{
    "9604/astro-ph9604054_arXiv.txt": {
        "abstract": "Significant 1.8 \\MeV\\ emission from the Carina region has been detected by COMPTEL. The emission is concentrated within 6 degrees or less near the Carina nebula NGC 3372, one of the brightest \\HII\\ regions known in our Galaxy. This region contains a wealth of extreme young open clusters whose massive stars possibly contributed to an enrichment of \\al26\\ in the ISM within the last few million years. The relation of these clusters and the peculiar object \\etacar\\ with the observed emission is discussed. The \\al26\\ yield of the clusters is estimated using current theoretical nucleosynthesis models. ",
        "introduction": "\\al26\\ is the first radioactive isotope which was detected by its \\gray\\ line emission in the interstellar medium (\\cite{rf:mahoney}). With its short lifetime of $\\sim~10^6$ yrs, it is a clear tracer of ongoing nucleosynthesis in our Galaxy. \\al26\\ is assumed to be produced in various sites, such as core-collapse supernovae (SNe), Wolf-Rayet (WR) stars, O-Ne-Mg novae, and asymptotic giant-branch (AGB) stars (see review of \\cite{rf:pd95}). After two years of observations, the imaging telescope COMPTEL aboard the CGRO satellite revealed the first map of the Milky Way in the light of the \\al26\\ decay line at 1.809 \\MeV\\ (\\cite{rf:diehl95a}). The sky map clearly shows that the emission is confined to the Galactic plane. Besides the important emission from the central Galactic radian, 1.8 \\MeV\\ emission is seen in the direction of Cygnus, Vela, and near the anticentre. One of the most prominent features is found at $(l,b)=(286.5\\dg,0.5\\dg)$ in the direction of Carina. Among all features in the map it appears to be the most concentrated one, and additionally it lies in a nearly emission free region of the plane. We present here a detailed study of this region and discuss the potential source candidates. ",
        "conclusions": "Prominent 1.8 \\MeV\\ \\gray\\ line emission which is attributed to the radioactive decay of \\al26\\ was detected from the Carina region by the imaging telescope COMPTEL. We find a remarkable correlation of the 1.8 \\MeV\\ feature with a concentration of open clusters younger than 20 Myr which suggests that WR stars and core collapse SNe are the origin of the observed emission. The presence of evolved objects in the Carina clusters clearly indicates recent \\al26\\ ejection to the ISM. This region demonstrates that OB-associations and young open clusters could be a natural explanation for some of the emission features and `hot-spots' in the COMPTEL 1.8 \\MeV\\ sky-map. Especially noncoeval star-formation can produce veritable starbursts (\\cite{rf:ss95}) which would result in substantial concentrations of \\al26\\ within a localized region. If such a burst occurred within the last million years it could be observable through its strongly peaked 1.8 \\MeV\\ emission. We will investigate further regions of peaked 1.8 \\MeV\\ emission to verify this hypothesis."
    },
    "9604/astro-ph9604112_arXiv.txt": {
        "abstract": "In the present paper we discuss the modifications introduced into the first-order Fermi shock acceleration process due to a finite extent of diffusive regions near the shock or due to boundary conditions leading to an increased particle escape upstream and/or downstream the shock. In the considered simple example of the planar shock wave we idealize the escape phenomenon by imposing a particle escape boundary at some distance from the shock. Presence of such a boundary (or boundaries) leads to coupled steepening of the accelerated particle spectrum and decreasing of the acceleration time scale. It allows for a semi-quantitative evaluation and, in some specific cases, also for modelling of the observed steep particle spectra as a result of the first-order Fermi shock acceleration. We also note that the particles close to the upper energy cut-off are younger than the estimate based on the respective acceleration time scale. In Appendix A we present a new time-dependent solution for infinite diffusive regions near the shock allowing for different constant diffusion coefficients upstream and downstream the shock. ",
        "introduction": "In the test particle approximation, the first-order Fermi shock acceleration with infinitely extended diffusive regions upstream and downstream the shock leads to a power-law particle spectrum with the spectral index  $$\\alpha = {3R \\over R-1} \\qquad , \\eqno(1.1)$$  \\noindent and the acceleration time scale  $$T_{acc} = {3 \\over U_1-U_2} \\, \\left( {\\kappa_1 \\over U_1} + {\\kappa_2 \\over U_2} \\right) \\qquad , \\eqno(1.2)$$  \\noindent where the index \"$1$\" (\"$2$\") indicates respectively the upstream (downstream) quantity, $U$ is the shock velocity, $R \\equiv U_1/U_2$, $\\kappa$ is the spatial diffusion coefficient. For a review of the results referring to the diffusive acceleration mechanism one should consult some of the numerous review papers (e.g. Drury 1983; Blandford \\& Eichler 1987; Berezhko et al. 1988; Jones \\& Ellison 1991).  In discussions of the acceleration process, the background conditions are usually considered with particle diffusion coefficients changing at most moderately with the distance from the shock. As a result, there are infinite diffusive regions for cosmic ray particles. Then, particles are removed from the acceleration region near the shock only due to advection with the general plasma flow far downstream. However, if the waves responsible for particle scattering are created due to the process of cosmic ray streaming instability upstream the shock, e.g., the finite amplitude waves resonant with particles near the spectrum cut-off energy could be created only close to the shock within a finite time available for the acceleration process. The analogous situation is the case if the shock propagates through the finite volume of turbulent plasma. Then, the energetic particles diffusing far from the shock will encounter the conditions enabling them to permanently escape from the shock. A finite extent of the shock wave to the sides and some particular boundary conditions may also allow for such escape. The situation can be qualitatively modelled by introducing the upstream (Berezhko et al. 1988) and/or downstream boundary for the energetic particle escape.  \\begin{figure} \\vspace{6cm} \\caption{The configuration of the shock with the {\\bf u}pstream {\\bf f}ree {\\bf e}scape {\\bf b}oundary (\"u.f.e.b.\") and the {\\bf d}ownstream {\\bf f}ree {\\bf e}scape {\\bf b}oundary (\"d.f.e.b.\"). The situation is presented as seen in the shock rest frame. Between the shock and a given escape boundary there is a diffusive layer with some finite diffusion coefficient $\\kappa$. Behind the boundary one assumes infinite diffusion coefficient allowing for a free streaming to infinity (escape) of particles hitting the boundary.} \\end{figure}  The conditions mentioned above are sometimes discovered in analysing {\\it in situ} measurements near heliospheric shock waves. Then, an attempt to analyse the data within the standard approach, involving equations (1.1-2) to describe the first-order Fermi acceleration, may fail. For example, basing on such an analysis Bialk \\& Dr\\\"oge (1993) rejected the possibility of the first-order acceleration at the considered shock wave and suggested the second-order acceleration downstream the shock to play a role, instead. Unfortunately, to date there is no theory available describing effects of the enhanced particle escape at the acceleration process and discussions of the 'difficult' cases have to be based on numerical methods. As a step forward, in the present paper we develop a simple analytic theory allowing for evaluation of measurements in terms of the diffusive length scales for energetic particles. It can become a starting point for more elaborate computations or numerical modelling. The theory describes modifications introduced into the acceleration process due to finite extent of diffusive regions near the shock (Figure~1). Below, in Section 2 we discuss a simple time-dependent solution of the diffusion equation for cosmic ray particles at the shock with the upstream and/or the downstream escape boundary. To obtain analytic solutions, we consider a simple case with constant diffusion coefficients, leading in the stationary situation to power-law cosmic ray spectra. It enables us to discuss the relation between the spectrum inclination and other parameters of the acceleration process.  In Section 3,  we derive the particle spectral index and the acceleration time scale as a function of the boundary distance from the shock. The introduced particle sinks at escape boundaries lead to steepening of the spectrum accompanied by a substantial -- a factor of two or three for reasonable spectral indices -- decrease of the acceleration time scale. We discuss the dependence of these quantities on the upstream and downstream boundary distance from the shock. Finally, in Section 4, we shortly summarize the results. The presented theory allows one to interpret the above mentioned Bialk \\& Dr\\\"oge (1993) data in terms of the first-order acceleration. We also note that the particles close to the upper spectrum energy cut-off are younger than the estimate based on the respective acceleration time scale valid at infinite times. The new time-dependent solution for the particle distribution function in the case of infinite boundary distance is presented in Appendix A. It is a generalization of the Toptygin (1980) solution for the case of different constant diffusion coefficients upstream and downstream the shock. ",
        "conclusions": "The presented results allow one to model a wider range of cosmic ray spectra as the output of the first-order Fermi acceleration at shock waves, including distributions with a spectral index steeper than the value given by Equation~(1.1). One should note that accelerating particles with a steep spectrum due to finite escape boundaries allows to obtain spectra extended to somewhat higher energies, due to decrease of the acceleration time scale with respect to the one given in Equation~(1.2). Another point to be mentioned here is an apparent asymmetry of the acceleration process with respect to imposing the escape boundary upstream or downstream the shock. We would like to note, however, that the respective influence of these boundaries discussed in the previous section can change if the downstream particle diffusion coefficient is much smaller than the upstream one. In general, the boundary allowing for quick particle escape will always preferentially act to increase the particle spectrum inclination, while the other one will limit in a higher degree the acceleration time scale. Let us also note the fact that a downstream boundary further away than approximately two diffusive length scales (cf. Figures~3) has almost no influence on the particle spectral index at the shock. Therefore physical processes and/or conditions behind that distance are not expected to modify the particle distribution at the shock.  The realistic conditions near astrophysical shocks are expected to involve the diffusion coefficients depending on particle momentum. The same will hold for the particle escape probability, defined by the boundary distance in our model. Therefore, the present results can be used only to a rather general evaluation of the acceleration conditions and compatibility of the observed (possibly non-power-law) spectra and time scales with the shock dynamics. A detailed modelling requires numerical methods and not-frequently available information about the local physical conditions and the boundary conditions.  Another fact of interest for modelling the generation of highest energy particles near the spectrum cut-off should be mentioned in this place. The first such particles to appear at the shock are those which have not spent too much time diffusing far from the shock, in analogy to the case with escape boundaries. Therefore, at a given energy, the time for these particles to appear can be substantially shorter than the respective acceleration time scale given in Equation~(1.2) or (2.24), the one valid for the steady state particle spectrum. This fact is visible in analytic time-dependent solutions (e.g. Drury 1991; our solution in Appendix A).  \\thanks{We are grateful to the anonymous referee for valuable remarks. This work was partly done during the visit of MO to Max-Planck- Institut f\\\"{u}r Radioastronomie in Bonn. He is grateful to Prof. R. Wielebinski and other colleagues from the Institute for their hospitality and valuable discussions. The work of MO was supported from the KBN grant PB 1117/P3/94/06. RS acknowledges partial support by the Deutsche Forschungsgemeinschaft (Schl 201/8-2). }"
    },
    "9604/astro-ph9604168_arXiv.txt": {
        "abstract": "We present the $V$-band globular cluster luminosity functions (GCLFs) of the Fornax Cluster galaxies NGC~1344, NGC~1380, NGC~1399, and NGC~1404. Our observations reach to $V = 24.8$, roughly one magnitude beyond the GCLF turnover $m^0_V$, with $\\sim90\\%$ completeness.  From the amplitude of the galaxy surface brightness fluctuations, we also estimate the number of globular clusters fainter than this cutoff magnitude. The GCLFs of these galaxies are well fitted by Gaussians; the weighted means of their turnover magnitudes and dispersions are $\\langle m^0_V\\rangle = 23.88 \\pm 0.10$ mag and $\\langle\\sigma\\rangle = 1.35\\pm 0.07$ mag.  The assumption of a universal value for the absolute magnitude of the turnover $M^0_V$ places the Fornax cluster $0.13 \\pm 0.11$~mag more distant than Virgo.  However, in light of recent Cepheid and other high-precision distance measurements, as well as ongoing HST observations of GCLFs for the purpose of determining the extra-galactic distance scale, we choose to re-examine the universal GCLF hypothesis. Based on data from groups and clusters of galaxies, we find evidence that $M^0_V$ becomes fainter as the local density of galaxies increases. We speculate on the possible cause of this trend; if it is confirmed, GCLF observations will be less useful for determining distances, but may provide important information for constraining theories of star formation in primordial galaxy halos. ",
        "introduction": "The globular cluster luminosity function\\break (GCLF) is often employed as a standard-candle distance indicator based on the assumption a universal value for its mean, or turnover, magnitude $M^0$ (see \\cite{j92} for a review of the method). Until recently, it was impossible to apply the GCLF method to determine the distances of galaxies further away than Virgo.  Now, with HST and improvements in ground-based seeing and instrumentation, it becomes potentially much more powerful for determining the extra-galactic distance scale.  Furthermore, new Cepheid and other high-precision local distance measures would allow for a firm calibration of the method.  Baum et al.\\ (1995a,b) have used HST to observe the the globular clusters (GCs) of the Coma galaxies NGC 4881 and IC~4051 down to $V=27.6$ and $V=28.4$, respectively.  They derive values of the Hubble constant $H_0$ near 60 km/s/Mpc. Also with HST, Whitmore et al. (1995) studied the GCs of the extremely rich M87 system to two magnitudes beyond $M^0_V$ (the $V$-band GCLF turnover) and derived $H_0 = 78 \\pm 11$ km/s/Mpc. Only a small part of the discrepancy in derived $H_0$ values can be accounted for by the different calibrations used by the two groups. This situation leads one to suspect that the GCLFs themselves may be intrinsically different, especially as there remains no firmly established physical basis for assuming a universal $M^0$.  Previously, there have been suspicions that $M^0$ is different for spirals and ellipticals (\\cite{sh93}), with the root cause of this difference being metallicity variations (Ashman, Conti, \\& Zepf 1995), but $M^0_V$ was assumed not to vary among large ellipticals. As Whitmore et al.\\ candidly remark, this ``crucial assumption'' of a universal GCLF is ``a hypothesis that needs further verification.''  In this Letter, we examine the current state of the universal GCLF hypothesis. First, we present new observations of GCs around four Fornax galaxies. Fornax is an important cluster for testing distance determination methods, as it is spatially much more concentrated than Virgo while being at nearly the same distance (e.g. \\cite{t91}; Ciardullo, Jacoby, \\& Tonry 1993). We find that the GCLF exhibits remarkably little variation for these Fornax galaxies. Next, we use independent distance measurements to galaxies and galaxy groups to compare derived $M^0_V$ values in different environments. We find somewhat startling % evidence that $M^0_V$ becomes fainter as the local density of galaxies increases. Further verification is once again needed, but if the observed trend proves real, it would have major implications for the GCLF method of distance measurement as well as for theories of GC formation. We conclude with a discussion of these implications, in particular how the local galaxy environment may govern the properties of GC populations.~~~~~ ",
        "conclusions": "We have found that the GCLF is remarkably constant within the Fornax cluster, but, as Figure~2 shows, varies with environment. Ashman et al.\\ (1995) suggested that metallicity differences result in $M^0_V$ values which are systematically brighter by $\\sim0.15$~mag for spirals. Large ellipticals usually do have higher metallicity GC populations; however, NGC 4881 in Coma has a GC color/metallicity distribution similar to that of the MW (\\cite{baum1}), yet its $M^0_V$ is very faint.  In addition, the Leo elliptical NGC 3379, with its relatively high metallicity GC population (\\cite{abt94}), has an exceedingly bright $M^0_V$, though with a large uncertainty (\\cite{h90}). Finally, we note that the magnitude of the environmental effect we propose is a factor of 3-6 larger than the Ashman et al. $M^0_V$ metallicity shift.  \\begin{figure} \\epsscale{0.75} \\plotone{fig2.eps} \\caption{\\tenrm\\baselineskip 11pt The GCLF turnover magnitude $M^0_V$ plotted against the velocity dispersion of the host galaxy's environment, used as a measure of the local density.  See text for details. \\label{fig2}} \\end{figure}  \\begin{figure} \\epsscale{0.85} \\plotone{fig3.eps} \\caption{ \\tenrm\\baselineskip 11pt The effects of variable GC creation/de-struction mechanisms. In part (a) the GC creation process, shown as a power-law growing to smaller mass, is universal (dark solid line), and the destruction process, a power law which wipes out low-mass objects, varies with environment (dashed lines).  Since the total number of GCs is the integral under the intersecting creation/destruction lines, this situation results in an anti-correlation between GC specific frequency $S_N$ and GC mean logarithmic luminosity $-M^0$. In part (b) the creation process (dashed lines) varies with environment, while the destruction process (dark solid line) is universal. The result here is a positive correlation between $S_N$ and $-M^0$. \\label{fig3}} \\end{figure}  We suggest that the most straightforward way to produce the present-day near-Gaussian GCLF is to assume that two simultaneous and competing effects were operating when GCs formed: a ``creation'' process which preferentially created low-mass GCs, cutting the mass function off at the high end, and a ``destruction'' process which inhibited the formation of, or quickly destroyed, low-mass GCs. If each process operated in a manner which was independent of the details of the environment, then the  final mass (luminosity) function would be universal, but if one depended more sensitively on environment than did the other, the final mass function would vary. This situation is schematically illustrated in Figure~3, where we use $S_N$ for the GC ``specific frequency'' (number of GCs per unit luminosity of the host galaxy). In Figure~3a, we assume that the creation process is relatively universal, but that the destruction process varies.  This leads to $M^0$ being variable, and predicts an inverse correlation between the number of GCs formed and their mean brightness.  On the other hand, if the destruction process is constant and the creation process is more variable we will again get a variable $M^0$, but with a direct correlation between the number of GCs and their luminosity, as shown in Figure~3b.  Empirically, we think we see evidence for the latter sort of behavior among ``coeval'' galaxies, i.e., those located within the same physical association.  In Virgo, for instance, M87 has a very large $S_N$ and a slightly brighter $M^0$ than its close neighbors, and similarly in Fornax for NGC~1399. In this context, the GCLF would depend on the extent to which the host galaxy dominated its local environment. On the other hand, the main point of this paper is that we see the former behavior among very ``heterogeneous'' systems of galaxies.  Young groups dominated by spirals have fewer GCs than galaxy clusters such as Virgo, which in turn may have fewer than rich clusters such as Coma, and we find that the central luminosity of the GCLF is declining along this sequence.  As an example of how such an interplay of opposing processes might work in practice, we consider the common picture of structure formation through gravitational instability. Here, the ``creation'' process is the primordial spectrum of density fluctuations which favors low-mass clusters, and the ``destruction'' process is the inhibition of the collapse of low-mass objects resulting from the Jeans mass. In this picture, the Jeans mass can be a very rapidly growing function of time (Tegmark et al.\\ 1996), and the densest systems of galaxies, forming first, would have experienced a less restrictive low-mass cutoff and hence have more, and fainter, GCs. This is precisely the case depicted in Figure~3a.  Harris \\& Pudritz (1994) have proposed a detailed astrophysical theory of GC formation which is perhaps more illustrative of the ``coeval'' case of Figure~3b.  They suggest that the ``creation'' process is made more efficient by the larger external pressures of dense environments. Their ``destruction'' process is the tidal disruption and evaporation of low-mass GCs, and this might be less sensitive to environment (although see Murali \\& Weinberg 1996). Of course, if the cutoff is very abrupt (a steep ``destruction'' line), then $M^0$ will not correlate very strongly with $S_N$.  We do not have better ``creation'' and ``destruction'' processes to offer than have been advanced elsewhere, but we believe that this description is a profitable way to frame the discussion. Until now, the assumed constancy of $M^0$ has been a serious obstacle to reasonable models for GC formation, so we conclude by re-emphasizing our primary point. The GCLF apparently {\\it does} depend on environment, with $M^0_V$ being fainter in denser regions, although it may be remarkably constant within a single group of galaxies. This dependence will present challenges for the use of the GCLF as a distance indicator.   On the other hand, it opens the door for correlations between $M^0$ and $S_N$, and $M^0$ and environment, which may yield valuable insights into the conditions and processes which prevailed at the time of GC/galaxy formation."
    },
    "9604/astro-ph9604106_arXiv.txt": {
        "abstract": "Quasar absorption lines provide detailed information on the chemical, kinematic, and ionization conditions in galaxies and their environments, and provide a means for studying the evolution of these conditions back to the epoch of the first quasars. Among the collection of absorbing structures along the lines of sight to quasars there is an evolutionary sequence of galaxies that represent predecessors of the Milky Way and provide a direct view of its history. Absorption spectra of lines of sight through the Milky Way and through nearby galaxies reveal a variety of chemical species, ionization conditions, and kinematic substructures.  These absorption profiles are produced by low density gas distributed in rotating disks, high velocity halo clouds, satellite galaxies and their debris, superbubbles, and other sub--galactic gaseous fragments. Guided by knowledge gained by studying nearby galaxies, we are developing interpretations of the variety of observed absorption signatures.  Images of $z\\sim1$ galaxies responsible for {\\Mg} absorption also allow us to explore the statistical connections between the galaxy properties and their gaseous content. Quasar absorption lines are fast becoming a powerful evolutionary probe of gaseous conditions in the Milky Way. ",
        "introduction": "Quasar absorption line (QAL) studies have evolved past the stage of simply counting the number of absorption lines due to a particular ion as a function of equivalent width and redshift.  No longer do we merely study disjoint classes of objects called {\\Mg}, {\\Lya}, and {\\C} absorbers; we are now poised at the dawn of an era in which we can explore the structures, metallicities, ionization conditions, and kinematics within the individual galaxies that produce the absorption lines.  Once the absorption signatures of Milky Way--like galaxies are recognized, QALs can be used to trace the detailed evolution of galactic gas from $z=5$ to $z=0$, the full redshift range over which quasars have been observed.  Historically, QAL systems were divided into various categories based upon how the samples were selected: 1) by {\\Lya} (classified as damped if the neutral column density $N_H > 10^{20.3}$~{\\cm2}, Lyman limit if $10^{17.2} < N_H < 10^{20.3}$~{\\cm2}, and forest if $N_H < 10^{17.2}$~{\\cm2}), 2) by {\\Mg}, and 3) by {\\C}.  Evolution in the co--moving number per unit redshift\\footnote{The co--moving number is parameterized by $N(z) \\propto (1+z)^{\\gamma}$.  For no evolution, $\\gamma = 0.5$ in a $q_0=0.5$ universe and $\\gamma = 1.0$ in a $q_0=0.0$ universe.} can be measured for each of these ``populations'' for a given detection threshold in equivalent width (column density in the case of {\\Lya}).  For $z>2$, the number of {\\Lya} forest systems is observed to rapidly decrease with time with $\\gamma \\sim 1.9$ (Bechtold 1994), and for $z<2$ is observed to slowly decrease with $\\gamma$ consistent with no evolution (Bahcall \\etal 1996).   In the redshift interval $0.4 < z < 2.2$, the population of {\\Mg} absorbers ($W_0 > 0.3$~{\\AA}) is consistent with no evolution, though the subgroup of stronger absorbers ($W_0 > 1$~{\\AA}) evolve away with time very strongly ($\\gamma \\sim 2.2$). The opposite trend, an increasing number with decreasing redshift, is seen for {\\C} systems ($W_0 > 0.15$~{\\AA}) in the redshift interval $1.2 < z < 3.7$ (Steidel 1990).  Overall, the statistical evolution of these various populations is due to some combination of evolution in the metallicities, ionization conditions, and kinematic structuring of the absorbing clouds.  A separation of these effects awaits detailed study of the individual absorption line systems.  To better appreciate the inferences about the physical conditions of the gas based upon absorption lines, we would like to know what type of galaxy and what part of the galaxy is being probed by the QSO light path. Specifically relevant to this proceedings is the question of when we are looking through part of a Milky Way--like disk, high velocity cloud, LMC--like satellite, and/or something like the LMC Stream.  A good strategy for sorting this out is to study nearby galaxies and the Milky Way itself, to draw inferences on the kinematic, chemical, and spatial distribution of the clouds from the absorption lines they produce, and to ascertain what fraction of each absorber ``population'' is presented by a galaxy of known type and morphology.  Deep imaging of $z\\leq1$ galaxies allows us to go further back in time and to determine what luminosity and type of galaxy is associated with particular QALs.  With the addition of theoretical modeling of the expected absorption conditions arising from various processes (eg.~winds, fountains, superbubbles, and infalling material), the evolution of strengths, overall shapes, and subcomponent structures of the absorption line profiles will allow us to infer a time sequence in the Milky Way's history.  This article begins with a review of the insights that are likely to be gained through absorption studies of the Milky Way and of nearby galaxies.  The third section is a discussion of observations of {\\Mg} absorbing galaxies at $0.3 < z < 1.0$ that also focuses on studies designed to discern the overall geometric cross sections of low ionization galactic gas.  In the fourth section, we present HIRES (Vogt \\etal 1994) spectra of a number of {\\Mg} absorbers (Churchill 1996a) and discuss what we hope to learn from the variety of substructures that are observed.  The use of {\\C} systems at high $z$ in developing a picture of the history of the early Milky Way galaxy is discussed in section five, as is the relevance of damped {\\Lya} systems for understanding the evolution of the Milky Way disk.  We conclude with speculations on the physical nature of the objects that give rise to various types of QALs. ",
        "conclusions": ""
    },
    "9604/astro-ph9604127_arXiv.txt": {
        "abstract": "HIRES/Keck spectra of Mg II $\\lambda 2796$ absorption arising in the ``halos'' of 15 identified $0.4 < z < 0.9$ galaxies are presented. Comparison of the galaxy and absorbing gas properties reveal that the spatial distribution of galactic/halo gas does not follow a smooth galactocentric dependence. The kinematics of absorbing gas in $z\\sim 1$ galaxies are not suggestive of a {\\it single}\\/ systematic velocity field (i.e.~rotation or radial flow) and show little dependence on the QSO--galaxy impact parameter. From the full HIRES dataset of 41 systems ($0.4 < z < 1.7$), strong redshift evolution in the cloud--cloud velocity dispersion is measured. Direct evidence for turbulent or bulk motion in ``high velocity'' clouds is found by comparing Fe~II and Mg II Doppler parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604061_arXiv.txt": {
        "abstract": "We consider the effect on matter-enhanced neutrino flavor transformation of a randomly fluctuating, delta-correlated matter density.  The fluctuations will produce a distribution of neutrino survival probabilities.  We find the mean and variance of the distribution for the case of solar neutrinos, and discuss the possibility of placing a limit on solar density fluctuations using neutrino data. ",
        "introduction": "\\label{sec:intro}  Matter-enhanced neutrino oscillations, especially in connection to the solar neutrino problem \\cite{baha}, have been extensively studied in the recent years. More recently some interest has developed in the problem of neutrino flavor transformations via the Mikheyev-Smirnov-Wolfenstein (MSW) effect in a randomly fluctuating matter density.  A general approach to neutrino oscillations in such inhomogeneous matter was developed in Ref.~\\cite{sawyer}.  A Boltzmann-like collision integral with blocking factors, describing the decoherence of neutrinos in matter, was given in Ref.~\\cite{raffelt}.  Matter fluctuations which are not random, but harmonic \\cite{koonin, haxton} or changing stepwise \\cite{ks-step} were also considered.  Redfield equations for a neutrino traveling in a region with delta-correlated Gaussian noise were recently developed in Ref.~\\cite{us1} and applied to two-neutrino flavor transformations in the post-core bounce supernova environment in Ref.~\\cite{us2}. In parallel to these papers, an analytical procedure to calculate the survival probability was described in Ref.~\\cite{lujan}, and further implications of solar matter density random noise upon resonant neutrino conversion were studied in Ref.~\\cite{nrsv}.  The aim of this paper is to expand the analysis of Ref.~\\cite{us1} to investigate the mean and variance of the distribution of neutrinos when a randomly-fluctuating, delta-correlated electron density is present in the sun.  A general treatment of fluctuations is presented in Section \\ref{sec:fluc}.  Mean survival probabilities and the variances of the survival probability distribution are given in Sections \\ref{sec:mean} and \\ref{sec:higher}, respectively.  In Section 5, we discuss the results and present conclusions. ",
        "conclusions": "\\label{sec:disc}  The variance is a potentially important tool for the exploration of solar density fluctuations on some time scales.  Fluctuations on the time scale of a radiochemical experiment's run would broaden the distribution of count rates, since different runs would have different survival probabilities.  None of the experiments currently operating has noted a broader distribution of rates than expected \\cite{sage, gallex, homestake}, suggesting that the neutrino data will probably limit, rather than measure, such fluctuations.  For the favored, small-angle solution, the variance is strongly peaked when neutrinos are produced near the resonance.  This suggests that the finite radial distribution of neutrino production will have an important effect, likely extending the peak to lower values of $\\Delta m^2 / E$, since those correspond to resonance farther out in the sun.  Further, it means that density fluctuations will affect pp neutrinos more strongly than other neutrinos, so that gallium experiments may put the strongest limit on fluctuations.  To develop a very rough estimate of the limit on fluctuations, we note that their effect should become noticeable when the ratio $\\sigma / \\langle P_e \\rangle$ becomes comparable to the relative $1\\sigma$ experimental uncertainty.  To simplify the argument, we assume that the signal observed at the gallium experiment consists only of $\\sim 0.3$ MeV pp neutrinos.  (Standard MSW analyses indicate the near-complete suppression of other neutrinos.)  As an example, we will consider the GALLEX results.  GALLEX has a $1\\sigma$ uncertainty of approximately 13\\% and an energy-averaged survival probability of about 60\\% \\cite{gallex}.  A 1\\% density fluctuation on the time scale of a GALLEX run, 20 to 28 days, with $\\sin^2 2\\theta=0.01$ and $\\Delta m^2 / E = 1.63 \\times 10^{-5}$ eV$^2$/MeV (chosen to give a 60\\% survival probability), has $\\sigma/\\langle P_e \\rangle = 0.15$. It is likely, then, that a careful study could rule out fluctuations on that level.  In the future, real-time, high-statistics detectors such as HELLAZ \\cite{hellaz} and Borexino \\cite{borexino} could be used to investigate fluctuations on shorter time scales.  In particular, a helioseismological g-mode oscillation could leave a signature in the neutrino data.  Unfortunately, the current high-statistics experiments SNO \\cite{sno} and Super-Kamiokande \\cite{superk} probably cannot be used to probe fluctuations in this way, as they are not sensitive to low-energy neutrinos.  A proper averaging over neutrino production location would be extremely cpu-intensive.  In this regard, approximate analytic techniques to compute the moments of the density matrix, such as that of Ref.~\\cite{lujan}, may be very useful in analyzing data."
    },
    "9604/astro-ph9604075_arXiv.txt": {
        "abstract": "We show that the process of photoionizing a gas of atomic hydrogen and helium by line radiation whose energy is slightly above the helium single-ionization threshold is unstable if the helium fraction by number is less than approximately one half. However, in the two scenarios we consider here, based on the Decaying Dark Matter (DDM) model of cosmological reionization, there is no significant growth. In the first scenario we consider ionization and recombination to be approximately in equilibrium.  This is relevant to high photon flux rates and early reionization, but in that case the heating is balanced by Compton cooling, which is very stabilizing. In the second scenario we ignore recombination.  This is relevant to low photon flux rates or to the last stage of the reionization. In that case there is too little growth on a cosmological time scale to be significant. ",
        "introduction": "The hydrogen in the intergalactic medium (IGM) is highly ionized and has been at least since $z=4.3$ \\cite{Steidel:Sargent,GP4.7}. Recent observations indicate that the helium in the IGM is mostly at least singly ionized, at least at large redshifts \\cite{Reimers&Vogel:HeGP,Jakobsen:HeGP,Miralda-Escude:HeGP}. Though the IGM might be ionized mainly by quasars \\cite{Meiksin:Madau:Q2,Madau:Q1} or an early generation of stars \\cite{FK:Reheating}, additional ionization by photons emitted by decaying dark matter (DDM) has many appealing features \\cite{Flannery:Press,Sciama:overview,Sciama:ddm_book,Scott:Rees:Sciama}.  In particular, if a generation of neutrinos has a mass sufficient to close the universe, and if such a particle decays into a much lighter particle and a photon, the photon's energy will be near the ionization energies of hydrogen and helium. Specifically, \\begin{equation} E_\\gamma\\approx{1 \\over 2} m_\\nu = {1 \\over 2} 91.5 h^2 \\,{\\rm eV} \\end{equation} where $h$ is the Hubble constant in units of 100 km/s/Mpc. Thus, for $h > .55$, the decay photons would have sufficient energy to ionize hydrogen and for $h > .74$, they could (singly) ionize helium.  The decay photons would also provide a source of out-of-equilibrium energy which might lead to small-scale structure formation. \\citeN{Hogan:ion_inst} proposed a mechanism of small-scale structure formation through the ionization of a pure-hydrogen plasma, but \\citeN{me:no_ion_inst} showed that this mechanism does not work. We have continued to seek possible instabilities of the system, since a neutrino-dominated universe requires a non-gravitational mechanism for spawning structure on galactic scale and below and there is no fundamental reason why this system should be stable. The free energy per atom is sufficient for many e-foldings of instability growth ($E_\\gamma \\gg k T$), and the scale of instability can be comparable to proto-galactic structure. A hydrodynamical instability would mimic many of the desirable features found in the gas component of a Cold Dark Matter (CDM) model, as described for example in \\citeN{Miralda-Escude:Mini-Halo}.  \\subsection{Helium-Ionization Instability Mechanism}  \\citeN{Meiksin} suggested that there is an instability if DDM photons have sufficient energy to ionize helium.  The instability is thermal, i.e. it is based on low-density regions being preferentially heated. Low-density regions receive more heat because they have a larger ratio of neutral hydrogen to neutral helium and a hydrogen ionization deposits more energy in the plasma than a helium ionization (because there is a larger difference between the photon energy and the ionization threshold).  The ratio of neutral hydrogen to neutral helium is larger in low density regions because they are more ionized and the higher ionization cross-section of helium (for photons just above its ionization threshold) ensures that it will absorb photons out of proportion to its abundance.  The reason that low-density regions are more ionized is different in the two scenarios considered below. When recombination equilibrium holds, it is due to the fact that recombination happens in proportion to the square of the density (because a positive ion and an electron need to find each other) while the ionization rate is proportional to the density (the rate of production of photons is assumed to be unrelated to the density). When recombination equilibrium doesn't hold, the low density regions are more ionized because there are more photons per atom.  Though any source of uniformly distributed photons will do for this instability, we concentrate on decaying neutrinos and parameterize the photon production rate by the neutrino lifetime in units of $10^{24}$ seconds, $\\tau_{24}$. ",
        "conclusions": "Though there is an interesting instability involving helium ionization by line radiation, it doesn't seem to have a cosmological application in the DDM scenario. As other candidate radiation sources presume the existence of small-scale structure, we conclude that this class of instability is unlikely to play a role in the initial formation of structure from smooth cosmic gas."
    },
    "9604/astro-ph9604133_arXiv.txt": {
        "abstract": "We present rotation curves of the Galaxy based on the space-velocities of 197 OB stars and 144 classical cepheids, respectively, which range over a galactocentric distance interval of about 6 to 12\\,kpc. No significant differences between these rotation curves and rotation curves based solely on radial velocities assuming circular rotation are found. We derive an angular velocity of the LSR of $\\Omega_0 = 5.5 \\pm 0.4$\\,mas/a (OB stars) and $\\Omega_0 = 5.4 \\pm 0.5$\\,mas/a (cepheids), which is in agreement with the IAU 1985 value of $\\Omega_0 = 5.5$\\,mas/a. If we correct for probable rotations of the FK5 system, the corresponding angular velocities are $\\Omega_0 = 6.0$\\,mas/a (OB stars) and $\\Omega_0 = 6.2$\\,mas/a (cepheids). These values agree better with the value of $\\Omega_0 = 6.4$\\,mas/a derived from the VLA measurement of the proper motion of Sgr\\,A$^{*}$. ",
        "introduction": "The galactic rotation curve has been determined for the inner parts of the galactic disk, interior to the solar annulus, from H\\,{\\sc i}-measurements using the tangential point method (Burton \\& Gordon 1978), whereas the outer rotation curve has been determined using radial velocities of objects with individually known distances, e.\\,g.\\ OB stars (Fich et al.\\ 1989), planetary nebulae (Schneider \\& Terzian 1983), young open clusters (Hron 1987) and carbon stars (Metzger \\& Schechter 1994). An alternative method is based on the vertical thickness of the galactic H\\,{\\sc i}-layer (Merrifield 1992). Recently Brand \\& Blitz (1993) have rederived the outer rotation curve from CO radial velocities of OB stars associated with H\\,{\\sc ii} regions, and Pont et al.\\ (1994) have used new radial velocity measurements of classical cepheids for this purpose. All these methods rely on the assumption of circular orbits around the galactic centre, so that radial velocities can be converted to circular velocities. Obviously proper motions of the objects may provide independent information on the rotation curve. The PPM catalogue which has recently been compiled at the Astronomisches Rechen-Institut (R\\\"oser \\& Bastian 1991, Bastian et al.\\ 1993, R\\\"oser et al.\\ 1994) is well suited as a broad data base of proper motions of high accuracy for such purposes. ",
        "conclusions": "\\subsection{Rotation curves} Once the velocities have been corrected for the solar motion, it is straightforward to determine the circular velocity $v_{\\mbox{\\tiny c}}(R)$ for according to the formula \\begin{equation} v_{\\mbox{\\tiny c}}=(U-\\Omega_0 Y)\\sin \\varphi+(V+\\Omega_0 X)\\cos \\varphi + \\Omega_0 R\\,\\,, \\end{equation} where $R$ is the distance of the star from the galactic centre. It is now no longer necessary to assume that the stars move on circular orbits around the galactic centre, because the full space velocities are used.  Alternatively, assuming again circular orbits, one may derive rotation curves based solely on radial or tangential velocity components of the stars, \\begin{eqnarray} \\label{rotrad} \\Omega_0(R) &=& \\frac{v_{\\mbox{\\tiny rad}}}{R_0 \\sin l \\cos b} + \\Omega_0 \\qquad \\\\ \\label{rotpm} \\Omega_0(R) &=& \\frac{(\\mu_l + \\Omega_0) r \\cos b}{R \\cos (\\varphi + l)} + \\Omega_0 \\quad, \\end{eqnarray} where $r$ denotes the distance of the star from the Sun, $v_{\\mbox{\\tiny rad}}$ the radial velocity, and $\\mu_l$ the proper motion in the direction of galactic longitude $l$. Both velocities have to be corrected again for the solar motion. $r$ is the distance of the star from the Sun.  The resulting rotation curves determined using either the full space velocities or deprojected radial or tangential velocity components are shown in Figs.\\,\\ref{ob1} to \\ref{cep2}. They agree very closely and are all consistent with a flat shape of the rotation curve $\\Omega(R)=\\Omega_0 R_0 / R$.  Stars with unsuitable projection angles onto the supposed circular velocities have been excluded. In the case where only radial velocity data are used stars with $| \\sin l\\,| < 0.3$ have been rejected. If only proper motion data are used stars with $| \\cos(\\varphi +l)\\,| < 0.7$ have been rejected. Therefore the overlap between these two samples is small and we cannot correlate directly the rotation curve derived from radial velocities with the rotation curve based on proper motions or both.  Instead, we show in Fig.\\,\\ref{diff} the deviations of circular velocities derived from the space velocities from circular velocities based solely on deprojected radial velocities of the stars illustrating again that both methods give consistent results.  \\subsection{Residual velocities} The orientations of the velocity residuals after subtraction of the systematic velocity components due to a flat rotation curve from the space velocities of the stars are shown in Fig.\\,\\ref{illu} projected onto the galactic plane. The velocity residuals are dominated by the errors of the proper motions. Most of the velocity residuals can be shown to be randomly orientated with the exception of the Perseus spiral arm. There is a trend of coherent motion along the spiral arm with the effect that the OB stars in the spiral arm tend to lag behind the general rotation of the disk. Exactly such a behaviour is predicted by the density wave theory of spiral structure for stars recently born in the shock front of the interstellar gas included by the spiral density wave (Shu et al.\\ 1972). This is not observed for the cepheids. They are about 100 times older than the OB stars and their systematic flow pattern is already dissolved (Wielen 1979)."
