neutron star pair formed via a completely unremarkable scenario.  Because of the small eccentricity in the system, however, our Monte Carlo kick study clearly indicates that it is unlikely that a substantial kick was imparted to the neutron star at birth.  However, one should be cautious about drawing far reaching conclusions based on this one system.  We point out that there are at least two other relatively wide X-ray binaries that may exhibit a similarly small orbital eccentricity: 2S 1553--54 and GS 0834--43 (see \\cite{bil97} and references therein).  The orbit of the transient X-ray source 2S 1553--54 was determined on only one occasion during a single orbital cycle (\\cite{kra83}).  The system was found to have a period of 29 days and an eccentricity $e < 0.09$ (2$\\sigma$ limit). The companion star has never been optically identified.  Nonetheless this system seems like another reasonable candidate for a Be/X-ray binary with only a small eccentricity.  The transient source GS 0834--43 (\\cite{wil97}) was found to have an orbital period of 106 days, but the orbital parameters were not uniquely determined due to effects of accretion torque noise.  However, the best fitting solutions yield an orbital eccentricity of $\\sim$\\,0.17.  The optical counterpart of GS~0834--43 has recently been identified with a Be star (\\cite{isr00}).  While it is always conceivable that this system was produced via a highly improbable event, it seems more reasonable in the case of X Per to draw the conclusion that not all neutron stars receive kicks of a magnitude consistent with the Hansen \\& Phinney (1997) distribution of single pulsar velocities.  Therefore we adopt the somewhat unconventional view that perhaps a substantial fraction of all neutron stars do {\\em not} suffer significant kicks at birth (see also \\cite{ibe96,cor98})."
    },
    "0004/astro-ph0004128_arXiv.txt": {
        "abstract": "A sample of 4208 objects with magnitude $15 < g^* < 22$ and colors of main sequence A stars have been selected from 370 square degrees of Sloan Digital Sky Survey (SDSS) commissioning observations.  The data is from two long, narrow stripes, each with an opening angle of greater than 60$^\\circ$, at Galactic latitudes $36^\\circ < \\vert b\\vert < 63^\\circ$ on the celestial equator.  Relative photometric calibrations good to 2\\% and consistent absolute photometry allows this uniform sample to be treated statistically over the large area. An examination of the sample's distribution shows that these stars trace considerable substructure in the halo.  Large overdensities of A-colored stars in the North at $\\rm (l,b,R) = (350,50,46\\> kpc)$ and in the South at $\\rm (157,-58, 33\\> kpc)$ and extending over tens of degrees are present in the halo of the Milky Way.  \\citet{ietal00} has detected the northern structure from a sample of RR Lyraes in the SDSS.  Using photometry to separate the stars by surface gravity, both structures are shown to contain a sequence of low surface gravity stars consistent with identification as a blue horizontal branch (BHB).  Both structures also contain a population of high surface gravity stars two magnitudes fainter than the BHB stars, consistent with their identification as blue stragglers (BSs).  The majority of the high surface gravity stars in the Galactic halo may be blue straggler stars like these.  A population of F stars associated with the A star excess in the southern structure is detected (the F stars in the northern structure at 46 kpc would be too faint for the SDSS to detect).  From the numbers of detected BHB stars, lower limits to the implied mass of the structures are $6\\times 10^6 M_\\odot$ and $2\\times 10^6 M_\\odot$, though one does not yet know the full spatial extent of the structures.  The fact that two such large clumps have been detected in a survey of only 1\\% of the sky indicates that such structures are not uncommon in the halo.  Simple spheroidal parameters are fit to a complete sample of the remaining unclumped BHB stars and yield (at $r<40$ kpc) a fit to a halo distribution with flattening ($c/a = 0.65\\pm 0.2$) and a density falloff exponent of $\\alpha = -3.2 \\pm 0.3$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0004/astro-ph0004285_arXiv.txt": {
        "abstract": "First, we review the current status of the detection of strong `external' variability in the CLASS gravitational B1600+434, focusing on the 1998 VLA 8.5--GHz and 1998/9 WSRT multi-frequency observations. We show that this data can best be explained in terms of {\\it radio-microlensing}.  We then proceed to show some preliminary results from our new multi-frequency VLA monitoring program, in particular the detection of a strong feature ($\\sim$30\\%) in the light curve of the lensed image which passes predominantly through the dark-matter halo of the lens galaxy. We tentatively interpret this event, which lasted for several weeks, as a {\\it radio-microlensing caustic crossing}, i.e. the superluminal motion of a $\\mu$as-scale jet-component in the lensed source over a single caustic in the magnification pattern, that has been created by massive compact objects along the line-of-sight to the lensed image. ",
        "introduction": "Optical microlensing has unambiguously been detected both in our own Galaxy (e.g. the EROS, MACHO and OGLE collaborations; see these proceedings), as well as in external galaxies (e.g. Q2237+0305; Irwin et al. 1989). However, no convincing detections of microlensing in other wavebands have been claimed thus far. In our Galaxy this is mostly due to the very low surface number density of bright sources, other than stars, that are compact enough (i.e. $\\la$few mas) to be significantly microlensed. For cosmological microlensing there are similar arguments. Because the angular Einstein radius of a point-mass scales approximately as $\\propto$$D^{-1/2}$, where $D$ is the distance to the lens, the lensed source must be $\\la$few\\,$\\mu$as in angular size to by microlensed by compact stellar-mass lenses at cosmological distances (few Gpc). It was thought until recently that only the optical to $\\gamma$-ray emitting regions of quasars and AGNs had these extremely small angular scales and would still be observable over cosmological distances.  In this proceeding we will show, however, that `radio-microlensing' has most likely been detected in the CLASS gravitational lens B1600+434 (Jackson et al. 1995) and that it promises to be a new and exciting technique for the study of massive compact objects in galaxies at intermediate redshifts. ",
        "conclusions": "The CLASS gravitational lens B1600+434 shows strong `external' variability in the VLA 5 and 8.5-GHz light curves of lensed image A, which passes predominantly through the halo of the lens galaxy. Neither Galactic scintillation nor Extreme Scattering Events (ESE) can satisfactorily explain all the observations. Only `radio microlensing' can explain the observed variability time-scales and amplitudes, as well as its frequency dependence, without invoking extreme assumptions. It also offers a natural explanation for the strong ($\\sim$30\\%) event that we more recently observed, as being a `radio-microlensing caustic crossing'.  Based on the 1998 VLA and 1998/9 WSRT data, we already attempted to place constraints on the mass-function of compact objects in the lens-galaxy halo, finding a lower limit of $\\ga$0.5\\,M$_\\odot$ (KdB00).  However, these constraints are still weak and depend strongly on assumptions about scatter-broadening in image B.  The way foreward is to use {\\sl only} the lightcurve of image A to put constraints on the mass function in its line-of-sight.  With the detection of distinct radio-microlensing events, this task has become a realistic possibility. We have started to work on this (Koopmans et al. in prep) and hopefully we will soon be able to place significantly better constraints on the mass function of compact objects in the lens-galaxy halo, bringing cosmological radio-microlensing up to Galactic standards!"
    },
    "0004/hep-ph0004159_arXiv.txt": {
        "abstract": "Critical slowing down of the relaxation of the order parameter has phenomenological consequences in early universe cosmology and in ultrarelativistic heavy ion collisions. We study the relaxation rate of long-wavelength fluctuations of the order parameter in an $O(N)$ scalar theory near the critical point to model the non-equilibrium dynamics of critical fluctuations near  the chiral phase transition. A lowest order perturbative calculation (two loops in the coupling $\\lambda$) reveals the breakdown of perturbation theory for long-wavelength fluctuations in the critical region and the emergence of a hierarchy of scales with hard $q\\geq T$, semisoft $T\\gg q \\gg \\lambda T$ and soft $\\lambda T\\gg q$ loop momenta which are widely separated in the weak coupling limit. A non-perturbative resummation is implemented to leading order in the large $ N $ limit which reveals the infrared renormalization of the static scattering amplitude and the crossover to an effective three dimensional theory for the soft loop momenta near the critical point. The effective three dimensional coupling is driven to the Wilson-Fisher three dimensional fixed point in the soft limit. This resummation provides an infrared screening and for critical fluctuations of the order parameter with wave-vectors $ \\lambda T\\gg k \\gg k_{us} $ or near the critical temperature $\\lambda T \\gg m_T \\gg k_{us}$ with the {\\bf ultrasoft} scale $k_{us} = \\frac{\\lambda T}{4\\pi}\\exp{[-\\frac{4\\pi}{\\lambda}]}$ the relaxation rate is dominated by {\\em classical} semisoft loop momentum leading to $ \\Gamma(k,T)= \\lambda T/(2\\pi N) $. For wavectors $ k\\ll k_{us} $ the damping rate is dominated by hard loop momenta and given by $ \\Gamma(k,T)= 4\\pi T/\\left[3N \\ln(T/k)\\right] $. Analogously, for homogeneous fluctuations in the ultracritical region $ m_T \\ll k_{us} $ the damping rate is given by $  \\Gamma_0(m_T,T)=4\\pi T/\\left[3N \\ln(T/m_T)\\right] $. Thus critical slowing down emerges for ultrasoft fluctuations. In such regime the rate is independent of the coupling $\\lambda$ and both perturbation theory and the classical approximation within the large $N$ limit break down. The strong coupling regime and the shortcomings of the quasiparticle interpretation are discussed. ",
        "introduction": "The program of relativistic heavy ion collisions both at Brookhaven and at Cern seeks to understand the phase diagram of QCD in conditions of temperatures that were achieved during the first 10$\\mu$s after the Big Bang or densities several times that of nuclear matter which could exist at the center of neutron stars. Current theoretical understanding\\cite{QCD} leads to the conclusion that QCD could undergo two phase transitions: a confinement-deconfinement (or hadronization) and the  chiral phase transition. Current lattice data seem to suggest that both occur at about the same temperature $T_c \\approx 160 \\mbox{MeV}$\\cite{QCD}. The consensus emerging in the field is that several types of observables will have to be studied simultaneously and event-by-event analysis of data will have to be carried out to extract unambiguous signals both hadronic and electromagnetic to reveal the presence of a Quark-Gluon Plasma phase. Recent results reported from CERN-SPS\\cite{cern}  seem to indicate a strong evidence for the existence of the  QGP in Pb-Pb collisions, and RHIC at Brookhaven will begin operation soon with Au-Au collisions with four dedicated detectors capable of event-by-event analysis of hadronic and electromagnetic observables.  For QCD with only two flavors of massless quarks (u,d) it has been argued\\cite{o4pis,o4raja} that the chiral phase transition at finite temperature but vanishing baryon number density is of second order and described by the universality class of O(4) Heisenberg ferromagnets. It has also been suggested recently that at finite baryon density there is a second order critical point described by the Ising universality class\\cite{ising}. Second order critical points are characterized by strong critical long-wavelength fluctuations and a diverging correlation length that  could lead to important experimental signatures\\cite{QCDPT}. These signatures would be akin to  critical opalescence near the critical point in binary fluids\\cite{QCDPT} and could be observed in an event-by-event analysis of the fluctuations of the charged particle transverse momentum distribution (mainly pions)\\cite{QCDPT}. These fluctuations are characterized by the typical correlation length of the order parameter and it has been  suggested that the phenomenon of critical slowing down, ubiquitous near the critical point of second order phase transitions, can lead to strong departures from equilibrium that will determine the value of the correlation length when long-wavelength fluctuations freeze out\\cite{rajaslow}. Critical slowing down of long-wavelength fluctuations near a second order critical point is the statement that the long-wavelength Fourier components of the order parameter relax very slowly towards equilibrium\\cite{MA}. In mean-field theory in {\\em classical critical phenomena},  the relaxation time diverges proportional to the susceptibility near the critical temperature but thermal fluctuations renormalize the relaxation time to be of the form $\\tau(\\vec k=0) \\propto \\xi^z$ with  $\\xi$ the correlation length, or at critical point $\\tau(k) \\propto k^z$ with $z$ a dynamical critical exponent  \\cite{MA,hoh}. Another similar manifestation of an anomalously slow relaxation of long-wavelength fluctuations arises in weakly first order phase transitions when the system enters into the mixed phase where the (isothermal) speed of sound, which determines the velocity of propagation of long-wavelength pressure waves, becomes anomalously small resulting in the softest point of the equation of state. In this case there have also been suggestions that there are experimental consequences of this softening in relativistic heavy ion collisions in observables  related to collective flow and the transverse momentum distributions of particles at freezeout\\cite{shuryak}.  In classical normal fluids near the critical point  the vanishing of the (isothermal) speed of sound, critical opalescence (strong scattering of light by long-wavelength fluctuations) and critical slowing down are all related\\cite{forster}, and in ferromagnets  the spin diffusion constant vanishes near the critical point again signaling critical slowing down\\cite{forster}.  The softening of the equation of state near the critical point of QCD could also have important cosmological implications. When the the QGP enters the mixed phase with hadrons, the speed  of sound becomes anomalously small and the time scale for propagation of pressure waves over a given critical wavelength becomes longer than the free fall time for gravitational collapse which is then unhindered by the pressure of the hadronic gas. This could lead to the formation of primordial black holes\\cite{jedamzik} with a possible imprint in the acoustic peaks in the cosmic microwave background\\cite{widerin}. Other possible cosmological relics from the QCD phase transition with a mixed phase had been predicted, from strange quark nuggets to MACHO's\\cite{raja,mau}.  