    },
    "9604/astro-ph9604149_arXiv.txt": {
        "abstract": "\\asca\\ data are used to obtain two-dimensional gas temperature maps of the hot non-cooling flow clusters A2256, A2319, A2163 and A665. In all four clusters, the temperature decreases significantly at off-center distances of $\\sim 1\\,h^{-1}$~Mpc ($H_0\\equiv 100\\,h$\\kmsmpc). Central regions of the two nearer clusters A2256 and A2319 are resolved by \\asca\\ and appear largely isothermal except for the cooler spots coincident with the subunits in their X-ray surface brightness. Although the existence of this substructure may suggest ongoing merger activity, no asymmetric features in the temperature distribution resembling those in the hydrodynamic merger simulations (e.g., Schindler \\& \\muller\\ 1993) are apparent. In the outer parts of the clusters, the temperature declines symmetrically with radius. In A2256 and A2319, it follows a polytropic slope with $\\gamma\\simeq 1.3-1.5$. This is somewhat steeper than the simulations predict for a flat CDM universe and is closer to the open universe predictions (Evrard et al.\\ 1996). The temperature drop is more prominent in the outer regions of A2163 and A665 and appears even steeper than adiabatic (although not inconsistent with it). If the gas in the outskirts of these two clusters is indeed as cool as we measure, the cluster atmospheres should be convectively unstable and transient. Also, such a steep temperature profile could not possibly emerge if the gas was heated only via the release of its own gravitational energy during infall. This may indicate the presence of an additional heat source in the inner cluster, such as merger shocks transferring energy from the dark matter to the gas. The results suggest that A2256 and A2319 are pre-merger systems and A2163 and A665 are ongoing or post-mergers. ",
        "introduction": "Spatially-resolved measurements of the cluster gas temperature are necessary for such an extensive and important problem as determining cluster masses (e.g., Fabricant et al.\\ 1984; White et al.\\ 1993). Beyond that, the cluster temperature structure can provide information on the dynamical history of these systems. Rich and massive clusters should be just forming now in the hierarchical clustering scenarios with a high matter density parameter $\\Omega$ (e.g., Blumenthal et al.\\ 1984). On the other hand, in an open universe most present-day clusters should be old, because their formation is inhibited after $t\\sim \\Omega t_0$, $t_0$ denoting the present epoch (White \\& Rees 1978). Hydrodynamic simulations of cluster growth (e.g., Navarro et al.\\ 1995; Evrard et al.\\ 1996, hereafter EMN) predict a largely constant temperature profile in the inner part and its decline in the outer regions for the relaxed clusters, with a steeper decline in the open universe models. Young clusters which have recently undergone a merger should retain a complex temperature structure (e.g., Schindler \\& \\muller\\ 1993). However, while there is a wealth of cluster simulations in various cosmological scenarios, until recently, direct spatially-resolved temperature measurements have been possible with only a limited accuracy, especially for the hotter, more massive systems (e.g., Hughes 1991; Eyles et al.\\ 1991; Miyaji et al.\\ 1993; Briel \\& Henry 1994, hereafter BH; Henry \\& Briel 1995). \\asca\\ with its broad energy coverage combined with imaging capability (Tanaka et al.\\ 1994) is set to significantly improve the situation. Some results have already appeared (e.g., Arnaud et al.\\ 1994; Markevitch et al.\\ 1994, 1996, hereafter M94 and M96; Ikebe et al.\\ 1996).  In this {\\em Letter}, we use \\asca\\ data to derive temperature maps of nearby A2256 ($z=0.058$), A2319 ($z=0.056$), and distant A2163 ($z=0.201$) and A665 ($z=0.18$). All four are hot, lack cooling flows and are probably not fully relaxed, which is suggested by either the substructure in their X-ray images or by galaxy velocities (e.g., Briel et al.\\ 1991; Elbaz et al.\\ 1995). For A2256, a temperature map was earlier presented by BH who used \\rosat\\ PSPC, and our results are compared with theirs. \\asca\\ results on the temperature structure near the center of A2163 were reported in M94. In M96, a steep radial temperature decrease was found in this cluster. Interestingly, a recent measurement of the Sunyaev-Zeldovich effect toward A2163 by Holzapfel et al.\\ (1996) independently suggests a similar decrease, although with marginal significance. Below, a less model-dependent, two-dimensional approach to the \\asca\\ data is employed to confirm the result of M96 and find similar phenomena in other three clusters. ",
        "conclusions": "Although A2256 and A2319 clearly exhibit substructure in their \\rosat\\ X-ray images, no large-scale merger signatures, such as those predicted by hydrodynamic simulations (e.g., Schindler \\& \\muller 1993; EMN), are seen in the temperature maps of their central parts. This may indicate that the current mergers have not proceeded far enough to disturb the bulk of gas. For example, Roettiger et al.\\ (1995) specifically simulated A2256 and found that the cluster image, galactic velocities and absence of the cluster-scale temperature variations are consistent with an epoch of about 0.2~Gyr before core passage of an infalling subunit. The observed relative symmetry of the temperatures in A2256 and probably A2319 suggests that their outer parts have been undisturbed by major mergers for the past few Gyr, making these clusters good candidates for an accurate mass measurement, which will be made in a future paper.  Apart from the subgroups, the temperature profiles of these two clusters are qualitatively similar to those predicted by the simulations of Navarro et al.\\ and EMN for clusters in equilibrium. Interestingly, the observed temperature decline starts at smaller radii than EMN predict for the flat universe models without galactic winds, and is closer to their open universe model, in which clusters are expected to have steeper density and temperature profiles (Hoffman \\& Shaham 1985; Crone et al.\\ 1994; Jing et al.\\ 1995).  However, the published simulations including gas are limited to the CDM initial perturbations spectrum, and our sample is limited to just a couple of rather specific clusters (e.g., lacking cooling flows unlike most of the clusters). A study of several other clusters is underway with \\asca, which will show how common this phenomenon is.  The temperature falls even steeper in A2163 and perhaps in A665. However, as was noted in M96, the low outer values may in fact not be representative of the mean gas temperature at those radii. For example, the measured electron temperature may be lower than that of ions heated by shock waves, because the timescale of electron-ion equipartition via collisions becomes non-negligible at such low plasma densities. Other possibilities involve cold gas clumps or point sources which cannot be localized by either \\asca\\ or \\rosat\\ but significantly contribute to the flux. On the other hand, if the outer gas temperatures are indeed as low as measured, the observed steep profiles would have interesting implications for the physical conditions of the gas (although it may be premature to speculate using such poor data constraints.) Firstly, the outer cluster parts with a steeper than adiabatic temperature decline should be convectively unstable, and convection should develop on a timescale of the order of the free-fall time (a few Gyr) and erase the gradient. Thus, existence of a steep gradient implies that the cluster cannot have remained in its present state for a longer time than this. Another interesting problem is how such a temperature distribution may have emerged. Early simulations of infall of the cold gas into the cluster and its heating via the release of its potential energy (e.g., Bertschinger 1985) predict that such a process should form shallower temperature distributions. A steeper slope may therefore indicate that the gas in the central part has accumulated additional energy from another source. Hydrodynamic merger simulations predict (Pearce et al.\\ 1994) that during a merger, energy is transferred from the dark matter to the gas, increasing its entropy in the center. Thus the observed profiles may independently indicate that these clusters have experienced major mergers.  There is some evidence of the asymmetric temperature variations in the central part of A2163 (M94) and recent weak lensing analysis reveals two mass peaks near its center (Squires et al.\\ 1996). A665 has a markedly asymmetric X-ray image which may remain from a merger. Schindler \\& \\muller\\ (1993) predict that a merger shock wave would manifest itself at certain stages as a sharp projected temperature gradient in the cluster outer part, not accompanied by a similarly noticeable feature in the wide-band surface brightness, which is what we may be observing in the two more distant clusters."
    },
    "9604/astro-ph9604023_arXiv.txt": {
        "abstract": "We study the production of Na and Al around the hydrogen shell of two red-giant sequences of different metallicity in order to explain the abundance variations seen in globular cluster stars in a mixing scenario. Using detailed stellar models together with an extensive nuclear reaction network, we have calculated the distribution of the various isotopic abundances around the hydrogen shell at numerous points along the red-giant branch. These calculations allow for the variation in both temperature and density in the shell region  as well as the timescale of the nuclear processing, as governed by the outward movement of the hydrogen shell. The reaction network uses updated rates over those of Caughlin \\& Fowler (1988). We find evidence for the production of Na and Al occurring in the NeNa and MgAl cycles.  In particular, Na is significantly enhanced throughout the region above the hydrogen shell. The use of the newer reaction rates causes a substantial increase in the production of $^{27}$Al above the hydrogen shell through heavy leakage from the NeNa cycle and should have an important effect on the predicted surface abundances.  We also find that the nuclear processing is considerably more extensive at lower metallicities. ",
        "introduction": "Briley et al. \\markcite{r5}(1994) and Kraft \\markcite{r6}(1994) have reviewed the observational data on variations in the abundances of C, N, O, Na, and Al in globular cluster red-giant-branch (RGB) stars. The variations in Na and Al are particularly significant, since they have long been regarded as one of the principal arguments in favor of a primordial origin for the abundance anomalies (see e.g., Cottrell \\& Da Costa \\markcite{r19}1981). Star-to-star variations of Na were first observed by Cohen \\markcite{r1} (1978) and Peterson \\markcite{r2}(1980) in M13 and M3, while similar variations of Al, which are correlated with the CN band strength, were found by Norris et al. \\markcite{r3}(1981) in NGC 6752. Since these original observations, numerous other groups have confirmed the general existence of Na and Al vs. N correlations and Na and Al vs. O anticorrelations in globular cluster red giants (Drake et al. \\markcite{r24}1992; Kraft et al. \\markcite{r25}1992, \\markcite{r26}1993). Recently Norris \\& Da Costa \\markcite{r7}(1995) have concluded that Na variations exist in all clusters, while Al variations are greater in the more metal-poor clusters.  Except for the modest alterations due to the first dredge-up, canonical stellar evolution models predict no changes in the C, N, O, Na, and Al surface abundances during the RGB phase. In an attempt to explain this discrepancy between the predicted and observed abundance variations, Sweigart \\& Mengel \\markcite{r8}(1979; hereafter, SM79) suggested that meridional circulation currents, driven by internal rotation, might be able to mix material across the radiative zone that separates the top of the hydrogen shell (H shell) from the base of the convective envelope in canonical RGB models. They found that there was a region of significant extent just above the H shell within which the CN cycle has processed C into N and, somewhat closer to the shell, there was a region within which the ON cycle has processed O into N. The mixing was postulated to begin at the point along the RGB where the H shell burns through the hydrogen discontinuity that was previously produced by the deep penetration of the convective envelope during the first dredge-up. A progressive depletion of carbon is seen along the giant branch in M92, M15, and NGC 6397 (Bell, Dickens, \\& Gustafsson \\markcite{r23}1979; Carbon et al. \\markcite{r27}1982; Trefzger et al. \\markcite{r28}1983; Briley et al. \\markcite{r29}1990), although the luminosity at which this depletion begins is uncertain, for observational reasons. If Na and Al were also manufactured in the CN- and ON-processed regions, then mixing might also explain the observed Na and Al variations. Complimenting the SM79 hypothesis, Denisenkov \\& Denisenkova \\markcite{r9} (1990; hereafter, DD90) have suggested $^{23}$Na can be produced from proton captures on $^{22}$Ne in the ON-processed region and that the rotation rate required by SM79 is sufficient to reconcile the observations with theory.  Expanding on the concepts of SM79 and DD90, Langer, Hoffman, \\& Sneden \\markcite{r10}(1993; hereafter, LHS93; see also Langer \\& Hoffman \\markcite{r31} 1995 and Denissenkov \\& Weiss \\markcite{r32} (1996)) examine the production of Na and Al in RGB stars by following the reactions of the relevant nuclei in a simplified model with a low metallicity (Z = 0.0001 or [Fe/H] = -2.3) and a constant temperature (T$_{9}$ = 0.040 where T$_{9}$ = T/10$^{9}$K) and density ($\\rho$ = 44.7 g cm$^{-3}$). Their study shows a significant enhancement of $^{23}$Na in the ON-processed region which derives from a series of proton captures on $^{20}$Ne in the NeNa cycle. Furthermore, their model yields an increase in $^{27}$Al which is made from proton captures on both $^{25}$Mg and $^{26}$Mg via the reactions $^{25}$Mg$(p,\\gamma)^{26}$Al$(\\beta^{+})^{26}$Mg$(p,\\gamma)^{27}$Al. Their study also shows that a MgAl cycle is set up after 1.7 Myr. This cycle begins with a proton capture on $^{24}$Mg, which is itself enhanced through leakage from the NeNa cycle by a $(p,\\gamma)$ reaction with $^{23}$Na, and is completed with $^{27}$Al$(p,\\alpha)^{24}$Mg. Both the NeNa and MgAl cycles are depicted graphically in Figure 1.   \\begin{figure} \\plotone{cavallo1.eps} \\caption{The reactions involved in the NeNa and MgAl cycles. The half-lives and reaction types are given parenthetically. In the case of competing decay paths, the solid lines show the stronger path, derived with the rates used in our code for the relevant temperature range. The dashed lines show the weaker decays which lead to a depletion in the total abundance of each cycle (\"leakage\").} \\end{figure}  The purpose of this ${\\it Letter}$ is to develop the work of LHS93 further by using detailed stellar evolutionary sequences to examine the production of Na and Al around the H shell in RGB stars. This represents an improvement over LHS93 because: 1) we incorporate the latest reaction rates in our burning code, 2) we explore a wider parameter space in both metallicity and luminosity, and 3) our more realistic RGB models take into account the variation in temperature and density around the H shell and incorporate the timescale for the nuclear processing to produce Na and Al, as set by the rate at which the H shell moves outward in mass.  The Na and Al produced in this fashion can then be mixed outward into the envelope over the relevant timescale for the mixing process. Section 2 summarizes our numerical techniques and input physics while section 3 presents the results of our calculations for some representative cases.  We conclude with a brief discussion of our results in section 4. ",
        "conclusions": "By combining our realistic stellar models with an updated nuclear reaction network, we are able to follow the nucleosynthesis of C, N, O, Ne, Na, Mg, and Al around the H shell of two sequences of differing metallicity as they evolve up the RGB. In qualitative agreement with the observations, our results show an increase in $^{23}$Na above the H shell throughout the entire giant branch, independent of metallicity.  Furthermore, we produce sizable enhancements of $^{27}$Al for the low metallicity sequence as it approaches the tip of the RGB, without having to increase the initial $^{25,26}$Mg abundance or the $^{26}$Mg proton capture rate as in Langer \\& Hoffman \\markcite{r31} (1995). Thus, our results can potentially reconcile the Na and Al abundance anomalies observed in globular cluster giants with the mixing of these elements from the stellar interior.  Although our work is based on up-to-date reaction rates, those rates can still be quite uncertain in some cases. For example, our results deviate from the observations in the case of $^{24}$Mg. Contrary to the findings of Shetrone \\markcite{r30}(1996), who observes a decrease in $^{24}$Mg with increasing luminosity in M13 giants, we show an enhancement of $^{24}$Mg with continuing evolution for the low metallicity sequence.  However, the high metallicity sequence shows only a marginal increase  in the $^{24}$Mg abundance. Thus, our present inability to qualitatively reproduce the $^{24}$Mg observations might be due to the current uncertainty in the nuclear cross-sections. Perhaps a faster rate for the $^{24}$Mg$(p,\\gamma)^{25}$Al reaction or a slower rate for the $^{27}$Al$(p,\\alpha)^{24}$Mg reaction would decrease the $^{24}$Mg abundance.  Figures 2 and 4 show how the $^{24}$Mg abundance can depend on the choice of reaction rates.  In addition to depending on the accuracy of the nuclear reaction rates, any final predictions of surface abundances will also depend strongly on the choice of a mixing mechanism.  The mixing timescale, the depth of the mixing, and the dependence on metallicity will all affect the quantitative results for the surface abundances.  In a future paper, we will explore these aspects of the mixing mechanism and will compare the predicted variations of the surface abundances along the RGB with the observational constraints.   In order to better understand the effects of metallicity, we will also extend the present work to a  Population I red giant. Furthermore, we will use the same set of sequences to study the CNO isotopes. Finally, we intend to follow the production and destruction of $^{3}$He along the RGB in order to determine better how low mass stars affect the chemical evolution of this isotope in the Galaxy."
    },
    "9604/astro-ph9604171_arXiv.txt": {
        "abstract": "It is generally supposed that when the ``compactness\" $l \\equiv L\\sigma_T/(r m_e c^3)$ in photons above the pair-production threshold is large, few $\\gamma$-rays can escape.  We demonstrate that even when $l \\gg 1$, if the high energy and low energy photons are produced in geometrically-separated regions, many of the $\\gamma$-rays can, in fact, escape. Pair-production along a thin surface separating the two sources creates enough Compton optical depth to deflect most of the low energy photons away from the high energy ones. Those few low-energy photons which penetrate the shielding surface are reduced in opacity by advection to large distance and small density, by relativistic beaming along the inner edge of the surface, and by Compton upscattering to higher energies inside the surface. The pairs in this surface flow outward relativistically, forming a structure resembling a pair-dominated mildly relativistic jet. ",
        "introduction": "One of the most basic concepts in the study of astrophysical $\\gamma$-ray sources is that of ``compactness\" (see, {\\it e.g.}, the review by Svensson 1986).  Its importance stems from many causes, but one of its central implications has to do with the optical depth of a $\\gamma$-ray source to pair production.  Two photons with energies $x_1$ and $x_2$ (in units of $m_e c^2$) may react to produce an $e^{\\pm}$ pair when their energies satisfy the relation $x_1 x_2 \\geq 2/(1 - \\cos\\theta)$, where $\\theta$ is the angle between their directions of motion.  If a luminosity $L$ in photons of energy just above threshold is made isotropically within a source of size $R$, then the optical depth to pair production $\\tau_{\\gamma\\gamma}$ is greater than or of order unity when \\begin{equation} l \\equiv {L \\sigma_T \\over R m_e c^3} > 4\\pi, \\end{equation} where $\\sigma_T$ is the Thomson cross section. The only caveat attaching to this statement is an order unity correction dependent on the details of the spectrum.  In virtually all work since, it has been assumed that when $l \\gg 1$, few high energy photons can escape the source region, although subsequent pair annihilation can partially restore them ({\\it e.g.}, Guilbert, Fabian \\& Rees 1983; Svensson 1987; Blandford 1990).  This argument has been used in many contexts because sources rich in $\\gamma$-rays are very often copious sources of softer photons as well, and are also often variable enough to indicate causality bounds on the source size which result in large estimated compactnesses.  However, it is not always true that $l \\gg 1$ implies $\\tau_{\\gamma\\gamma} \\gg 1$.  In particular, when the sources of high and low energy photons are geometrically separated, the very pair production which has been thought to prevent $\\gamma$-ray escape can instead ensure it.  Because high compactness generally also entails a significant Compton optical depth (Guilbert {\\it et al.} 1983), an optically thick Compton scattering layer can form between the two source regions.  High energy photons incident upon this surface from one side are either absorbed in pair production reactions, or else lose much of their energy by Compton recoil; low energy photons, which strike it from the other side, are scattered away, so that only a few cross the surface onto the side where the $\\gamma$-ray source is located.  The net result is that those high energy photons directed away from the surface can escape freely, while the energy of the $\\gamma$-rays directed into the surface is transformed into a relativistic outflow of mixed $e^{\\pm}$ pairs and photons.  The remainder of this paper is devoted to working out a ``toy-model\" version of this idea so as to illustrate its qualitative features.  It would be premature at this stage to use these ideas as the basis for a detailed model of any particular source or class of sources; our goal instead is to work out a number of gross features and scaling properties of this class of model to provide guidance for future, more realistic applications.  Hence the following description is necessarily highly idealized. ",
        "conclusions": "We have shown that when $L_\\gamma/L_x > 0.08$, the pairs that are created by $\\gamma$-ray -- X-ray reactions are squeezed into a screening surface shaped like an hour-glass at relatively small distances, but which becomes asymptotically conical at large distances.  As equation 7 demonstrates, when $l_{\\gamma} \\gg 1$ and $l_x$ is large enough to make the surface optically thick to pair production, the screening surface acquires a significant Compton optical depth at least to heights $z_s \\sim a$.  When that occurs, most of the incident X-rays are scattered back outside the surface, although the relativistic bulk motion will beam them within an angle $1/\\Gamma$ of the local surface tangent.  As a result, very few X-rays penetrate into the central region, and the $\\gamma$-rays directed within the surface are free to escape.  This conclusion differs from that of Bednarek (1993), who found that a compact X-ray ring would lead to $\\gamma$-ray absorption, because he neglected the Compton scattering of X-rays by the pairs.  On the other hand, most of the energy of the $\\gamma$-rays initially directed into the screening surface is absorbed.  Thus, this mechanism has the net effect of using a fraction of the $\\gamma$-rays (the covering fraction of the surface around the origin) to preserve the remainder. When $L_{\\gamma}/L_x \\sim 1$, the two fractions are comparable.  Both the escaping $\\gamma$-rays and the X-rays are collimated by this process.  The $\\gamma$-rays can only escape in those directions not covered by the screening surface.  On the other hand, those X-rays initially directed toward the surface are focussed into directions nearly tangent to the surface, and also inverse Compton scattered to higher energies. Therefore, when our line of sight lies nearly parallel to the asymptotic direction of the screening surface, we should see an X-ray spectrum which is both quite strong and quite hard.  Thus, objects in which this process operates can be expected to have strikingly different high energy spectra when viewed from different directions.  Because the asymptotic direction of the screening surface changes with $L_{\\gamma}/L_x$, it is possible for fluctuations in $L_\\gamma/L_x$ to strongly modulate the observed X-rays.  When $L_\\gamma/L_x$ is small, our line of sight is likely to lie outside the screening surface, so we see the unmodified X-ray spectrum; when $L_{\\gamma}/L_x$ is large, the screening surface swings outward, cutting off our view of the X-rays; for some range of intermediate values, our line of sight will lie close enough to the direction of the surface that we will see the beamed and upscattered X-rays.  It is possible that these effects have been observed in the BL Lac object PKS 2155-304.  Ordinarily its spectral index above a few keV is in the range 1 -- 2 (Sembay \\etal 1993), but in one observation (Urry \\& Mushotzky 1982) the spectral index from 15 keV up to the sensitivity limit of the instrument, $\\simeq 40$ keV, was $\\simeq -1.5$! Such behavior might be explained if an excursion to especially large $L_\\gamma/L_x$ opened the screening surface wide enough for our line of sight to be nearly aligned with it.  In many of the most powerful high energy $\\gamma$-ray sources known there is independent evidence strongly suggesting that much of the radiation comes from relativistic jets (Hartman \\etal 1994; Dermer, Schlickeiser \\& Mastichiadis 1992; Blandford 1993; Sikora, Begelman, \\& Rees 1994). Some have argued that the existence of strong high energy $\\gamma$-radiation in these objects is itself evidence for relativistic motion ({\\it e.g.} Zdziarski and Krolik 1993; Dondi \\& Ghisellini 1995).  While the mechanism we described here certainly does not argue against relativistic motion in the source, it does create a loophole in the arguments for relativistic motion solely on the basis of high energy compactness.  Our point of view is that in more realistic pictures of these sources, {\\it both} relativistic motion of the source and optical depth effects such as the ones we have discussed in this paper may operate. One might easily imagine, for example, that the toy geometry we have explored here should, in real sources, be modified to take into account a relative velocity between the $\\gamma$-ray source and the X-ray source which is very likely relativistic, and may or may not be aligned with the symmetry axis.  Such relative motion could beam the $\\gamma$-rays away from the X-ray source, so that a smaller fraction of the $\\gamma$-rays are used to produce a shielding wall which is still sufficiently optically thick to be effective.  This effect could help solve an otherwise troubling problem in jet models of $\\gamma$-ray production in AGN: that if $\\gamma$-rays are produced too close to the center of the system, pair production on X-rays would lead to a saturated pair cascade and the production of a much larger X-ray luminosity (through inverse Compton scattering by the pairs that are produced in the cascade) than is observed (Ghisellini \\& Madau 1995).  These effects have another significant consequence: there is a collimated outflow in the screening surface whose luminosity is equal to the absorbed $\\gamma$-ray luminosity, which is likely to be an interesting fraction of $L_\\gamma$.  This outflow, composed of a mixture of electron-positron pairs and photons of comparable energy, is automatically mildly relativistic, with a modest bulk $\\Gamma$.  Particularly when $L_\\gamma/L_x < 1$ so that the opening angle is comparatively small, these outflows have many of the characteristics of the relativistic jets thought to be responsible for much of the lower-frequency radiation in these objects.  \\centerline"
    },
    "9604/astro-ph9604037_arXiv.txt": {
        "abstract": "A cosmological multidimensional hydrodynamic code is described and tested. This code is based on modern {\\it high-resolution shock-capturing} techniques. It can make use of a linear or a parabolic cell reconstruction as well as an approximate Riemann solver. The code has been specifically designed for cosmological applications. Two tests including shocks have been considered: the first one is  a standard shock tube and the second test involves a  spherically symmetric shock. Various additional cosmological tests are also presented. In this way, the performance of the code is proved. The usefulness of the code is discussed; in particular, this powerful tool is expected to be useful in order to study the evolution of the hot gas component located inside nonsymmetric cosmological structures. ",
        "introduction": "In this paper, a new hydrodynamical code is presented and tested. This code is the multidimensional extension of a previous one described by Quilis et al.(1994); it has been designed with the essential aim of simulating the evolution of a nonsymmetric three-dimensional (3D) distribution of gas. The features of this distribution are similar to those of the hot rarified gas located inside cosmological structures. Further applications of the multidimensional code will be presented elsewhere.  Even if the numerical techniques, the physical principles and the equations involved in a code are appropriate, the hydrodynamical code must be tested in order to discover unexpected failures and possible limitations. Six tests have been selected spanning a significant set of situations which  allow us to check all the relevant ingredients of the code. The chosen tests correspond to six problems with known analytical or numerical solutions. These solutions are compared with those obtained using our multidimensional code. Results are encouraging.  Two types of methods can be used in order to study structure formation in Cosmology: Analytical and numerical methods. Among the analytical methods, several Eulerian and Lagragian hydrodynamical approaches have been proposed. In these approaches, the Universe is considered as a fluid (see Bertschinger 1991). Let us mention, as examples, the {\\it Zel'dovich solution} (Zel'dovich 1970), the {\\it Adhesion model} (Gurbatov et al. 1989), and the {\\it frozen flow approximation} (Matarrese et al. 1992). These approaches apply beyond the linear regime, but they have limitations. Zel'dovich's 3D solution only works properly up to the mildly nonlinear regime and both the frozen flow and the adhesion models introduce fictitious forces in order to avoid caustic formation. On account of these limitations, it seems that the use of other techniques such as the numerical ones is appropriate. N-body simulations are used in the pressureless case and very robust hydrodynamical codes have been proposed for studying collisional matter. Let us mention the code developed by Cen (1992), which uses {\\it artificial viscosity}, and the smooth particle hydrodynamic codes originally developed by Gingold \\& Monaghan (1977) and Lucy (1977), independently. This paper is a preliminary study, in which modern numerical techniques are implemented in order to design a code for future applications to structure formation in Cosmology.   We have built up a multidimensional cosmological hydrodynamical code based on {\\it modern high-resolution shock-capturing} (HRSC)  techniques. These HRSC methods were specifically designed for solving hyperbolic systems of conservation laws and have two main features: they are at least second order accurate on the smooth part of the flow and they give well resolved nonoscillatory discontinuities (LeVeque 1992). By construction HRSC schemes avoid using numerical artifacts,  such as the artificial viscosity, in order to smear shocks. With HRSC techniques strong shocks are sharply solved, typically, in two or three numerical cells, and they are free of spurious oscillations due to the Gibbs phenomenon. This last property could be of crucial importance in 3D calculations where the numerical grid is constrained for obvious technical reasons and has a poor resolution.   In recent years, HRSC methods have been applied widely in many astrophysical fields: interacting stellar winds (see, e.g., Mellema et al. 1991), type II supernovae explosions (see, e.g., M\\\"uller 1994) , relativistic jets (Mart\\'{\\i} et al., 1995), etc. More recently, some codes --with cosmological applications-- based on HRSC techniques have been presented. A very recent multidimensional cosmological hydro-code built up using these techniques is the one described by Ryu et al. (1993). In Quilis et al. (1994) we analyzed the main features of a one-dimensional code, and in the present paper we are concerned with the multidimensional version of this code.  Hereafter, $t$ stands for the cosmological time, $t_0$ is the age of the Universe, $a(t)$ is the scale factor of a flat background. $\\dot{X}$ stands for the derivative of the function $X$ with respect to the cosmological time. Function $\\dot{a}/a$ is denoted by $H$. Hubble constant is the present value of $H$; its value in units of $100 \\ Km \\ s^{-1} \\ Mpc^{-1}$ is $h$. Velocities are given in units of the speed of light. $\\rho$ and $\\rho_{_{B}}$ stand for mass density and  background mass density, respectively. The density contrast is $\\delta=(\\rho-\\rho_{_{B}})/\\rho_{_{B}}$. The background is flat. $p$ and $\\epsilon$ stand for pressure and internal energy per unit mass, respectively.   The plan of this paper is as follows: In Section 2, our numerical code is described. In Section 3, the results of several tests are shown. Finally, a general discussion is presented in Section 4. ",
        "conclusions": "In this paper, we have numerically  solved the full multidimensional system of hydrodynamical equations describing the evolution of a gas --including gravity and an expanding cosmological background-- taking advantage of the fact that they are a system of conservation laws with sources. This property is crucial for using modern HRSC techniques.  A multidimensional hydrodynamic code based on these techniques has been built up. This code includes the possibility of choosing between two spatial reconstructions in order to get better resolutions, i.e. a linear reconstruction (MUSCL) or a parabolic reconstruction (PPM).  An algorithm for solving Poisson's equation at each time step is included as well. This Poisson solver is based on the FFT. The use of this transform is appropriated because the code described in Section. 2 gives the density at each time step and this density is the only ingredient required by the FFT in order to compute the gravitational potential; in other words, if the FFT is used, no unknown boundary conditions for the gravitational potential are required as inputs. Finally, the time elapsed by the FFT increases very slowly as the number of points per edge --of the elemental cube-- increases. Although the use of the FFT has important advantages, this technique only leads to admissible simulations in the central part of the elemental box. This is an unavoidable limitation attached to the use of the FFT. The elemental cube must be carefully chosen in each case.   Our code has passed successfully a battery of six severe tests. Four of them (shock tube problem, spherical shock reflection, Zel'dovich's solution and the Newtonian pressureless spherically symmetric solution) are in fact considered as standard bed-tests in classical and cosmological hydrodynamics. We have implemented two more numerical tests from previous numerical solutions.  The behaviour of our code is good in all six cases. The tests show that the code works in the presence of shocks, rarefactions, contact discontinuities, cosmological expansion, gravity and pressure. In the spherical cosmological test, it has been seen that high density contrasts of the order of $10^{2}$ can be reached by using spatial grids with $64 \\times 64 \\times 64$ points (see Fig. 7). Larger density contrasts would require a greater number of points and, consequently, they would have a greater computational cost. Grids having  $128 \\times 128 \\times 128$ nodes should allow the description of the hot gas component located inside clusters up to density contrasts between $10^{2}$ and $10^{3}$ (rich clusters).  In cosmological applications and as Ryu et al. (1993) pointed out, it might happen that regions having large kinetical --compared with the thermal energy-- appear. As we have discussed before, a conservative algorithm could lead to important numerical difficulties. In practice we have not noticed this problem in any of the tests presented in this paper, overall in the shock reflexion test where $\\frac{E_{th}}{E} \\ge 10^{-4}$ and the analytical solution was recovered quite well.  Important perspectives arise in the case of several cosmological problems, specially, if the code presented in this paper is coupled to a N-body one (see Bertschinger \\& Gelb 1991 for a description of this kind of codes) describing the evolution of the pressureless matter. In order to design this coupling, it should be taken into account that the hot gas and the pressureless component are gravitationally coupled. Let us point out that the resulting coupled code could lead to very realistic simulations of rich clusters, in which, the observed features of the baryonic component would play an important role. See Quilis et al.(1995) for a discussion about this point.      Here, we summarize the steps in order to build up our hydro-code. In one of them, the spectral decompositions of the Jacobian matrices described in the Appendix A, are needed. For more details see  Section 2.  The scheme for the numerical procedure in updating vector $\\vec u^n$ to $\\vec u^{n+1}$, is the following: \\begin{enumerate} \\item{}The unknowns $\\vec u$ are known at the center of the numerical cells at the time step $n$, i.e. $\\vec u^n$. \\item{}Reconstruction procedure allows to compute the unknowns at the interface between a cell and its neighbours. It must be done for each direction. For instance, in the $x$ direction is \\begin{eqnarray} \\vec u_{i-1,j,k}^n\\, ,\\vec u_{i,j,k}^n\\, ,\\vec u_{i+1,j,k}^n \\Longrightarrow \\left\\{ \\begin{array}{c} {\\vec u}_{i+{1\\over 2},j,k}^R \\\\ {\\vec u}_{i+{1\\over 2},j,k}^L \\end{array} \\right . \\end{eqnarray}  Reconstruction , in our code, can be linear or parabolic (PPM).  \\item{}Numerical fluxes are computed using Roe prescription, Eq.(22). For example in the $x$ direction, \\begin{eqnarray} {\\widehat {{\\vec f}}}^n_{i+{1\\over 2},j,k} = \\frac{1}{2}\\left( ( {\\vec f}({\\vec u}_{i+{1\\over 2},j,k}^{L})^n + {\\vec f}({\\vec u}_{i+{1\\over 2},j,k}^{R})^n -\\sum_{{\\eta} = 1}^{5} \\mid \\widetilde{\\lambda}_{\\eta}^x\\mid \\Delta \\widetilde {\\omega}_{\\eta} {\\widetilde {\\vec R}}^x_{\\eta})\\right) \\end{eqnarray} where \\begin{eqnarray} \\vec {L}^x_{\\eta} \\cdot ({\\vec u}^{R}_{i+{1\\over 2},j,k} - {\\vec u}^{L}_{i+{1\\over 2},j,k} )= \\Delta\\widetilde{\\omega}_{\\eta} \\, \\, \\, \\, \\, \\, , \\eta=1,...,5 \\end{eqnarray} symbol ($\\, \\tilde{} \\, $) refers to mean values in the interface. The fluxes in directions $y,z$ , $\\vec g$ and $\\vec h$, are obtained analogously. \\item{}Poisson's equation is solved by using FFT. \\item{}Sources are obtained at each cell, $\\vec s (\\vec u_{i,j,k})$. \\item{} Advancing in time, \\begin{eqnarray} {\\vec u}_{i,j,k}^{n+1}={\\vec u}_{i,j,k}^n  - \\Delta t \\cal L (\\vec u^n_{i,j,k}) \\end{eqnarray} where $\\Delta t=t^{n+1} - t^n$, and $\\cal L$ is the operator \\begin{eqnarray} \\cal L(\\vec u_{i,j,k}) &=& \\frac{\\hat{\\vec f}(u_{i+{1\\over 2},j,k}) - \\hat{\\vec f}(u_{i-{1\\over 2},j,k})} {\\Delta x_i} +\\frac{\\hat{\\vec g}(u_{i,j+{1\\over 2},k}) - \\hat{\\vec g}(u_{i,j-{1\\over 2},k})} {\\Delta y_j} \\nonumber\\\\ &+&  \\frac{\\hat{\\vec h}(u_{i,j,k+{1\\over 2})} - \\hat{\\vec h}(u_{i,j,k-{1\\over 2}})} {\\Delta z_k} + {\\vec s}(u_{i,j,k}) \\end{eqnarray}  A third order Runge-Kutta, proposed by Shu and Osher (1988), has been chosen in order to solve Eq. (B4). The expressions corresponding to this Runge-Kutta like solver are: \\begin{eqnarray*} \\vec u^{1}=\\vec u^{n} + \\Delta t \\cal L(\\vec u^{n}) \\end{eqnarray*} \\begin{eqnarray*} \\vec u^{2}={3 \\over 4}\\vec u^{n} + {1\\over 4}\\vec u^{1}+ {1\\over 4}\\Delta t \\cal L(\\vec u^1) \\end{eqnarray*} \\begin{eqnarray} \\vec u^{n+1}={1 \\over 3}\\vec u^{n} + {2\\over 3}\\vec u^{2}+ {2\\over 3}\\Delta t \\cal L(\\vec u^{2}) \\end{eqnarray} and $\\vec u^{1}$, $\\vec u^{2}$ being two intermediate states.   \\item{}The unknowns $\\vec u$ are known at the center of the numerical cells at the time step $n+1$, i.e. $\\vec u^{n+1}$.  \\end{enumerate}   \\newpage"
    },
    "9604/astro-ph9604084_arXiv.txt": {
        "abstract": "We present a useful analytic approximation to the solution of the Lane-Emden equation for infinite polytropic index - the isothermal sphere. The optimized expression obtained for the density profile can be accurate to within 0.04\\% within 5 core-radii and to 0.1\\% within 10 core-radii. ",
        "introduction": "We construct an analytic approximation to the full non-singular isothermal sphere. The approximations currently in use provide a good fit either to the inner regions ($r\\,\\leq\\,2\\,{r_{\\rm core}}$) or to the asymptotic behavior. The Lane-Emden equation for the gaseous polytrope with polytropic index ${n\\rightarrow{\\infty}}$ is identical to that of a self-gravitating isothermal sphere. All polytropes with $n\\geq$ 5 are infinite and hence no analytic solutions exist. In a recent paper, \\citeN{liu} has exhaustively examined approximate analytic solutions for polytropes with general index, where he obtains a solution for the isothermal case to within $<$ 1\\%. In this brief note, we report a simpler analytic form, accurate to within 0.04\\% within 5 core-radii. ",
        "conclusions": "The analytic approximation presented above is potentially useful in the context of many physical problems and is particularly useful since the projected quantities have simple analytic forms. We also point out that within 5 core-radii, while the analytic approximation to the isothermal sphere currently in use systematically over-estimates the mass enclosed our formula is accurate to within 0.04\\%."
    },
    "9604/astro-ph9604150_arXiv.txt": {
        "abstract": "We have analyzed spatially resolved spectra of the A754 cluster of galaxies obtained with ASCA. Through earlier observations with HEAO-1, Einstein, and ROSAT as well as optical studies, A754 has been established as the prototype system for a merger in progress. The combination of spectral and spatial resolution over a broad energy band provided by ASCA has set unprecedented constraints on the hydrodynamical effects of a cluster merger. We find significant gas temperature variations over the cluster face, indicating shock heating of the atmosphere during the merger. The hottest region, $>$ 12 keV (90\\% confidence), is located in the region of the Northwest Galaxy clump though the entire region along the cluster axis appears to be hotter than the mean cluster temperature ($\\sim$9 keV). The cool, $\\le$5 keV, gas originally found with the HEAO1-A2 experiment, resides in the exterior of the cluster atmosphere and in plume of gas we identify with a stripped cool atmosphere of the infalling subcluster. We have also attempted to reconstruct an iron abundance map of this merging system. Though poorly constrained, no significant deviations of abundance from the mean value are apparent in the individual regions.  A754 is the only cluster so far which shows the significant temperature pattern expected in a subcluster merger, in both the ROSAT (Henry \\& Briel 1995) and ASCA data, providing the first possibility to compare it with theoretical predictions. The cluster does not feature a hot peak accompanied by two hot lobes perpendicular to the cluster axis, predicted by hydrodynamic simulations of a head-on merger. The observed temperature and surface brightness maps suggest that the two colliding subunits have missed each other by about 1~Mpc, and are now moving perpendicular to the cluster axis in the image plane (as, e.g., in the simulations by Evrard et al.\\ 1996). ",
        "introduction": "A754, a rich hot cluster of galaxies at $z=0.0541$ (Bird 1994) has become the prototype of a merging cluster. It has been observed by every modern X-ray observatory, and a series of papers in the last few years, based on these data have established that there is evidence of a merger in progress. Fabricant (1986) used the X-ray imaging data obtained by the Einstein IPC and interpreted the elongated shape of the surface brightness distribution as merging subclusters. Henriksen (1986) fit a $\\beta$-model to the Einstein imaging data to obtain the gas density profile and utilized a polytropic relationship to predict gas temperature profiles which were then constrained by the HEAO1-A2 spectra.  These data required non-isothermality in the gas and were consistent with a monotonically decreasing temperature profile with the hottest material in the center and the coolest on the outside. In 1993, Henriksen analyzed HEAO1-A2 spectra along with Einstein IPC and SSS spectra and found further evidence of non-isothermality in the gas and that the data were also consistent with a simpler two-component model including a very hot and a cooler gas. The lack of combined spatial and spectral resolution in the data made the location of the thermal components ambiguous. Henry and Briel (1995, hereafter HB) analyzed ROSAT PSPC data in the energy band 0.5--2~keV, which provided a high quality image of the cluster as well as spectral constraints. These authors found that the hot gas is in the general vicinity of the North-West (NW) galaxy clump identified by Zabludoff and Zaritsky (1995), while the cluster brightness peak, in the vicinity of the South-East (SE) galaxy group, has a lower temperature as does the outer region of the cluster atmosphere.  Hydrodynamical simulations of the effect of the subcluster merger on the intracluster medium predict the existence of relatively long-lived spatial temperature variations in the post-merger cluster gas (e.g., Schindler \\& \\muller\\ 1993; Roettiger, Burns \\& Loken 1993; Evrard et al.\\ 1996). ASCA data have spectral and spatial resolution combined with a broad energy band which is sufficient to test these predictions and determine the evolutionary stage of clusters, and put constraints on the physics of cluster mergers. In this Letter, we present an analysis of spatially resolved spectra obtained by ASCA of A754, which have provided temperature and abundance maps of the cluster. ",
        "conclusions": "Summarizing the findings of ROSAT and ASCA, A754 is significantly non-isothermal. There is a ridge of hot gas located approximately alongside the cluster elongation axis, with the highest temperatures in the area of the NW galaxy concentration. The X-ray brightness peak, elongated in the direction perpendicular to the cluster axis, contains gas which is significantly cooler than the cluster average, as are the cluster outskirts toward North and East. The question which we would like to address using our derived temperature distribution is the evolutionary history of A754. Detailed hydrodynamic simulations of cluster mergers (e.g., Schindler \\& \\muller\\ 1993; Roettiger, Burns \\& Loken 1993; Pearce et al.\\ 1994) predict that as a result of a recent or ongoing merger, the cluster should have a strongly peaked temperature distribution with the highest temperature at the site of the subcluster collision, accompanied by hot lobes perpendicular to the collision axis. Briel \\& Henry (1994) reported on the detection of such hot lobes in another cluster, A2256, using ROSAT PSPC.  However, their existence was not confirmed by ASCA (Markevitch 1996), who also failed to detect the central temperature peak in that cluster. Using ASCA data, Arnaud et al.\\ (1994) and Markevitch et al.\\ (1994) reported irregular temperature structure in Perseus and in A2163, respectively, attributing it to the merger effects, although the former have ignored the PSF scattering while the statistical significance of the latter result was marginal. A754 remains the only cluster for which there is a significant and unambiguous indication of the irregular temperature structure expected in a merger, providing the first possibility to compare it with theoretical predictions.  If one assumes, in line with all previous studies of A754, that the merger in this cluster proceeds in the direction of the cluster elongation, one expects to find a temperature structure quite different from what is observed.  We offer two possible explanations for this. One is that unlike in the simulations which considered highly supersonic ($\\sim$ 3000 km~s$^{-1}$) mergers, the subcluster collision velocity is low and the post-shock gas temperature, proportional to the upstream shock velocity squared, is not significantly high. Analysis of optical data indicates that mergers may likely occur at a lower velocity, for example, $\\sim 500$ km~s$^{-1}$ for A98 (Beers, Geller \\& Huchra 1982) and $\\sim$ 1500 km~s$^{-1}$ for A3395 (Henriksen \\& Jones 1996). The line-of-sight velocity difference of the galaxy clumps in A754 is very low, about 100 km~s$^{-1}$ (Zabludoff \\& Zaritsky 1995) though application of a simple dynamical model by these authors results in a post-merger, relative velocity of $\\sim2000$ km~s$^{-1}$. However, the observed heating along the cluster axis remains to be explained in this scenario.  A more plausible explanation is a merger with a non-zero impact parameter. Evrard et al.\\ (1996) show a simulation of a merger which is reprinted here as our Fig.~2 and is strikingly similar to A754.  In this scenario, the merger proceeds in the image plane with the subunits infalling from North and South with the impact parameter $\\sim 0.5\\,h^{-1}$ Mpc, as shown by the gas velocity arrows in the left panel of Fig.~2 (gas is expected to drag behind the dark matter during a merger.)  The hottest region is at the site of the NW group penetrating the larger subunit, and the brightness elongation to the South is the tail of this group.  The plume of cool gas at the X-ray brightness peak is the stripped atmosphere, and perhaps even a cooling flow, belonged to the subunit associated with the SE galaxy group, its elongation (more clearly seen in the full resolution ROSAT image presented in HB) pointing in the direction of that group's infall.  The fact that the cold gas still exists in the cluster may indicate that the merger has not proceeded very far and this is the first encounter of the subclusters. It is therefore interesting to see if any large-scale elemental abundance differences exist, indicating that the gas belonged to different subclusters and still retains its identity (if the subunit's abundances have been different). We have averaged the abundances over the North-East (regions 2, 3, 6, 10, 13 in Fig.~1) and the South-West (regions 4, 7, 8, 11, 12) areas of the cluster associated with the two infalling subunits (excluding the regions 1, 5, 9), and found the abundances of $0.45 \\pm 0.12$ and $0.29 \\pm 0.13$ (68\\% intervals), respectively. They are consistent within the statistical errors, although there is a hint of difference and it would be interesting to measure them with a better accuracy.  To conclude, using ASCA data, we have found a prominent structure in the temperature map of A754 and put some constraints on the abundances in the different cluster regions. All of the evidence, including the X-ray image and the optical galaxy distribution, is consistent with a non-head-on merger proceeding with high velocity in the image plane. With such favorable viewing angle and the merger stage, A754 is an ideal laboratory to test various evolutionary scenarios."
    },
    "9604/astro-ph9604078_arXiv.txt": {
        "abstract": "We discuss the non-linear evolution of the angular momentum $\\bfL$ acquired by protostructures, like protogalaxies and protoclusters, due to tidal interactions with the surrounding matter inhomogeneities. The primordial density distribution is assumed to be Gaussian and the non-linear dynamics of the collisionless mass fluid is followed using Lagrangian perturbation theory. For a Cold Dark Matter spectrum, the inclusion of the leading-order Lagrangian correction terms results in a value of the rms ensemble average $\\lan\\bfL^2\\ran^{1/2}$ which is only a factor of 1.3 higher than the corresponding linear estimate, irrespective of the scale. Consequently, the predictions of linear theory are rather accurate in quantifying the evolution of the angular momentum of protostructures before collapse sets in. In the Einstein-de Sitter universe, the {\\it initial} torque is a good estimate for the tidal torque over the whole period during which the object is spun up. ",
        "introduction": "The problem of the acquisition of angular momentum by protostructures in the universe is of considerable interest in theories of galaxy and cluster formation. A widely accepted view is that the present luminous structures acquired their spin via gravitational tidal interactions with the surrounding matter inhomogeneities (Hoyle 1949; Peebles 1969; Doroshkevich 1970; White 1984; Barnes \\& Efstathiou 1987; Hoffman 1986, 1988; Heavens \\& Peacock 1988; Ryden 1988; Quinn \\& Binney 1992; Eisenstein \\& Loeb 1995; Catelan \\& Theuns 1996).  So far, however, the theoretical analysis of the growth of the tidal galaxy angular momentum has been essentially limited to the linear regime, during which the galaxy spin grows proportionally to the cosmic time $t$ (Doroshkevich 1970; White 1984).  In this paper we examine analytically, for the first time, the question of how the galaxy tidal angular momentum evolves during the mildly non-linear regime. Previous attempts in this direction may be found in Peebles (1969) and White (1984). However, in contrast to their approach, we employ actual perturbative solutions of the dynamical equations that describe the motion of the fluid. More in detail, we apply Lagrangian perturbation theory.  The Lagrangian approach shows to be $ideal\\,$ in treating the evolution of the galaxy spin, because it is powerful in describing the non-linear growth of the mass-density fluctuations on one hand (Zel'dovich 1970a, b; Buchert 1992; Bouchet \\etal 1992; Catelan 1995); and on the other hand the usual difficulty of inverting the mapping from Lagrangian coordinates $\\bfq$ to Eulerian coordinates $\\bfx$ is completely by-passed. This is because the angular momentum $\\bfL$ is $invariant\\,$ with respect to the Eulerian or Lagrangian description.  The layout of this paper is as follows: in the next section we first briefly review the basics of the Lagrangian theory and the perturbative solutions of the Lagrangian fluid equations. Next, we compute within this framework the perturbative corrections to the linear tidal angular momentum $\\bfL^{(1)}$ acquired by a protoobject. The resulting expressions are then simplified by calculating their averages over the ensemble of realisations of the linear Gaussian gravitational potential $\\psi^{(1)}$ for objects with given inertia tensor. We then compare the non-linear spin corrections with the results of the linear analysis recently performed by Catelan \\& Theuns (1996). In the main text we restrict ourselves mainly to the case of a flat universe, leaving the more involved treatment of closed and open universes to Appendices. ",
        "conclusions": "In this paper we analysed the growth of the tidal angular momentum $\\bfL$ acquired by a protoobject (protogalaxy or protocluster) during the mildly non-linear evolution of the matter density perturbations, assuming the latter to be Gaussian distributed. The dynamics of the collisionless matter fluid is described using the Lagrangian approach in the formulation given by Catelan (1995). This formulation is very suitable to study the problem at hand, because the Lagrangian expressions are considerably simpler than their Eulerian counterparts, yet the protogalaxy's tidal spin is a vector {\\it invariant} under the change of Eulerian to Lagrangian spatial coordinates, $\\bfx$ and $\\bfq$ respectively. Specifically, the difficult problem of inverting the mapping $\\bfx=\\bfq+\\bfS$, where $\\bfS$ is the displacement vector, in order to recover the Eulerian quantities from the Lagrangian ones, is completely avoided.  The strategy we follow is straightforward. The non-linear spin corrections $\\bfL^{(h)}$, where $\\bfL^{(1)}$ is the linear angular momentum, are calculated approximating the fluid elements' trajectories $\\bfS$ by the perturbative solutions $\\bfS_h$ of the Lagrangian fluid equations~(8) and (9). This leads to the expression~(20). Since we are interested in computing the lowest-order perturbative corrections to the ensemble average $\\lan\\bfL^{(1)2}\\ran$ for objects with given inertia tensor, we need to calculate corrections to $\\bfL$ up to third-order. This has the added advantage that we take account of the full physical content of equations~(8) and (9), since the latter are cubic in the displacement. The calculation is summarised as follows: from the knowledge of $\\bfS=\\bfS_1+\\bfS_2+\\bfS_3$ (where $\\bfS_1$ corresponds to the displacement in Zel'dovich approximation), we deduce the corresponding corrections $\\bfL=\\bfL^{(1)}+\\bfL^{(2)}+\\bfL^{(3)}$ and finally get the perturbative expansion $\\lan\\bfL^2\\ran= \\lan\\bfL^{(1)2}\\ran+\\lan\\bfL^{(2)2}\\ran+2\\,\\lan\\bfL^{(1)}\\!\\cdot\\bfL^{(3)}\\ran$. The term $\\lan\\bfL^{(1)}\\!\\cdot\\bfL^{(2)}\\ran$ is zero for an underlying Gaussian matter distribution, but it should be taken into account in the framework of more general non-Gaussian statistics (work in progress). Assuming Gaussian statistics here, we disregard it. In sections 2.1 and 2.2 (for the Einstein-de Sitter universe; in Appendix A for a more general Friedmann universe) we reviewed the Lagrangian theory and the perturbative solutions $\\bfS_1, \\bfS_2$ and $\\bfS_3$ of the Lagrangian fluid equations. Using these results we calculate the corrections $\\lan\\bfL^{(2)2}\\ran$ and $\\lan\\bfL^{(1)}\\!\\cdot\\bfL^{(3)}\\ran$ in section 3, after summarising the results of linear theory (i.e., the term $\\lan\\bfL^{(1)2}\\ran$). The final expressions are rather cumbersome (the details of the calculations have been deferred to Appendix B), but we can summarise the main features of our results as follows: for an Einstein-de Sitter universe, \\begin{itemize} \\item $\\lan\\bfL^{(1)2}\\ran^{1/2} \\propto \\tau^{-3} \\propto t \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\left[\\propto \\dot{D}(\\tau)\\right]\\;;$\\\\ \\item $\\lan\\bfL^{(2)2}\\ran^{1/2} \\propto \\tau^{-5} \\propto t^{5/3} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; \\;\\;\\;\\;\\;\\;\\;\\left[\\propto \\dot{E}(\\tau)\\right]\\;;$\\\\ \\item $\\lan\\bfL^{(1)}\\!\\cdot\\bfL_h^{(3)}\\ran^{1/2} \\propto \\tau^{-5} \\propto t^{5/3} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\left[\\propto \\left(\\dot{D}(\\tau)\\dot{F_h}(\\tau)\\right)^{1/2}\\right]\\;;$\\\\ \\item $\\lan\\bfL^{(1)}\\!\\cdot\\bfL^{(12)}\\ran^{1/2} \\propto \\tau^{-5} \\propto t^{5/3} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\left[\\propto \\left(\\dot{D}(\\tau)[\\dot D(\\tau)E(\\tau)-D(\\tau)\\dot E(\\tau)]\\right)^{1/2}\\right]\\;,$ \\end{itemize} \\noindent where $D$ is the growth factor of the density perturbations, $E$ and $F_h$ ($h=a,b,c$) are the growing modes of the second- and third-order Lagrangian displacements, respectively. We see that the perturbative corrections to $\\lan\\bfL^{(1)2}\\ran^{1/2}$ grow proportionally to $t^{5/3}$ in the Einstein-de Sitter universe, in agreement with Peebles (1969). The expressions between square brackets give the generalisations of the results for a generic Friedmann universe. Interestingly, all the corrections we have analysed are proportional to the same invariant of the inertia tensor $\\calJ$ of the matter contained in the homogeneous Lagrangian volume $\\Gamma$, a result we can express as \\begin{itemize} \\item $\\lan\\bfL^{(1)2}\\ran \\propto \\lan\\bfL^{(2)2}\\ran \\propto \\lan\\bfL^{(1)}\\!\\cdot\\bfL^{(3)}\\ran \\propto \\mu_1^2 - 3\\mu_2\\;$, \\end{itemize} \\indent where $\\mu_1$ and $\\mu_2$ are the first and the second invariant of the inertia tensor (see equations~(40) and (41)). This invariant $\\mu_1^2 - 3\\mu_2$ has been thoroughly investigated in Catelan \\& Theuns (1996). As a consequence of this factorisation we have been able to express the order of magnitude of the non-linear corrections to $\\lan\\bfL^2\\ran$ in terms of the linear contribution, $\\lan\\bfL^2\\ran = (1+\\Upsilon)\\,\\lan\\bfL^{(1)2}\\ran$ (equation~(54)), where $\\Upsilon\\approx 0.6$ for the standard CDM spectrum at galactic scales. Taking into account that the non-linear correction is small, we conclude that linear theory gives a good description of the angular momentum up to maximum expansion. Since in addition linear theory predicts, in the Einstein-de Sitter model, a growth rate $\\bfL \\propto t$, it follows that the {\\it initial} torque is a good estimate for the tidal torque over the whole period during which the object is spun up: $d\\bfL(t)/dt\\approx d\\bfL(0)/dt$.  Finally, as is the case with almost any analytic calculation, comparison of these results against observations is hampered by the fact that the very final stages of galaxy formation are likely to be highly non-linear and in addition dissipative processes may play an important role as well. Analytic investigations are not able to take such highly complex phenomena into account."
    },
    "9604/astro-ph9604002_arXiv.txt": {
        "abstract": "We present a simple toy model of the distribution of objects responsible for gravitational microlensing.  We use Monte Carlo simulations to demonstrate how difficult it is to determine the parameters of the lens mass distribution on the basis of the observed distribution of event time scales.  A robust determination requires $ \\sim 100 $ events, or more, even if the geometry of lens distribution, and the lens kinematics are known. ",
        "introduction": "One of the main objectives of the searches for microlensing events (Paczy\\'nski 1996, and references therein) is to determine the mass function of the lensing objects. The determination of typical lens masses, or even their distribution function, were attempted by Alcock et al. (1993), Udalski et al. (1994), Zhao et al. (1995, 1996), Han \\& Gould (1996), and many others.  We know that some lenses must be ordinary stars, but some may be brown dwarfs, and perhaps also planetary mass objects, and even more exotic things like black holes.  It is recognized that even if all lensing objects were of the same mass $ M $, they would give rise to a broad range of time scales $ t_0 $ of the observed microlensing events. Although this is known to a relatively small number of people in the field this fact is not generally appreciated. Considering the broad interest in the nature of dark matter, and the relevance of microlensing searches to this problem, we feel it is justified to present a very simple model of the lens distribution and kinematics which allows some results to be obtained analytically, and the rest can be calculated with one-dimensional numerical integrals.  This model retains some of the most important generic characteristics: a very broad distribution of $ t_0 $ with power law tails towards very short and very long time scales for a delta function distribution of lens masses. Using Monte Carlo simulations we demonstrate how difficult it is to obtain accurate information about the mass function even within the framework of our simple model. In reality one is not sure what is the correct space distribution and kinematics of lensing masses, making the interpretation even more difficult.  The model is described in the next section.  The results of the Monte Carlo simulations are presented in section 3, and the discussion of the results is given in the last section. The technical details of the model can be found in the appendix. ",
        "conclusions": "The Monte Carlo simulations presented in this paper demonstrate that even if we have a full knowledge of the space distribution and kinematics of the lensing objects, or equivalently, if we know the relation between the lens mass and the distribution of event time scales (see eq. 4 for an example), it may still be difficult to obtain the parameters of the lens mass function. The problem is particularly serious when the width of the mass function is narrow ($\\beta \\la 1$). In addition, if the mass function power law is quite different from $-1.5$, then it will be difficult to probe the high mass end when $\\alpha \\ll -1.5$, and the low mass end when $\\alpha \\gg -1.5$ (cf. eq. 8).  Recently, Han \\& Gould (1996) obtained a rather tight limit on the mass function of the Galactic disk using about 50 microlensing events toward the Galactic bulge. In their analysis, the upper mass limit is fixed at $10 M_{\\odot} $ in a rather ad hoc manner. As we have shown in Fig. 2, such a fixed upper mass limit will make the determination of the other parameters appear more accurate. A more realistic approach is to estimate the upper mass limit using the maximum likelihood in order to avoid carrying our a priori assumption into a limit on the physical parameters. In general, if at least three separate parameters are to be estimated, for example $ M_0$, $ \\alpha $, and $ \\beta $, then any method must effectively use the information contained in at least the first three moments.  However, the higher the moment the more uncertain is its value as estimated from a relatively small sample of events. Therefore, the accuracy of the determination of any distribution parameter is lower in the 3-parameter determination than in a 2-parameter determination. Unfortunately, the smaller errors in a 2-parameter determination could be misleading, unless we have an independent and reliable information about the value of the third parameter.  In our second example, very broad mass functions are sampled through a realistic detectability window (cf. Fig. 4). The current published sensitivity windows allow one to probe about three to four decades of mass range and we find that the ``amplitude'' $ A $ and the local slope of the lens mass function $ \\alpha $ can be determined reasonably well with $ \\sim 100 $ events. However, the estimate of the range of masses within the detectability window, $ \\beta _{\\mathrm{min}} $ or $ \\beta _{\\mathrm{max}} $, may be very uncertain with as many as $ n = 100 $ events, as shown in Fig. 4.  If the mass function is broader than the detectability window then the total event rate or the total optical depth (or total mass) cannot be measured as the majority of events are outside the detectability window (Fig. 3), having either too short or too long $ t_0 $.  This points to the necessity of broadening the detectability window, a task easy to accomplish in the near future. Some first attempts have already been made by the EROS (Ansari et al. 1995) and MACHO collaborations (Bennett et al. 1996).  How can we be sure that we have reached the low and the high mass end of the lens distribution?  In principle this is simple: we have to broaden the detectability window and we should detect so many events that the generic power law tails in the distribution (cf. eq. 3) become apparent.  This may call for well over 100 events. Such a high number of events is within reach for the Galactic Bulge, where the rate appears to be very high (Udalski et al. 1994, Alcock et al. 1995b,c). The determinations of any parameters of the distribution of lens masses based on as few as 10 events are subject to very large uncertainties, as shown with the dotted contours in Fig. 2.  In reality we do not know what the geometry of the lens distribution is, and what is the lens kinematics.  While looking towards the Galactic Bulge the majority of lenses may be in the Bulge itself (Kiraga \\& Paczy\\'nski 1994, Udalski et al. 1994, Paczy\\'nski et al. 1994, Zhao et al. 1995, 1996), or they may be in the disk (Alcock et al. 1995b,c).  While looking towards the LMC the majority of lenses may be in our galaxy (Alcock et al. 1995a) or in the LMC (Sahu 1994; Wu 1994).  Depending on the location and kinematics the relation between the lens masses and the event time scale $ t_0 $ may be very different, and this leads to an additional major uncertainty on top of purely statistical uncertainty discussed in this paper. It seems rather difficult to disentangle all parameters one needs to describe the distribution and kinematics of the lenses {\\it and} their mass function on the basis of the observed distribution of event time scale.  Fortunately, the future studies of the variation of the optical depth with the location in the sky will help to identify the lens location. If lensing of the Galactic Bulge stars is dominated by the Bulge lenses then the optical depth should vary rapidly with the galactic longitude (Kiraga \\& Paczy\\'nski 1994).  If such a variation is not present this will indicate that the lenses are located mostly in the galactic disk. If lensing towards the LMC is dominated by the LMC lenses then the optical depth should increase strongly towards the LMC center (Sahu 1994).  If the optical depth is uniform over the LMC then the lenses must be mostly in our galaxy.  This is very simple in principle, but it will take hundreds of events to establish beyond reasonable doubt.  Such a large number will be readily collected towards the Galactic Bulge, but it will take a very long time for the LMC as the rate in that direction is very low (Alcock et al. 1995a).  Once the geometry of the lens distribution is established it will be possible to develop reliable models of their kinematics, and infer a reasonable statistical relation between the lens masses and the event time scales $ t_0 $.  \\vskip 0.5cm  This project was supported by the NSF grant AST93-13620 and by the ``Sonderforschungsbereich 375-95 f\\\"ur Astro-Teilchenphysik'' der Deutschen Forschungsgemeinschaft. We are very grateful to Dr. Peter Schneider for a critical reading of the manuscript.  \\newpage \\appendix"
    },
    "9604/astro-ph9604144_arXiv.txt": {
        "abstract": "The parallax effect in ground-based microlensing (ML) observations consists of a distortion to the standard ML light curve arising from the Earth's orbital motion. This can be used to partially remove the degeneracy among the system parameters in the event timescale, $t_0$. In most cases, the resolution in current ML surveys is not accurate enough to observe this effect, but parallax could conceivably be detected with frequent followup observations of ML events in progress, providing the photometric errors are small enough. We calculate the expected fraction of ML events where the shape distortions will be observable by such followup observations, adopting Galactic models for the lens and source distributions which are consistent with observed microlensing timescale distributions. We study the dependence of the rates for parallax-shifted events on the frequency of followup observations and on the precision of the photometry. For example, we find that for hourly observations with typical photometric errors of 0.01 mag, 6\\% of events where the lens is in the bulge, and 31\\% of events where the lens is in the disk, (or $\\approx 10$\\% of events overall) will give rise to a measurable parallax shift at the 95\\% confidence level. These fractions may be increased by improved photometric accuracy and increased sampling frequency. While long-duration events are favored, the surveys would be effective in picking out such distortions in events with timescales as low as $t_0 \\approx 20$ days. We study the dependence of these fractions on the assumed disk mass function, and find that a higher parallax incidence is favored by mass functions with higher mean masses. Parallax measurements yield the reduced transverse speed, $\\tilde{v}$, which gives both the relative transverse speed and lens mass as a function of distance. We give examples of the accuracies with which $\\tilde{v}$ may be measured in typical parallax events. Fitting ML light curves which may be shape-distorted (e.g., by parallax, blending, etc.) with only the 3 standard ML parameters can result in inferred values for these quantities which are significantly in error. Using our model, we study the effects of such systematic errors and find that, due primarily to blending, the inferred timescales from such fits, for events with disk lenses, tend to shift the event duration distribution by $\\approx 10$\\% towards shorter $t_0$. Events where the lens resides in the bulge are essentially unaffected. In both cases, the impact-parameter distribution is depressed slightly at both the low and high ends. ",
        "introduction": "The recent observations by the MACHO (Alcock et al.~1995a, Alcock et al.~1996), EROS (Aubourg et al.~1993), OGLE (Udalski et al.~1994a), and DUO (Alard et al.~1995) collaborations of gravitational microlensing (ML) of stars in the LMC and Galactic bulge have generated tremendous excitement in astrophysics. These surveys provide a new probe of Galactic structure and low-mass stellar populations. ML observations of the LMC allow measurements of the dark-matter content of the Galactic halo (Bennett et al.~1996), placing important constraints on dark-matter theories. Also intriguing have been observations of ML events towards the Galactic center, which probe the inner disk and bulge of our Galaxy. The roughly 100 events observed to date imply an optical depth towards the bulge of ($3.3 \\pm 1.2$) $\\times 10^{-6}$, three times that predicted by theoretical estimates (Griest et al.~1991), and reveal the need for a better understanding of Galactic structure.  One explanation for the observed excess of bulge events would be the presence of  a hitherto undiscovered population of compact, sub-stellar objects, implying the existence of more mass in the disk than previously believed (Alcock et al.~1994; Gould, Miralda-Escude, \\& Bahcall 1994). This would require an upturn in the stellar mass function (MF) below the hydrogen-burning limit. Another possibility is that the lenses comprise an ordinary stellar population, and the enhancement of ML events is due to non-axisymmetric structure, such as a bar, in the Galactic bulge. The optical depth calculated {}from a self-consistent bar model, matching both kinematic observations of the bulge and the COBE image of the Galaxy, is consistent with the MACHO and OGLE results (Kiraga \\& Paczy\\'{n}ski 1994; Zhao, Spergel, \\& Rich 1995; Zhou, Rich, \\& Spergel 1995). Furthermore, the best-fit models to the observed ML duration distribution have a median lens mass of about $0.2\\ \\Msolar$, {\\em above} the hydrogen-burning limit, and consistent with ordinary stellar populations.  Determining the nature and location of the lenses will enable us to learn whether the enhancement of bulge events is due to lenses in the bar or disk, and probe a region of the stellar MF about which little is known. However, while numerous cases of ML have been observed, essential information such as the mass, distance, and velocity distributions of the lensing population can only be disentangled in a statistical manner, due to the degeneracy of microlensing light curves with respect to these parameters. For an ideal ML event, where the source and lens are assumed to be point-like, the lens is assumed to be dark, and the velocities of the observer, source, and lens are constant, the amplification is given by \\begin{equation} A[u(t)] = \\frac {u^2 +2} {u(u^2 +4)^{1/2}};\\qquad  u(t)= \\left[  \\left( \\frac {t-t_{\\max}} {t_0} \\right)^2 + u_{\\min}^2\\right]^{1/2}, \\label{standardamplification} \\end{equation} where $u$ is the impact parameter from the observer-source line to the lens in units of the Einstein radius, ${t_0} = R_{e}/v$ is the event timescale, $v$ is the transverse speed of the lens relative to the source-star line of sight, and $t_{\\max}$ is the time at which peak amplification, $A_\\max = A(u_{\\min})$, occurs. The Einstein radius, $R_{e}$, is determined by the geometry of the event and the lens mass, and is given by \\begin{equation} R_e = \\left[ {4G\\mass_l D_{ol} (D_{os} - D_{ol}) \\over c^2 D_{os}} \\right]^{1/2}, \\label{einsteinradius} \\end{equation} where $\\mass_l$ is the lens mass and $D_{ol}$ and $D_{os}$ denote the observer-lens and observer-source distances, respectively. We see from Eq.~(\\ref{standardamplification}) that the ML light curve can be fit by the three parameters $t_0$, $t_{\\max}$, and $u_{\\min}$. The latter two, however, only tell when the lens passed nearest to the line of sight and how close it came, revealing little useful information about the event. All the parameters of interest, namely  $\\mass_l$, $v$, $D_{ol}$, and $D_{os}$, are folded into the single parameter $t_0$. (In practice, $D_{os}$ may be obtained to reasonable accuracy.)  Several techniques for breaking the degeneracy have been proposed, many of which involve fitting for distortions to the generic ML light curve which may be detected in some fraction of observed events (Nemiroff \\& Wickramasinghe 1994; Witt \\& Mao 1994; Maoz \\& Gould 1994; Loeb \\& Sasselov 1995; Han \\& Gould 1995; Gaudi \\& Gould 1996). These distortions arise from violations of the standard ML assumptions. For example, if the lenses consist partly of ordinary, low-mass dwarfs in the Galactic disk and bulge, the contribution of the lens light would distort the observed light curve differently in different wavebands, violating the assumption of achromaticity, and distorting the shape, [c.f. Eq.~(\\ref{standardamplification})] of the light curve (Kamionkowski 1995; Buchalter, Kamionkowski, \\& Rich 1995, hereafter BKR). Ground-based observations of this color-shift effect in two or more wavebands can remove entirely the degeneracy in $t_0$. If the lenses are not completely dark, then color shifting should, in principle, affect every event and the observed incidence of this effect is simply a function of resolution of the data (namely sampling frequency and level of photometric error). The MACHO collaboration has already observed another type of distortion, where the time symmetry of the light curve is broken by the Earth's orbital motion (Alcock et al.~1995b). This parallax effect allows one to compare the projected Einstein ring with the size of the Earth's orbit and thereby obtain an additional constraint relating $\\mass_l$, $v$, and $D_{ol}$, effectively removing one degree of degeneracy. Strictly speaking, this effect is also present in every event, but is not expected to be frequently observed by any of the existing ML surveys; the light curves are sampled too infrequently and the photometric errors are too large. The single confirmed parallax event detected to date was fortuitous in that it was both long enough ($t_0 \\approx 110$ days) for the light curve to be heavily sampled, and well-situated during the MACHO observing season so that the asymmetry could be fit along the entire light curve. In addition, the effect is more dramatic for longer events, during which the Earth can move through an appreciable fraction of its orbit. It is only for these rarer long events that a parallax shift may be observed by the low-resolution ML surveys. However, with the early-warning alert systems developed by both the MACHO and OGLE groups (Stubbs et al. 1994; Udalski et al. 1994b), it is conceivable that a program of followup observations with frequent and precise measurements could measure the light curves with sufficient accuracy to detect the parallax effect in a larger fraction of events. The PLANET (Albrow et al.~1995) and GMAN (Pratt et al.~1996) collaborations are currently performing such observations, with the primary purpose of detecting planets around lenses, and Tytler et al. (1996) are proposing another such search. Planetary masses would give rise to smaller event timescales and thus produce narrow spikes on the lensing light curve which likely go undetected in current surveys. These spikes could be resolved with observations by dedicated telescopes performing rapid sampling of events in progress with high photometric precision. Such ground-based surveys are well-suited to pick out the parallax effect for shorter-term events, as well as distortions due to unlensed light from the lens.  In this paper, we determine the fraction of ML events toward the Galactic bulge which should show a measurable parallax shift in such followup programs. We calculate the fractions which will arise if the lenses are all in the bulge and if the lenses are all in the disk, using realistic models for the lens distributions which are consistent with the currently observed ML timescale distribution. A Monte Carlo technique is used to simulate events and generate parallax-distorted light curves for each. It is then determined for each event whether the distorted light curve can be distinguished at the 95\\% confidence level from a standard ML light curve. The calculation is performed for several values of the sampling frequency and for several values of photometric accuracies that may be attainable. Our results indicate that, depending on the survey sensitivity, the expected fraction of parallax-shifted events ($F_{PSE}$) ranges from 0.06 to 0.31 for bulge self-lensing events, 0.31 to 0.77 for disk-lensing-bulge events, and 0.49 to 0.83 for (rare) disk self-lensing events, with current-generation followup surveys favoring the lower values. Since the parallax effect is particularly important for the case of disk lenses, we also examine the dependence of this effect on the adopted mass function (MF) for the disk, and find higher incidences from MFs with higher mean masses. We further calculate the dependence of $F_{PSE}$ on event timescale, finding that the intensive followup programs should measure a substantial contribution to $F_{PSE}$ from events with $t_0$ as low as 20 days, and we also examine what may be learned from a typical PSE. We emphasize that, like color-shift analysis, parallax analysis can be directly applied to all events, so that information about the mass, velocity and spatial distributions of the lenses becomes available on an event by event---rather than on a statistical---basis.  Although parallax and blending effects may be present to some degree in a large fraction of events, they are expected to be small in most cases, so that a standard three-parameter fit to a slightly distorted light curve may be deemed adequate. However, fitting the standard ML amplification function to a light curve distorted by these effects can result in systematic errors in the inferred fit parameters, most notably $t_0$ and $u_{\\min}$. These inaccuracies may lead to a systematic miscalculation of the duration and amplification distributions, and of the overall optical depth. Thus, we perform another simulation to generate shape-distorted light curves including both parallax and color-shift effects (i.e., blending due to the lens). A standard three-parameter fit to these light curves is then applied, and the inferred parameters are used to compare the resulting amplification and event duration distributions with their actual values. We find for the latter that the errors incurred by three-parameter fits are negligible for bulge lenses, but may be quite large ($\\geq 10\\%$) if the lenses are in the disk, particularly if the disk MF declines at the low-mass end. We also find that errors in the inferred minimum impact parameter lead to a non-uniform distribution in peak amplification.  The plan of the paper is as follows: In Section 2, we review the characteristics of parallax-distorted events. In Section 3, we summarize the models and techniques used in the calculation and present the results of our calculations of expected fraction of PSEs. In Section 4, we examine the impact of parameterizing shape-distorted light curves with the standard ML fit on the inferred ML parameters, and in Section 5 we make some concluding remarks. ",
        "conclusions": "While many prospects for learning more about ML events involve intensive followup observations aimed at resolving distortions to the standard ML light curve, few detailed calculations of rates and features of such distortions exist for realistic Galactic models. We have employed a detailed model to calculate the fraction of ML events which will exhibit a significant parallax distortion, as seen by the followup monitoring programs. We find that with frequent and precise observations of events in progress, from 10\\% up to 39\\% of all events arising from various lens-source pairings are expected to exhibit a parallax shift, including some events with $t_0 \\lesssim 2$ months. Such events can be used to place an additional constraint among the parameters of interest (namely $\\mass_l,D_{ol},\\mbox{ and }v$) with reasonable precision, allowing one to express $\\mass_l$ as a function of $D_{ol}$ and thus better determine the characteristics of the lensing population. Failing to account for ML distortions, particularly blending, can lead to errors in the inferred parameters (most notably $t_0$) which propagate into the inferred duration distribution and impact-parameter distribution and impact the assumed overall optical depth.  The most important application of analyzing such distortions is in discerning whether observed bulge events are due to structure in the bulge, or excess mass in the form of sub-stellar objects in the disk. However, the significance of such work extends to many areas. Determining the origin of the excess events will lead to a better understanding of Galactic structure and stellar populations, particularly the MF of low-mass stars in the disk and bulge. The latter relates directly to our understanding of the process of star formation. Moreover, if the bulge of our Galaxy is representative of relaxed stellar systems such as elliptical galaxies, then knowledge of the bulge luminosity function and dynamical structure will have implications for galaxy formation and evolution. Precise knowledge of the mass of the disk and bulge also constrains the halo mass and places strong limits on the dark-matter content. If it can be determined that there is more mass in the bulge and/or disk, this implies less halo dark matter and has important consequences for the predicted event rates in direct and indirect searches for exotic dark matter (Jungman et al.~1996)."