A familiar argument is typically invoked to state that while the QCD phase transition in the Early Universe occurred in local thermodynamic equilibrium (LTE) this {\\em may} not be the case in Relativistic Heavy Ion Collisions. The argument compares the typical collisional relaxation time scale obtained from a  strong interaction process $\\tau_{coll} \\sim 10^{-22}~\\mbox{secs}$ to the time scale for cooling near the critical temperature $\\sim 160 \\mbox{MeV}$, i.e. ${T/ \\dot{T}} \\sim H^{-1} \\sim 10^{-5}~\\mbox{secs}$. The argument is that since $\\tau_{coll} \\ll   H^{-1}$ the phase transition occurs in LTE in cosmology, whereas in relativistic heavy ion collisions at RHIC and LHC energies these time scales will be comparable. However this argument completely neglects the possibility that long-wavelength fluctuations could undergo critical slowing down and freeze out, i.e. fall out of local thermal equilibrium, even {\\em before} the phase transition. The freeze-out of long-wavelength fluctuations during the phase transition could  result in important non-equilibrium effects on the size and distribution of primordial black holes or any other cosmological relic just as they could lead to important observables in the momentum distributions of charged pions in relativistic heavy ion collisions\\cite{QCDPT,rajaslow,shuryak}.  Indeed there are simpler experimental situations where this is the case, in typical normal fluids a collisional relaxation time (away from the critical point) is of the order of $10^{-9}~\\mbox{secs}$ while near the critical point (even at $10\\%$ of the critical temperature) critical slowing down becomes very dramatic and thermalization time scales become of the order of minutes if not hours\\cite{forster,stanley,boyleshouches}.  Thus the phenomenological importance of critical slowing down for the QCD phase transition both in Relativistic Heavy Ion Collisions as well as in Early Universe Cosmology motivates us to study this phenomenon in a model Quantum Field Theory that bears on the low energy (chiral) phenomenology of QCD, the $O(N)$ linear sigma model. Furthermore the study of critical slowing down is the precursor to a more complete program to understand transport phenomena  and the relaxation of hydrodynamic modes at or near a critical point.  {\\bf Goal:} Our goal is to provide a consistent microscopic description of critical slowing down at or near criticality directly from an underlying quantum field theory that is at least phenomenologically motivated to study the QCD phase transitions. This will be a first step in a program that seeks to offer a consistent description of transport near critical points that eventually may be merged with a hydrodynamic description to obtain a more reliable picture of critical phenomena near the deconfinement and chiral phase transitions and an assessment of the potential phenomenological observables both in early universe cosmology as well as in relativistic heavy ion collisions.  We begin this program in this article by focusing on the relaxation rate of long-wavelength fluctuations of the order parameter  at and near  the critical point in a consistent non-perturbative framework.  {\\bf Strategy:} We begin our study of critical slowing down by analyzing the relaxation rate of long-wavelength fluctuations of the order parameter at and near the critical point in  an $O(N)$ scalar field theory, which is a phenomenological arena to study the relaxation of sigma mesons and pions. Our first step is to obtain the relaxation rate to lowest order in perturbation theory (two loops). This calculation reveals clearly the breakdown of the perturbative expansion for long-wavelength fluctuations at or near the critical point as a result of the strong infrared behavior for soft loop momentum and the necessity for a non-perturbative treatment. We then implement a non-perturbative resummation of bubble-type diagrams via the large $ N $ approximation to obtain the damping rate in the next-to-leading order in the large $ N $ limit. The resummation implied by the large $ N $ limit to order $1/N$ is akin to that obtained via the renormalization group with the one loop beta function and reveals the softening of the scattering amplitude and the  crossover to an effective three dimensional theory for momenta $q   \\ll  \\lambda T$ with $\\lambda$ the quartic coupling.  {\\bf Summary of main results:}  We have obtained the relaxation rate for long-wavelength fluctuations of the order parameter at the critical point and for homogeneous fluctuations near criticality both to lowest order in perturbation theory (two loops) and near the critical point  to next to leading order in the large $N$ limit.  The two-loop results for the relaxation rate for a fluctuation of wavevector $\\vec k$ of the order parameter at the critical point is found to be $\\Gamma(k,T) \\propto \\lambda^2 T^2/k$ whereas near the critical point, homogeneous fluctuations (with $\\vec k=0$) relax with a rate $ \\Gamma_0(m_T,T) \\propto \\lambda^2 T^2/m_T $. Here $ m_T \\propto |T-T_c|^{1/2}\\ll T_c$ is the effective thermal mass. These results clearly reveal the breakdown of the perturbative expansion in the long wavelength limit $ k \\rightarrow 0 $ at $ T = T_c $ and for $ T \\to T_c $ and $ k = 0 $.  A detailed analysis of the different contributions to these results for the relaxation rate shows that the rate is dominated by very soft  loop momentum $q   \\ll  \\lambda T$ which in the weak coupling limit $\\lambda    \\ll  1$ are classical. The implementation of a non-perturbative resummation via the large $ N $ limit explicitly leads to  an effective scattering amplitude that vanishes in the static long-wavelength limit as a consequence of the crossover to a three dimensional theory for loop momenta $q   \\ll  \\lambda T$. This effective scattering amplitude allows us to recognize  that the effective  {\\em three dimensional} coupling for soft momenta approaches the three dimensional non-trivial (Wilson-Fisher) fixed point in the long-wavelength limit near the critical point. The large $N $ resummation for the relaxation rate incorporates this effective three dimensional coupling in the spectral density that determines the imaginary part of the retarded self-energy for the order parameter. Since the effective three dimensional coupling is driven to its fixed point at long wavelength,  the contribution from very soft loop momenta $ q \\ll \\lambda T$ which give the strongest infrared behavior in lowest order in perturbation theory is effectively screened by this renormalization of the coupling. Consequently the most important contribution to the relaxation rate arises both from the {\\em semisoft classical} region of loop momentum $ T   \\gg   q   \\gg   \\lambda T $ and also from the {\\em hard} region $ q\\geq T $. A detailed analysis of the contribution from the loop momenta reveals a non-perturbative {\\bf ultrasoft scale} $$ k_{us}\\simeq{\\lambda \\; T \\over 4 \\pi }e^{-{4 \\pi\\over \\lambda}} \\; . $$ We find that for soft momenta $ k \\gg k_{us} $ the damping rate is dominated by classical semisoft loop momenta and given by \\be \\label{gamacli} \\Gamma(k,T) \\buildrel{k \\gg k_{us}}\\over= { \\lambda T \\over 2\\pi N } \\left[ 1+ {\\cal O}\\left( {1 \\over \\ln{\\lambda T \\over k}}\\right)\\right]\\; . \\ee For $ k \\,, m_T \\ll k_{us} $ the classical approximation breaks down and the damping rate at the critical point $ m_T=0 $  for $ k \\ll k_{us} $ is given by \\be \\label{gamaint} \\Gamma(k,T) \\buildrel{k \\ll k_{us}}\\over= { 4\\,\\pi \\, T \\over 3\\, N \\ln{T\\over k}}\\left[1+{\\cal O}\\left({1\\over \\ln{T\\over k}}\\right) \\right] \\; . \\ee For homogeneous  fluctuations near the critical point ($ k=0,~~m_T \\propto |T-T_c|^{1/2}\\neq 0$) the damping rate is given by \\be \\label{gamaint2} \\Gamma_0(m_T,T) \\buildrel{m_T \\ll k_{us}}\\over={ 4\\,\\pi \\, T \\over 3\\, N \\ln{T\\over m_T}}\\left[1+{\\cal O}\\left({1\\over \\ln{T\\over m_T}}\\right) \\right]\\; . \\ee Thus critical slowing down, i.e,  the vanishing of the quasiparticle width $ \\Gamma $ for long-wavelengths  emerges in the ultrasoft limit $ k\\ll k_{us}  $  or very near the critical point $ m_T \\ll k_{us} $ where it vanishes logarithmically  slow in the $ k, m_T \\rightarrow 0 $ limit to this order in $ 1/N $. Notice that in such regimes the rate is independent of the coupling $ \\lambda $.  The large $ N $ approximation is not limited to weak coupling and our results apply just as well to a strong coupling case $ \\lambda \\geq 1 $ wherein we find that the  relaxation rate is given by eq.(\\ref{gamaint})-(\\ref{gamaint2}). However this analysis clearly reveals that for weak coupling there emerges a {\\em hierarchy} of widely separated scales for loop momenta: from hard $ q\\geq T $ to semisoft $ T   \\gg   q   \\gg   \\lambda T $, and soft $ \\lambda T   \\gg   q $ that lead to different contributions to the relaxation rate. Which is the relevant scale for the damping rate is determined by the wavevector of the fluctuation of the order parameter and the proximity to the critical temperature. For $ k~,~ m_T \\gg k_{us} $ the classical approximation does apply and the damping rate is dominated by the soft and semisoft classical loop momenta [with the result (\\ref{gamacli})], whereas for $ k~,~m_T \\ll k_{us} $ the classical approximation breaks down and the damping rate is dominated by hard loop momenta $ q\\geq T $ [with the results (\\ref{gamaint})-(\\ref{gamaint2})].  A similar hierarchy exists in non-abelian plasmas\\cite{lebellac}-\\cite{yaffe} and we compare and contrast our results in the scalar theory with those in the hard thermal loop approximation in abelian and non-abelian plasmas\\cite{lebellac}-\\cite{yaffe}.  This article is organized as follows: in section II we introduce the model, obtain the real-time equation of motion for the order parameter and describe the strategy  followed to obtain the relaxation rate. In section III we carry out a perturbative analysis of the relaxation rate to two loops order, recognize the breakdown of perturbation theory and compare to the case of the hard thermal loop resummation program in gauge theories. In sections IV and V we introduce the large $ N $ limit, obtain the effective static scattering amplitude in leading order in the large $ N $ and discuss the dimensional crossover for soft momenta and the effective three dimensional coupling being driven to the three dimensional fixed point. We then use these results to obtain the relaxation rate and  near criticality to order ${\\cal O}(1/N)$ in the large $ N $ limit and explicitly discuss the screening of the soft loop momenta. The contribution from classical soft and semisoft momenta and that of hard loop momenta are analyzed separately to highlight the important differences. In this section we discuss further the validity of a quasiparticle interpretation of the collective long-wavelength fluctuations of the order parameter. In section VI we summarize our conclusions and results and discuss the next step of the program. In appendix A the equations of motion in the large $ N $ limit are derived and in appendix B the polarization integral is computed. ",
        "conclusions": "In this article we have began the program of studying transport and relaxation at and near the critical point in second order phase transitions.  The focus here is to provide a systematic study of critical slowing down from first principles in a phenomenologically motivated quantum field theory. The ultimate  goal of this program  is to assess the potential experimental signatures associated with the critical slowing down of long wavelength fluctuations at the chiral phase transition. Obtaining a robust understanding of such phenomena will have important implications in the QGP and or chiral phase transitions in early universe cosmology and ultrarelativistic heavy ion collisions.  Our study reveals novel phenomena that require a non-perturbative framework for their consistent and sistematic treatment which is different from the hard thermal loop resummation program used in gauge theories.  Whereas critical slowing down has been studied thoroughly in {\\em classical} critical phenomena\\cite{MA,hoh} and these results were used for  preliminary estimates of the correlation length at freezeout in heavy ion collisions\\cite{rajaslow}, we are  aware of only one prior attempt\\cite{pietroni} to study critical slowing down in a full relativistic quantum field theory and a similar recent analysis\\cite{newrenor}. In references\\cite{pietroni,newrenor}  the Wilson renormalization group was used to explore the relaxation of the $ k=0 $ mode of the order parameter slightly away from criticality working at $ N=1 $. The final results of \\cite{pietroni,newrenor} are that for $ t =(T-T_c)/T_c $ approaching the critical limit the damping rate for homogenous configurations vanishes as $ t^\\nu\\log t$ with $\\nu\\simeq 0.5-0.6 $. The analysis of\\cite{pietroni,newrenor}, relies on a truncation of the exact renormalization group equations and their numerical evolution. In our opinion there are two very important limitations in this approach: i) in refs.\\cite{pietroni,newrenor} the absorptive parts were treated in a rather simplistic manner and associated with the scattering vertex rather than the self-energy, but more importantly, this simplified treatment does not include consistently the Landau damping and multiple particle thresholds that are the important ingredients in a consistent and sistematic description of damping and relaxation. ii) the method relies on the numerical evolution of a set of equations that had been truncated without a clear control or understanding of the errors induced by this truncation. In particular, it is not clear if the result $ \\Gamma(t) \\sim t^\\nu\\log t $ is stable with respect to different truncations.  In our study we have systematically focused on the important aspects associated with Landau damping, many-particle threshold effects and  a consistent study of real-time phenomena at finite temperature. As is evident in our study of absorptive parts of the self-energy in sections II and III  a simplified treatment that does not include consistently these can hardly reveal the rich hierarchy of scales and the different physics associated with these: the soft scale $ k_s \\simeq \\lambda T/(4\\pi) $ and the ultrasoft scale $  k_{us} \\simeq\\lambda T\\exp[-4\\pi/(\\lambda T)] $. Consistently to next to leading order in the large $N$ limit  we see that  slowing down of relaxation of long-wavelength fluctuations only begins to emerge at the ultra-soft scale and in contrast to the results obtained in\\cite{pietroni,newrenor}, we obtain $ \\Gamma \\sim T/[N\\log t] $. Thus, although there is agreement on the statement that the relaxation rate vanishes at criticality the consistent large $N$ resummation leads to a very different behavior of the relaxation rate.  \\bigskip  Our main results can be summarized as follows: a consistent treatment of critical slowing down and of transport phenomena at or near a critical point requires a non-perturbative framework to resum the contributions from {\\em soft} loop momenta which is different from the hard thermal loop program of abelian and non-abelian gauge plasmas. A perturbative two loop calculation reveals clearly the emergence of a hierarchy of loop momentum scales from hard $ q\\geq T $ to semisoft $ T \\gg q \\gg \\lambda T$ and soft $ \\lambda T \\gg q $,  for weak coupling $ \\lambda \\ll 1 $ the scales in this hierarchy are widely separated and the semisoft and soft scales are classical. Recognizing the shortcomings of a perturbative treatment for long-wavelength fluctuations, we implemented a non-perturbative resummation via the large $N$ limit to next to leading order. The large $ N $ limit provides a consistent non-perturbative framework for resummation of infrared contributions. It clearly displays the infrared renormalization of the scattering amplitude in the static limit at or near the critical point and the crossover to three dimensional physics for soft loop momentum. The resummation of the scattering amplitude leads to an effective three dimensional coupling that interpolates between the bare coupling for loop momenta $ q \\gg \\lambda T $ and the three dimensional Wilson-Fisher fixed point for $ q \\ll \\lambda T $.  The infrared renormalization of the effective coupling screens the contribution from soft loop momentum to the self-energy and the relaxation rate, which is now dominated by the contribution of {\\em semisoft and hard loop momenta} $T \\gg q \\gg \\lambda T$. Furthermore a new {\\em ultrasoft} (in the weak coupling limit) non-perturbative scale emerges $ k_{us} \\simeq { \\lambda \\; T \\over 4 \\pi }e^{-{4 \\pi\\over \\lambda}} $ that signals the breakdown of the classical approximation and the dominance of hard loop momentum modes.  For $ k\\;, m_T  \\gg k_{us} $ the damping rate is dominated by the classical semisoft scales and given by $ \\Gamma(k,T) = \\frac{\\lambda T}{2N\\pi} $ whereas for $ k \\leq  k_{us} $ the hard loop momenta region dominates and leads to the damping rate $ \\Gamma(k,T) = { 4\\,\\pi \\, T \\over 3\\, N \\ln{T \\over k}} $ at criticality or $ \\Gamma_0(m_T,T) = { 4\\,\\pi \\, T \\over 3\\, N \\ln{T \\over m_T}} $  near criticality for homogeneous fluctuations,  which reveal the slowing down of relaxation of critical ultrasoft fluctuations with a damping rate that is {\\em independent of the coupling}.  As discussed above, however, these results and those found in\\cite{pietroni,newrenor} seem to indicate a breadown of the quasiparticle picture of collective excitations of the order parameter because the resonance parameter $ \\Gamma(k,T)/\\omega_p(k) \\gg 1 $ in the long-wavelength limit and the excitation decays on time scales much shorter than the natural oscillation time $ \\omega^{-1}_p(k) $. At this stage it is not clear if this feature is a true physical manifestation of relaxation of collective excitations at or near the critical point or that further resummation  and other contributions that are beyond the leading order in the large $ N $  must be accounted for. We are currently studying this possibility by introducing the renormalization group at finite temperature and analyzing in detail the contribution from vertex and wave function renormalizations and expect to report on further understanding on these issues in a forthcoming article. At this stage our study has revealed a wealth of new phenomena and a hierarchy of scales which will require a deeper understanding for a complete and consistent treatment of transport and eventually hydrodynamics near or at the critical point. Only a thorough understanding of these phenomena can lead to an unambiguous assessment of the phenomenological implications of critical fluctuations either in the formation of cosmological relics in the early universe or in experimental observables in ultrarelativistic heavy ion collisions thus motivating and justifying their study.  \\bigskip  {\\large\\bf Acknowledgements}\\\\  D. B. thanks the N.S.F for partial support through grant awards: PHY-9605186 and INT-9815064 and LPTHE (University of Paris VI and VII) for warm hospitality and partial support, he also thanks S. Raja for interesting conversations.  H. J. de Vega thanks the Dept. of Physics at the Univ. of Pittsburgh for hospitality.  We thank the CNRS-NSF cooperation programme  for partial support. The early stages of the work of M.S. were supported by a grant from Padova University.  \\appendix"
    },
    "0004/astro-ph0004236_arXiv.txt": {
        "abstract": "We report the discovery of quasi-periodic oscillations (QPOs) at roughly 187 Hz and 150 Hz in the X-ray intensity of X-ray nova XTE J1859+226. The source was observed during a recent outburst with RXTE. Besides these high-frequency QPOs, we have also detected QPOs (and sometimes their harmonics) at 6--7 Hz, and significant broad-band variability at low frequencies. These properties, as well as the observed hard X-ray spectrum, make XTE J1859+226 a black hole candidate (BHC).  The detection of QPOs at two distinct frequencies $\\gtrsim 100$ Hz in two observations separated by about 4 hours provide additional insights into the high-frequency QPO phenomenon in BHCs. The importance lies in the proposed interpretations which invariably involve the effects of strong gravity near a black hole. We compare our results to those of other BHCs, and discuss the impact of the observational data on the models in a global context. ",
        "introduction": "A new X-ray nova, designated as XTE J1859+226, was discovered by the ASM/RXTE on October 9, 1999 (Wood et al. 1999). Subsequent observations (2-60 keV) of the source with the PCA/RXTE revealed a hard X-ray spectrum roughly of power-law shape (Markwardt et al. 1999). The source was detected up to 200 keV (McCollough \\& Wilson 1999; Dal Fiume et al. 1999). It is interesting to note that the hard X-ray flux of the source (as measured by the BATSE/CGRO; McCollough \\& Wilson 1999) began to decrease while the soft X-ray flux (as measured by the ASM) was still rising. This is remarkably similar to another X-ray nova XTE J1550-564 (Cui et al. 1999) and may be common for such objects.  Model fits to the X-ray spectrum of XTE J1859+226 implied a column density about 3--8 $\\times 10^{21}$ $cm^{-2}$ along the line of sight (Markwardt et al. 1999; Dal Fiume et al. 1999). The relatively low foreground extinction facilitated rapid optical follow-up observations. The optical counterpart (Garnavich et al. 1999), as well as the radio counterpart (Pooley \\& Hjellming 1999), was identified soon after the discovery. The optical spectrum of XTE J1859+226 showed characteristics typical of soft X-ray transients during an X-ray outburst (Wagner,~R.~M. et al. 1999; Hynes et al. 1999). The optical/UV continuum spectrum from the HST/STIS observations imply a binary period $<$ 1 day (Hynes et al. 1999). Subsequent optical observations indeed detected a possible period of 9.15 hours (Garnavich \\& Quinn 2000), which are likely associated with the orbital motion.  XTE J1859+226 showed a strong quasi-periodic oscillation (QPO) in X-ray intensity during the rising phase of the outburst. The QPO was first detected at 0.45 Hz (Markwardt et al. 1999) and moved up to $\\sim$5.5 Hz as the outburst proceeded (Dal Fiume et al. 1999), similar to other X-ray novae (Fox \\& Lewin 1998; Cui et al. 1998, 1999; Dieters et al. 2000). The combination of the hard X-ray spectrum and QPOs makes XTE J1859+226 a black hole candidate (BHC). Here we report the detection of QPOs at much higher frequencies. ",