    },
    "9604/gr-qc9604020_arXiv.txt": {
        "abstract": "\\baselineskip .15in We study the motion of a spinning test particle in Schwarzschild spacetime, analyzing the Poincar\\'e map and the Lyapunov exponent. We find chaotic behavior for a particle with spin higher than some critical value (e.g. $S_{cr} \\sim 0.64 \\mu M$ for the total angular momentum $J=4 \\mu M$),  where $\\mu$ and $M$ are the masses of a particle and of a black hole, respectively.  The inverse of the Lyapunov exponent in the most chaotic case is about three orbital periods, which suggests that chaos of a spinning particle may become important in some relativistic astrophysical phenomena.  The ``effective potential\" analysis enables  us to classify the particle orbits into four  types as follows. When the total angular momentum $J$ is large, some orbits are bounded and the ``effective potential\"s are classified into two types: (B1) one saddle point (unstable circular orbit) and one minimal point (stable circular orbit) on the equatorial plane exist for small spin;  and (B2) two saddle points bifurcate from the equatorial plane and one minimal point remains on the equatorial plane for large spin. When $J$ is small, no bound orbits exist and the potentials are classified into another two types: (U1) no extremal point is found for small spin; and (U2)  one saddle point appears on the equatorial plane, which is unstable in the direction perpendicular to the equatorial plane, for large spin.   The types (B1) and (U1) are the same as those for a spinless particle, but the potentials (B2) and (U2) are new types caused by spin-orbit coupling.  The chaotic behavior is found only in the type (B2) potential. The ``heteroclinic orbit'', which could cause chaos, is also observed in type (B2). ",
        "introduction": "\\label{sec1}\\setcounter{equation}{0} Chaos is now one of the most important ideas used to explain various non-linear phenomena in nature.  Since the research on the three body problem by Poincar\\'e, many studies about chaos in celestial mechanics and astrophysics have been done and revealed the important role of chaos in the Universe\\cite{moser},\\cite{wisdom}. Although we know many features of  chaos in Newtonian dynamics, we do not know, so far, so much  about those in general relativity.   If gravity is strong, e.g., a close binary system or a particle near a black hole, we have to use Einstein's theory of gravitation. Because the gravitational field  in general relativity is non-linear, we may find a new type of chaotic behavior in strong gravitational fields, which do not appear in Newtonian dynamics\\cite{misner}-\\cite{barrow}. In a previous paper\\cite{sota}, we studied a criterion for chaos of a test particle motion around an $N$-black hole system (or an $N$-naked singularity system) and found that a local instability determined by the Riemann curvature tensor provides us a sufficient test for chaos.  We also found that the existence of an unstable circular orbit, which guarantees the existence of a homoclinic or heteroclinic orbit, plays a crucial role for chaos.   However, the relativistic systems  analyzed  so far by several authors\\cite{sota}-\\cite{moeckel}, in which chaotic behavior of a test particle is found, are rather unrealistic\\cite{sota}-\\cite{varvoglis}, except for the perturbed spacetimes of the Schwarzschild black hole solution\\cite{bombelli},\\cite{moeckel}.  As for the other interesting cases of systems, for example, the $N$-extreme black hole system\\cite{contopoulos}-\\cite{yustsever} is unstable, the existence of a strong uniform magnetic field around a black hole\\cite{karas} is not likely, and naked singularities\\cite{sota} may not exist.  We may wonder whether any realistic relativistic system can be chaotic and when chaos may play  an important role in such a relativistic astrophysical phenomena.  In astrophysics, rotation of a system plays a quite important role.  The angular momentum or spin may completely change the evolution of the system. In a dynamical system, rotation or spin is one of the most important elements and it may sometime cause chaotic behavior.  For example, the H\\'enon-Heiles system, which describes the motion of a star moving in the potential of a rotating galaxy, is chaotic\\cite{henon}. We also know that some spin-orbit interaction induces chaos in Newtonian gravity.  This may also be true in a relativistic system such as the evolution of a binary system. The motion of coalescing binary systems of neutron stars and/or black holes is very important to  study because they are promising sources of gravitational waves, which we are planning to  detect by large-scale laser interferometric gravitational observatories, such as US LIGO\\cite{abromovici}. If we will detect the signal of gravitational waves emitted from these systems and compare it with theoretical templates, we may be able to determine a variety of astrophysical parameters of the sources such as their direction, distance, masses, spin, and so on\\cite{cutler}. In order to extract exact information about such sources from the observed signal we need the exact theoretical templates of the gravitational waveforms. To  make such templates, it is very important to know the exact motion of sources. Hence, the equations of motion in the post-Newtonian expansion in terms of a small parameter $\\epsilon\\approx (v/c)^2 \\sim GM/r$ have been studied  by many authors\\cite{damour}. Those can be written schematically as \\begin{equation} \\frac{d^2\\mbox{\\boldmath$x$}}{dt^2}=\\mbox{\\boldmath$a$}_{N} +\\mbox{\\boldmath$a$}_{PN}^{(1)}+\\mbox{\\boldmath$a$}_{SO}^{(3/2)} +\\mbox{\\boldmath$a$}_{2PN}^{(2)}+\\mbox{\\boldmath$a$}_{SS}^{(2)} +\\mbox{\\boldmath$a$}_{RR}^{(5/2)}+O(\\mbox{\\boldmath$a$}^{(3)}), \\label{1.1} \\end{equation} where the subscripts $N, PN, SO, 2PN, SS$ and $RR$ denote Newtonian, post-Newtonian, spin-orbit coupling, 2nd post Newtonian, spin-spin coupling, and radiation reaction terms, respectively\\cite{will}; and the superscript corresponds to the order of expansion in $\\epsilon$. To make sufficient templates, we may need at least the 3rd-order post Newtonian contribution to obtain  the S/N ratio required from the observation\\cite{3rd1}-\\cite{3rd3}.  But this is still under investigation in the world.  The spin effect is also important. The spin terms in Eq.(\\ref{1.1}), $\\mbox{\\boldmath$a$}_{SO}$ and $\\mbox{\\boldmath$a$}_{SS}$, induce a precession of the orbital plane through the spin-orbit or spin-spin coupling, resulting in modulation of the gravitational waveforms \\cite{kidder},\\cite{apostlatos}. In \\cite{apostlatos}, it is also shown that the orbital plane may behave  very strangely  due to the spin effects. We  cannot verify whether or not any chaotic behavior occurs  in their system. But from  the studies on spin effects in Newtonian dynamics, we know that a spin effect can make a motion chaotic. We then expect that a relativistic system such as  a coalescing binary pulsar may also show the similar non-linear phenomena. The gravitational waveform from the system with chaotic motion will be different from a system with a regular motion, for example, a regular precession of the orbital plane as shown in \\cite{kidder}. The chaos might be too strong to make a complete template of gravitational waves, or rather it might give us new information about astrophysical parameters from a time series of the observed waveforms.  We will discuss this problem for a coalescing binary system with highly spinning bodies elsewhere.  Thus we believe that a study about spin effects on the orbital evolution of a relativistic system and its gravitational waveform is very important from the viewpoint of observations as well as of academic interest. In this paper, to clarify the spin effect on the orbital motion, especially the spin-orbit interaction, we study the motion of a spinning test particle around a Schwarzschild black hole. So far, studies about a spinning test particle in relativistic spacetime have been done by many authors since the basic equations were derived by Papapetrou\\cite{papa} and reformulated by Dixon\\cite{dixon}. Corinaldesi and Papapetrou already discussed a spinning test particle in Schwarzschild spacetime\\cite{papa2}. But, apart from the supplementary condition, from which they adopted a different equation from the present standard one,  they presented the basic equations and discussed some terms with physical interpretations. They did not analyze the orbits in detail from a viewpoint of the dynamical system.  Kerr or Kerr-Newman spacetime was also analyzed by several authors\\cite{Rasband}-\\cite{Rudiger}. In \\cite{Rasband},\\cite{Tod} and \\cite{hojman},  the effective potential of the spinning particle is given and the spin effects on the binding energy are discussed. In \\cite{mino} and \\cite{mino2},  the gravitational waves produced by a spinning particle falling into a Kerr black hole or moving circularly around it is discussed and the energy emission rate from those systems is calculated. But in those papers they discussed only the case of the orbit in  the equatorial plane or on the symmetric axis of the black hole. Since we  are interested in chaotic motion induced by spin-orbit coupling here, we have to discuss the most generic situation, i.e., the orbital motion off the equatorial plane.  This paper is organized as follows. In section 2 we shall briefly review the basic equations, i.e., the equations of motion for a spinning test particle in relativistic spacetime, a supplementary condition and some constants of motion.  We specify the background spacetime to be a Schwarzschild black hole, then we write down those equations and introduce a sort of  ``effective potential\", which enables us to classify the particle behavior. In section 3, performing numerical integrations, we show that chaos occurs for a highly spinning test particle. Summary and some remarks follow in section 4.  Throughout this paper we use units $c=G=1$.  We define the signature of the metric as $(-,+,+,+)$. ",
        "conclusions": "\\label{sec5}\\setcounter{equation}{0} In this paper, using the pole-dipole approximation,  we study the motion of a spinning test particle near a Schwarzschild black hole to clarify its dynamical properties such as a chaos. We find the motion of the particle can be chaotic under some appropriate conditions. Because the motion of a spinless particle in this spacetime is never  chaotic because of its integrability, this chaotic behavior is purely induced by the spin-orbit interaction. The ``effective potential\" of the particle is also introduced to classify the dynamical behaviors. The ``effective potential\"s are classified into  four different types depending on the total angular momentum  $J$ and the spin $S$. When $J$ is large, some orbits are bounded and the ``effective potential\"s are classified into two types: for type (B1) one saddle point (unstable circular orbit) and one minimal point (stable circular orbit) on the equatorial plane exist for small spin;  and for (B2) two saddle points bifurcate from the equatorial plane and one minimal point remains on the equatorial plane for large spin. If $J$ is small, no bound orbits exist and the potentials are classified into another two types: for type (U1) no extremal point is found for small spin; and   for type (U2)  one saddle point appears on the equatorial plane, which is unstable in the direction perpendicular to the equatorial plane, for large spin. The types (B1) and (U1) are the same as those for a spinless particle, but the types (B2) and (U2) are new potentials which appear through a spin-orbit coupling.  The chaotic behavior is found only in the type (B2) potential. We believe that the appearance of saddle points is important, because we find chaos only for the orbits which approach the saddle points.  The critical value of the spin beyond which chaos will occur is $S \\sim 0.64 \\mu M$ for $J= 4 \\mu M$. We also present the Lyapunov exponent, which increases rapidly after the spin $S$ gets larger than the above critical value.  This supports the use of the ``effective potential'' as a criterion for chaos. The typical value of the Lyapunov exponent is about several orbital periods of the particle, which may become important in some relativistic astrophysical phenomena.    In a real astrophysical system such as  a binary system, the symmetry of the system is  lower than the present case. There may be  other important effects in addition to the  spin-orbit interaction, which make the motion more complicated. Then we may expect that chaos occurs even in the real system, or that the other effects stabilize the system and chaos will never be found. We need further analysis taking into account the other effects such as the spin-spin interaction or a force due to  multipole moments, even if we adopt a test particle analysis\\cite{comment2}. As for the spin-spin interaction, we can analyze a spinning test particle in  a rotating Kerr black hole. In our preliminary analysis by the ``effective potential\" near the equatorial plane, the critical value of the spin for chaos gets smaller as the angular momentum of the black hole becomes larger.  The detailed analysis is under investigation.  Another important point is whether or not such a chaotic behavior, if it exists, affects any realistic astrophysical or physical phenomena. One of the important targets for such an investigation is a coalescing binary system, where we need general relativity, in particular relativistic dynamics of compact objects as we mentioned in section 1. To examine it, since we have studied here only the condition for occurrence of chaos, then we next have to know the evolution of the system including emission of gravitational waves.  The particle traveling around a black hole emits the gravitational waves, extracting  the energy and the angular momentum from the system. This will tell us whether or not  the evolutionary path will get into the region of parameter space where chaos will take place. We then have to calculate the emission rates of the energy and the angular momentum, $\\dot{E}$ and $\\dot{J}$ for the present system\\cite{mino},\\cite{mino2} and follow the evolution.  \\vskip 1cm We would like to thank M. Sasaki, M. Shibata, Y. Sota, H. Tagoshi and T. Tanaka for useful discussions and thank P. Haines for reading the paper carefully. SS also acknowledges K. Imafuku a stimulating discussion. This work was supported partially by the Grant-in-Aid for Scientific Research  Fund of the Ministry of Education, Science and Culture  (No. 06302021 and No. 06640412),   and by the Waseda University Grant for Special Research Projects.  \\newpage \\renewcommand{\\theequation}{\\mbox{$A .\\arabic{equation}$}} \\setcounter{equation}{0} \\begin{flushleft} {\\bf Appendix A: The Basic Equations by Use of Spin Vector} \\end{flushleft} \\vspace{.5 cm}  In the text, we have used  a spin tensor $S^{\\mu\\nu}$. However, it may be sometimes more convenient or more intuitive to describe the basic equations by use of a spin vector $S_\\mu$, which is defined by \\begin{equation} S_\\mu = - \\frac{1}{2}\\epsilon_{\\mu\\nu\\rho\\sigma} u^{\\nu}S^{\\rho\\sigma}, \\label{def:spin} \\end{equation} which gives the following constraint: \\begin{equation} p^\\mu S_\\mu = 0. \\end{equation}   The equations of motion (\\ref{eqn:xdot}), (\\ref{eqn:pdot}), and (\\ref{eqn:sdot}) are now \\begin{eqnarray} \\frac{dx^{\\mu}}{d\\tau}&=&v^{\\mu}, \\label{eqn:xdot2} \\\\ \\frac{Du^{\\mu}}{D\\tau}&=& -\\frac{1}{\\mu} R^{*\\mu}_{~~\\nu\\rho\\sigma}v^{\\nu}u^{\\rho}S^{\\sigma}, \\label{eqn:pdot2} \\\\ \\frac{DS^{\\mu}}{D\\tau}&=&-\\frac{1}{\\mu} u^{\\mu} (R^{*}_{~\\alpha\\beta\\gamma\\delta}S^\\alpha v^{\\beta}u^{\\gamma}S^{\\delta}), \\label{eqn:sdot2}, \\end{eqnarray} where \\begin{equation} R^{*\\alpha\\beta}_{~~~\\gamma\\delta} \\equiv \\frac{1}{2} R^{\\alpha\\beta\\rho\\sigma} \\epsilon_{\\rho\\sigma\\gamma\\delta}. \\end{equation} The relation between the 4-velocity and the momentum is \\begin{equation} v^{\\mu} = N \\left(u^{\\mu} - \\frac{1}{\\mu^2} ~^{*}R^{*\\mu\\alpha}_{~~~\\beta\\gamma} S_\\alpha u^\\beta S^\\gamma \\right), \\label{eqn:v-u2} \\end{equation} where \\begin{equation} ~^{*}R^{*}_{~\\mu\\nu\\rho\\sigma} \\equiv \\frac{1}{2} \\epsilon_{\\mu\\nu\\alpha\\beta} R^{*\\alpha\\beta}_{~~~\\rho\\sigma} = \\frac{1}{4} \\epsilon_{\\mu\\nu\\alpha\\beta} R^{\\alpha\\beta\\gamma\\delta} \\epsilon_{\\gamma\\delta\\rho\\sigma}. \\end{equation} and $N$ is the normalization constant determined from $v_\\mu v^\\mu =-1$. The supplementary condition to fix the center of mass is \\begin{equation} v_\\mu S^\\mu =0. \\end{equation} The spin vector is perpendicular to the 4-velocity as we expected. Note that $v_\\mu S^{\\mu\\nu} \\neq 0$. Equation (\\ref{eqn:spin}) is just \\begin{equation} S^2 = S^\\mu S_\\mu . \\label{eqn:spin2} \\end{equation}  In what follow, we assume a Schwarzschild black hole as the background spacetime. We write down explicitly  the relation (\\ref{def:spin}) between  $S^{\\mu\\nu}$ and  $S_{\\mu}$ as \\begin{eqnarray} S_t &= &-r^2 \\sin \\theta \\left[u^r S^{\\theta\\phi} +u^\\theta S^{\\phi r} + u^\\phi S^{r \\theta}\\right] \\label{eqn:sst}\\\\ S_r &= &r^2 \\sin \\theta \\left[u^t S^{\\theta\\phi} -u^\\theta S^{t \\phi} + u^\\phi S^{t \\theta}\\right] \\label{eqn:ssr}\\\\ S_\\theta &= & r^2 \\sin \\theta  \\left[u^t S^{\\phi r} -u^\\phi S^{t r} + u^r S^{t \\phi}\\right] \\label{eqn:ssth}\\\\ S_\\phi &= & r^2 \\sin \\theta  \\left[u^t S^{r \\theta} -u^r S^{t \\theta} + u^\\theta S^{t r}\\right] \\label{eqn:ssph} \\end{eqnarray} Conversely, Eqs. (\\ref{eqn:sst}) $\\sim$ (\\ref{eqn:ssph}) with Eqs.(\\ref{eqn:ps0}), (\\ref{eqn:sst})  give \\begin{eqnarray} S^{\\theta\\phi}&= & -\\frac{u_t}{r^2 \\sin\\theta}\\tilde{S}_r \\label{eqn:ssr3}\\\\ S^{\\phi r} &= & -\\frac{u_t}{r^2 \\sin\\theta}\\tilde{S}_\\theta \\label{eqn:ssth3}\\\\ S^{r \\theta} &= & -\\frac{u_t}{r^2 \\sin\\theta}\\tilde{S}_\\phi , \\label{eqn:ssph3} \\end{eqnarray} where \\begin{equation} \\tilde{S}_\\mu \\equiv S_\\mu -\\frac{S_t}{u_t}u_\\mu . \\end{equation}  Using Eqs.(\\ref{eqn:jx2})$\\sim$(\\ref{eqn:stph}), Eqs.(\\ref{eqn:sst})$\\sim$(\\ref{eqn:ssph}) are now \\begin{eqnarray} \\tilde{S}_r &= &\\frac{J}{u_t}\\left[-\\left(1+fu_r^2\\right)\\cos\\theta \\right] \\label{eqn:ssr2}\\\\ \\tilde{S}_\\theta &= &\\frac{r}{u_t}\\left[J\\left( \\sin\\theta - fu_r \\frac{u_\\theta}{r}\\cos\\theta\\right) -\\frac{p_\\phi}{\\sin\\theta}\\right] \\label{eqn:ssth2}\\\\ \\tilde{S}_\\phi &= &\\frac{r}{u_t}\\left(p_\\theta\\sin\\theta -J f u_r\\frac{u_\\phi}{r} \\cos\\theta\\right) \\label{eqn:ssph2}\\\\ S_t &= &J\\left(-f u_r \\cos\\theta +\\frac{u_\\theta}{r}\\sin\\theta\\right) \\label{eqn:sst2} \\end{eqnarray} This expression already includes the constants of motion for the angular momentum.  The initial direction of the spin vector at $r=r_0,\\;\\theta =\\pi/2,\\;\\phi=0$ is given by the equations \\begin{eqnarray} \\tilde{S}_r&= & 0 \\label{eqn:sr3} \\\\ \\tilde{S}_\\theta &= &\\frac{r}{u_t}\\left( J -p_\\phi\\right) \\label{eqn:sth3} \\\\ \\tilde{S}_\\phi &= &\\frac{r}{u_t}p_\\theta . \\label{eqn:sph3} \\end{eqnarray}  As for the angle of the spin $\\alpha$, we find that \\begin{equation} \\alpha = \\tan^{-1} \\frac{ \\tilde{S}_\\phi}{\\tilde{S}_\\theta} . \\end{equation} Using this definition with (\\ref{eqn:spin2}), we have \\begin{eqnarray} \\tilde{S}_\\theta &= &\\frac{\\mu Sr_0^2}{\\sqrt{\\mu^2 r_0^2 + J^2 \\sin^2\\alpha}}\\cos\\alpha \\\\ \\tilde{S}_\\phi &= &\\frac{\\mu Sr_0^2}{\\sqrt{\\mu^2 r_0^2 + J^2 \\sin^2\\alpha}}\\sin\\alpha . \\end{eqnarray} $u_\\theta, u_\\phi$ are then some functions of $u_t$ through Eqs. (\\ref{eqn:sth3}) and (\\ref{eqn:sph3}).  Inserting them into the energy conservation equation (\\ref{eqn:energy2}), we obtain the quadratic equation for  $u_t$.  We also have Eq. (\\ref{eqn:mass}) for $u_r$.  Solving them, we can finally set up the initial data of $u^\\mu$ and $S^\\mu$.    Setting $p_r=p_\\theta=0$ and using Eqs. (\\ref{eqn:mass}) and (\\ref{eqn:spin2}), we find the ``effective potential\" (\\ref{eqn:effpot}). After giving $J,S$ and $E$ and setting initial data, we can solve the dynamical equations  (\\ref{eqn:xdot2}) and  (\\ref{eqn:pdot2}) with algebraic equations (\\ref{eqn:v-u2}), (\\ref{eqn:ssr2}) $\\sim$ (\\ref{eqn:sst2}), although here we have solved for all variables using the dynamical equations in order to estimate the accuracy by the constraint equations.   .\\newpage \\baselineskip .15in"
    },
    "9604/hep-th9604028_arXiv.txt": {
        "abstract": "Photon splitting in a very strong magnetic field is analyzed for energy $\\omega < 2m$. The amplitude obtained on the base of operator-diagram technique is used. It is shown that in a magnetic field much higher than critical one the splitting amplitude is independent on the field. Our calculation is in a good agreement with previous results of Adler and in a strong contradiction with recent paper of Mentzel et al. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604135_arXiv.txt": {
        "abstract": "Using the OH-airglow suppressor spectrograph at the University of Hawaii 2.2m telescope and the CGS4 spectrometer at the United Kingdom Infrared Telescope, we have found exceptionally large Balmer decrements in two unusual high-$z$ QSOs, Hawaii 167 ($z=2.36$, $\\ha / \\hb = 13$) and Q0059-2735 ($z=1.59$, $\\ha / \\hb = 7.6$), the latter being a so-called low-ionization broad absorption line QSO (BALQSO). We argue that these objects are young QSOs heavily enshrouded by dust. In fact, the internal reddening might be so large as to completely extinguish the QSO light in the restframe UV, allowing us to see the underlying stellar population.  Our possible detection of the 4000 \\AA\\ break in Hawaii 167 supports this idea.  Its small amplitude indicates a very young age for the population, $\\sim$ 15 Myrs.  To explain the properties of these QSOs, we propose a model in which a young QSO is surrounded by a shell of young massive stars mixed with significant amounts of dust.  We predict that as the QSO emerges from this dust cocoon, it will eventually take on the appearance of a normal BALQSO. ",
        "introduction": "In the Hawaii deep $K$-band survey (\\cite{Cowie94a}; \\cite{Songaila94}), one extraordinary object was found at a redshift% \\footnote{The redshift was originally reported as 2.35, but our reanalysis of the \\ha\\ line indicates that it is closer to 2.36.}of 2.36, the highest by far in the sample.  This object, named Hawaii 167, has strange characteristics: its high redshift, compact morphology, and broad ($\\sim$ 5000 km s$^{-1}$) Balmer emission lines are sure signs of this object's being a QSO while its restframe UV spectrum resembles those of starburst galaxies, showing strong metal absorption lines with no obvious emission features (\\cite{Cowie94b}).  Cowie et al.\\ (1994b) suggested that Hawaii 167 was either a very exotic broad absorption line QSO (BALQSO), a starburst galaxy, or most probably a mixture of both.  If Hawaii 167 is in any way related to starbursts, it may be of considerable importance: the small area coverage of the survey ($\\sim$ 77 arcmin$^{-2}$) suggests that this type of object might in fact be quite common in faint near-IR samples.  The key to understanding Hawaii 167 may lie in the QSO called Q0059-2735 (\\cite{Hazard87}), the only previously known object with similar characteristics.  This object is the most extreme member of the low-ionization BALQSOs (also often called \\mgii\\ BALQSOs) comprising $\\sim$ 15\\% of BALQSOs, which themselves represent $\\sim$ 10\\% of all QSOs (\\cite{Weymann91}).  This class of QSOs shows broad absorption lines (BALs) of low-ionization ions (\\mgii , \\aliii ) as well as those of the high-ionization species (\\civ , \\siiv ) regularly seen in BALQSOs.  What separates low-ionization BALQSOs from the normal BALQSOs seems to be their richness in dust: they are known to be more common in IRAS-selected samples (\\cite{Low89}), and their UV continuum shows signs of moderate reddening (\\cite{Sprayberry92}). Also, the absence or the extreme weakness of \\oiii\\ in these objects can be understood if dust is preventing ionizing radiation from reaching the outer low-density regions where the formation of the forbidden line is possible (\\cite{Boroson92}).  Voit et al.\\ (1993) suggested that low-ionization BALQSOs are probably ``young quasars in the act of casting off their cocoons of gas and dust,'' which implies some kind of connection between these objects and the ultraluminous IRAS galaxies (\\cite{Sanders88}).  Indeed, one of the ultraluminous IRAS galaxies, IRAS 07598+6508, was found to be a low-ionization BALQSO (\\cite{Lipari94}).  As already mentioned, one interpretation for Hawaii 167 put forth by Cowie et al. (1994b) was that Hawaii 167 is a mixture of both a QSO and a starburst galaxy: its internal reddening is so large that its QSO light, dominating the restframe optical, is completely extinguished in the restframe UV, leaving only the light of the surrounding starbursting galaxy.  This idea is supported by the lack of any UV broad emission lines (BELs), and the large Balmer decrement ($> 8$) inferred from the non-detection of \\hb.  In this paper, we will investigate this possibility by examining the near-IR spectra of Hawaii 167 and Q0059-2735, using the Balmer decrement as an indicator of internal reddening. ",
        "conclusions": "Although any conclusion drawn from such a small data set is necessarily tentative, at this point we think that the star-light hypothesis is the most likely explanation for the spectra of these QSOs.  It is very difficult to imagine that the intrinsic Balmer decrement could be close to, or even larger than 10 as in the case for the reddened-QSO-light hypothesis; nor is it easy to imagine that any dust grains could have UV albedos close to unity as in the case for the scattered-QSO-light hypothesis.  Furthermore, neither of these could explain the small observed equivalent widths of the \\mgii\\ and \\hb\\ BELs.  Therefore, we conclude that these objects are in fact heavily dust-enshrouded young QSOs whose internal reddening is so large as to completely extinguish the QSO light in the restframe UV, allowing us to see the underlying stellar population.  A possible counter argument to this explanation, nevertheless, is to attribute the large Balmer decrements and the small equivalent widths to the scatter in the intrinsic QSO properties.  This argument is especially strong for Q0059-2735, whose Balmer decrement and \\hb\\ equivalent width are deviant but might still be within the range of their natural distribution, which itself has a large scatter.  Our response to such an argument is that low-ionization BALQSOs are very likely to be dusty from other lines of evidence (see Introduction), and therefore it is very natural to explain the peculiar properties of these objects as the effect of dust.  Probably most of the low-ionization BALQSOs discovered so far are not as heavily reddened as the two objects studied here because the determination of the low internal reddening by Sprayberry \\& Foltz (1992) stems not only from the continuum shape but also from the emission line strength.  Our claim, however, is that there are objects which are extreme versions of these optically selected low-ionization BALQSOs, and they are probably as dusty as the ultraluminous IRAS galaxies.  The identification of one of the ultraluminous IRAS galaxies, PC07598+6508, as a low-ionization BALQSO (\\cite{Lipari94}) strongly supports this idea.  In this sense, these two classes of objects may well be the same things looked at from different angles or at different times.  A model of a young QSO surrounded by a shell of starbursting population mixed with dust seems to explain (at least qualitatively) the main characteristics of these QSOs, and possibly suggest the evolutionary connection between the low-ionization BALQSOs and normal BALQSOs.  This starburst in the shell surrounding a QSO could be identified as spheroid formation (\\cite{Cowie94b}; \\cite{Kormendy92}) though this is a speculative guess at this point.  This model of course has its own problems, and we mention a few here: First, it is unclear whether we could really explain the BALQSO phenomena with supernova explosions in this surrounding shell.  Second, a great deal of fine-tuning seems to be required if we are to produce an almost perfectly power-law SED out of two completely separate light components, a QSO and its underlying stellar population.  Finally, it is uncertain whether the use of the Bruzual-Charlot model is justified for these objects.  The starburst might be limited to a class of massive stars within a small mass range; also, the light from supernovae might be dominating the SEDs.  We also mention another rather technical problem: that is, the assessment of effects due to \\feii\\ emission features.  It is known that this type of QSO shows strong \\feii\\ emission, and that these features are especially conspicuous around \\hb.  Therefore, it is possible that the small equivalent width of \\hb\\ is simply due to a large contribution of \\feii\\ emission to the continuum.  Also, our measurement of the 4000 \\AA\\ break amplitude might be affected by the \\feii\\ emission features.  Although these effects are potentially serious, we are not able to assess the effects of \\feii\\ emission here because our limited spectral coverage and resolution prevent us from performing detailed modelling of \\feii\\ emission features.  Despite the various uncertainties, these objects present two very interesting possibilities: that is, 1) heavy internal reddening of QSOs could result in a separation of their stellar light and AGN light in spectral space, and 2) if such heavily reddened QSOs exist at high redshifts, they could easily have escaped our detection until now.  The first point leads us to the possibility that we may be able to study the underlying stellar populations of QSOs by using this type of object.  In this paper, we have presented an example of such analyses, which suggests that the stellar population in Hawaii 167 looks very young, probably around 15 Myrs old if we take the fit of the Bruzual-Charlot model at its face value.  The great importance of such a spectral analysis is that we do not have to spatially resolve the QSO: at high redshifts where such a spatial separation becomes difficult, spectroscopic studies of this class of QSOs may be the only way to understand what high-$z$ QSOs really are.  Regarding the second point, if this type of object exists in significant numbers, it is very likely that many of them would not have been picked up in the previous optical surveys due to their faint UV continuum and absence of strong UV emission lines.  In fact, Hawaii 167-like objects could be completely dark in the observer's optical band if starbursts have not started in the host galaxy.  This implies that there might be a whole population of dust-enshrouded young QSOs at high redshifts which has so far eluded our detection.  We were able to identify Hawaii 167 simply because it has barely enough UV continuum for our optical spectroscopy with a 4-m class telescope, and this might explain the flatness of Hawaii 167's SED as a selection effect.  In any case, if these objects are in fact young QSOs emerging from their dust cocoons, they are probably much more abundant at $z > 2$, where the comoving space density of QSOs is rapidly increasing. The redshift of Hawaii 167 is indeed just above 2.  Wide-field sensitive near-IR surveys should eventually tell us if such a population of objects exists.  We thank S.\\ Charlot for providing the isochrone synthesis model, and E.\\ M.\\ Hu for helpful comments on the manuscript.  This work was partly supported by the Grant-in-Aid of the Ministry of Education, Japan (07044080)."
    },
    "9604/astro-ph9604028_arXiv.txt": {
        "abstract": "We present a spectroscopic study of the $138$ field galaxies to a redshift $z=0.3$ from the $I$-selected Canada-France Redshift Survey. $117$ (85\\%) spectra exhibit at least $H\\alpha$ in emission, the remaining $21$ (15\\%) are purely absorption-line spectra.  We focus our analysis on spectra with H$\\alpha$ and H$\\beta$ in emission, accounting for about half of this low-$z$ sample, which we classify using emission-line ratio diagrams. Using photoionization models, we determine the extreme boundaries of H II galaxies in these diagnostic diagrams, and demonstrate that the emission-line ratios of a significant fraction of galaxies require harder photoionization sources than massive O stars.  We find that about $17 \\%$ of the field galaxies have emission-line ratios consistent with active galaxies, {\\it e.g.}  Seyfert 2 or LINERs.  After correcting for stellar absorption under the Balmer lines, we conclude that the fraction of such galaxies is at least $8\\%$ of the field galaxy population at $z\\leq 0.3$. ",
        "introduction": "Studying emission-line spectra leads to a better under\\-standing of the nature of the galaxies in deep surveys; line ratio diagrams can separate narrow emission-line galaxies such as H II galaxies and active galaxies.  In this paper, we adopt the terminology of ``H II galaxies'' for galaxies with line ratios that can be explained by OB stars as the main io\\-ni\\-za\\-tion source of the nebula gas, as for example starburst ga\\-la\\-xies (SBG), blue compact galaxies (BCG), H II region-like galaxies.  Galaxies with line ratios requiring harder io\\-ni\\-za\\-tion sources than hot main sequence stars, are gathered together under the name ``active galaxies'', as Seyfert 2 and Low Ionization Nuclear Emitting Regions (LINER) gala\\-xies.  By observing several emission lines and considering the appropriate line ratios, a narrow emission-line object can be classified re\\-lia\\-bly (see {\\it e.g.}  Veilleux \\& Osterbrock 1987, hereafter VO). \\footnotetext[2]{Visiting observer with the Canada-France-Hawaii Telescope, operated by the NRC of Canada, the CNRS of France and the University of Hawaii.}  The wide spectral range of the $I$-selected Canada-France Redshift Survey ($4500-8500\\ \\AA$) offers significant advantages for a systematic investigation of the emission-line ratios of the field galaxies at low redshifts. The survey is described in a series of earlier papers (Lilly {\\it et al.} 1995a (CFRS-I), Le F\\`evre {\\it et al.} 1995a (CFRS-II), Lilly {\\it et al.} 1995b (CFRS-III), Hammer {\\it et al.} 1995a (CFRS-IV) and Crampton {\\it et al.} 1995a (CFRS-V)).  An analysis of a few low-$z$ emission-line galaxies already showed the presence of very strong emission-line objects from a survey preliminary to the CFRS (Tresse {\\it et al.} 1993).  At $z <0.3$, forbidden emission lines such as [O II]$\\lambda 3727$, [O III]$\\lambda 4959$, [O III]$\\lambda 5007$, [S II]$\\lambda\\lambda 6717,6731$ (hereafter [S II]$\\lambda6725$) can be measured, and their ratio to the appropriate Balmer line strength allows a classification according to the main ionization source responsible for the emission lines.  Our study can shed new light on the nature of the po\\-pulation of faint blue galaxies at low redshifts. Indeed, an excess of these galaxies relative to local counts is seen in the CFRS luminosity function (Lilly {\\it et al.} 1995c (see \\mbox{CFRS-VI})) for absolute magnitudes $M(B)> -20$ mag$^{1}$.\\footnotetext{$^1$Absolute magnitudes are expressed assuming $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ throughout the paper.} A similar excess has been found previously from $B$-selected deep surveys (Broadhurst {\\it et al.} 1988, $b_{J}=21.5$; Colless {\\it et al.} 1990, $b_{J}=22.5$).  These surveys do not cover all the optical wavelength range, hence emission-line studies have been res\\-tric\\-ted mainly to the single [O II]$\\lambda3727$ line. Our systematic study of line ratios is a major step forward from previous analyses of these galaxies, yielding a better understanding of the evolution of this population with cosmic epoch.  The plan of this paper is as follows. In Section 2, we review some key points concerning the CFRS observations and selection criteria which are important for our analysis, and we describe the classification of all the CFRS spectra up to $z=0.3$.  In Section 3 we focus on the emission-line galaxies, describing our line measurements. In Section 4, we use two line ratio diagrams to classify these galaxies. In Section 5, we construct photoionization models to determine the upper limit of the H II galaxy locus in these diagnostic diagrams, and we compare our data to this limit. In Section~6, we assess the effects of the stellar absorption underlying the Balmer lines. Section 7 presents a statistical analysis of our sample.  In Section 8, we investigate the photometric characteristics of our low-$z$ sample. Section 9 discusses our results. ",
        "conclusions": ""
    },
    "9604/astro-ph9604114_arXiv.txt": {
        "abstract": "We have undertaken a pilot project to measure the rotation velocities of spiral galaxies in the redshift range 0.18 $\\le z \\le$ 0.4 using high dispersion long slit spectroscopy obtained with the Palomar 5m telescope. One field galaxy and three cluster objects known to have strong emission lines were observed over wavelength ranges covering the redshifted lines of [OII], CaII K, H$\\beta$, and [OIII]. Two of the objects show extended line emission that allows the tracing of the rotation curve in one or more lines.  A line width similar to that obtained with single dish telescopes for the 21--cm HI line observed in lower redshift galaxies can be derived from the observed H$\\beta$, [OII], and [OIII] emission by measuring a characteristic width from the velocity histogram.  These moderately distant galaxies have much stronger emission lines than typical low--redshift spirals but they appear to be kinematically similar.  Application of the Tully-Fisher relation suggests that the two galaxies with rotation curves are intrinsically brighter at R-band than nearby galaxies. ",
        "introduction": "A major goal of cosmology is to determine the temporal history and fate of the universe.  Galaxies provide the primary means to accomplish this; thus it is critically important to understand their formation and evolution.  Between the present time and the epoch corresponding to a redshift $z \\simeq$ 0.4, significant evolution of the cluster population results in a higher fraction of blue galaxies at the earlier epoch. Most recent high resolution imaging of the blue cluster members confirms their spiral nature (Lavery \\etal 1992; Dressler \\etal 1993). A number of authors conclude that galaxy--galaxy interactions play a primary role in producing an enhancement in the blue galaxy fraction in distant clusters (Thompson 1988; Lavery and Henry 1988; Lavery \\etal 1992) while others suggest that the excess activity in blue galaxies is associated with the interaction between hot intracluster gas and infalling galaxies (Dressler \\etal 1985; Bothun and Dressler 1986). The detailed study of the kinematics of the blue cluster galaxies and their counterparts in the field at similar redshifts will place constraints on both the formation and evolutionary processes.  In recent years, the HI Tully--Fisher (TF) relation and its optical analog have been used by a growing community in attempts to map out the local flow field.  These peculiar motions hinder determination of the Hubble constant from nearby galaxies.  However, at distances larger than $z \\simeq$ 0.04, peculiar motions should be negligible with respect to the overall Hubble flow and the main source of error is the intrinsic scatter in the TF relation (about 0.3 mag) and the local calibration.  Thus, if the TF method can be properly understood and applied, it may permit the determination of redshift--independent distances and evaluation of the Hubble constant, at intermediate redshift.  The prospects for high redshift application of the TF method using the H$\\alpha$ line have been discussed by van der Kruit and Pickles (1988).  Detection of the H$\\alpha$ line as the distance to the emitting object increases is hampered not only by the radial fall-off in emissivity but also because the redshifting of the line from its rest wavelength places the H$\\alpha$ line in the night--sky--contaminated red portion of the spectrum until it finally moves out of the optical window at $z \\ge$ 0.4.  An alternative approach, discussed in the context of star formation indicators by Kennicutt (1992a,b), relies on lines whose rest wavelengths are significantly lower than that of H$\\alpha$, specifically, the emission lines of [OII], H$\\beta$, and [OIII]. All three lines are often (though {\\it not always}) seen in the spectra of nearby spiral galaxies, the same objects in which H$\\alpha$ is detected.  In this {\\it Letter}, we report the results of a pilot project to conduct high resolution, high--sensitivity long--slit spectroscopic observations of galaxies in the redshift range 0.18 $\\le z \\le$ 0.4.  In \\S \\ 2 we discuss our observational strategy.  A summary of the results, including the extended rotation curves of two galaxies with $z \\simeq$ 0.2, is presented in \\S \\ 3. Finally, we discuss the degree of success of our experiment and the potential for using this technique to derive an estimate of the Hubble constant out to such distances. ",
        "conclusions": "The results of this pilot effort demonstrate the potential of investigating the kinematics of galaxies at intermediate redshift via long--slit spectroscopy.  Given very good seeing ($<$ 1\\arcsec) and adequate weather conditions, this project confirms the efficacy of the measurement of rotation curves in distant galaxies.  The rotation curves shown in figures 1 and 3 are well-defined and appear to flatten, characteristics seen in nearby normal spiral galaxies.  At the same time the presence of strong emission in the objects makes them quite unusual compared to low redshift galaxies.  Although both galaxies showing extended emission appear to have nearby companions, there is no evidence of interaction in their rotation curves.  In order to compare the extent of the rotation curve derived from the various lines, we obtained similar high resolution long slit spectra for a small sample of relatively nearby, normal galaxies ($z \\sim$ 0.02) over wavelength ranges covering all of the relevant lines including H$\\alpha$. Although the current sample is still small, primarily because of weather limitations, we note that the extent of the rotation curve traced in all lines is similar within the constraint that the signal--to--noise ratio is much higher for the H$\\alpha$ emission.  The velocity widths measured from different lines are consistent within the errors.  In the nearby objects, the oxygen lines are far weaker than the H$\\alpha$ line, but unlike our high $z$ candidates these galaxies had not been pre--selected for strong oxygen emission.  Measurements of the velocity widths allow estimation of galaxy properties and/or distances via application of the TF relation. Since most TF studies of nearby galaxies have been calibrated using radio 21--cm line widths, the optical widths must be transformed to a system appropriate to the radio TF relation.  This is particularly critical for distant studies where the velocity width rather than a detailed rotation curve may be all that is available.  A relative calibration between optical and HI line widths has been derived from a sample of low redshift galaxies, most of them in nearby Abell clusters.  Using 145 objects with both high quality HI line widths and optical velocity widths derived in the same manner employed here, we have derived a transformation relation between the optical width W and the 21--cm line width as used by Pierce and Tully (1988).  Using this transformation, we derive an equivalent corrected (for turbulent broadening) 21--cm line width $W_c^{21}$ of 185 $\\pm$ 15 \\kms for Galaxy 2545.3 and 264 $\\pm$ 11 \\kms for BOW 85. A more complete discussion of the comparison of radio and optical width measurements will be presented elsewhere (Vogt \\etal 1993). The corrected 21--cm width derived for 2545.3 is rather small, particularly relative to the mean expected for a large sample of low redshift objects (Roberts and Haynes 1993). It is possible, given the uncertainties, that this object may be less inclined than the current estimate, or that the rotation curve continues to rise significantly, beyond the radius at which line emission is detected.  Using the TF relation, the intrinsic brightness of distant galaxies can be compared, albeit crudely, to those nearby. Using the R-band calibration of the TF relation by Pierce and Tully (1988), a Hubble constant of 75 km s$^{-1}$ Mpc$^{-1}$, and the Scd galaxy energy distribution of Pence (1976) for the K-correction, we estimate that BOW 85 and 2545.3 are 0.9$\\pm$0.4 and 3.1$\\pm$0.7 magnitudes brighter respectively than predicted from the nearby galaxy sample.  (Another 0.3 magnitude error should be added in quadrature with the quoted errors to account for the intrinsic scatter in the TF relation).   Given the various uncertainties and complications in comparing the distant galaxies to the nearby ones, this increase in the brightness of distant galaxies relative to nearby ones of the same velocity width should be considered preliminary. Similar results have been reported by Franx (1993) who applied the luminosity--velocity dispersion relation to E+A galaxies in the cluster A665 at z=0.18.  Because of our galaxy selection process and the potential for galaxy evolution, use of the TF relation to derive a Hubble constant is not justified at this time.  In the future a better statistical sample will allow a comparison of the slope of the TF relation at high and low redshift and the evaluation of the importance of merging in the blue galaxy population.  High resolution images such as those obtained recently with HST by Dressler \\etal (1993) will give better estimates of the galaxy morphologies and inclination angles.  Multiband application of the TF relation will help to investigate galaxy evolution since long wavelengths should be progressively less affected by it. Alternatively, if the effects of galaxy evolution can be disentangled and the local calibration of the TF relation verified, the Hubble constant can be determined."
    },
    "9604/astro-ph9604100_arXiv.txt": {
        "abstract": "We present orbital-phase resolved I and K-band spectroscopy of Cygnus~X-3.  All spectra show emission lines characteristic of Wolf-Rayet stars of the WN subclass.  On time scales longer than about one day, the line strengths show large changes, both in flux and in equivalent width.  In addition, the line ratios change, corresponding to a variation in spectral subtype of WN6/7 to~WN4/5.  We confirm the finding that at times when the emission lines are weak, they shift in wavelength as a function of orbital phase, with maximum blueshift coinciding with infrared and \\Xray\\ minimum, and maximum redshift half an orbit later.  Furthermore, we confirm the prediction -- made on the basis of previous observations -- that at times when the emission lines are strong, no clear wavelength shifts are observed.  We describe a simplified, but detailed model for the system, in which the companion of the \\Xray\\ source is a Wolf-Rayet star whose wind is at times ionised by the \\Xray\\ source, except for the part in the star's shadow.  With this model, the observed spectral variations can be reproduced with only a small number of free parameters.  We discuss and verify the ramifications of this model, and find that, in general, the observed properties can be understood.  We conclude that \\CygX3 is a Wolf-Rayet/\\Xray\\ binary. ",
        "introduction": "}  Cygnus X-3 is a bright \\Xray\\ source that is peculiar among \\Xray\\ binaries.  It has huge radio outbursts, mildly relativistic ($\\beta\\simeq0.3$) jets, smooth 4.8\\,hour orbital modulation of its \\Xray\\ light curve, and a rapid increase of the orbital period on a time scale of 850\\,000 years.  It also has a very strong iron line in its \\Xray\\ spectrum, is very bright in the infrared, and has been claimed to be detected at very high energies (for reviews of its properties, see, e.g., Bonnet-Bidaud \\& Chardin \\cite{bonnc:88}; Van der Klis \\cite{vdkl:93}).  One of the first models for \\CygX3 was put forward by Van den Heuvel \\& De Loore (\\cite{vdhedl:73}).  They suggested that the system is composed of a compact object and a helium star of several solar masses, and that it represents a late evolutionary stage of massive \\Xray\\ binaries (the so-called `second Wolf-Rayet phase'; Van den Heuvel \\cite{vdhe:76}).  Massive helium stars have been identified observationally with the group of `classical', or population~I Wolf-Rayet stars (Van der Hucht et al.\\ \\cite{vdhu&a:81}).  Such stars have strong winds, which in \\CygX3 would be the underlying cause for the \\Xray\\ modulation (due to scattering of the X~rays), the increase of the orbital period (due to the loss of angular momentum) and the brightness in the infrared (due to free-free emission in the wind). In the course of further evolution, the helium star is likely to explode as a supernova.  If the system is not disrupted, a binary such as the Hulse-Taylor pulsar PSR\\,1913+16 could be formed (Flannery \\& Van den Heuvel \\cite{flanvdh:75}).  In this model, it is predicted that the optical/infrared counterpart shows Wolf-Rayet features in its spectrum.  This prediction was confirmed by Van Kerkwijk et al.\\ (\\cite{vker&a:92}, hereafter Paper~I), who found strong, broad emission lines of \\HeI\\ and \\HeII\\ -- but no evidence for hydrogen -- in I and K-band spectra of \\CygX3, as expected for a Wolf-Rayet star of spectral type WN7.  In subsequent observations, it was found (Van Kerkwijk \\cite{vker:93a}, hereafter Paper~II) that large changes in the absolute and relative strengths of the emission lines had occurred.  Furthermore, the spectra showed orbital-phase dependent wavelength shifts of the emission lines, with maximum blueshift occurring at the time of infrared and \\Xray\\ minimum, and maximum redshift half an orbit later.  It was shown that these wavelength shifts could be understood if the Wolf-Rayet wind were almost completely ionised by the \\Xray\\ source at the time of the observations, except in the part shadowed by the helium star.  It was found that both the wavelength shifts and the modulation of the infrared continuum could be reproduced with a detailed model (with only a small number of free parameters).  Based on the model, it was predicted that at times when strong emission lines were present in the infrared spectra, there would be little modulation of the lines and the continuum, whereas at times when the emission lines were weak, there would be a clear modulation of the lines and continuum.  For the latter case, it was expected that at high resolution the line profile would be resolved in two components.  Based on the idea that a stronger \\Xray\\ source would ionise more of the wind, it was also predicted that the \\Xray\\ source should be in its low state (low flux, hard spectrum) when the emission lines were strong, and in its high state (high flux, soft spectrum) when they were weak.  Kitamoto et al.\\ (\\cite{kita&a:94}), however, found that \\Xray\\ and radio data indicated that the source was in its high state on 1991 June 21, when the infrared lines were strong, while radio data indicated it was in its low state on 1992 May 29, when the lines were weak.  They suggested a modification of the model presented in Paper~II, in which the source's state at all wavelengths was a function of the mass-loss rate of the Wolf-Rayet star only.  In this case, a high mass-loss rate would lead to increased infrared and \\Xray\\ fluxes, as well as to radio outbursts.  At the same time, the wind would become optically thick to X~rays, leading to a lower degree of ionisation, and hence stronger infrared emission lines.  In this paper, we discuss the model for the infrared continuum and lines in detail, and compare it to the available observations.  In Sect.~\\ref{sec:obsred}, the procedures used for making and reducing the observations are described, both for the observations presented in Papers~I and~II, and for a number of additional observations.  We present the observations in Sect.~\\ref{sec:spectra}, and point out the characteristic similarities and differences shown by the spectra.  In Sect.~\\ref{sec:model}, we describe our model for the system, and use it to calculate light curves and line profiles as a function of orbital phase.  Furthermore, we qualitatively interpret the long-term changes, and verify the predictions for the lines made in Paper~II. In Sect.~\\ref{sec:massloss}, we estimate the velocity in the wind, and discuss different estimates of the mass-loss rate.  We discuss the ramifications expected for a more realistic treatment of the wind in Sect.~\\ref{sec:ramifications}.  In Sect.~\\ref{sec:fcont}, we estimate the infrared flux distribution of \\CygX3, and compare it with the one predicted for our model, and with the ones observed for other Wolf-Rayet stars.  We draw conclusions about the nature of \\CygX3 in Sect.~\\ref{sec:bigdisc}. ",
        "conclusions": "}  The first infrared spectroscopic observations of \\CygX3 (Paper~I) showed the presence of Wolf-Rayet emission features in the infrared spectrum of \\CygX3, which indicates that a strong, dense, helium-rich wind is present in the system, in which both the infrared lines and continuum originate.  In follow-up observations, the emission lines were much weaker, and they shifted in wavelength as a function of orbital phase (Paper~II).  It was argued that this could be understood in the context of the helium-star model, if the Wolf-Rayet wind is sometimes strongly ionised and heated by the X~rays originating from the compact object.  Furthermore, it was argued that if this were the case, the observed modulation of the infrared continuum would follow as a natural consequence.  In support of this argument, it was shown that observed variations could be reproduced with a simple model of the system, in which the wind has a two-temperature structure.  In this paper, we have presented a number of additional observations. These observations confirm the prediction made in Paper~II that the emission lines show wavelength shifts as a function of orbital phase if they are weak, but not if they are strong.  We have described the model used in Paper~II in detail, and presented model continuum light curves and model line profiles for various combinations of the parameters.  From these results, we have found that from the assumption of a two-temperature wind it follows naturally not only that the lines shift in wavelength, but also that the profiles at maximum redshift are stronger and narrower than at maximum blueshift.  A problem we have encountered, is that the estimate of mass-loss rate inferred from the infrared flux is an order of magnitude larger than the estimate inferred from the rate of increase of the orbital period. We have also found that the observed flux distribution is not consistent with the model.  However, since similar deviations are found for normal Wolf-Rayet stars, we believe that these discrepancies reflect the simplifications made for our model, and that they do not invalidate our main conclusions, viz., that the continuum and the lines arise in the wind of the helium star, and that the modulation of both lines and continuum is due to the wind being highly ionised by the \\Xray\\ source, except in the helium star's shadow.  An open question that remains is what the underlying cause is of the changes in state, in the infrared as well as in \\Xray\\ and radio. Watanabe et al.\\ (\\cite{wata&a:94}) have found that the radio and \\Xray\\ states are correlated: large radio outbursts occur when the \\Xray\\ source is in its high state (high flux, soft spectrum), but not when it is in its low state (low flux, hard spectrum).  Kitamoto et al.\\ (\\cite{kita&a:94}) determined that the radio/\\Xray\\ source was in its high state in June 1991, when the infrared spectra showed strong lines, and in its low state in May 1992, when the lines were weak. They suggested that the state at all wavelengths was a function of the mass-loss rate of the Wolf-Rayet star only.  When high, the infrared continuum would be stronger and the accretion rate would be high, leading to a strong \\Xray\\ source and radio outbursts (the latter are presumably due to the ejection of relativistic jets; e.g., Geldzahler et al.\\ \\cite{geld&a:83}).  The wind would be dense enough to be optically thick to low-energy X~rays, leading to a low state of ionisation.  When the mass-loss rate was lower, the infrared continuum flux and the accretion rate would be low, leading to \\Xray\\ low state and the absence of radio outbursts.  They suggested that the wind would become optically thin to low-energy X~rays, so that it could be ionised despite the lower \\Xray\\ luminosity.  At present, the hypothesis of a variable wind mass-loss rate as the underlying cause of the changes in state is consistent with the observations.  Other Wolf-Rayet stars of the WN subclass, however, do not show evidence for large changes in mass-loss rate, showing, e.g., constant flux to within a few percent on all observed time scales (Moffat \\& Robert \\cite{moffr:91}).  It might be that the changes in accretion rate and optical depth that are required are due to processes near the \\Xray\\ source and its accretion disk, perhaps similar to those underlying the changes in state observed in black-hole candidates like \\CygX1 (e.g., Oda \\cite{oda:77}).  One possibly major problem that has been addressed only briefly so far (in Paper~I) is that the radii of Wolf-Rayet stars derived from model fits to Wolf-Rayet spectra (e.g., Schmutz et al.\\ \\cite{schm&a:89}), are much larger than the radius of the Roche lobe in \\CygX3 (Paper~I; Conti \\cite{cont:92}).  Schmutz (\\cite{schm:93}) has tried to model the 1991 I and K-band spectra of \\CygX3 using the Wolf-Rayet wind model described by Hamann \\& Schmutz (\\cite{hamas:87}) and Wessolowski et al.\\ (\\cite{wess&a:88}).  He found that the stellar parameters were typical for a Wolf-Rayet star of spectral type WN7, but that given the absolute K magnitude, the photospheric radius was larger than the orbital separation by a factor~3.  In order to resolve this discrepancy, he suggested that the distance to \\CygX3 was much smaller than 10\\,kpc, so that the luminosity, and thus the mass and radius, were much smaller.  He suggested that the 21\\,cm absorption features, on which the 10\\,kpc distance is based, might be due to circumstellar shells.  Although the suggestion of a lower distance is tempting in that it would also resolve the discrepancy between the different estimates of the mass-loss rate, we believe that it is unlikely that there are circumstellar shells at such velocities and with such strengths that they produce a 21\\,cm absorption profile which corresponds closely with the combined emission profile of the Local, Perseus and Outer arms (e.g., Dickey \\cite{dick:83}).  Instead, we believe it more likely that the estimate of the radius derived from the model-atmosphere fits is wrong.  We stress that these estimates are based on an {\\em{}assumed} velocity law, which is used to extrapolate inward from the regions where the continuum and lines are formed, to the so-called `zero-velocity' radius.  Usually, for this velocity law, a $\\beta$-law is used ($v=v_\\infty(1-R_*/r)^\\beta$).  This velocity law has been found to be a reasonable approximation for O stars, but there is no reason to assume it is valid for Wolf-Rayet stars as well, especially close to the star (for more thorough discussions, we refer to Kudritzki \\& Hummer \\cite{kudrh:90}; Hillier \\cite{hill:95}; Schmutz \\cite{schm:95}).  An observational indication that the radii of Wolf-Rayet stars are actually quite small comes from studies of Wolf-Rayet stars in eclipsing binaries.  For instance, Cherepashchuk et al.\\ (\\cite{cher&a:84}) found that the eclipse light curve of V444~Cyg indicated a core radius (electron-scattering optical depth unity) of only 2.8\\,\\Rsun.  Hamann \\& Schwartz (\\cite{hamas:92}) argue that this estimate is not unique, and that the radius of the Wolf-Rayet star could be much larger.  From polarisation measurements, however, St-Louis et al.\\ (\\cite{stlo&a:93}) find that radii larger than $\\sim\\!4\\,\\Rsun$ are excluded.  A similar conclusion is reached by Cherepashcuk et al.\\ (\\cite{cher&a:95}) based on a spectroscopic estimate of the luminosity ratio.  Using this and other observational evidence, Moffat \\& Marchenko (\\cite{moffm:96}) conclude that the radii of Wolf-Rayet stars are not unlike the radii expected from stellar model calculations (e.g., Langer \\cite{lang:89}).  For such radii, the helium star would fit well inside the Roche lobe in \\CygX3 (Paper~I).  We conclude from our data that the idea that the companion of \\CygX3 is a massive Wolf-Rayet star is entirely plausible.  It seems to us that \\CygX3 is a very interesting system, in which a Wolf-Rayet wind can be probed through the combined effects of occultation and variable ionisation by X~rays."
    },
    "9604/astro-ph9604010_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{Intro4}  For many years broadband colors have been used to obtain a basic insight into the contents of galaxies.  Broadband photometry is relatively easy to obtain and gives an immediate impression of the spectral energy distribution (SED) of an object.  Broadband colors are particularly efficient when used for statistical investigations such as this one. Colors have been used to estimate the stellar populations of galaxies (e.g.\\ Searle et al.~\\cite{Sea73}; Tinsley~\\cite{Tin80}; Frogel~\\cite{Fro85}; Peletier~\\cite{PelPhD}; Silva \\& Elston~\\cite{SilEls94}) and it has been suggested that colors can give information about the dust content of galaxies (Evans~\\cite{Eva94}; Peletier et al.~\\cite{Pel94}, \\cite{Pel95}).  In this paper I use radial color profiles to investigate the stellar and dust content of galaxies.  The problem of determining the stellar content of galaxies from integrated SEDs has been approached from two sides, often called the empirical and the evolutionary approach (for a review, see O'Connell~\\cite{OCon87}).  In the first method, stellar SEDs are fitted to the observed galaxy SEDs (Pickles~\\cite{Pic85}; Peletier~\\cite{PelPhD}).  This method works only if one has spectral (line) information.  Generally, the broadband colors of a galaxy can be explained by a combination of the SEDs of two or three types of stars (Aaronson~\\cite{Aar78}; Bershady~\\cite{Ber93}).  In the second, more theoretical approach, stellar SEDs are combined, using some knowledge of initial conditions and evolutionary time scales of different stellar populations, to produce evolutionary stellar population synthesis models (for reviews Tinsley~\\cite{Tin80}; Renzini \\& Buzzoni~\\cite{RenBuz86}; more recent models are e.g.\\ Buzzoni~\\cite{Buz89}; Bruzual \\& Charlot~\\cite{BruCha96}; Worthey~\\cite{Wor94}).  The papers of Disney et al.~(\\cite{DDP89}) and Valentijn~(\\cite{Val90}) have renewed the debate on whether spiral galaxies are optically thick or thin. Broadband colors of galaxies can be used to examine this problem, because the dependence of dust extinction on wavelength causes reddening. This can be used to measure extinction at a certain point through the disk using a galaxy or another object behind it (Andredakis \\& van der Kruit~\\cite{AndKru92}) or to measure extinction within a galaxy, for instance across a spiral arm dust lane (Rix \\& Rieke~\\cite{RixRie93}; Block et al.~\\cite{Blo94}). To measure the global dust properties of a galaxy by reddening one can use the color profile. If one assumes that dust is more concentrated towards the center (just like the stars), the higher extinction in the center produces a color gradient that makes galaxies redder inwards (Evans~\\cite{Eva94}; Byun et al.~\\cite{Byun94}).  Most previous studies investigating stellar population and dust properties of galaxies used their integrated broadband colors.  The use of surface photometry colors is less common, as it is easier to compare integrated photometry than surface photometry for large samples of galaxies.  Integrated photometry samples the bulk properties of galaxies, but because the light distribution of galaxies is strongly concentrated, one effectively measures the colors of the inner regions of galaxies.  The half total light radius of an exponential disk is $\\sim$1.7 scalelengths, while luminosity profiles are easily traced out to 4-6 scalelengths.  Therefore, half of the light in integrated colors comes from an area that is less than $1/5$ of the area commonly observed in galaxies (say within $D_{25}$).  Our knowledge of the star formation history (SFH) and the dust content of galaxies improves when we start looking at local colors instead of integrated colors.  A first improvement is obtained by using the radial color distribution (i.e.\\ the color profile) of a galaxy.  This has been common practice for elliptical galaxies (e.g.\\ Peletier et al.~\\cite{Pel90a}; Goudfrooij et al.~\\cite{Gou94}), but not for spiral galaxies, because elliptical galaxies are assumed to have a simple SFH and low dust content (but see Goudfrooij~\\cite{GouPhD}) opposed to spirals.  Color studies of spiral galaxies have been concentrated on edge-on systems, in the hope to be better able to separate the dust and stellar population effects (e.g.~Just et al.~\\cite{Jus96}).  Even more detailed information about galaxies can be obtained by the use of azimuthal profiles (Schweizer~\\cite{Sch76}; Wevers et al.~\\cite{Wev86}) and color maps, but these techniques require high resolution, high signal-to-noise observations and are hard to parameterise to global scales, which means that they cannot be used in statistical studies.  Due to the large variety of galaxies, statistical studies of galaxies require large samples.  The introduction of CCDs into astronomy made it possible to obtain for large samples of galaxies accurate optical surface photometry in reasonable observing times (Kent~\\cite{Kent84}). Very large data sets of CCD surface photometry have recently become available (Cornell~\\cite{Cor87}; Han~\\cite{Han92}; Mathewson et al.~\\cite{Mat92}; Giovanelli et al.~\\cite{Gio94}).  Unfortunately, most of these samples are observed in only one or two passbands. Furthermore, the surface photometry is often reduced to integrated magnitudes and isophotal diameters to study extinction effects with an inclination test or to study the Tully-Fisher relation (Tully \\& Fisher \\cite{TulFis77}, hereafter TF-relation).  Fast plate measuring machines have also produced surface photometry of large sets of galaxies (e.g.~Lauberts \\& Valentijn~\\cite{ESO-LV}, hereafter ESO-LV), but again only in one or two passbands.  Since near-infrared (near-IR) arrays have become available only in the late eighties, near-IR surface photometry is available for only a few somewhat larger samples of spiral galaxies.  Terndrup et al.~(\\cite{Ter94}) observed 43 galaxies in $J$ and $K$, which was complemented with $r$ passband photometry of Kent~(\\cite{Kent84}, \\cite{Kent86}, \\cite{Kent87}).  They explained the observed colors mainly by population synthesis and invoked dust only for the reddest galaxies.  Peletier et al.~\\cite{Pel94} imaged 37 galaxies in the $K$-passband and combined their data with the photometry of the ESO-LV catalog. They explained their surface photometry predominantly in terms of dust distributions and concluded that spiral galaxies are optically thick in the center in the $B$ passband, under the assumption that there are no population gradients across the disk.  There are two sets of observations that allow a direct physical interpretation of color gradients in spiral galaxies: 1) The current star formation rate (SFR) as measured by the \\ha\\ flux has a larger scalelength than the underlying older stellar population (Ryder \\& Dopita~\\cite{RydDop94}). There are relatively more young stars in the outer regions of spiral galaxies than in the central regions. This will be reflected in broadband colors of spiral galaxies. 2) From metallicity measurements of \\hii\\ regions it is known that there are clear metallicity differences in the gas among different galaxies and that there are metallicity gradients as function of radius within galaxies (Villa-Costas \\& Edmunds~\\cite{VilEdm92}; Zaritsky et al.~\\cite{Zar94}). If the metallicity gradients in the gas are also (partly) present in the stellar components, the effects might be observable in the broadband colors. Stellar population synthesis models incorporating both age and metallicity effects are needed in the comparison with observations of radial color gradients.  Broadband photometry is often assumed to trace baryonic mass, and the transformation from light to mass is performed by postulating a mass-to-light ratio ($M/L_\\lambda$).  Both dust extinction and differences in stellar populations will influence $M/L_\\lambda$ ratios, most notably in the bluer optical passbands.  Color differences, among galaxies and locally within galaxies, will translate in different $M/L_\\lambda$ values; one can expect that this will influence studies involving rotation curve fitting and the TF-relation.  In this paper I concentrate on the use of color profiles as a diagnostic tool to investigate dust and stellar content of spiral galaxies.  Other processes that may contribute to the broadband colors (e.g.\\ emission from hot dust in the $K$ passband) are ignored as they are expected to be small in most cases.  The structure of this paper is as follows.  In Sect.~\\ref{data4} the data set is described and the color profiles of the 86 spiral galaxies using the $B, V, R, I$ and $K$ passband data are presented.  Section~\\ref{colprofsec} describes the extinction models and the stellar population models used in this paper and then compares these models to the data.  In Sect.~\\ref{struccol}, I investigate the relation between the color properties of the galaxies and the structural galaxy parameters derived in the previous papers of this series.  Implications of the current measurements are discussed in Sect.~\\ref{discus4} and the paper is summarized in Sect.~\\ref{concl4}. ",
        "conclusions": "\\label{concl4}  The stellar and the dust content of a large sample of galaxies was investigated using the color profiles of these galaxies.  Data in four optical and two near-IR passbands were combined simultaneously to derive structural properties of the sample as a whole, rather than for individual galaxies.  The main conclusions are:  \\begin{itemize} \\item Almost all spiral galaxies become bluer with increasing radius. \\item The colors of galaxies correlate strongly with surface brightness, both within and among galaxies. The morphological type is an additional parameter in this relationship, because at the same surface brightness late-type galaxies are bluer than early-type galaxies. \\item Realistic 3D radiative transfer modeling indicates that reddening due to dust extinction cannot be the major cause of the color gradients in face-on galaxies. The predicted color vectors in color--color space are not compatible with the data, unless the assumed scattering properties of the dust are entirely wrong. \\item The color gradients in the galaxies are best explained by differences in SFH as function of radius, with the outer parts of galaxies being on average much younger than the central regions. This implies that the stellar scalelength of galaxies is still growing. The central stellar populations in a galaxy must have a range in metallicities to explain the red central colors of the galaxies. \\item A consequence of the population changes implied by the color differences in and among galaxies is that there are large changes in $M/L$ values in and among galaxies.  These changes in $M/L$ make the missing light problem in spiral galaxies as derived from rotation curve fitting even more severe. \\item The $H$ and $K$ passbands are recommended for standard candle work and for studies depending on $M/L$ ratios in galaxies. \\end{itemize}"
    },
    "9604/astro-ph9604156_arXiv.txt": {
        "abstract": "We investigate the Kelvin-Helmholtz instability of stratified jets. The internal component (core) is made of a relativistic gas moving with a relativistic bulk speed. The second component (sheath or envelope) flows between the core and external gas with a nonrelativistic speed.  Such a two-component jet describes a variety of possible astrophysical jet configurations like e.g. (1) a relativistic electron-positron beam penetrating a classical electron-proton disc wind or (2) a beam-cocoon structure. We perform a linear stability analysis of such a configuration in the hydrodynamic, plane-parallel, vortex-sheet approximation.  The obtained solutions of the dispersion relation show very apparent differences with respect to the single-jet solutions.  Due to the reflection of sound waves at the boundary between sheet and external gas, the growth rate as a function of wavenumber presents a specific oscillation pattern. Overdense sheets can slow down the growth rate and contribute to stabilize the configuration. Moreover, we obtain the result that even for relatively small sheet widths the properties of sheet start to dominate the jet dynamics. Such effects could have important astrophysical implications, for instance on the origin of the dichotomy between radio-loud and radio-quiet objects. ",