        "conclusions": "The most important result from this work is the detection of QPOs at high frequencies ($\\gtrsim 100$ Hz) in XTE J1859+226. These QPOs most probably originate from a single periodic process that varies in its period. It is remarkable that the period of the process can change on hour timescales. Our detections are corroborated by Markwardt et al. (2000), who found QPOs at $\\sim$174 Hz from their own observations. Very similar behavior is seen in XTE J1550-564: a QPO was first discovered at $\\sim$184 Hz during a giant X-ray flare near the peak of the outburst (Remillard et al. 1999b) and another was later detected at $\\sim$284 Hz at a much lower flux (Homan, Wijnands, \\& van der Klis 1999). These QPOs are, however, in sharp contrast to the high-frequency QPOs seen in microquasars, GRS 1915+105 (Morgan, Remillard, \\& Greiner 1997) and GRO J1655-40 (Remillard et al. 1999a), which were thought to maintain a constant frequency. This implies that there could be intrinsic physical differences underlying such QPOs in BHCs. Alternatively, the different behaviors may merely result from insufficient statistics or sporadic coverage. For XTE J1859+226, the QPOs are stronger at higher energies, in terms of their fractional rms amplitudes, which is typical of nearly all QPOs in BHCs (review by Cui 1999a).  Several models have been proposed to explain the high-frequency QPOs in BHCs, all of which invoke strong gravity near the central black hole (review by Cui, Chen, \\& Zhang 1998). It was first suggested that the QPOs might be associated with the Keplerian motion at the last stable orbit around Schwarzschild black holes (Morgan, Remillard, \\& Greiner 1997). Perhaps the motion of ``blobs'' or ``hot spots'' in an accretion disk extending all the way to the last stable orbit somehow causes the modulation of X-ray emission. However, the emissivity of the disk vanishes at the last stable orbit, due to the torque-free boundary condition, so it is perhaps more sensible to associate the QPOs with orbiting hot spots located in a region where the disk emissivity peaks. While appealing for its simplicity the model is incompatible with the spectral results (Cui, Chen, \\& Zhang 1998). Nowak et al. (1997) subsequently proposed that the QPOs are associated with the oscillation modes in the accretion disk. Not only can such modes be supported by the disk, they also become trapped in the innermost portion of the disk, purely due to relativistic effects (Kato \\& Fukue 1980; Nowak \\& Wagoner 1992, 1993). Finally, Cui et al. (1998) argued that the QPOs could be due to the Lense-Thirring precession of the tilted circular orbits of blobs or hot spots at the inner edge of the accretion disk, motivated by observational evidence for the presence of rapidly rotating black holes in microquasars (Zhang, Cui, \\& Chen 1997).  All the models were initially proposed to explain the ``constant frequency QPOs'' in microquasars. Consequently, the accretion disk was assumed to extend to the last stable orbit in the models. Since the high-frequency QPOs now appear to be common among BHCs, an obvious question to ask now is whether the models can accommodate the QPOs that vary in frequency, as in the case of XTE J1550-564 and XTE J1859+226. For those models that involve Keplerian motion or Lense-Thirring precession, ``variable frequency QPOs'' can be explained if the inner edge of the accretion disk is allowed to move. Because the precession frequency in the Lense-Thirring model depends strongly on the radius of the orbit (roughly, $f \\propto r^{-3}$), a $\\sim$35\\% change in the QPO frequency (as in the extreme case of XTE J1550-564) requires that the inner edge of the disk moves radially outward only by $\\sim$15\\%. Therefore, the variable frequency QPOs would pose least serious challenges to this model. In comparison, the dependence of the Keplerian frequency on the orbital radius is not as steep; the inner edge of the disk would be required to move about twice as far in the Keplerian motion model. The problem is far more serious (if not fatal) for the disk oscillation model, because the oscillation modes tend to be fairly well localized in the accretion disk. While much improved statistics are clearly needed to ultimately establish the physical origin of the high-frequency QPOs in BHCs, it is clear that the discovery of variable frequency QPOs has begun to seriously constrain the models.  An additional similarity between XTE J1859+226 and XTE J1550-564 is the presence of a prominent and evolving QPO at a few Hz during the rising phase of the outburst. However, the QPO is almost certainly of a different physical origin for the two sources, because the phase lag associated with the fundamental component is of an {\\em opposite} sign (cf. Cui, Zhang, \\& Chen 2000). While the complicated QPO lags in XTE J1550-564 are likely the same phenomenon as was first discovered in a subset of the QPOs of GRS 1915+105 (Cui 1999b), the hard QPO lag in XTE J1859+226 makes the QPO here similar to a second type of the low-frequency QPOs in XTE J1550-564, which were detected at a later stage of the outburst and also showed a sub-harmonic component (Wijnands, Homan, \\& van der Klis 1999). The phase lag measurement has, therefore, added new puzzles to the already very intriguing QPO phenomenology in BHCs (see Cui 1999b for a discussion of possible implications of different QPO lags on the models).  It has been argued, based on the observed dependence of QPO properties on mass accretion rate (as indicated by X-ray flux) and photon energy, that the QPOs in BHCs form a heterogeneous class of phenomena, each of which may well be produced by a different physical mechanism (Cui 1999a; Cui et al. 1998). For a given source, the QPOs are known to come and go in an unpredictable manner, and a different set of QPOs may dominate the PDS at different times (see an extreme example offered by GRS 1915+105; Morgan, Remillard, \\& Greiner 1997). Here we have shown that such changes took place on time scales as short as roughly one hour in XTE J1859+226 (i.e., the time between Obs. 3 and Obs. 2 or 4). The source clearly behaved very differently during Obs. 3, both in terms of the QPOs and broad-band noise (see Fig.~2). It is interesting that the timing properties (low- and high-frequency QPOs plus broad-band variability) all seem to be correlated. A plausible model must, therefore, be able to account for such a global change."