        "introduction": "Kelvin-Helmholtz instability due to shearing induced by large velocity gradients is well known to develop in extragalactic jets. This has been shown by several analytical and numerical works such as the ones reviewed for instance by Birkinshaw, 1991. These studies often consider the jet as a single fluid with one bulk velocity and a single interface made by a sheared layer with the external medium.  However there exist both observational and theoretical pieces of evidence supporting the fact that components with different bulk velocities are present inside jets themselves.  Among the observational data showing the presence of different velocities inside jets we can mention (i) the detection in some compact sources of VLBI components with different apparent speed, (ii) the superluminal effect which proves that some jets are highly relativistic at VLBI scale while properties of the extended component rather suggest classical or mildly relativistic large scale jets, (iii) the presence in some sources of entrained gas which emits optical lines from which one can infer velocities likely smaller than the main bulk velocity of the jet.  The peculiar morphology of the wide angle tail radio sources (WATs) appears also very difficult to explain (Eilek et al, 1984; O'Donoghue et al, 1990 and 1993) except if one assumes that matter with two different speeds is present in the jet, namely one relativistic component radiating before the inner hot spot and one component with slow speed radiating after the inner hot spot (Leahy, 1984). Recent data on M87 provide the first direct measurement of apparent bulk velocity at the kiloparsec scale and show a complex velocity pattern with the presence of quite different bulk flow velocities along the jet (Biretta et al, 1995). In 3C273, the radio emitting jet appears longer and wider than the optical one. Bahcall et al (1995) suggest that there are actually two superposed components, a fast-moving inner jet surrounded by a slow-moving ``cocoon''.  Besides that, a Fanaroff-Riley type II radio source can also be considered as a stratified jet with two different layers corresponding to the inner real jet and to the backflow forming the surrounding cocoon.  A quite new way of looking at the radio data appears as well in favour of the existence of several components in radio sources. Starting with the use of color-color diagrams (Katz-Stone et al, 1993), Rudnick et al, 1994 and Katz-Stone and Rudnick (1994) show that it is reasonable to assume a distribution of radiating particles which is not a power-law throughout the sources and can then partially isolate the different parameters which contribute to the synchrotron brightness, namely the magnetic field B, the number of radiating particles $N_T$ and some fiducial Lorentz factor $\\gamma_0$ that characterizes their energy distribution. Adding some knowledge concerning the shape of the radiating particle distribution, they can produce frequency-independent maps, proportional to $N_TB$, $\\gamma^2_0B$ and $N_T/\\gamma_0^2$. This powerful method provides a completely changed view of the sources in which new features are discovered. For instance, in the eastern lobe of Cygnus A, they detect an edge-brightened channel girdled by rings, which corresponds to a real enhancement of the radiating particle density, likely located around the counterjet (Katz-Stone, Rudnick, 1994). Recent developments of their technics lead them to suggest that jets in both FRI and FRII sources may have coaxial sheaths a few times wider than the jets themselves. From their ``tomography'' analysis of radio data which combine maps at different frequencies, they detect two-component structures in 3C449 and 1231+674, with flat spectrum jets surrounded by sheaths of steep spectrum emission with different polarization properties (Rudnick, 1995). It is not yet known whether such sheaths emanate from the inner jets or have been directly ejected from the central engines. However their presence deeply emphasizes the likelihood of multi-component jet models.  Indeed several models of jet formation and propagation reach the concept of stratified jets.  First Chan and Henriksen, 1980, mentioned the ``multilevel'' structure as a basic modification to add to their self-similar jet model in order to explain why jets appear to have nozzles on very different scales. Smith and Raine, 1985, investigate a two-level configuration and explore the possibility to produce collimated outflows by the interaction of a nuclear wind from the very inner region of an active galactic nucleus with a Compton-heating-induced wind from an accretion disc. Baker et al, 1988, deal with another kind of two-component model while studying the radiative properties of a relativistic particle beam injected into an extragalactic jet and the possibility of collective emission of radio waves. A completely different approach by Melia and K\\\"onigl, 1989, studies the Compton-drag deceleration of ultrarelativistic nuclear jets.  It predicts the existence of a transverse gradient in the asymptotic bulk Lorentz factor distribution of the particles in the radio jets. K\\\"onigl and Kartje (1994) also consider highly stratified winds from accretion discs, with inner ionized gas and outer dusty neutral outflow. Sol, Pelletier, Ass\\'eo, 1989, propose a two-component model taking into account a relativistic electron-positron beam extracted from the funnel of an accretion disc and streaming through a classical electron-proton collimated wind coming out from all parts of the disc.  This model provides a simple explanation to the velocity dilemma if VLBI jets and superluminal motion are related to the relativistic beam while kiloparsec scale radio features are associated to the slower collimated wind.  Stability of such a configuration relatively to the excitation of Langmuir, Alfven and whistler waves has been found for electron-positron beams propagating along strong enough magnetic field with small enough Lorentz factor (Sol et al, 1989; Achatz et al, 1990; Pelletier and Sol, 1992;  Achatz and Schlickeiser, 1992). The condition for stability against excitation of Langmuir waves requires a strong longitudinal magnetic field $B$ such that the electron gyrofrequency is larger than the plasma frequency in the wind, namely $B>3.2\\times 10^{-3} \\sqrt{n_p}$ in CGS units where $n_p$ is the wind density.  The constraint on the Lorentz factor $\\gamma$ comes from the necessity to avoid excitation of Alfven waves which imposes $\\gamma < \\sqrt{m_p\\over m_e} \\simeq 43$.  Therefore it appears possible to quench the beam-plasma instability and to ensure a two-component configuration stable from the point of view of microphysics.  However, the question of large scale fluid instability of such two-component jet is still open. It is investigated in the present work. As a first approach to the study of the Kelvin-Helmholtz stability of stratified jets, we consider here a core-sheet structure. It is the simplest ``multilevel'' configuration which appears in some jet models (Smith and Raine, 1985;  Baker et al, 1988;  Sol et al, 1989) and provides an approximation to the general case of jets surrounded by cocoons or sheaths.  \\begin{figure}     % \\epsfxsize=\\hsize \\epsfbox{fig1.ps} \\caption[]{ Schema of a stratified jet with core, sheet and external components in the unperturbed state. } \\end{figure} ",
        "conclusions": ""
    },
    "9604/astro-ph9604142_arXiv.txt": {
        "abstract": " ",
        "introduction": "Redshift-space distortion of galaxy two-point correlation functions is known as a powerful tool in estimating the cosmological density parameter $\\Omega_0$; on nonlinear scales Davis \\& Peebles (1983) computed the relative peculiar velocity dispersions of pair of galaxies around $1h^{-1}$Mpc from the anisotropies in the correlation functions of the CfA1 galaxy redshift survey, and then concluded that $\\Omega_0=0.20 e^{\\pm0.4}$ (also see Mo, Jing \\& B\\\"{o}rner 1993; Suto 1993; Ratcliffe et al. 1996). In linear theory Kaiser (1987) showed that the peculiar velocity field systematically distorts the correlation function observed in redshift space; the averaged redshift-space correlation function $\\xi^{(s)}(x)$ of galaxies is related to its real-space counter part $\\xi^{(r)}(x)$ as \\begin{eqnarray} \\label{eq:kaiser} \\xi^{(s)}(x) &=& \\left( 1 + {2\\over 3}\\beta + {1\\over 5}\\beta^2 \\right) \\xi^{(r)}(x) , \\\\ \\label{eq:beta} \\beta &=& { f(z=0) \\over b} \\sim {1\\over b} \\left[ \\Omega_0^{0.6} + {\\lambda_0 \\over 70} \\left(1+ {\\Omega_0 \\over 2}\\right)\\right] , \\end{eqnarray} where $b$ is the bias parameter, $f(z=0)$ is the logarithmic derivative of the linear growth rate with respect to the scale factor $a$ (eq.[\\ref{eq:fz}] below) at $z=0$.  As is clear from the empirical fitting formula in the second line by Lahav et al. (1991), however, this formula depends on $\\Omega_0$ but is practically insensitive to the cosmological constant $\\lambda_0$.  At higher redshift $z \\simgt 1$, another anisotropy due to the geometrical effect of the spatial curvature becomes important. We derive a formula for {\\it the cosmological redshift-space distortion} in linear theory. This formula turns out to be a straightforward generalization of that derived by Hamilton (1992) for $z=0$. It provides a promising method to estimate both $\\Omega_0$ {\\it and} $\\lambda_0$ from future redshift surveys of galaxies and quasars including the Sloan Digital Sky Survey (SDSS) and the 2dF. A very similar idea was put forward independently by Ballinger, Peacock and Heavens (1996) in which they considered the anisotropy of the power spectrum while we developed a formulation in terms of two-point correlation functions.  We outline the derivation in \\S 2, and present some examples of results in cold dark matter (CDM) models and in power-law correlation function models. ",
        "conclusions": ""
    },
    "9604/astro-ph9604163_arXiv.txt": {
        "abstract": "Previously unseen profile components of the Crab pulsar have been discovered in a study of the frequency-dependent behavior of its average pulse profile between 0.33 and 8.4\\,GHz.  One new component, $36^\\circ$ ahead of the main pulse at 1.4\\,GHz, is not coincident with the position of the precursor at lower frequencies. Two additional, flat-spectrum components appear after the interpulse between 1.4 and 8.4\\,GHz.  The normal interpulse undergoes a transition in phase and spectrum by disappearing near 2.7\\,GHz, and reappearing $10^\\circ$ earlier in phase at 4.7 and 8.4\\,GHz with a new spectral index.  The radio frequency main pulse disappears for $f > 4.9$\\,GHz, even though it is seen at infrared, optical, and higher energies. The existence of the additional components at high frequency and the strange, frequency-dependent behavior is unlike anything seen in other pulsars, and cannot easily be explained by emission from a simple dipole field geometry. ",
        "introduction": "Since its discovery (\\markcite{Staelin \\& Reifenstein 1968}, \\markcite{Comella et al. 1969}), the Crab Nebula pulsar has been studied at many wavelengths, including optical, X-ray and $\\gamma$-rays, to understand the properties of the emission mechanism. Radio emission from the Crab pulsar is unusual because of its giant pulses, which are very powerful intensity fluctuations varying from 100 to 1000 times the mean intensity that occur at random intervals from one pulse to several thousand pulse periods.  But the pulsar's steep radio spectrum and the radio-bright Crab Nebula background make observations above 1\\,GHz difficult with single dish antennas. Average profiles formed at 4.7\\,GHz during the single successful observation in a decade of attempts at the Arecibo 305-m telescope, revealed that the main pulse (MP) to interpulse (IP) separation appeared to be smaller than at longer wavelengths, and that additional pulse profile components were present following the IP. To investigate this atypical frequency-dependent behavior of the pulse profile and to overcome the background contribution of the Nebula, the observing program was continued at the Very Large Array (VLA) of the National Radio Astronomy Observatory. ",
        "conclusions": "In a multifrequency study of the Crab pulsar, we have found new and unusual components that defy explanation by emission from a simple dipole field geometry.  Two of the new components have a flatter spectrum than the main pulse and interpulse, and so they survive to become visible at high radio frequencies, possibly even at infrared energies.  The interpulse is replaced between 2.7 and 4.7\\,GHz with a new component exhibiting a much different spectrum.  And a new component found $36^\\circ$ ahead of the main pulse may be a leading conal component to the core emission of the precursor. We are making polarimetric observations at frequencies higher than 1\\,GHz of the the new components to understand their emission and location in the Crab's magnetosphere.  We find that the occasional ``giant'' pulses from the Crab are emitted mostly at the phase of the MP and less often at the IP, but never at the phases of the other components.  Since there seem to be no ``giant'' pulses at optical or $\\gamma$-ray wavelengths where the emission is incoherent (\\markcite{Lundgren 1994}), the giant pulse phenomenon must be related to the degree of radio emission coherence, rather than energetic particle production. The emission region conditions must, therefore be similar at the MP and IP, but the LFC, HFCs, and precursor must be significantly different.  The VLA is part of the National Radio Astronomy Observatory, which is operated by Associated Universities Inc., under a cooperative agreement with the National Science Foundation. The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under contract with the National Science Foundation. This work has been conducted with partial support of NSF grant AST-9315285. DAM acknowledges the support of NRAO as a Junior Research Fellow, and also thanks D.~Nice of NRAO for software support and advice during the data reduction process of this project. We thank J.~Rankin for valuable background information and discussion of the results, and we thank J.~Gil for a provocative review.  \\vfill\\eject"
    },
    "9604/astro-ph9604096_arXiv.txt": {
        "abstract": "Spatially resolved velocity profiles are presented for nine faint field galaxies in the redshift range {0.1 $\\lesssim$ \\z\\ $\\lesssim$ 1}, based on moderate--resolution spectroscopy obtained with the Keck~10~m telescope.  These data were augmented with high--resolution {\\it Hubble Space Telescope} images from WFPC2, which provided V and I photometry, galaxy type, orientation, and inclination.  The effects of seeing, slit width, and slit misalignment with respect to galaxy major axis were modeled along with inclination for each source, in order to derive a maximum circular velocity from the observed rotation curve.  The lowest redshift galaxy, though highly elongated, shows a distorted low--amplitude rotation curve that suggests a merger in progress seen perpendicular to the collision path.  The remaining rotation curves appear similar to those of local galaxies in both form and amplitude, implying that some massive disks were in place at \\z\\ $\\sim 1$.  The key result is that the kinematics of these distant galaxies show evidence for only a {\\it modest} increase in luminosity ($\\Delta$M$_B \\lesssim 0.6$) compared to velocity--luminosity (Tully--Fisher) relations for local galaxies. ",
        "introduction": "The rotational velocity and luminosity of disk galaxies are found to be strongly correlated (Roberts \\etal~1975; Tully \\& Fisher 1977).  This scaling relation --- the ``Tully--Fisher (TF) relation'' --- provides a powerful tool to tackle such problems as deriving \\h0 (\\eg Pierce \\& Tully 1988), or mapping the local galaxy streaming motions (\\eg Aaronson \\etal~1986).  These studies have been confined to nearby galaxies by the use of single dish \\HI\\ radio observations (\\eg Haynes \\& Giovanelli 1991) or optical emission line spectra (\\eg Rubin \\etal~1985;  Mathewson \\etal~1992).  Extending velocity width studies to more distant galaxies would be particularly valuable to investigate galaxy evolution (Kron 1986; van der Kruit \\& Pickles 1988).  Current galaxy evolution models range from those with mild amounts of luminosity brightening in the past (\\eg Gronwall \\& Koo 1995) to those requiring more dramatic changes (\\eg Broadhurst \\etal~1988, Colless \\etal~1990, Glazebrook \\etal~1995) to explain the large numbers of faint blue galaxies.  By comparing a distant sample of rotation curves to local TF relations, we can directly constrain the global brightening of disk galaxies in the past.  The \\HI\\ Tully--Fisher method is limited by the sensitivity of current radio telescopes to \\z\\ $\\lesssim$ 0.1.  Beyond this, two approaches have been used. Vogt \\etal~(1993) derived rotation curves from strong optical emission lines, for two spirals at \\z~$\\sim$~0.2.  Without spatial resolution, an alternative measure of distant spiral kinematics may be extracted from the velocity widths of emission lines.  Forbes \\etal~(1996) measured velocity widths for a sample of 18 faint field galaxies with redshifts 0.2~$<$~\\z~$<$~0.84, while Colless \\etal~(1994) examined the equivalent widths of \\fOII\\ for a sample of 26 field galaxies with redshifts 0.1~$\\lesssim$~\\z~$\\lesssim$~0.7.  This paper presents rotation curves for nine field galaxies at redshifts {0.1~$\\lesssim$~\\z~$\\lesssim$~1}.  This project combines spatially--resolved spectra from the 10~m~Keck telescope with inclinations and position angles determined from the refurbished {\\it Hubble Space Telescope} (HST).  These data provide a first glimpse into the internal kinematics of galaxies out to a redshift of \\z~$\\sim$~1, or one--third to one--half (for $\\Omega_0$ = 1 - 0) of the age of the universe. ",
        "conclusions": "An immediate --- although not surprising --- result is that the shapes of the rotation curves of these high redshift galaxies are similar to those of local galaxies.  The high-redshift rotation curves are relatively symmetric, show a ``solid-body'' rise in the inner regions, and turnover to a relatively constant circular velocity in the outer parts. The maximum velocities (see Table~1) are comparable to those of local spirals.  Rough calculations yield masses between 1 and 5 $\\times$ $10^{11}$ $M_{\\odot}$, well within the range of masses found for nearby spiral galaxies.  This result, combined with the apparent disks which are obvious in the HST images, shows conclusively that some {\\it massive} disk systems were in place by \\z~$\\sim$ 1.  In Figure~\\ref{TF} we compare the data to local TF relations in the rest--frame $B$-band, which corresponds to \\V\\ at \\z~$\\sim$ 0.4 and to \\I\\ at \\z~$\\sim$ 0.8 (\\ie the $k$ corrections are small).  The comparison in Figure~\\ref{TF}$a$ is to the TF relation (inverse fit, \\ie \\Vmax\\ as a function of $M_B$) for 32 spiral galaxies in the Ursa Major cluster (Pierce \\& Tully 1988, 1992).  This relationship is based on \\HI\\ velocity width measurements (corrected for turbulent broadening), but we have not converted our optical--line terminal velocities to radio widths since this correction ($\\lesssim$~15~\\kms) is small compared to the optical error (\\eg Mathewson \\etal~1992; Giovanelli \\etal~1996).  The two sources with redshift \\z~$\\sim$~0.2 from Vogt \\etal~(1993) are plotted for comparison; with only ground--based imaging available, they are less well constrained in inclination.  Excluding these and the peculiar source 074$-$2262, we compute a weighted offset of 0.55~$\\pm$~0.15 mag relative to the $B$-band TF relation of Pierce \\& Tully, with a dispersion of 0.71 mag.  This dispersion matches the quadratic summation of errors (0.65) in the logarithmic velocity widths (0.47), the rest--frame $B$ magnitudes (0.2), and an assumed intrinsic scatter in the TF relation (0.4) (\\cf Willick \\etal~1996 and references therein), indicating that our error estimates are consistent.  An analogous calculation for the $I$-band relation yields an offset of 0.36~$\\pm$~0.18~mag.  Figure~\\ref{TF}$b$ shows the same galaxies and the TF relation (double regression fit) for local field galaxies of Hubble types Sa, Sb, and Sc as published in Rubin \\etal~(1985), corrected to \\Hconst{75}.  Since our galaxy types (see Table~1) are evenly distributed between Sb and Sc, we compare to a local relation midway between the Sb and Sc, and measure an offset of 0.38~$\\pm$~0.22~mag.  Several sample selection effects and assumptions work to make these offsets upper limits.  First,  at high redshift, our samples are biased toward intrinsically luminous galaxies, which will affect somewhat our results with respect to TF relations (\\cf Teerikorpi 1984).  Second, we have selected spatially-extended objects, which will bias our sample toward larger galaxies. Third, we have selected objects with detectable emission lines, which will bias the sample toward galaxies with stronger than average star formation --- and therefore higher luminosity.  Corrections for these biases would all {\\it reduce} the true offsets.  Moreover, if we have not traced the rotation curve to sufficiently large radii, we may have underestimated the maximum velocities, and if galaxies were less dusty at earlier epochs, we may have overcorrected for extinction.  These errors, if present, would also {\\it reduce} the true offsets. Finally, adopting \\q0\\ = 0.5 instead of \\q0\\ = 0 would decrease the restframe luminosities by 0.1 -- 0.4 magnitudes, and again reduce the true offset.  The issue of a representative local TF sample for comparison is also critical. Key factors include the photometry passband, the observations and analysis used to determine internal velocities, the selection biases due to catalog limits, the distortions due to peculiar velocities, internal extinction corrections, and the TF fitting technique (\\eg forward, inverse, or double regression). Note, however, that our sample lies in the central region ($-19.8 < M_B < 21.7$, $180 <$ \\Vmax\\ $< 290$ \\kms) of the range fit in the local samples, where the differences between various fits are minimized.  We assume an error due to these effects of 0.35 mag and combine this with the measured uncertainties in our offsets ($\\sim 0.2$) from local TF relations to arrive at an upper limit of $\\Delta$M$_B \\lesssim 0.55 \\pm 0.38$ mag.  In summary, we find an offset relative to local TF relations of $\\Delta$M$_B \\sim 0.6$ mag as a strong upper limit.  We do {\\it not} see an obvious trend with redshift or morphological type, but our sample is small.  Our result is compatible with other kinematic studies of field galaxies at intermediate redshifts, such as Rix \\etal~(1996), who find a brightening of 0.44 mag for blue field galaxies at \\z~$\\sim$ 0.3.  Forbes \\etal~(1996) concluded that galaxies near \\z~$\\sim$ 0.5 show a surface brightness increase of 0.6 $\\pm$ 0.1 mag, while Schade \\etal~(1995) find a 1.2 $\\pm$ 0.25 mag increase in disky galaxies at redshifts 0.5 $<$~\\z~$<$ 1.2. Since the average bulge--to--total ratio for the Schade \\etal\\ sample is $\\sim$ 0.1, any reductions due to a significant bulge would be small.  While the samples are not directly comparable, our data suggest somewhat less brightening and exclude any total (versus surface) brightening by more than 1.5 magnitudes at the 99\\% CL.  The results can be reconciled if disky galaxies of similar mass were not much brighter at \\z~$\\sim$ 1 than today but had slightly smaller disk scale lengths and thus higher surface brightness.  This Letter presents first results in the investigation of kinematics via rotation curves in high redshift galaxies.  We plan future observations to increase the sample size, and will constrain slits to lie within 30\\deg~of the major axes.  We will use this larger, higher--quality data set to focus on a better understanding of sample selection biases, and will further explore a more analogous local TF relation.  Nevertheless, the current data clearly indicate that optical rotation curves can be measured up to \\z~$\\sim$ 1 and will provide an important constraint on our understanding of the evolution of disk galaxies."
    },
    "9604/astro-ph9604119_arXiv.txt": {
        "abstract": "We present a color analysis of the galaxy populations of candidate clusters of galaxies from the Palomar Distant Cluster Survey (Postman et al.\\ 1996).  The survey was conducted in two broad band filters that closely match $V$ and $I$ and contains a total of 79 candidate clusters of galaxies, covering an estimated redshift range $0.2 \\simless z \\simless 1.2$.  We examine the color evolution in the 57 richest clusters from this survey, the largest statistical sample of distant clusters to date. The intermediate redshift ($0.2 \\simless z \\simless 0.4$) clusters show a distinct locus of galaxy colors in the color--magnitude diagram. This ridge line corresponds well with the expected no--evolution color of present--day elliptical galaxies at these redshifts. In clusters at redshifts of $z \\simgreat 0.5$, this red envelope has shifted bluewards compared to the ``no--evolution'' prediction. By $z \\sim 0.8$ there are only a few galaxies which are as red in their rest-frame as present--day ellipticals, consistent with recent claims on the basis of optical--infrared colors. The detected evolution is consistent with passive aging of stellar populations formed at redshifts of $z \\simgreat 2$.  Though the uncertainties are large, the Butcher--Oemler effect is observed in the Palomar clusters. The fraction of blue galaxies increases with the estimated redshift of the cluster at a 96.2\\% confidence level. The measured blue fractions of the intermediate redshift clusters ($f_{b} \\sim 0.2 - 0.3$) are consistent with those found previously by Butcher \\& Oemler (1984). The trend in the Palomar clusters suggests that $f_{b}$ can be greater than 0.4 in clusters of galaxies at redshifts of $z \\simgreat 0.6$. ",
        "introduction": "The study of the galaxy populations of rich clusters provides important constraints on the formation mechanisms of both clusters and galaxies. Present--day clusters show a distinct correlation between the structure of the cluster and the galaxy population. Irregular, open clusters, such as Virgo, are spiral--rich. These systems show no obvious condensations, though the galaxy surface density is at least five times as great as the surrounding field ($n_{\\rm gal} > 30~h^{3}~{\\rm galaxies~{Mpc}^{-3}}$). These clusters may be highly assymetric and have significant degrees of substructure.  Dense, centrally concentrated clusters, such as Coma, contain predominantly early--type galaxies in their cores (Abell 1958; Oemler 1974; Dressler 1980; Postman \\& Geller 1984).  These clusters have a single, outstanding concentration among the bright member galaxies and typically display a high--degree of spherical symmetry. Central densities can reach as high as $10^{4}~h^{3}~{\\rm galaxies~{Mpc}^{-3}}$. The galaxy content of clusters is part of the general morphology--density relation of galaxies; as the local density increases, the fraction of elliptical (E) and S0 galaxies increases, while the fraction of spiral galaxies decreases (Dressler 1980; Postman \\& Geller 1984).  One of the most intriguing results of the study of intermediate and high--redshift clusters of galaxies has been strong evolution in the galaxy population.  Butcher \\& Oemler (1984; hereafter BO) were the first to make a comprehensive study of intermediate redshift clusters with regard to their galaxy populations. They examined the fraction of blue galaxies ($f_{b}$) in 33 clusters at $z \\simless 0.5$. They found that $f_{b}$ is an increasing function of redshift in both open and compact clusters of galaxies, indicating that clusters at these redshifts are significantly bluer than their low--redshift counterparts. Recent HST image data (Dressler et al.\\ 1994; Couch et al.\\ 1994; Oemler, Dressler \\& Butcher 1996) reveal that many of these blue ($g-r < 1.2$) galaxies are either ``normal'' spirals or have peculiar morphologies, producing non--elliptical fractions which are 3 to 5 times higher than the average current epoch cluster.  Fading through cessation of star formation may play a role in this evolution. Gunn \\& Dressler (1988) find that the spectra of cluster galaxies with $z \\simgreat 0.6$ show, on average, smaller 4000$\\AA$ decrements and a higher frequency of post--starburst features (the ``E+A\" spectral class) than those at $z < 0.6$ (however, see Zabludoff et al.\\ 1996).  Detailed photometric observations of other intermediate redshift ($z \\simless 0.4$) clusters have confirmed the original results of BO. Even though these clusters show an increased fraction of blue galaxies, they still contain a population of E/S0s which distinguish itself by extremely red colors and a tight color--magnitude (CM) relation (a ``red envelope'').  Both the mean color and the CM relation is consistent with that of present--day ellipticals (Sandage 1972; Visvanathan \\& Sandage 1977; Butcher \\& Oemler 1978; Couch \\& Newell 1984; BO; Ellis et al.\\ 1985; Sandage, Bingelli \\& Tammann 1985; Millington \\& Peach 1990; Arag\\'on-Salamanca, Ellis \\& Sharples 1991; Luppino et al.\\ 1991; Dressler \\& Gunn 1992; Le Borgne, Pell\\'o \\& Sanahuja 1992; Dressler et al.\\ 1994; Molinari et al.\\ 1994; Smail, Ellis \\& Fitchett 1994; Stanford, Eisenhardt \\& Dickinson 1995).  Arag\\'on-Salamanca et al.\\ (1993; hereafter A93) have studied a small sample of 10 rich clusters at $0.5 < z < 0.9$ in the optical--infrared colors.  They observe an increase in the number of blue members which they interpret as the high--redshift extension to the Butcher--Oemler effect. This trend is also well studied by Rakos \\& Schombert (1995) in an intermediate and high--redshift sample of 17 rich clusters. They find the fraction of blue galaxies increases from 20\\% at $z = 0.4$ to 80\\% at $z = 0.8$ using Str\\\"omgren photometry in the cluster rest-frame.  In addition, the red envelope, presumably that of the early--type population, moves {\\it bluewards} with redshift (A93; Rakos \\& Schombert 1995; Oke, Gunn \\& Hoessel 1996). At $z \\sim 0.9$, there are few cluster members with colors as red as present--day ellipticals (see also Smail et al.\\ 1994). The color distribution of this high-redshift elliptical population is relatively narrow, and the trend is uniform from cluster to cluster; this suggests a homogeneous population which formed within a narrow time span (e.g.\\ Bower, Lucey \\& Ellis 1992a,b). Dickinson (1995) finds similar results in a cluster of galaxies which is associated with the $z = 1.206$ radio galaxy 3C 324. The galaxies exhibit a narrow, red locus in the CM magnitude diagram. This branch is $\\sim 0.6$ mag bluer than the expected ``no--evolution'' value, though the intrinsic rms color scatter is only 0.2 mag. The observed color trend for the red envelope of galaxies is consistent with passive evolution of an old stellar population formed by a single burst of star formation at redshifts of $z \\simgreat 2$.  The reasonably small color scatter would imply closely synchronized intra--cluster star formation (Bower et al.\\ 1992a,b; A93; Dickinson 1995).  It is critical to test the claims of significant color evolution in high--redshift clusters of galaxies against larger, more statistically complete samples. In this paper, we examine the galaxy colors of the largest, statistically complete sample of distant clusters presently available, the Palomar Distant Cluster Survey (hereafter PDCS; Postman et al.\\ 1996). This sample of 57 distant clusters will be used to trace the color evolution in the galaxy population and to investigate the Butcher--Oemler effect in clusters up to a redshift of $\\sim 1$. In \\S 2 of this paper, we briefly describe the PDCS, the original photometric reduction, and the cluster sample used in this analysis. The evolution in the color distributions of the cluster galaxies is presented in \\S 3. The Butcher--Oemler effect (the fraction of blue galaxies as a function of redshift) is presented in \\S 4. We summarize the results in \\S 5. ",
        "conclusions": "We have examined the color evolution of the galaxy populations of the richest candidate clusters of galaxies from the Palomar Distant Cluster Survey. The sample contains 57 intermediate and high--redshift clusters of galaxies, the largest statistical sample of distant clusters. Our principal conclusions are summarized below.  \\newcounter{discnt}  \\begin{list} {\\arabic{discnt}.}  {\\usecounter{discnt}}  \\item For intermediate redshift ($z \\simless 0.4$) clusters, we find that there is a distinct ridge line or ``red envelope'' in the color--magnitude diagram. The color of this locus corresponds well with the expected no--evolution color of present--day ellipticals at these redshifts. This result is consistent with previous optical studies of intermediate redshift clusters.  \\item At progressively higher redshifts, this red envelope, as characterized by the mean galaxy color, has shifted bluewards compared to the no--evolution prediction. By $z \\sim 0.9$, there are only a few cluster galaxies which are as red in their rest--frame as present-day ellipticals.  Through simple Monte--Carlo simulations of a synthetic non--evolving cluster population, we have confirmed that we are able to observe these very red members if they are present at these redshifts.  The blueing trend of the mean galaxy color by $\\sim 0.4$ mag at $z \\sim 0.6$ and over 1 mag at $z \\sim 0.8$ is consistent with that previously observed in the optical--infrared colors.  A comparison between the observed color evolution in the Palomar clusters with simple Bruzual \\& Charlot (1996) models of passive aging of an old stellar population formed in a single burst of star formation indicates formation epochs of $z \\simgreat 2$ for the early-type cluster galaxies.  \\item Though the uncertainties are large, the Butcher--Oemler effect is observed in the Palomar clusters. We find that the fraction of blue galaxies increases with the estimated redshift of the cluster at a 96.2\\% ($2 \\sigma$) confidence level. We observe blue fractions of $f_{b} \\sim 0.05 - 0.2$ in clusters at redshifts of $0.2 \\simless z \\simless 0.4$. These fractions are consistent with that found previously by Butcher \\& Oemler (1984). Our calculations of $f_{b}$ are only complete to redshifts of $z \\sim 0.6$; at these redshifts, $f_{b} \\sim 0.3 - 0.4$. The observed trend indicates that clusters at $z \\simgreat 0.6$ can have blue fractions which are greater than 0.4. This result is consistent with that found by Rakos \\& Schombert (1995). They have examined 17 clusters over a redshift range similar to ours and find that the blue fraction $f_{b}$ increases to 0.8 in clusters at $z \\sim 0.9$.  \\end{list}   \\vskip 0.3cm  The referee Jim Schombert is graciously thanked for his insightful comments and an extremely productive visit. It is also a pleasure to thank Marc Postman for his continual guidance and essential scientific contributions to this paper and Ian Smail for his special attention and thorough review of this text. Neta Bahcall, Jim Gunn, Michael Strauss, and Ed Turner are thanked for providing invaluable comments on a preliminary version. The following people are acknowledged for their generous contributions : Robert Lupton for his useful comments and aid in photometric conversions, Don Schneider for providing the spectral energy distributions and response functions, St\\'ephane Charlot and Ann Zabludoff for supplying and explaining the Bruzual \\& Charlot synthesis codes, and St\\'ephane Courteau for being available for scientific consultations. LML graciously acknowledges support from a Carnegie Fellowship and NASA contract NGT--51295.  \\vfill \\eject"
    },
    "9604/astro-ph9604082_arXiv.txt": {
        "abstract": "We have obtained near-infrared imaging and an optical spectrum of the proto-planetary nebula IRAS~07131$-$0147, a highly polarized bipolar reflection nebula believed to be in evolutionary transition from the asymptotic giant branch to the planetary nebula phase. Our images reveal point reflection symmetry in the lobes - a relatively rare morphological feature. We place an upper limit on the distance of 6.5 kpc. Utilizing numerical integrations of single grain scattering models we find the nebula to lie at an inclination angle of $i=20\\pm5$ degrees in the plane of the sky. We present a refined geometric interpretation of IRAS~07131$-$0147 consistent with our new observational data and argue that the central star is likely to have an as yet undetected binary companion. ",
        "introduction": "Proto-planetary nebulae (PPN) occur in the poorly understood transition between the asymptotic giant branch (AGB) and planetary nebula (PN) stages of stellar evolution.  As a star moves up the AGB, it loses mass at an increasing rate, culminating in a final cataclysmic superwind (\\cite{Wood 93}).  The end of mass-loss signals the transition to the beginning of the PPN phase (\\cite{Kwok 93}).  At this point, the central star with a typical core mass of $M_{c} = 0.6M_{\\odot}$ and temperature of $T_{*} \\approx 5000$ K (\\cite{Volk and Kwok 89}) evolves blueward in the HR diagram, its temperature increasing as it loses additional envelope mass through hydrogen shell burning, until it reaches $T_{*} \\approx 30,000$ K necessary to ionize the circumstellar shell and produce a PN. The ionization of the nebula defines the end of the PPN stage. The typical ages of PN, the large dust densities required for their formation, and the minimum central star temperature needed to ionize the ejected gas strongly constrain the duration of the PPN phase to a few thousand years (\\cite{Renzini 83}).  Some PPN have easily discernable bipolar morphologies reminiscent of ``butterfly'' PN, in which the central star illuminates two diametrically opposed, symmetric lobes of nebulosity (\\cite{Balick 87}). Butterfly PN structure is generally attributed to a density contrast characterized by the presence of a higher concentration of dust in an equatorial plane around the central star and an interacting wind, where a hot wind plows into a previously ejected, slow moving envelope (\\cite{Livio 93}). The morphological similarity between bipolar PPN and butterfly PN, and their evolutionary relation, allow us to interpret the objects in an analagous manner.  Morris (1981) investigated the dust distributions that would reproduce bipolar reflection nebulae and found that a density contrast was a necessary criterion. A waist of dust blocks the light from the star in the equatorial region, while starlight emerging into the lower density lobes of the polar regions experiences less extinction. His dust distributions reproduce the ``horns'' seen in some bipolar PPN; they are caused by changes in the optical depth structure of the lobes. The equatorial density enhancement appears well founded.  Observational evidence indicates that such dust distributions begin forming well before the PPN stage (\\cite{Johnson and Jones 91}).  However, full development of lobe structure, in agreement with the scattering models of Morris, and by analogy with the butterfly PN scenario of Balick, requires an interacting wind.  Herein lies one of the most puzzling aspects of PPN; namely, where does the second wind come from?  If we accept, for the moment, that some mechanism is capable of producing a fast wind and we account for the rapid evolution of the post-AGB central stars, then PPN may lie along a morphological sequence related to their age/evolutionary state. They evolve from young systems with totally obscured central stars surrounded by reflection nebulae (e.g. AFGL 2688, OH 0739$-$14), to middle-aged systems in which the circumstellar material has been sufficiently dispersed to reveal the central star (e.g. Frosty Leo), to old systems in which the central star has begun to photoionize the surrounding nebula, which is then seen in both reflection and emission (e.g. M2-9, Mz-3) (for examples see \\cite{Bujar 92} and references therein). This sequence is admittedly based on a small sample of transition objects whose ages are poorly known. Furthermore, some PPN have completely obscured central stars and photoionized central cavities (e.g. AFGL 618; \\cite{Kwok 93}); possibly evolving so quickly as to appear with both young and old features. Nevertheless, we proceed with our data interpretation in the context of this intuitive, albeit loosely defined, evolutionary sequence.   Wolstencroft, Scarrott, and Menzies (1989) discovered a bipolar reflection nebula at the position of the infrared object IRAS~07131$-$0147 (galactic coordinates l=217$^{\\circ}$, b=+4.7$^{\\circ}$) during a search for optical counterparts of IRAS sources with peak flux densities at $25~\\micron$.  They classified the nebula as a PPN of intermediate age based on the fact that the lobes are seen in reflection only and the central star is relatively unobscured.  The central star was found to be oxygen-rich and classified as spectral type M5 III based upon strong TiO bands in its optical spectrum and a $10~\\micron$ silicate feature in the IRAS $Low$ $Resolution$ $Spectrum$ (LRS) (Scarrot et al.~1990, hereafter S90). The IRAS LRS is consistent with PPN model spectra; furthermore, the rising flux from 10 to 20 $\\micron$ indicates that IRAS~07131$-$0147 is somewhat evolved (\\cite{Volk and Kwok 89}).  The latter feature is characteristic of an evacuated cavity created as the dust shell expands away from the central star and cools. The optical morphology and IRAS data suggest that the nebula has evolved considerably, yet the central star has not started moving toward hotter temperatures as single star evolutionary calculations predict.  In this paper, we present near-infrared imaging and an optical spectrum of IRAS~07131$-$0147 that allow us to refine our understanding of the physical nature of this interesting transitional object. ",
        "conclusions": "Morris (1981, 1987) proposed a binary star bipolar jet collimation mechanism which invokes a bound disk excreted from the red giant and a wind emanating from accretion onto the secondary. This model is capable of producing fast winds even for PPN with late type central stars. Multi-dimensional hydrodynamic studies have shown that dust distributions consistent with the formation of bipolar reflection nebulae are an inevitable consequence of AGB mass ejection in common envelope binary systems (\\cite{Livio 93}). In this scenario, the excretion of the disk results from a brief period of friction-induced ``spiraling-in'' during the common envelope evolution.  Accretion, and possibly jets, would follow this phase.  The suggestion by Livio (1993) that the secondary may retain high entropy material, expand adiabatically, and transfer material back onto the red giant could be responsible for increasing the mass of the hydrogen envelope over its critical value (\\cite{Kwok 93}) and postponing its blueward evolution though the PPN phase. The coupling of the orbital and rotational angular momenta of the stars in the system may cause the precession of the jets (\\cite{Morris and Reipurth 90}). Recent 3-dimensional hydrodynamic simulations of precessing jets in point-symmetric nebulae conclude that interacting binary systems are the most plausible jet production mechanism (\\cite{Cliffe 95}). The resolution of binary systems in several PPN through adaptive optics imaging (\\cite{Roddier et al. 95}) lends some observational support to the binary mechanism. Objects such as IRAS~07131$-$0147 and IRAS~09371+1212 (Frosty Leo), that display evidence of evolved bipolar nebular structure and late type (unevolved) central stars, are likely the typical examples of binary PPN.  We believe the observational data now available for IRAS~07131$-$0147 supports the binary mechanism of PPN formation.  The optical and near-infrared images suggest an evolved PPN morphology where the central star is no longer obscured by its equatorial waist. The IRAS LRS also indicates an evolved nebular morphology, where the dust has expanded away from the central star and cooled.  Yet the central star is a late M giant, virtually unevolved through the PPN phase. We find striking evidence for the precession of the fast dust-shaping wind in this object.  Our near-infrared images reveal point symmetry in the inner lobe regions.  Furthermore, a comparison with the inclination derived from optical polarization data, sensitive to the outer lobe regions, and our infrared-derived inclination, sensitive to the inner lobe regions, suggests that the fast wind also precessed around the equatorial axis.  This dynamic geometric interpretation is consistent with the observed lobe colors and the ridge-like brightening of the optical nebulosity several arcseconds away from the central star. No viable theory other than the binary mechanism has been proposed for the collimation or precession of a fast wind in PPN systems with late type central stars.  The binary central star hypothesis could be tested in several ways.  High signal to noise optical spectra of the reflection lobes could reveal a faint blue bump, emanating from the unseen binary companion.  This was suggested by Cohen et al.~(1985) to explain a blue bump in the lobe spectrum of the PPN, OH 0739$-$14; the binary companion was later resolved by adaptive optics imaging (\\cite{Roddier et al. 95}). The $Hubble$ $Space$ $Telescope$ may be able to resolve a companion in IRAS~07131$-$0147. Alternatively, if the binary system is eclipsing, a photometric monitoring campaign could confirm our hypothesis."
    },
    "9604/astro-ph9604031_arXiv.txt": {
        "abstract": "We show that the Einstein ring radius and transverse speed of a lens projected on the source plane, $\\hat{r}_{\\rm e}$ and  $\\hat{v}$, can be determined from the light curve of a binary-source event, followed by the spectroscopic determination of the orbital elements of the source stars. The determination makes use of the same principle that allows one to measure the Einstein ring radii from finite-source effects. For the case when the orbital period of the source stars is much longer than the Einstein time scale, $P\\gg t_{\\rm e}$, there exists a single two-fold degeneracy in determining $\\hat{r}_{\\rm e}$. However, when $P \\lesssim t_{\\rm e}$ the degeneracy can often be broken by making use of the binary-source system's orbital motion. For an identifiable 8\\% of all lensing events seen toward the Large Magellanic Cloud (LMC), one can unambiguously determine whether the lenses are Galactic, or whether they lie in the LMC itself. The required observations can be made after the event is over and could be carried out for the $\\sim 8$ events seen by Alcock et al.\\ and Aubourg et al.. In addition, we propose to include eclipsing binaries as sources for gravitational lensing experiments. ",
        "introduction": "There are many different effects that make a light curve of a microlensing event deviate from its characteristic achromatic and symmetric form: luminous lenses (Kamionkowski 1995; Buchalter, Kamionkowski, \\& Rich 1995), differential magnification during close encounters (Gould 1994; Nemiroff \\& Wickramasinghe 1994; Witt \\& Mao 1994; Witt 1995; Loeb \\& Sasselov 1995; Gould \\& Welch 1996), parallax effects caused by the Earth's orbital motion (Gould 1992; Alcock et al.\\ 1995), and finally binary-lens events (Mao et al.\\ 1994; Axerlod et al.\\ 1994; Udalski et al.\\ 1994; Mao \\& Di Stefano 1995; Alard, Mao, \\& Guibert 1995; Alcock et al.\\ 1996c). Whenever any of these distortions is detected, it provides information about the physical parameters of individual lenses: distance to the lens for luminous lens, lens proper motion, $\\mu=v/D_{\\rm ol}$, for differential magnification, observer-plane projected Einstein ring radius, $\\tilde{r}_{\\rm e} = (D_{\\rm os}/D_{\\rm ls})r_{\\rm e}$, for parallax, and the geometry of a lens binary system and sometimes the proper motion for binary lens events. Here, $v$ is the speed of the lens relative to the Earth-source line of sight, and the physical and angular Einstein ring radius are related to the physical parameters of the lens by $$ r_{\\rm e} = \\left( {4GM_{L} \\over c^2}{D_{\\rm ol}D_{\\rm ls} \\over D_{\\rm os}} \\right)^{1/2},\\qquad \\theta_{\\rm e} = {r_{\\rm e} \\over D_{\\rm ol}}, \\eqno(1.1) $$ where  $D_{\\rm ol}$, $D_{\\rm ls}$, $D_{\\rm os}$ are the distances between the observer, source, and lens, $r_{\\rm e}$ is the physical size of the Einstein ring, and $M_{L}$ is the mass of the lens.   The light curve can be also distorted when the source is composed of a binary system: binary-source event (Griest \\& Hu 1992). The binary-source event light curve distortions take various forms depending on many factors, e.g., the angular size of the projected separation between the source stars, source trajectories within the Einstein ring, the angular size of the Einstein ring projected onto the source plane, and the orbital motion. Griest \\& Hu (1992) concentrated on binary-source events for which the binary period is long compared to the event time scale. They briefly discussed the possibility of events with short-period binary sources and illustrated the dramatic oscillation which these could in principle generate. They noted, however, that for the expected lens parameters the amplitude of these oscillations would be extremely small.   In the present paper, by contrast, we concentrate on binary-source events where the binary stars move substantially, i.e., short-period binaries. We show that the Einstein ring radius and the transverse speed projected on the source plane, $\\hat{r}_{\\rm e}$ (`source-plane Einstein ring radius') and $\\hat{v}$ (`source-plane transverse speed'), can be determined from the light curve of a binary-source event, provided that the observations are followed by spectroscopic determination of the binary-source orbital elements. Throughout this paper, we use  a ``hat'' ($\\hat{\\ }$) to represent a quantity projected onto the source plane. The source-plane Einstein radius and transverse speed are defined by $$ \\hat{r}_{\\rm e} = D_{\\rm os}\\theta_{\\rm e} = r_{\\rm e}{D_{\\rm os}\\over D_{\\rm ol}}, \\ \\ \\hat{v} = {\\hat{r}_{\\rm e} \\over t_{\\rm e}}, \\eqno(1.2) $$ where $t_{\\rm e}=r_{\\rm e}/v$ is the Einstein ring crossing time. Note that when the source distance is known [e.g., for observations toward the Large Magellanic Cloud (LMC)], measuring $\\hat{r}_{\\rm e}$ and $\\hat{v}$ is equivalent to measuring $\\theta_{\\rm e}$ and the proper motion $\\mu$ since $\\theta_{\\rm e} = \\hat{r}_{\\rm e}/D_{\\rm os}$ and $\\mu = \\hat{v}/D_{\\rm os}$. The basic principle that makes this measurement possible is that a binary acts like an enormous finite source and therefore is much more susceptible to finite-source effects than are single stars. Once the proper motion is measured, one can uniquely separate Galactic versus LMC self-lensing events because of the large difference in the proper motions between the two populations of events (see \\S\\ 6).   Short-period binary sources are important for several reasons. First, it is these events that allow one to unambiguously measure the proper motion. For long period binaries, one can determine $\\hat{r}_{\\rm e}$, but with a two-fold degeneracy (see \\S\\ 3). Second, for lenses in the LMC the amplitude of the oscillations in the flux due to the binary-source effect is expected to be of order 10\\% and would be easily observable (see \\S\\ 5). Within the framework of the standard model, LMC events are expected to be relatively rare. However, Sahu (1994) has argued that essentially all the events currently detected toward the LMC are due to LMC lenses. This may seem unlikely in view of the large optical depth (Alcock et al.\\ 1996a), but it is nevertheless important to test this hypothesis. Since the alternative is that the halo is composed in large part of Massive Compact Halo Objects (MACHOs). As we show in \\S\\ 5, about 10\\% of LMC sources are short-period binaries. These binaries allow a direct test of the Sahu (1994) hypothesis. Third, while the effects are smaller for  Galactic lenses, they are not negligible. The significant advances now being made in rapid detection and follow-up observations (Pratt et al.\\ 1995; Albrow et al.\\ 1996; Alcock et al.\\ 1996b) open the possibility that even 1\\% or 2\\% oscillations caused by binary sources may soon be measurable. Finally, binary sources are competitive with other methods of measuring proper motion toward the Galactic bulge. ",
        "conclusions": ""
    },
    "9604/astro-ph9604177_arXiv.txt": {
        "abstract": "We report the discovery of two new R Coronae Borealis (RCB) stars in the Large Magellanic Cloud (LMC) using the MACHO project photometry database.  The identification of both stars has been confirmed spectroscopically.  One is a cool RCB star ($T_{eff}\\sim$ 5000 K) characterized by very strong Swan bands of $C_2$ and violet bands of CN, and weak or absent Balmer lines, G-band and $^{12}C^{13}C$ bands. The second star is an example of a hot RCB star of which only 3 were previously known to exist in the Galaxy and none in the LMC.  Its spectrum is characterized by several C II lines in emission.  Both stars have shown deep declines of $\\Delta V \\ge 4$ mag in brightness. The new stars are significantly fainter at maximum light than the three previously known LMC RCB stars. The amount of reddening toward these stars is somewhat uncertain but both seem to have absolute magnitudes, $M_V$, about half a magnitude fainter than the other three stars. Estimates of $M_{Bol}$ find that the hot RCB star lies in the range of the other three stars while the cool RCB star is fainter.  The two cool LMC RCB stars are the faintest at $M_{Bol}$. The discovery of these two new stars brings to five the number of known RCB stars in the LMC and demonstrates the utility of the MACHO photometric database for the discovery of new RCB stars. ",
        "introduction": "The R Coronae Borealis (RCB) stars represent a rare type of hydrogen-deficient carbon-rich supergiants which undergo very spectacular declines in visual brightness of up to 8 magnitudes at apparently irregular intervals (Clayton 1996). A cloud of carbon-rich dust forms along the line of sight to the RCB star eclipsing the photosphere, causing a severe drop in its brightness and the appearance of a rich emission-line spectrum.  As the dust cloud disperses, the star returns to maximum light. RCB stars have a wide range of temperatures but they can be divided simply into three groups, cool ($\\sim$5000 K), warm ($\\sim$7000 K) and hot ($\\sim$20,000 K).  Typical representatives of these groups are S Apodis, R Coronae Borealis and V348 Sagittarii, respectively.  Most RCB stars fall in the warm category. Hot RCB stars are quite rare with only 3 examples known. The typical warm RCB spectrum at maximum light looks like an F or G supergiant with a few important differences: the Balmer lines are very weak or absent; the spectrum contains many lines of neutral carbon, and bands of $C_2$ and CN. The cool RCB-type spectrum resembles the warm type but with much stronger molecular absorption bands. The hot RCB stars show similar lightcurve behavior to the cooler stars but their spectra are very different (Pollacco \\& Hill 1991).  The spectrum of V348 Sgr, the best studied hot-type star, shows strong emission lines of C II and He I as well as the Balmer lines, Ne I and various forbidden lines (Dahari \\& Osterbrock 1984). Most RCB stars in all three categories show excesses at near-IR and IRAS wavelengths.  The RCB Stars are very rare either because they form only in unusual circumstances or because they are a brief episode in stellar evolution. Only about 30 RCB stars are known in the Galaxy, and until now only 3 in the LMC despite their high intrinsic luminosities. Their evolutionary history remains very uncertain. Two major evolutionary models have been suggested for the origin of RCB stars, the Double Degenerate and the Final Helium Shell Flash conjectures (Sch\\\"{o}nberner 1986; Renzini 1990; Iben, Tutukov, \\& Yungelson  1996). Both involve expanding white dwarfs to the supergiant sizes assumed for RCB stars. A third model suggests that RCB stars are binaries in the second common envelope phase with a low mass companion orbiting inside the envelope (Whitney, Soker, \\& Clayton 1991). Recently, Iben et al. (1996) added the merger of a neutron star and a helium-rich star to the list of possible RCB star precursors.  An important input parameter to these models is stellar luminosity. This parameter can only be estimated when the distance to a star is known. However, there is no reliable distance estimate to any Galactic RCB star. Since they are not ``normal\" stars, their distances can only be estimated if they are associated with an object at a known distance or through other indirect methods. Previous estimates of Galactic RCB star luminosities are summarized in Table 1. In addition, a star in the cluster NGC 6231 was initially identified as an RCB star but turned out to be a normal reddened star (Bessel et al. 1970; Herbig 1972). In a similar manner to Doroshenko et al. (1978), Rosenbush (1981, 1982, 1989, 1995), using estimates of reddening along sightlines to RCB stars and the structure of the interstellar medium,  finds a wide range of absolute magnitudes,  $M_V$ =  -5 to +2.5. The RCB star,  V482 Cygni was identified with a quadruple star system containing a K5 III star based on proximity on the sky implying $M_V$ = -2.8 (Gaustad et al. 1988). This association was refuted by Rao \\& Lambert (1993) who find that V482 Cyg has significantly different radial velocities and interstellar columns than the K5 III star.  They estimate a larger distance consistent with an $M_V \\sim$ -4.6. Other RCB stars, including RY Sagittarii,  have close companions although none have been shown to be physical pairs (Andrews et al. 1967, Feast 1969; Milone 1995). Estimates of Galactic RCB star luminosities differ by factors of up to $10^3$.  Due to the absence of reliable distance estimates for the Galactic stars, the LMC RCB stars play a pivotal role in RCB star research. Absolute luminosities can be derived from the LMC RCB stars which are at a known distance. Using their apparent magnitudes and the known distance of the LMC (m-M = 18.6), an absolute magnitude of $M_V\\sim-4~to -5$ is derived.  However, this is based on only 3 stars (Feast 1972). This result, that RCB stars have supergiant size and luminosity, puts strong constraints on the evolutionary models outlined above.  One of the dividends from the search for Massive Compact Halo Objects (MACHO's) towards the LMC is the discovery of a large number of new variable stars.  Over 40,000 variables have been discovered so far (Cook et al. 1995). RCB candidates have been selected on the basis of their lightcurve behavior and confirmed spectroscopically. When only fragmentary lightcurve data are available, RCB stars may be confused with symbiotic, cataclysmic or semi-regular variables (Lawson \\& Cottrell 1990). ",
        "conclusions": "Figure 3 shows the lightcurve of MACHO*05:33:49.1-70:13:22.  Almost the entire 1200~d of coverage involves one deep decline of $\\Delta V \\ge$ 4 mag. The decline begins around JD 2448925 with a steep drop of $\\sim$4 mag in a few days. There is a slight recovery around JD 249000 followed by another fading and then a slow recovery to maximum light. This lightcurve is typical of an RCB star decline.  The final recovery to maximum light can be gradual as the dust cloud disperses and sometimes takes several years (e.g. Alexander et al. 1972). Figure 4 shows that lightcurve of  MACHO*05:32:13.3-69:55:59.  It is quite active showing 3 major declines around JD 2448900, 2449325 and 2449650.  There is a great variation in decline activity from star to star and also from time to time for an individual star (Clayton 1996).  The Galactic RCB star, V854 Cen, has shown similar activity to MACHO*05:32:13.3-69:55:59 in the last few years (Lawson et al. 1992). Both MACHO*05:33:49.1-70:13:22 and MACHO*05:32:13.3-69:55:59 show small amplitude variations at maximum light similar to other RCB stars.  Figures 3 and 4 also show the $(V-R)_{KC}$ color behavior. MACHO*05:33:49.1-70:13:22 becomes redder at the beginning of the decline and returns to its normal color as it returns to maximum light.  Early and late in a decline the star is reddened by a dust cloud which is not optically thick.  Deep in a decline the cloud may be optically thick so reddening may not be seen. MACHO*05:32:13.3-69:55:59 shows $(V-R)_{KC}$ colors which become bluer at the onset of the decline and then return to normal as the star recovers to maximum.  RCB stars experience red and blue declines (Cottrell, Lawson, \\& Buchhorn 1990).  The colors can vary from decline to decline depending on how much of the photosphere and the emission-line regions are obscured by dust, by the optical depth of the dust, and by the relative strength of the emission lines. Sometimes very early in a decline, the colors are unchanged at first and then become bluer. This can occur if the forming cloud is smaller than the photosphere and some unreddened starlight is still visible (Cottrell et al. 1990). Red declines occur if the forming cloud covers the entire photosphere. It is notable that MACHO*05:32:13.3-69:55:59 has had three blue declines in a row. Unfortunately, the data for Galactic RCB stars are sparse so the relative frequency of red and blue declines is not known. For the LMC RCB stars, another possibility is confusion with a blue star in the aperture although no such star is visible on the CCD frames.  The spectrum of MACHO*05:33:49.1-70:13:22, shown in Figure 7, is very similar to that of the hot RCB star, V348 Sgr (Dahari \\& Osterbrock 1984; Leuenhagen \\& Hamann 1994). The spectrum of V348 Sgr is classed as WC11 since it shows emission at C II but not C III (Leuenhagen \\& Hamann 1994). In addition, its lightcurve, IR excess and hydrogen deficiency distinguish it as an RCB star. The  MACHO*05:33:49.1-70:13:22 spectrum shows strong C II emission at 3919, 4267 and 4735 to 4747 \\AA. There is also possible C II emission seen at 4618 to 4630 \\AA~and near 4861 \\AA~blended with $H\\beta$. In addition, MACHO*05:33:49.1-70:13:22 seems to have been detected in the IRAS Serendipitous Survey (Kleinmann et al. 1986) at a level of 0.1 Jy at 12 \\micron. This is similar to the flux detected for HV 12842 and both are consistent with an extrapolation of flux levels measured for Galactic RCB stars. The spectrum, lightcurve and IR excess of MACHO*05:33:49.1-70:13:22 indicate that it is a hot RCB star, only the fourth known and the first discovered outside the Galaxy.  In addition, although the spectrum is low resolution the C II lines show a redshift of 259$\\pm$31 $km~s^{-1}$ which is appropriate for LMC membership.  The spectrum of MACHO*05:32:13.3-69:55:59 is shown in Figure 8.  This spectrum is a stereotypical cool RCB spectrum similar to S Aps and V517 Ophiuchi (Kilkenny et al. 1992). The spectrum shows deep Swan bands of $C_2$ with bandheads at 4382, 4737, and 5165 \\AA, and violet bands of CN with bandheads at 3883, 4216 and 4606 \\AA.  In addition, the spectrum is distinguished as an RCB star by the weak or absent Balmer lines, G-band and $^{12}C^{13}C$ bands (Lloyd Evans, Kilkenny, \\& van Wyk 1991). This is a result of severe hydrogen deficiency and a lack of $^{13}C$ typically seen in RCB stars (Clayton 1996).  MACHO*05:32:13.3-69:55:59 shows a $C_2$ (1,0) 4737 \\AA~band that dips about 80\\% below the continuum. HV 5637 shows a 72\\% depression and S Aps about 79\\% (Feast 1972). W Men and HV 12842 show much smaller dips of about 10\\% much like R CrB and RY Sgr. The (V-R$)_{KC}$ colors of S Aps and MACHO*05:32:13.3-69:55:59 are similar. The spectrum, color, and lightcurve of MACHO*05:32:13.3-69:55:59 show that it is a cool RCB star. The measured positions of the molecular bandheads show a redshift of 269$\\pm$21 $km~s^{-1}$ which is appropriate for LMC membership.  As mentioned in the introduction, the distance scale for the RCB stars depends entirely on the LMC members.  Using the recent photometry of Lawson et al. (1990), we find $V_{max}$ = 14.8, 13.8, and 13.7 for HV 5637, W Men and HV 12842, respectively. On the basis of this small sample, RCB stars are supergiants with a range of $\\sim$1 magnitude in absolute luminosity in the V-band.  Since HV 5637 is cooler than W Men and HV 12842, Feast (1979) suggested a relationship between temperature and luminosity. The observed range in $V_{max}$ is likely to be intrinsic rather than due to reddening differences. Estimates of foreground (circumstellar and interstellar) reddening are somewhat uncertain since RCB star colors are not known a priori. Older studies of the Galactic foreground find a fairly uniform screen of dust of E(B-V) $\\sim0.04 - 0.07$ (e.g. McNamara \\& Feltz 1980). Schwering \\& Israel (1991) re-examined the foreground reddening by comparing H I and IR observations. They find a small but significant variation in foreground reddening across the face of the LMC from E(B-V) = 0.07 to 0.17 mag. Any constant component of circumstellar dust around RCB stars is small. Using the observed B-V and the estimated $T_{eff}$ listed in Table 2, we calculate E(B-V) $\\sim$ 0.1-0.2 for HV 5637, W Men and HV 12842 (Johnson 1966). Goldsmith et al. (1990) estimate 0.08 and 0.10 for the circumstellar and interstellar components of the E(B-V) towards W Men.  Only W Men has a measured $(V-R)_{KC}$ (Eggen 1970). Converting this to Johnson $(V-R)_J$, a similar E(B-V) is obtained assuming a normal extinction curve with $R_v$=3.1 (Cousins 1980).  Therefore, the new stars presented here are very important.  The measured $V_{max}$ for these stars are 16.1 and 16.3, fainter by about 2 mag than the three known RCB stars.  Do these values really represent unreddened $V_{max}$? These stars have been followed photometrically for only about 3 years so it is possible that they have never fully recovered to maximum light. RCB stars go through very active phases where they are in decline for years (e.g. Mattei, Waagen \\& Foster 1991).  However, to remain in decline, dust must form regularly to compensate for the dispersal of previous dust clouds. Flat decline-lightcurve behavior only occurs deep in a decline when all direct starlight is extinguished and only scattered light is seen.  This kind of behavior is not seen when $\\Delta V \\sim$ 2 mag so it can't account for the observed lightcurve. Therefore, the recent lightcurve behavior of both stars, seen in Figures 3 and 4, is consistent with being at or near maximum light.  MACHO*05:33:49.1-70:13:22 seems to be approaching maximum light after a long decline.  This lightcurve shape is seen in many other RCB declines.  The MACHO*05:32:13.3-69:55:59 lightcurve is more complicated and the star is definitely in an active phase.  However, in the last 200 d during which time the spectrum was obtained, the star has been very constant with $\\Delta V \\sim$ 0.25 mag. There is no evidence in this time period for either dust formation or dispersal.  Another possibility is large interstellar extinction toward these stars.  The reddening can be estimated from the measured $(V-R)_{KC}$ colors which were converted to $(V-R)_J$ (Cousins 1980).  Assuming normal supergiant $(V-R)_J$ colors then E(B-V) $\\sim$0.3 for MACHO*05:33:49.1-70:13:22 and  $\\sim$0.4 for MACHO*05:32:13.3-69:55:59. This is slightly higher than the estimate of $\\sim$0.1-0.2 for the other LMC RCB stars. These values of E(B-V) are seen for many other LMC stars although most lie in or near the 30 Dor region.  Neither of the new LMC RCB stars lies in a visibly dusty region of the LMC. Estimates of $V_o$ and $M_V$ for all five stars are given in Table 2 using the estimated values of E(B-V).  The new stars fall $\\gtrsim$ 0.5 mag below the range of $M_V$ = -4.1 to -5.5 for the previously known LMC RCB stars. These estimates are somewhat uncertain because RCB colors may not be the same as normal supergiants. For instance, S Aps has the same maximum-light $(V-R)_{KC}$ color as MACHO*05:32:13.3-69:55:59 and a similar $T_{eff}$ (Lawson et al. 1990) yet its B-V colors imply a smaller reddening of E(B-V) $\\sim$0.1.  If the reddening of MACHO*05:32:13.3-69:55:59 is 0.1 then it has an even fainter $M_V \\sim$-2.6. Another indication of the uncertain reddening correction can be seen in Figures 5 and 6. The stars in the MACHO*05:33:49.1-70:13:22 field seem to be $\\sim$ 0.1 mag redder than the stars in the MACHO*05:32:13.3-69:55:59 field. Therefore, taking into account the uncertainties in the intrinsic and measured colors, and stellar $T_{eff}$, the uncertainty in E(B-V) is $\\sim$0.1-0.2.  So the uncertainty in $M_V$ is $\\sim$0.3-0.6.  A slightly smaller range of values is found when using $M_{Bol}$. The assumed effective temperatures are 5000 and 7000 K for cool and warm RCB stars, respectively. Using the values of $M_V$ calculated above and Bolometric corrections for normal supergiants, we get the values listed in Table 2 for $M_{Bol}$. A value of 20,000 K has been estimated for V348 Sgr (Sch\\\"{o}nberner 1986).  This value has been applied to MACHO*05:33:49.1-70:13:22. It's $M_{Bol}$ lies in the range of the warm RCB stars. The two cool RCB stars have the lowest Bolometric luminosities. The question of whether LMC and Galactic RCB stars are intrinsically different remains open.  There are abundance differences but they may lie within the range of variations seen in the Galactic RCB stars.  The slightly bluer UBV colors of the LMC RCB stars may also be an indication of abundance differences."
    },
    "9604/astro-ph9604088_arXiv.txt": {
        "abstract": "We examine models in which jets are responsible for the formation of the emission-line spiral structure in \\iras.  The kiloparsec-scale radio lobes in this active galaxy appear to be related to its extended emission-line spiral structure.  The radio structure consists mainly of extended symmetrically bent, FR\\,I-type lobes, which follow the emission-line spiral structure at their inner edge. In the central region of the galaxy a double radio source is observed with a separation of approximately 1\\,arcsec between its components, which are extremely well aligned with the hotspot from which the southern lobe expands outwards. Hill \\etal\\ (1988) suggested a model for the emission-line spiral structure invoking compressed interstellar matter, which is dragged away from the original jet path by the rotating ambient medium. From consideration of the propagation speed of the jets and the transverse ram pressure exerted by the rotating environment, we exclude this scenario as a possible origin of the spiral structure. We favour a model in which the jets themselves are bent by the rotating interstellar medium and possibly follows the emission-line spiral arms. We present fits of the model to the observed optical spiral structure. High sensitivity radio observations will be required to decide on the nature of the peculiar spiral structure in IRAS\\,04210+0400. ",
        "introduction": "The nature of the spiral structure of disk galaxies has been a widely discussed issue in extragalactic astronomy since their discovery. Several mechanisms have been put forward in  explanation, but to date no general consensus has been reached about a single process responsible for the spiral structure of galaxies. The most successful theories employ density waves (Lin \\& Shu, 1964) or galaxy interactions to produce the spiral arms (e.g. Toomre \\& Toomre, 1972). Another possibility is stochastic self propagating star formation considered by Gerola and Seiden (1978). Most probably, different mechanisms are at work in different galaxies, and these theories are complementary rather than competing.  On occasion, ejection phenomena have been suggested to be responsible for spiral arms in galaxies which did not fit into the conventional schemes. Van der Kruit \\etal\\ (1972) suggested that the anomalous radio and emission-line arms in NGC\\,4258 originated from an essentially instantaneous ejection of radio-emitting material from the centre of the galaxy in opposite directions. In their model, the observed radio arms are a transient phenomenon and represent the current position of these plasmons, which were ejected with a wide range of speeds. The trajectories of the individual parcels is determined by gravitational forces and ram pressure. Later on, several other similar models for the jets in NGC\\,4258 have been suggested which are consistent with the observation that the anomalous arms are bent in the same direction as the trailing spiral arms (e.g. Martin \\etal, 1989, and references therein).  Wilson \\& Ulvestad (1982) described a steady state model in which a two-sided jet propagates roughly in the rotational plane of a galaxy. The jet is then bent by the ram pressure exerted transverse to the jet by the rotating interstellar medium. They applied this model to the S-shaped radio structures found on a sub-kpc scale in the Seyfert galaxies NGC\\,1068, NGC\\,4151, and the radio galaxy 3C\\,293. The result of this model are leading spiral structures in the galaxy. Other mechanisms to obtain S-shaped jet pairs are pressure gradients in the hot gaseous halo of elliptical galaxies (Smith \\& Norman 1981) and precession of the jets (e.g. Gower \\etal\\ 1982).  Imaging, both ground based (Hill \\etal\\ 1988; Steffen \\etal\\ 1996a, 1996b) and by the Hubble Space Telescope (HST; Capetti \\etal\\ 1996) suggests that the spiral structure of the peculiar active galaxy \\iras\\ is due to a narrow ridge of emission line filaments. Its redshift is 0.0462 and was classified as a Seyfert~2 (Beichman \\etal, 1985), but the classification has been questioned by Hill \\etal\\ (1988). It shows an unsually extended and complex radio structure for this type of galaxies, with a total power of $2.4\\cdot10^{23}{\\,\\rm W\\,Hz^{-1}}$ at 20\\,cm. It has extended radio lobes, which are related to the optical spiral arms. Based on their observations, Hill \\etal\\ (1988) first suggested that the spiral structure could be due to the remnant of the interaction of the jets with the ambient medium. In this scenario the emission line filaments are produced during the passage of the jet through the interstellar medium. It is then carried away from its original location by the rotation of the galaxy forming a spiral structure.  In the present paper we analyze the suggestion by Hill \\etal\\ (1988) that the spiral structure in \\iras\\ is due to the interaction of the jets with the ambient medium. In Section \\ref{hill_mod.sec} we consider the model proposed by Hill \\etal\\ (1988). In Section \\ref{bent_jets.sec} we analyze the model proposed by Wilson \\& Ulvestad (1982) for other Seyfert galaxies. Section \\ref{discussion.sec} contains the discussion of our results and our conclusions are summarized in Section \\ref{conclusions.sec}. ",
        "conclusions": "\\label{conclusions.sec}  We have evaluated two different candidate models for the formation of the emission-line spiral structure in \\iras\\ involving jets ejected from the galactic nucleus. We exclude the suggestion by Hill \\etal\\ (1988) of straight jets which interacted with the ISM as a possible explanation for the observations.  We find that the second model of two jets bent by the rotating ambient medium is a better candidate.  Fitting this model to the optical spiral, we predict that the northern and the southern radio jets should be oriented at position angles PA$_n \\approx 45\\deg$ and PA$_s \\approx 135\\deg$ with an estimated error of $10\\deg$.  Detection of the jets on the intermediate scale between the inner and outer radio structures will be necessary to discriminate between the current models. High sensitivity radio observations will be required to search for the jets. \\\\"
    },
    "9604/astro-ph9604041_arXiv.txt": {
        "abstract": "The velocity dispersion of galaxies on scales of $r\\sim1h^{-1}$ Mpc, $\\sigma_{12}(r)$, may be estimated from the anisotropy of the galaxy-galaxy correlation function in redshift space. We present a reanalysis of the CfA1 survey, correct an error in the original analysis of Davis and Peebles (1983), and find that $\\sigma_{12}(r)$ is extremely sensitive to the details of how corrections for infall into the Virgo cluster are applied. We conclude that a robust value of $\\sigma_{12}$ cannot be obtained from this survey. We also discuss results from other redshift surveys, including the effect of removing clusters. ",
        "introduction": "Davis and Peebles (1983, hereafter DP83) calculated the velocity dispersion of galaxies, $\\sigma_{12}(r)$, on scales of $r \\sim 1 - 5 h^{-1} Mpc$ for the CfA1 redshift survey, a survey containing 1840 redshifts covering 1.83 steradians in the North galactic hemisphere (Huchra et al.  1983). Their result, $\\sigma_{12}(1) = 340\\pm40$ km/s on scales of $1 h^{-1}$ Mpc, became the standard by which N-body simulations were judged for perhaps ten years, and a primary argument against the Cold Dark Matter scenario for structure formation with the assumption that galaxies trace the mass fluctuations in an unbiased way, which yields much higher velocities on this scale (e.g., Davis et al. 1985, Gelb \\& Bertschinger 1994). The same calculation was done on the Southern Sky Redshift Survey (SSRS1, da Costa et al. 1991), with results of $\\sigma_{12}(1)\\sim 300$ km/s (Davis 1988), in apparent agreement with the CfA1 result. It is only recently that there have been attempts to reproduce the results of DP83 (Mo, Jing, \\& Borner 1993, Zurek et al. 1994), and to perform this analysis on new, larger redshift surveys (Fisher et al. 1994a, 1994b, Marzke et al. 1995, Guzzo et al. 1995). It is now apparent that there is a large variation in $\\sigma_{12}$ between different surveys (see Table 1). In addition, different workers (Mo et al. 1993, Zurek et al. 1994) obtain very different results (from DP83 and from each other) when they analyze the CfA1 survey. In this paper we clarify some details of the original calculation of DP83 which were not spelled out in the original paper, and present a reanalysis which shows why the results are so unstable. We also reproduce the earlier results for SSRS1 (Davis 1988), and investigate how removing clusters affects $\\sigma_{12}$. ",
        "conclusions": "We hope that this paper has removed the confusion regarding the value of the velocity dispersion in the CfA1 survey. We have shown that the value of the velocity dispersion is extremely sensitive to the way in which corrections for infall into the Virgo cluster are applied. This is because the Virgo cluster is in the foreground of the CfA1 survey and contains many intrinsically faint galaxies in a thermally hot region.  Increasing the distance to these galaxies by a small amount results in inclusion of fewer of these galaxies in the volume limited sample. Because $\\sigma_{12}$ is pair weighted, including or leaving out even $\\sim10$ galaxies from the Virgo cluster can change $\\sigma_{12}$ by $\\sim 100-200$ km/s. In addition, including corrections for cluster infall in the calculation of the correlation function in redshift space, $\\xi(r_{p}, \\pi)$, effectively removes part of the ``finger of god'' and reduces the velocity dispersion of the cluster.  Once again the pair-weighted nature of the statistic means that this will result in a significant reduction in the overall value of $\\sigma_{12}$ for the sample.  However, no infall corrections were used in recent calculations of $\\sigma_{12}$ for other redshift surveys and yet a wide range of values for $\\sigma_{12}$ is obtained for different surveys. We have argued that this is because $\\sigma_{12}$ is extremely sensitive to clusters, and existing redshift surveys do not sample a large enough volume of space to represent a fair sample of these relatively rare objects. One approach to solving this problem is to remove the clusters from the sample before calculating $\\sigma_{12}$; however, this reduces the ability of the statistic to discriminate between cosmological models. In addition, our work suggests that the results are likely to be sensitive to the details of how the clusters are identified and removed. Analysis of larger volume redshift surveys will be necessary in order to obtain a robust value of $\\sigma_{12}$ which is useful in discriminating between cosmological models or for estimating $\\Omega_{0}$. In the meantime, modified velocity statistics such as the galaxy-weighted velocity dispersion (Miller et al. 1996), a density dependent version of the pairwise velocity dispersion (Strauss 1996), or the median velocity of groups (Nolthenius, Klypin \\& Primack 1996) have been designed to be less sensitive to clusters and may be promising alternatives.  \\clearpage  \\begin{center} Acknowledgements \\end{center}"
    },
    "9604/astro-ph9604107_arXiv.txt": {
        "abstract": "We have made an attempt to compile all currently available data on optically identified QSO absorber systems (Lindner {\\it et al.} 1996) to establish the status quo of absorber galaxy data as a basis for the investigation of galaxy evolution. We present a first comparison with results from our galaxy evolutionary synthesis models to demonstrate the potential power of this kind of approach and to guide future observations to identify absorber galaxies. ",
        "introduction": "Our chemical and spectral synthesis model (Fritze -- von Alvensleben {\\it et al.} 1994) describes the evolution of various types of galaxies (E to Sd) with appropriate star formation histories and supplies us with time dependent values of luminosities from UV to NIR, colors and metallicities of model galaxies. Adopting any cosmological model characterized by $H_0$, $\\Omega_0$, $\\Lambda_0$ and the redshift of galaxy formation z$_{\\small form}$, these results can be transformed into redshift dependent quantities. Apparent R magnitudes as a function of redshift z calculated with this model are plotted in Fig.~1.  Pioneering work in the optical identification of QSO absorption systems was done by Bergeron \\& \\Boisse\\ (1991) and up to now, there are more than a dozen of publications reporting on photometric data, equivalent widths and impact parameters for absorbing galaxies. All available data on apparent R magnitudes are plotted in Fig.~1. Different symbols correspond to different authors. Absorbing galaxies with spectroscopically confirmed redshifts (i.e. \\zgal\\ = \\zabs) are marked by filled symbols, whereas open symbols indicate absorber candidates.  Accounting for the luminosity ranges from the luminosity functions of the various galaxy types (cf. $2 \\sigma_R$ bars in the lower right corner of Fig.~1) we can state that virtually all observational data points fall between the curves for E-- and Sd--models and, accordingly, we can establish global agreement between our galaxy evolution models and observational data up to $z \\approx 2$.  Most of the absorber galaxies appear to be early through intermediate type spirals (Sa--Sbc) but many ellipticals and some late type spirals seem to be present, too. The presence of intermediate and late type galaxies among QSO absorbers would imply that a considerable fraction of these galaxies do have extended gaseous halos metal rich enough to cause detectable absorption.  \\begin{figure} \\epsfysize=14.5cm \\vskip -6.5truecm {\\epsffile{QSOabs_fig.ps}} \\vskip -0.3truecm \\caption{Apparent R magnitude as a function of redshift for observational data and results from our galaxy evolution model using \\zform\\ $= 5$ and cosmological parameters \\H0{50}, $\\Omega_0 = 1$, and $\\Lambda_0 = 0$.} \\end{figure}  Varying the cosmological parameters we find that e.g. model galaxies for \\H0{50} and $\\Omega_0 = 0.1$ are much too faint as compared to observations leading us to exclude this compination of cosmological parameters. ",
        "conclusions": ""
    },
    "9604/astro-ph9604113_arXiv.txt": {
        "abstract": "We report results based on 35 new spectroscopic redshifts obtained with the Keck Telescope for field galaxies that also have photometry and morphology from survey images taken by the refurbished Hubble Space Telescope.  A sample of 24 redshifts for galaxies fainter than {\\it I} $= 22$ has a median redshift of $z \\sim 0.81$.  This result is inconsistent with the lower median redshift of $z \\sim 0.6$ predicted by the ``maximal merger models'' of Carlberg (1996), which otherwise fit existing data.  The data match an extrapolation of the Canada France Redshift Survey (CFRS), as well as  predictions of certain mild luminosity-evolution models.  Nearly half of the redshifts lie in two structures at $z \\simeq 0.81$ and $z \\simeq 1.0$, showing the presence of high density concentrations spanning scales of $\\sim 1 h^{-1}$ Mpc, i.e., the size of groups.  We find emission lines or the presence of possible  neighbors in 7 of 9 otherwise luminous galaxies with red central regions at redshifts beyond $z \\sim 0.7$. We also note a diversity of morphological types among blue galaxies at $z \\sim 1$, including small compact galaxies, ``chains,'' and ``blue nucleated galaxies.'' These morphologies are found among local, but generally less luminous, galaxies. Distant blue galaxies also include apparently normal late-type spirals.  These findings could imply modest bursts of star formation caused by mergers or interactions of small, gas-rich galaxies with each other or with larger, well-formed galaxies. This first glimpse of very faint $z \\sim 1$ field galaxies of diverse colors and morphologies suggests that a mixture of physical processes is at work in the formation and evolution of faint field galaxies. ",
        "introduction": "To date, we have yet to understand the origin of the large density of very faint, blue field galaxies despite enormous observational progress (e.g., \\cite{Lil95a}).  Deep  images to $I \\sim 24$ or fainter with the refurbished Hubble Space Telescope (HST)  show a predominance of late-type or unusual galaxy morphologies (e.g., \\cite{Gri94,For94,Gla95b,Dri95}), as well as small sizes that suggest a large fraction of low-luminosity dwarf galaxies (\\cite{Dri95}).  The  redshifts of individual  galaxies observed with HST remain largely unknown.  Existing redshift samples include 17 galaxies to $I \\leq 22$ from Forbes \\etal (1996), 32 to $I \\leq 22$ from Schade \\etal (1995), and 34 to $B \\leq 24.5$ (roughly $I \\leq 22.5$) from Cowie \\etal (1995) but do not yet probe fainter than  $I \\sim 22$.  The redshift distribution is, however, well established statistically to $I \\leq 22$ from the CFRS (Lilly \\etal 1995b).  Redshifts are crucial for determining the intrinsic properties of galaxies visible in deep HST images, and good spectra can also yield rotation velocities, velocity dispersions, stellar population age indices, and metallicities.  A program to obtain deep spectra using the 10-m Keck Telescope is now underway as a new initiative called the Deep Extragalactic Evolutionary Probe, or DEEP (\\cite{Mou93,Koo95a}).  This paper reports on initial results of a new DEEP survey that utilizes redshifts from the Keck Telescope, supplemented by photometry, colors, and morphologies from  images taken by Groth \\etal (1996) with the refurbished HST.  The target sample contains 230 galaxies, but poor weather permitted only 18\\% to be observed during the run.  The acquired redshift sample of 35 galaxies is, however, generally representative of the target sample. The magnitudes are so faint (median $I \\gtrsim 22$; median $B \\sim 25$) and the redshifts so high (median $z \\sim 0.8$), that the current data provide a first glimpse of the nature of faint, distant field galaxies of $\\sim L^*$ luminosities at an epoch of roughly half the Hubble age.  \\footnote{We adopt $h = 0.75$, $q_o = 0$, and $\\Lambda = 0$, i.e. $\\Omega = 0$. At redshift $z \\sim 1$, $I \\sim 22.6$ for a blue ($B-V \\sim 0.65$) galaxy of $M_B \\sim -20.4$ ($M^*$), the lookback time is about 6.5 Gyr for a Hubble age of 13 Gyr, and 1 arcsec corresponds to 7 kpc.} ",