    },
    "0004/gr-qc0004047_arXiv.txt": {
        "abstract": "\\noindent On  the mass neutrino phase calculations along both the particle geodesic line  and the photon null line, there exists a double counting, factor of 2,  when comparing the geodesic phase with the null phase. Moreover,  we compare the phase calculations among the same energy description, the same momentum description by means of the Minkowski diagram, and obtain the practical equivalence of these two descriptions. On the same velocity description, although it does not correspond to a reality of physical process, we still indicate its phase calculation in the Minkowski diagram,  which has the same result as those of same energy and same momentum cases. Further, in the curved spacetime, we also prove the existence of the double counting of the geodesic phase to the null phase. Our conclusions are same as others' results by the different methods. ",
        "introduction": "There have been many topics issued  on the quantum mechanics of the neutrino oscillation in the flat spacetime ~\\cite{zub98,bgg98,pak98,bp78,boe92} and in the curved spacetime \\cite{ful96,for96,koj96,gro96,bha99,zcm20,zcm01,zcm21}, most of them centered on the understanding of the quantum coherence condition between the various mass-eigenstates, as well as on the calculations  of the phase factor in the spacetime. The mass neutrino oscillation problem is connected basically to the phase difference between different mass-eigenstate, a property intimately related to the basic principles of quantum mechanics. In the standard treatment of the neutrino oscillation, the condition of same-energy (same-momentum) with different momenta (energies) of the mass-eigenstates is introduced~\\cite{bgg98,bp78,boe92,dod97,oku00,tsu01}, which arises the same practically applied massive neutrino oscillation phase, however it is indicated that the unambiguous theoretical description of the massive neutrino oscillation phase will be involved in the wave packet formalism and not in the plane wave approximation \\c{dod99}. Further, it is shown that the same energy or same momentum description is just the arbitrariness of the choice of the Lorentz frame in which it is valid,  nothing to do with the physical argument \\c{giu01}. The change of flavors is described by the oscillation length, and it is usually claimed that both conditions (same-energy and same-momentum) present practically the same neutrino oscillation results. Sometimes, a source-dependent condition is added~\\cite{kay81,win82,gol96}, which implies in giving up either the same-energy or the same-momentum prescription. This calculation of the neutrino oscillation phase, however, yields the same result of the standard treatment. Furthermore, the velocity difference of the various mass-eigenstates results in a spacetime separation for neutrinos of different flavor~\\cite{gro96,lip95,nie96}, which is another source of confusion on the interference condition for the different neutrino. The same velocity description has been paid much attention \\c{tak98,leo99,tak00}, but this prescription is pointed out to be forbidden kinematically \\c{oku007}. Actually, considerable confusion has arisen in the description, interpretation and understanding of the neutrino oscillation, a problem which involves the fundamental principles of both quantum mechanics and  special relativity, such as the uncertainty principle, the superposition principle, and simultaneity problems because the mass neutrinos are high energy quantum objects. Moreover it is often noted the factor of 2  of the neutrino phase calculations in the flat spacetime \\cite{lip95,lip98} and in the curved spacetime \\cite{gro96,bha99}, which is believed  to be the consideration of the space  phase or the time phase \\cite{gro96}, as well as the arrival time difference of the two mass neutrinos \\cite{bha99}.   In the Minkowski diagram of flat spacetime, we discuss how  the same-energy,  the same-momentum  of the neutrino phase present the same practical result, and the relation between the arrival time difference and   the double counting of the phase,  and further, we extend our discussions  to the curved spacetime. We also calculate the phase factor in the case of the same velocity description although this process is not physically realized. Here we stress that all our calculations are based on the plane wave treatment of the neutrino, otherwise the wave packet should be considered \\c{giu91}. We set $\\hbar = c = 1$ throughout this article. ",
        "conclusions": "As mentioned above, the same-energy and the same-momentum prescriptions in calculating the mass neutrino interference phase will obtain the same result if the arrival time difference is taken into account, which will be, mathematically,  equivalent to the  phase of a   treatment by using the same-velocity  condition. However, we remark that  generally neither energy nor momentum are equal in the factual physical process \\c{dod99,giu01}, the standard treatment of the same-energy and  the same-momentum prescriptions of massive neutrino phases is just a mathematical simplification in some sense when applying the plane wave approximation. Then our analysis  of the  massive neutrino standard phases in the Minkowski diagram indicates that the same phase factor will be obtained if the two massive neutrinos follow the null line  or possess the same velocity. If the phase calculated along the particle world line, the realtive phase will produce a factor of 2 \\cite{lip98}, the reason is that we despise the arrival time difference of the two mass neutrinos, which results in a double counting effect. This conclusion is correct in both flat spacetime and the curved spacetime. Further, for the better understanding of the mass neutrino interference, we should consider the neutrino mixing state as a  wave-packet \\c{dod99,giu91}, or a relativistic quantum ball. According to this scheme, two physical aspects must be considered. On one hand, there is the classical trajectory of the ball which is defined by giving the initial conditions in the classical sense. This motion is relativistic, leading thus to the relativity of the simultaneity as the velocities of the different mass neutrinos are not the same. The classical trajectory, therefore, is defined in terms of macroscopic quantities of the packet, the quantum average of quantities, in accordance with the correspondence principle of quantum mechanics.  In the classical sense, the average mass and velocity of the packet ball define its orbit. The distance from the source to the detector is usually measured accurately, which is a classical measurement with a macroscopic precision. On the other hand, there is the quantum dynamics of the packet, which is a two-state system like a neutral kaon or a B meson. This fuzzy dynamics takes place in the microscopic scale $d$, the characteristic dimension of the ball, which is small compared to the dimension $L$ of the macroscopic trajectory.  This quantum dynamics, that is, the internal evolution of the packet, is well described and understood in the particle physics context.   As a final comment, we would like to argue that, if the individual mass-eigenstate is supposed to have a well defined classical velocity, its orbit would be well defined at any spacetime position. This, however, is not in accordance with the quantum interference description, which consequently would never occur. This means that the quantum fuzzy is necessary for the wave-packet, a well defined orbit being valid only for the classical (macroscopic) quantum average. According to this point of view, two basic points concerns the description of the neutrino oscillation. For the kinematics, we use the plane-wave description with a Lorentz invariant phase. For the dynamics, we use the wave-packet and fuzzy quantum path. We use, therefore, two concepts for the neutrinos, particle and wave, classical and quantum. The neutrino is described in terms of a spinor wave-packet, and the spacetime translation induces the spinor to precess in the state vector-space. In other words, we take the mixing state as a classical object propagating along the classical world-line with a well defined velocity --- the group velocity. The individual mass-eigenstates propagate with a phase velocity, and only the relative phase velocity will produce the realistic interference phenomenon."
    },
    "0004/astro-ph0004146_arXiv.txt": {
        "abstract": "We study the stellar population of galaxies with active star formation, determining ages of the stellar components by means of spectral population synthesis of their absorption spectra. The data consist of optical spectra of 185 nearby ($z \\leq 0.075$) emission line galaxies (Terlevich et al., 1991). They are mostly H{\\sevensize II} galaxies, but we also include some Starbursts and Seyfert 2's, for comparison purposes. They were grouped into 19 high signal-to-noise ratio template spectra, according to their continuum distribution, absorption and emission line characteristics. The templates were then synthesized with a star cluster spectral base.  The synthesis results indicate that H{\\sevensize II} galaxies are typically age-composite stellar systems, presenting important contribution from generations up to as old as 500 Myr. We detect a significant contribution of populations with ages older than 1 Gyr in two groups of H{\\sevensize II} galaxies. The age distributions of stellar populations among Starbursts can vary considerably despite similarities in the emission line spectra. In the case of Seyfert 2 groups we obtain important contributions of old population, consistent with a bulge.  From the diversity of star formation histories, we conclude that typical H{\\sevensize II} galaxies in the local universe {\\it are not} systems presently forming their first stellar generation. ",
        "introduction": "H{\\sevensize II} galaxies are gas-rich dwarf galaxies undergoing violent star formation characterized by an important population of massive stars with strong narrow emission lines (Melnick et al., 1985, for a review see Melnick, 1992). Their study began with the discovery by Sargent and Searle (1970) that some compact galaxies in Zwicky's Catalogue (1971) had spectra comparable to those of giant H{\\sevensize II} regions.  Many of these objects have a compact morphology with a dominant starbursting region, but some have double or multiple components (Telles et al., 1997). Terlevich et al. (1991) built a catalogue of H{\\sevensize II} galaxy spectra and analysed the properties of the individual objects, such as distributions of redshift, Balmer decrement, H$\\beta$ luminosity and equivalent width. In particular, they derived $-1.4<$[O/H]$<-0.4$, showing that H{\\sevensize II} galaxies are typically metal deficient.  An important question is whether H{\\sevensize II} galaxies are young systems that are forming their first generation of stars, or systems that have multiple stellar generations, including several Gyr old ones. Several studies have searched for an underlying old stellar component but the results were not conclusive (Thuan, 1983; Campbell and Terlevich, 1984; Melnick et al., 1985; Farnell et al., 1988). It would be also worth probing whether non-ionizing blue stellar populations (20--500 Myr) occur and in which proportions.  Brinks and Klein (1988) found that one H{\\sevensize II} galaxy (II Zw 40) may be indeed a young system, recently formed by interaction of two H{\\sevensize I} clouds. On the other hand, Schulte-Ladbeck and Crone (1998) resolved red giant stars in the blue compact dwarf galaxy VII Zw 403 with HST, and concluded that it is not a young galaxy forming its first generation of stars, having an older underlying component. Indeed such red giants could be related to an intermediate or old age component since the colour-magnitude diagram could not distinguish such low turnoffs.  The spectral study of stellar populations requires a high enough signal-to-noise (S/N) ratio to clearly establish the continuum distribution and absorption lines. Averaging spectra of similar properties has proven to be a powerful tool for stellar population analyses of star clusters and galaxies (e.g. Bica, 1988; Bonatto et al., 1995; Santos et al., 1995).  The population synthesis of galaxies using a base of star cluster spectra (Bica, 1988) has provided information on population components in normal galaxies. The method employs a grid of star cluster spectral properties built as a function of age and metallicity (Bica and Alloin, 1986a, 1986b). The synthesis provides flux fractions for the successive generations of star formation and corresponding chemical enrichment. The method was also applied to the dusty metal-rich Starburst galaxy NGC6240 (Schmitt et al., 1996). Recently the algorithm was adapted to the Ultraviolet range (Bonatto et al., 1998, 1999).  In this Paper we use the star cluster synthesis method to investigate the stellar populations of H{\\sevensize II} galaxies and to compare them with those of more luminous emission-line galaxies such as Starbursts and Seyfert 2's. We intend to shed light on the question about H{\\sevensize II} galaxies as single or multigeneration systems. We present the sample and spectral group construction in Section 2. The population syntheses are carried out and discussed in Section 3. The concluding remarks are given in Section 4. We defer consideration of the emission line spectra to a forthcoming paper (Raimann et al., 1999, hereafter Paper II). ",
        "conclusions": "The population syntheses indicate that H{\\sevensize II} galaxies in the local universe as a rule are not forming their first generation of stellar population. The diversity of histories of star formation within H{\\sevensize II} galaxy groups shows that they are typically age-composite stellar systems, presenting important contributions from generations up to as old as 500 Myr. These galaxies may still be evolved objects in a cosmological sense if old components at low levels could be detected.  We find evidence of an old underlying stellar population in the H{\\sevensize II} galaxy groups G\\_UM504 and G\\_UM71. These two groups are probably genuine evolved galaxies.  Finally, we point out that H{\\sevensize II} galaxies with stellar populations up to 500 Myr may have had as many as 100 stellar generations, assuming a constant star formation rate and that a typical stellar generation lasts 5 Myr. This result is basically consistent with the gas enrichment, because H{\\sevensize II} galaxies, although metal deficient, are typically not extremely so, presenting a range of metallicities (Terlevich et al., 1991 and Paper II). Indeed the metals seen in H{\\sevensize II} galaxies must come from previous generations of stars since mixing times of the metals ejected by the current or a few previous generations are far too long to be seen now."
    }
}