        "conclusions": "\\subsection{Redshift Distribution}  What cosmological implications can we extract from this  faint sample of galaxies? Given the relatively small numbers, median redshifts are a good starting point for comparison to theory. First, the 100\\% complete sample of 9 galaxies  with $19.5 < I < 22$ has a median redshift of $z = 0.81$, with a possible range from 0.37 to 1.00 at the 95\\% confidence level; this is consistent with the median of 0.56 for CFRS based on a much larger sample. Second, the $I \\ge 22$ sample with 24 galaxies has the same median of 0.81 but a tighter range of $z \\sim 0.81$ to 1.00 at the 95\\% confidence level. This high median is an important finding, since Songaila \\etal (1994) found the median to remain at $z \\lesssim 0.6$ for galaxies fainter than $K \\sim 18$ (roughly $I \\sim 20.5$), as predicted by some merger models.  For example, the ``maximal merger model'' of Carlberg (1996), which matches observations to $I \\sim 22$, predicts that the median will stabilize at $z \\sim 0.6$ from $I \\sim 20$ all the way to $I \\sim 28$, contrary to what we find in this sample at $I > 22$.  The high median is, however, consistent with at least two other scenarios.  First, with no further evolution than that already found in the CFRS survey, Lilly \\etal (1995c) predict that the median should continue to rise to $z \\sim 0.9$ for $22 < I < 23$ and to $z \\gtrsim 1.0$ for $23 < I < 24$. Second, the predictions of models that do include mild luminosity evolution but no mergers, such as those of Gronwall \\& Koo (1995), are virtually identical  to those predicted by CFRS.  At this time, we have no basis to exclude other models, such as those of Cole \\etal (1994), that analytically track various other physical processes such as supernova gas removal and star-formation, as well as merging of dark matter halos.  The width of the distribution is also a potentially powerful discriminant. For example, the fiducial bursting dwarf model of Babul \\& Ferguson (1996) predicts a narrow redshift distribution peaked at $z \\la 1$ for a $B < 26$ redshift sample, with a cutoff in numbers at $z = 1$, chosen by them by fiat, when dwarfs first form.  Our data, with the bulk of the redshifts $0.7 \\la z \\la 1$, lend some support for this picture.  On the other hand, several very blue galaxies are seen beyond $z = 1$, while those very near $z = 1$ span a wide range of colors rather than being dominated by the very blue colors expected from an ongoing burst of star formation (see Figure 7).  Moreover, the diverse morphologies and diffuseness of many of the blue galaxies, discussed below, may be difficult to explain.  More faint redshifts with stricter selection criteria will be needed to provide more definitive tests of the bursting model.  Redshifts also probe clustering. CFRS already paved the way with the discovery of 12 redshifts in a 0.016 redshift interval, indicating a five to ten times overdensity of galaxies at $z = 0.985$ (Le F\\`evre \\etal 1994).  This possible supercluster structure was found in a 10'x10' field that coincidentally overlaps our field.  As seen in Figure~6a, we find a density enhancement at $z = 0.995 \\pm 0.004$ containing five galaxies, which may be the same structure, plus another stronger one at $z = 0.811 \\pm 0.003$ with ten galaxies (standard deviations are measured within $\\delta z = \\pm 0.02$ of the peak).  Since we see no visual evidence for any rich cluster of galaxies in our field and the velocity dispersion is relatively low ($\\sim 450$~\\kms), the $z = 0.81$ enhancement is probably a rich group.  \\subsection {Morphologies, Colors, and Galaxy Evolution}  We close with some discussion and speculations based on  HST morphologies and colors.  Figure~1 shows 9 galaxies with very red colors within their central one arcsec diameters.  The galaxies in the top two rows of Figure~1 appear morphologically to be elliptical or S0 galaxies (though Phillips \\etal [1996] find that 073-2675 is better fit with an exponential, rather than an $r^{1/4}$ profile).  Combined with very red colors, these galaxies indicate the presence of luminous ellipticals {in the field} that are already quite old by $z \\sim 1$.  This presumes that the red colors are not due to dust extinction, which remains to be determined.  Based on the $z = 2$ curve in Figure~7, the last significant star formation event presumably occurred at redshifts $z \\sim 2$ or greater.  Such very red galaxies have also been found among distant radio galaxies (e.g., McCarthy 1993) and cluster galaxies (e.g., Dickinson 1996).  Yet several of these galaxies are not quite ``normal.'' Galaxy 093-2470 appears to have four very blue compact objects imbedded in a symmetrical pattern within its halo, a pattern proposed to form a ``quad-lens'' system (Ratnatunga \\etal 1995, Broadhurst \\etal 1996, Crampton \\etal 1996).  Although the ``maximal merger model'' of Carlberg (1996) appears to be invalidated by the new redshifts, we do find some evidence for minor  mergers among several of these distant red galaxies.  Except for 073-2675, the remaining five galaxies in the first two rows of Figure~1 show close projected neighbors or tidal and merger ``tails,'' and some of this neighboring material is blue. Perhaps we are watching the infall of dwarf galaxies, some possibly quite gas-rich, if the blue colors are due to active star formation.  The implication would be that elliptical galaxies form early and yet can be built up by minor mergers over a much longer time period.  This scenario could reconcile  the apparent lack of luminosity or density evolution among red galaxies seen in the CFRS (Lilly \\etal 1995c) or the Mg II absorber sample (Steidel, Dickinson, \\& Persson 1994) with the presence of well-formed, red field galaxies at an early epoch.  The last row of Figure~1 presents three very red {\\it disk} galaxies.  With a peak-to-peak rotation curve velocity of $\\sim$600~\\kms ~(Vogt \\etal 1996), the conspicuous edge-on system (104-4024) demonstrates that some massive, thin-disk, dusty systems already exist at $z \\sim 0.8$, roughly half a Hubble age ago.  This single object, evidently similar to our Milky Way, supports what we believe to be the early formation of our own old disk. There are hints of very faint blue satellites that might eventually settle into this distant spiral, an evolutionary path that has also been proposed for the Milky Way (Majewski 1993).  Another disk-like system, 094-2210, shows a very red, bulge-like core surrounded by numerous blue ``blobs'' or arms.  If this galaxy is a disk or proto-disk system with little dust extinction, it suggests that some bulges are already quite old by $z \\sim 1$, with their most recent star formation occurring perhaps at $z \\ga 2$ (see Figure~7).  Object 103-2074 appears to be another disky system of early type (S0 or Sa) with a very red bulge.  Figure~2 shows two sets of galaxies: the bluer, high-redshift galaxies and those without redshifts. Those with measured redshifts $z \\gtrsim 0.75$ are organized into 7 rows divided into morphological groups.  We argue that all six of the galaxies without redshifts are likely to be at high redshifts $z \\ge 0.75$.  Since blue galaxies normally have strong emission lines of [OII] $\\lambda$ 3727 that should be detectable to $z \\approx 1.4$ (or especially of H${\\alpha}$ for redshifts lower than $z \\sim 0.4$), we argue that the five blue galaxies in our Keck sample without measured redshifts are probably at higher redshifts.  In support of this conjecture, we note that the bluest galaxy in our entire Keck sample (source 084-1720) has two definite absorption lines separated by $\\sim$35\\AA\\ and appearing around 6770\\AA\\ (see Figure 5); we tentatively identify this doublet as Fe II $\\lambda$2587 and $\\lambda$2600 that would yield a redshift of $z = 1.60$.  One relatively bright, very red galaxy (source 083-3138) also has no redshift, but based on the similarity of its color to the galaxies in Figure 1 with redshifts (see Figure 7), we infer that it is likely to be at high redshift.  The targets without redshifts are all presented in the last two rows of Figure~2.  The large diversity of morphologies displayed in Figure~2 is striking, ranging from compact galaxies shown in the first and second rows, normal disk-like systems in the third and fourth rows, to irregular, late-type systems with multiple blue star-formation sites in the fifth and sixth rows. This last group could  be related to a proposed new class of ``blue-nucleated'' galaxies (Schade \\etal 1995).  The seventh row shows galaxies that are perhaps more elongated versions of the above galaxies or that may belong to the new class of ``chain'' galaxies defined by Cowie \\etal (1995).  Comparably complex and diverse morphologies can, however,  be found even among local very blue galaxies such as clumpy irregulars, Markarian galaxies, extragalactic HII regions, merging galaxies, or other peculiar classes (c.f. images of HII galaxies presented by Melnick 1987).  These various local blue systems often show similar, multiple concentrations of intense star formation.  The new Keck redshifts and HST morphologies combine to  counter the  view that low redshift, very blue, low luminosity dwarf galaxies dominate the late-type, peculiar systems detected in deep HST images (e.g., Driver \\etal 1995). Instead, most of the Keck galaxies with these morphologies are at high redshifts and are thus relatively luminous, though typically less luminous than $L^*$.  The apparent sizes of many of these distant galaxies are $\\sim 1$ arcsec, however,  which imply metric sizes of several kiloparsecs, quite typical of local dwarf galaxies.  After corrections for cosmological dimming (about 2.3 magnitudes at $z \\sim 1$) and K-corrections (from $\\sim$ 0 mag for very blue galaxies  to over 4 mag for very red galaxies at $z \\sim 1$), the resultant restframe $B$ band surface brightnesses are higher than those of most local spirals and star-forming irregulars (Schade \\etal 1995, Phillips \\etal 1996).  Thus, the faint blue  galaxies  are not explained by a class of very low surface brightness galaxies that have been missed by local surveys (McGaugh 1994).  Whether the high redshift  galaxies are indeed low-mass systems (which  may or may not have any direct relationship to local dwarfs) or are the luminous tips of more massive systems will be decided by results of kinematic surveys.  So far, only small samples of very blue, distant, field galaxies have been observed, many quite compact and as distant as $z \\sim 0.8$, and these have yielded some luminous galaxies with very-low, emission-line velocity widths of $\\sigma \\la 70$~\\kms (Koo \\etal 1995, Colless 1995, Guzm\\'an \\etal 1996, Forbes \\etal 1996).  We conclude from the present sample that no single physical process (e.g., mergers or bursting dwarfs) dominates the evolution of faint galaxies at redshifts $z \\sim 1$.  We find a diversity of morphologies from normal to peculiar, with no firm evidence for entirely new classes of galaxies for which local counterparts cannot be found, a conclusion also reached by Forbes \\etal (1996).  We also find bulge and elliptical systems that are well formed and red.  Equally tantalizing are the hints that a large fraction of galaxies are participating in further agglomeration and continued star-formation due to interactions, major mergers, and infalling satellites. Taken together, these early Keck results for distant HST galaxies imply that much larger samples will be needed to disentangle and understand this exciting, complex, and very important early history of galaxies."
    },
    "9604/astro-ph9604055_arXiv.txt": {
        "abstract": "Using  a new physical model for star formation (Padoan 1995) we have tested the possibility that globular clusters (GCs) are formed from primordial mass fluctuations, whose mass scale ($10^8$ - $10^9$ M$_{\\odot}$) is selected out of a CDM spectrum by the mechanism of non-equilibrium formation of $H_2$.  We show that such clouds are able to convert about 0.003 of their total mass into a bound system (GC) and about 0.02 into halo stars. The metal enriched gas is dispersed away from the GC by supernova explosions and forms the galactic disk.  These mass ratios between GCs, halo and disk depend on the predicted IMF which is a consequence of the universal statistics of fluid turbulence. They also depend on the ratio of baryonic over non-baryonic mass ,$X_b$, and are comparable with the values observed in typical spiral galaxies for $X_b \\approx 0.1-0.2$.  The computed mass and radius for a GC ( $5\\times 10^5$ M$_{\\odot}$ and 30 pc) are in good agreement with the average values in the Galaxy.  The model predicts an exponential cut off in the stellar IMF below 0.1 M$_{\\odot}$ in GCs and 0.6 M$_{\\odot}$ in the halo. The quite massive star formation in primordial clouds leads to a large number of supernovae and to a high blue luminosity during the first two Gyr of the life of every galaxy. ",
        "introduction": "Globular clusters are the fossil record of galaxy formation. They are among the oldest known objects in the Universe.  Since the first colour-magnitude diagrams (CMD) for GCs were obtained by Arp, Baum and Sandage (1952) in the early 50's, GCs became the natural laboratory where the theory of stellar evolution was tested. On the other hand, the stellar evolution theory has allowed the determination of the age of GCs (e.g., Sandage 1962, Iben \\& Renzini 1984, Vandenberg et al. 1992, Jimenez et al. 1996), that is today considered as the best estimate for the age of the Universe.  Each galaxy contains hundreds of GCs, whose properties are surprisingly similar in the Universe (Harris and Racine 1979, Harris 1991), suggesting the presence of a common physical mechanism in the early stages of galaxy formation. The study of GCs should therefore provide important clues for the development of a theory for the origin of galaxies.  Sites of present day star formation do not form bound stellar systems as massive as GCs. Therefore modelling the formation of GCs could be a critical test for any theory of star formation.  It is known that stars are formed with an efficiency of a few percent in giant molecular clouds of the Galactic disk (Duerr, Imhoff \\& Lada 1982, Myers et al. 1986, Mooney and Solomon 1988), and bound open clusters with an efficiency ten times smaller (Larson 1986). Therefore it is not surprising if GCs contain only 0.1\\% of the luminous mass of the parent galaxy, since, from the point of view of the efficiency of present day star formation, they can only emerge from protoglobular clouds 1000 times more massive (as in Searle 1977, Searle \\& Zinn 1978, Harris \\& Pudritz 1994).  Nevertheless models of GCs formation have identified objects of just $10^6$ M$_{\\odot}$ with protoglobular clouds. Such models place the GC formation period either before the galaxy is formed or during its formation. The first model for GCs as primordial objects was proposed by Peebles \\& Dicke 1968, and was later revised by Peebles (1984) and Rosenblatt, Faber \\& Blumenhal (1988). The secondary formation scenario, in which the mass of the protoglobular clouds are determined by the detailed mechanism of cooling in the protogalactic cloud, was first proposed by Fall \\& Rees (1985) and later improved by Kang et al. (1990), Ashman (1990), Brown, Burkert \\& Truran (1991), Brown, Burkert \\& Truran (1995), Vietri \\& Pesce (1995).  An alternative scenario has been proposed to explain the formation of some of the GCs in large shocks, e.g. in merging and interacting galaxies (Ashman \\& Zepf (1992), Kumai, Basu \\& Fujimoto (1993)).  In this paper we explore the possibility that GCs are formed in large clouds identified with primordial mass fluctuations in a cold dark matter (CDM) spectrum. The progenitors of GCs are clouds as massive as a few $10^8$ M$_{\\odot}$ of baryons (Searle \\& Zinn 1978, Zinnecker et al. 1988 and Larson 1990 have suggested that GCs form in the core of very massive clouds). Such mass-scale is selected out of the standard CDM spectrum by the mechanism of non-equilibrium H$_{2}$ formation (see section 2) and, as mentioned above, it is required by the low efficiency of the process of star formation on large scales in order to form a bound stellar system.  The protoglobular cloud is efficiently cooled to approximately 100 K by H$_{2}$ collisional excitation. The cooling time is short enough that the gas can be considered isothermal. During the dissipative collapse (isothermal shocks radiate away most of the kinetic energy) of the baryonic gas (see section 2) star formation occurs until supernova explosions disperse away the gas (section 4).  We suggest that this star formation process in very massive primordial clouds can be responsible for the formation of the halo GCs (GCs with metallicity below 0.1 the solar value) and of a significant fraction of the halo stars. The dispersed metal-enriched gas is further processed in following star formation episodes.  The paper is organised in the following way. In section 2 we show how the mass-scale of protoglobular clouds is selected. Section 3 is a brief review of the statistical model of star formation (Padoan 1995) that is used in this work. Results are presented in sections 4 and 5. The last two sections of the paper contain the discussion and the conclusions. ",
        "conclusions": "The main conclusions of this paper are the following:  \\begin{itemize}  \\item GCs can be formed out of the nucleus of a  primordial cloud of a few $10^8$ M$_{\\odot}$ of baryons.  \\item The mechanism of $H_2$ formation selects this typical mass scale out of a CDM spectrum.  \\item Most halo stars in the galaxy (bulge excluded) can be formed together with the (halo) GCs in the same clouds; the halo and the GCs masses are sensitive to the ratio of baryonic over non baryonic matter, X$_{b}$, in the cloud.  \\item A ratio $0.1<X_{b}<0.5$ produces realistic values for the mass of GCs; if $X_{b}<0.1$ too large GCs are formed.  \\item A ratio $X_{b}<0.2$ is needed in order to originate most of the halo stars in the protoglobular clouds during the formation of GCs.  \\item The gas that is blown away from the disrupted protoglobular cloud by supernova explosions is used in the following star formation episodes; therefore the initial disk metallicity (minimum stellar metallicity in the disk) is [O/Fe]$=-1$.  \\item The IMF in GCs has a cut off below 0.1 M$_{\\odot}$ while the halo IMF has the cut off below 0.6 M$_{\\odot}$.  \\item During the first 2 Gyr of its life each galaxy has higher blue luminosity than normal disks. This is due to the fact that star formation  (halo stars) in the protoglobular cloud produces a significant number of massive stars, which also means an epoch of large number of supernova explosions.   \\end{itemize}  We have shown that this scenario is consistent with several observations.  The star formation rate and the stellar IMF are essential ingredients of the model. In order to predict them we have used a new physical model of star formation. In this model the functional form of the IMF is shown to be universal because it is based on the statistics of turbulence which are universal properties of flows.  Future work should concentrate in the attempt of simulating the collapse of CDM primordial protoglobular clouds. Cosmological simulations on such small scales are at the moment beyond the capabilities of available computer resources, since they require a very large dynamical range."
    },
    "9604/gr-qc9604033_arXiv.txt": {
        "abstract": " ",
        "introduction": "The design and construction of a number of new and more sensitive detectors of gravitational radiation is currently underway.  These include the LIGO detector being built in the United States by a joint Caltech/MIT collaboration \\cite{science92}, the VIRGO detector being built near Pisa by an Italian/French collaboration (Caron {\\sl et al.} 1995), the GEO-600 detector being built in Hannover by an Anglo/German collaboration (Danzmann {\\sl et al.} 1995), and the TAMA-300 detector being built near Tokyo \\cite{tama300}.  There are also several resonant bar detectors currently in operation, and several more refined bar and interferometric detectors presently in the planning and proposal stages.  It's not clear when the first sources will be detected -- this may happen soon after the detectors go ``on-line\" or it may require decades of further work to increase the sensitivity of the instruments.  But it's clear to me that eventually, when the sensitivity passes above some threshold value, the gravity-wave detectors will find sources.  The operation of these detectors will have a major impact on the field of gravitational physics.  For the first time, there will be a significant amount of experimental data to be analyzed, and the ``ivory tower\" relativists will be forced to interact with a broad range of experimenters and data analysts to extract the interesting physics from the data stream.  Being an optimist, I think that this ``interesting physics\" will include sources of gravitational radiation which we have not hitherto expected or even conceived of.  It promises to be an exciting time.  There are at least three major types of ``known\" sources (Schutz 1991; Thorne 1995).  These are (1) coalescing binary systems, composed of neutron stars and/or black holes, (2) pulsars (and other periodic sources), and (3) supernovae (and other transient or burst sources).  These are listed, roughly speaking, in increasing order of ``detection difficulty\".  For example, the waveforms (chirps) of the coalescing binary systems are already known to a high enough degree of precision (Blanchet {\\sl et al.} 1995; Apostolatos {\\sl et al.} 1994; Tagoshi \\& Sasaki 1994) to enable reliable detection, using the classical signal analysis technique of matched filtering (Dhurandhar \\& Sathyaprakash 1994; Apostolatos 1995).  Some order-of-magnitude studies have shown that the amount of required processing can be done with a dozen or so high-performance desk-top workstations \\cite{owen}.  The major unknown for this type of source is their number density, which is uncertain by more than an order of magnitude.  Hence the rate of events which occur close enough to us to be observable can not be estimated very precisely.  The second type of source is pulsars.  Once again, the form of the signal is known very precisely:  it's just a sine wave in the solar system barycenter, a coordinate system at rest with respect to the Sun \\cite{pulsars}.  In this case however, the signal processing problem is far more difficult, because the orbital motion of the earth around the sun and the rotational motion of the earth around its axis modulates the pulsar frequency.  This means one must separately analyze the signal for each of $\\approx 10^{14}$ separate patches on the celestial sphere, each of which would have a distinctive pattern of frequency modulation.  The required processing speeds are currently beyond the limits of even the most powerful computers.  Of course, it may be just a question of waiting until a better search algorithm is developed, or faster computers become available!  The third source which I've listed, supernovae, is perhaps the most difficult to detect.  There are two reasons.  First, we do not have precise predictions of the gravitational waveform such an event would produce -- and it's hard to design a data analysis algorithm to look for something unknown.  Furthermore, these explosions are expected to be rare events and because certain kinds of instrument noise might look quite similar to supernovae, I chose to list them last, and categorize them as the most difficult sources to detect.  The subject of these lectures is a fourth type of source, quite different in character from the three listed above.  These are the ``stochastic\" or ``background\" sources (Michelson 1987; Christensen 1992; Flanagan 1993). Roughly speaking, these are ``random\" sources, typically arising from an extremely large number of ``unresolved\" independent and uncorrelated events.  This type of background could be the result of processes that take place very shortly after the big bang, but since we know very little about the state of the universe at that time, it's impossible to say with any certainty.  Such a background might also arise from processes that take place fairly recently (say within the past several billion years) and this more recent contribution might overwhelm the parts of the background which contain information about the state of the early universe.  These sources are ``unresolved\" in the following sense.  If we study an optical source, somewhere in the sky, using a telescope with a certain angular resolution, then details of the source can be ``resolved\" if the angular resolution of the telescope is smaller than the angular size of the features or objects being studied.  In the case of the LIGO experiment, and similar detectors, the angular size of the antenna pattern is of order $90^\\circ$.  Hence almost any source is ``unresolved\" in that it makes a significant contribution to the detector output for almost any orientation of the detector and the source.  When many such sources are present, even if they are pointlike, the resulting signal has a stochastic nature.  My motivation for studying these stochastic sources is two-fold.  The first reason is a rather hopeful one.  Because the gravitational force is the weakest of the four known forces, the small-scale perturbations of the gravitational field decouple from the evolution of the rest of the universe at very early times.  Currently, our most detailed view of the early universe comes from the microwave background radiation, which decoupled from matter about $10^5$ years after the big bang, and gives us an accurate picture of the universe at this rather early time.  Some rather simple estimates (which I'll elaborate later) show that if the current crop of gravity-wave detectors do detect a background of cosmological origin, then it will carry with it a picture of the universe as it was about $10^{-22}$ seconds after the big bang.  This would represent a tremendous step forward in our knowledge, and is the main reason for my interest.  The second reason for my interest in stochastic sources is rather more practical.  One often hears it said that searching for signals in the output of a gravitational wave antenna is like searching for a needle in a haystack.  And indeed, for the first three types of source listed above, this is true.  One has to search through tons of rock (the data stream) in order to find the one precious gem (say, a binary chirp). Most of the rock is barren, and the challenge is to isolate the one tiny volume containing the material of interest.  As I have already discussed, for gravitational wave sources, the amount of computational power required for this careful search can be very large.  However the situation is rather different when one is searching for a stochastic background.  The analogy in this case is mining aluminum, where on average every ton of ore contains a certain number of kilograms of aluminum.  The situation is analogous for stochastic gravity-wave sources.  As you will see, whenever two detectors are operating simultaneously, even if only for a few seconds, we get a little bit more data and information about the stochastic background.  And as you will also see, it is easy to analyze this data.  The ``signal\" in this case is a very low bandwidth one, so the essential part of the data analysis for a stochastic background can be done on a garden-variety personal computer.  For example, in the case of the LIGO detectors, the part of the signal carrying a significant amount of information lies below a few hundred Hz (see Section~\\ref{s:optimal} and Fig.~\\ref{f:filters} for details).   Thus the rate at which information needs to be processed is only a few hundred data points/second; a very manageable rate.  These lectures are organized as follows.  In Section~\\ref{s:two}, I discuss some of the general properties that a stochastic background of gravitational radiation might have.  I show how such radiation is characterized by a spectral function, discuss some of its statistical properties, and show during which cosmological epoch the radiation which falls into the bandwidth of the ground- and space-based detectors was produced.  In Section~\\ref{s:three} I show how one can combine data from two or more gravity-wave detectors to either put limits on the amplitude of a stochastic background, or to actually detect it.  I do this in several steps, first giving a crude argument which demonstrates the main detection strategy, then discussing the reduction in sensitivity which comes about from the siting of the detectors and their relative orientations on the earth.  This is followed by a rigorous derivation of the optimal signal processing strategy, and a calculation of the expected signal-to-noise ratio and the minimum detectable energy-density in a stochastic background.  In Section~\\ref{s:what} I discuss the observational facts: what we actually know about the stochastic background.  There are strong limits on the spectrum of a stochastic background coming at very long wavelengths from observations of the isotropy of the Cosmic Microwave Background Radiation, and limits at higher frequencies from observations of timing residuals in millisecond pulsars.  Finally, there is a limit on the integrated spectrum arising from the standard model of big bang nucleosynthesis.  The last part of these lectures is far more speculative.  In Section~\\ref{s:sources} I discuss some of the potential ways in which a stochastic background might arise from processes that take place early in the history of the universe.  I discuss in some detail three particular models (inflation, cosmic strings, and bubble formation in a first-order phase transition) each of which gives rise to a different spectrum of radiation, and examine in some detail the mechanisms at work in each case. In each case I indicate if and when the experiments being built might detect these potential sources. This is followed by a short conclusion, and an appendix containing a few calculational details. ",
        "conclusions": "\\setcounter{equation}0 \\label{s:conc} I am often asked, after a series of lectures like this, ``Do you really think that a stochastic background is there?\" There is only one possible honest answer -- I do not know!  The problem is that we have very little information about what really happened in the very early universe.  And the scenarios that we have given as examples are just that, examples.  Personally, I think it is unlikely that any of them will turn out to be true.  However they do have {\\it illustrative} value.  They show that, under reasonable assumptions about what happened in the early universe, it {\\it is} possible to produce a stochastic background which is easily detected with the instruments now under construction.  The possibility of such detection has been anticipated for some time (Grishchuk 1976; 1988), and there are good hopes that they will be realized in the next decade.  So while the answer to the question is unknown, that does not matter. In the next decade we will probably know the answer in a very direct way, with real observational data.  And if we do detect a stochastic background, a new era in our understanding of the early universe will begin."
    },
    "9604/astro-ph9604074_arXiv.txt": {
        "abstract": "The solar neutrino anomaly, measurements discrepant from predictions of the Standard Solar Model, has existed for over 30 years.  Multiple experiments measuring fluxes from several reactions in the hydrogen fusion chain have added to the puzzle.  Each of the several elements of the enigma are resolved by recognition of measurements establishing that most of the sun's fusion must occur near the surface rather than the core. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604148_arXiv.txt": {
        "abstract": "The FIRST radio survey provides a new resource for constructing a large quasar sample.  With source positions accurate to better than 1\\arcsec\\ and a point source sensitivity limit of 1 mJy, it reaches 50 times deeper than previous radio catalogs.  We report here on the results of the pilot phase for a FIRST Bright Quasar Survey (FBQS). Based on matching the radio catalog from the initial 300 deg$^2$ of FIRST coverage with the optical catalog from the Automated Plate Machine (APM) digitization of Palomar Sky Survey plates, we have defined a sample of 219 quasar candidates brighter than E = 17.50.  We have obtained optical spectroscopy for 151 of these and classified 25 others from the literature, yielding 69 quasars or Seyfert~1 galaxies, of which 51 are new identifications.  The brightest new quasar has an E magnitude of 14.6 and z = 0.91; four others are brighter than E = 16.  The redshifts range from z=0.12 to 3.42.  Half of the detected objects are radio quiet with L$_{\\rm 21cm} < 10^{32.5}$ ergs/s.  We use the results of this pilot survey to establish criteria for the FBQS that will produce a quasar search program which will be 70\\% efficient and 95\\% complete to a 21-cm flux density limit of 1.0 mJy. ",
        "introduction": " ",
        "conclusions": "We have shown that the FIRST radio survey can be used in conjunction with the POSS plate material as the basis for an efficient QSO survey, highly complete for radio loud sources.  With the 1994 survey data, the current FIRST catalog covers a total 1550 deg$^2$ and includes over 138,000 radio sources.  Using our revised screening criteria, we have prepared a list of $\\sim 400$ additional QSO candidates for spectroscopic follow-up.  With this larger sample, selected in a uniform fashion and covering a large contiguous area of sky, we will begin to be able to address important questions regarding quasar radio and optical luminosity functions and their evolution, the large scale distribution of quasars, and the differences between radio loud and radio quiet objects."
    },
    "9604/astro-ph9604170_arXiv.txt": {
        "abstract": "We present simultaneous H$\\alpha$, Na~{\\sc i} D$_{1}$, D$_{2}$, and He~{\\sc i} D$_{3}$ spectroscopic observations on the RS CVn system UX Ari. We have found a dramatic increase in the excess H$\\alpha$ emission equivalent width by a factor of 2 in an interval of 1 day that indicates the beginning of a strong flare in this system. The presence of the He~{\\sc i} D$_{3}$ in emission in coincidence with the enhancement of the H$\\alpha$ emission confirms the detection of a flare. The application of the spectral subtraction technique reveals that the core of the Na~{\\sc i} D$_{1}$ and D$_{2}$ lines are also filled-in by chromospheric emission. ",
        "introduction": "The RS CVn system UX Ari (HD 21242) is a non-eclipsing double-lined spectroscopic binary with spectral type (G5V/K0IV) and orbital period of 6.43791 days. It is one of the few RS CVn systems, along with V711 Tau, II~Peg, and DM UMa, that shows H$\\alpha$ consistently in emission. The extensive Ca~{\\sc ii} H \\& K and H$\\alpha$ observations of UX Ari (Carlos \\& Popper 1971;  Bopp \\& Talcott 1978; Nations \\& Ramsey 1986; Huenemoerder et al. 1989; Frasca \\& Catalano 1994; Raveendran \\& Mohin 1995; Montes et al. 1995a,b,~c) show that the chromospheric emission belongs to the cooler component of the system, however some authors (Simon \\& Linsky 1980; Huenemoerder et al. 1989) also found a small contribution from the G5V secondary component.  Simon et al. (1980) observed a UV flare on UX Ari and found that chromospheric lines were enhanced 2.5 times their quiescent level, while transition region lines displayed an enhancement by a factor of up to 5.5. They proposed an activity scenario for long-lived RS CVn flares in which the component stars  have  large  corotating  flux  tubes  which occasionally interact. The resulting magnetic  reconnection  leads to flares. Recently, Henry \\& Newsom (1996) have detected, with photometric observations, an optical flare in this system.  In this paper, we report the detection of an optical flare in UX Ari through simultaneous observations of the chromospheric H$\\alpha$ and He~{\\sc i} D$_{3}$ lines. It is the first time that the He{\\sc i} D$_{3}$ line is observed in emission in UX Ari, moreover, a filled-in in the Na~{\\sc i} D$_{1}$ and D$_{2}$ lines have been found. In Sect. 2 we give the details of our observations and data reduction and we describe the method of spectral subtraction to obtain the active-chromosphere contribution. In Sect. 3 we describe the behaviour of the chromospheric activity indicators and Sect. 4 gives the conclusions.   \\begin{figure*} \\psscalefirst \\vspace{-0.5cm} \\hspace{-1.5cm} {\\psfig{figure=aal_fig1.ps,height=27.2cm,width=17.6cm,rheight=18.25cm,rwidth=6cm,angle=270}} \\caption[ ]{H$\\alpha$, Na~{\\sc i} D$_{1}$, D$_{2}$, and He~{\\sc i} D$_{3}$ spectra of UX Ari For each system we plot the observed spectrum (solid-line), the synthesized spectrum (dashed-line), the subtracted spectrum, additively offset to better display (dotted line), and the two components fit to the subtracted spectrum (dotted-dashed line). \\label{fig:plot_paphahe} } \\end{figure*} ",
        "conclusions": "We have shown here that the UX Ari binary system exhibits a strong increase in the H$\\alpha$ EW at 0.74 orbital phase. In the same phase the He~{\\sc i} D$_{3}$ appears in emission, and the Na~{\\sc i} D$_{1}$ and  D$_{2}$ lines show a larger filled-in than in the other phases. This behaviour yields us to the conclusion that we have detected a flare which released about 1.7~10$^{31}$~erg~s$^{-1}$ in the H$\\alpha$ line.   According to the magnetic loop model for UX Ari described by Simon et al. (1980) in which the component stars have  large  corotating  flux  tubes  which occasionally interact, the resulting magnetic  reconnection  leads to flares. When the emitting regions are seen in projection against the stellar disk the He~{\\sc i} D$_{3}$ line appears in absorption, when the regions are observed off the stellar limb He~{\\sc i} D$_{3}$ can be seen in emission. The orbital phase (0.74) of our UX Ari observation and its symmetrical position are the most favourable to observe the He~{\\sc i} D$_{3}$ line in emission. The other He~{\\sc i} D$_{3}$ observations in emission in II~Peg are also seen at these orbital phases.  This fact together with the strong increase in the H$\\alpha$ EW and the larger filled-in of the Na{\\sc i} D$_{1}$ and D$_{2}$ lines at 0.74 orbital phase support the presence of a flare in the UX Ari binary system.  Photometric observations of this system indicate that the light curve minimum (due to cool spots) normally occurred between phase 0.6-0.9 (Raveendran \\& Mohin 1995 and references therein), on the other hand Rodon\\`{o} \\& Cutispoto (1992) observed an anticorrelation between the V light curve and the U-B and B-V colours variations that is not consistent with the expected variations produced by the presence of cool spots alone, and might be due to continuous flaring activity  or bright plage-like regions associated with the spots. Our flare observation occurs at phase 0.74, so the flare eruption could be produced (as have been observed in the largest solar flares) by the interaction of new emerging magnetic flux with old magnetic structures associated to the spot groups present in the star at this phase."
    },
    "9604/astro-ph9604145_arXiv.txt": {
        "abstract": "We analyze compact groups of galaxies appearing in a galaxy formation simulation dominated by cold dark matter ($\\Omega_0$=1, $\\Omega_{baryon}$=0.1).  The simulation uses an N-body code to model the behavior of the non-baryonic matter and smoothed particle hydrodynamics to model the baryons.  One run includes gas dynamics alone, and the other incorporates star formation as well.  Groups identified as physically compact at $z$=0 form originally along filaments and become compact in the final $\\sim$20\\% of the simulation; they contain x-ray--luminous diffuse gas well before becoming compact.  The component masses, baryon fractions, and gas-to-galaxy mass fractions of the simulated groups are roughly similar to those of Hickson compact groups (HCGs), although they tend to be somewhat more gas-rich and have x-ray temperatures a factor of 3 lower than those seen in HCGs.  These discrepancies may be alleviated by adding the effects of energy input into the diffuse group gas by star formation and supernovae. ",
        "introduction": "Most galaxies are neither completely isolated nor in rich clusters, but instead are found in galaxy groups that contain only a handful of members.  Although galaxy groups are very common, they have not been extensively studied due to the difficulty of determining whether a galaxy density enhancement on the plane of the sky is a true group or simply a superposition of field galaxies.  Thus, many investigations of groups have focused on ``compact groups of galaxies'', particularly those catalogued by Hickson (1982).  Hickson compact groups (HCGs) were selected to have high average surface brightness and to be relatively isolated, conditions that increase the probability that the groups are physically compact. Subsequent redshift measurements showed that 92 of the original 100 HCGs contain at least three galaxies with concordant redshifts, strengthening this argument (\\cite{hic92}).  HCGs have been observed throughout the electromagnetic spectrum for signs of galaxy-galaxy interaction and other indications of physical compactness (see Pildis, Bregman, \\& Evrard 1995, hereafter Paper I, for an overview).  The results have been contradictory at best:  strong hints of interaction (e.g., HI deficits in HCG spiral galaxies, a probable diffuse x-ray--emitting intragroup medium in a subset of groups) coexist with the lack of other interaction signs (e.g., no obviously merged HCG remnants, few to no HCG ellipticals with signs of recent mergers).  Furthermore, the analysis of the x-ray observations of these groups is complicated by systematic differences between the reduction procedures of different groups.  Depending on the choice of background level and plasma metallicity, as well as other free parameters, the same data can be interpreted to give gas and total masses that differ by up to an order of magnitude (e.g., Paper I, \\cite{sar95,mul96}).  While preceding papers in this series (Paper I and Pildis, Bregman, \\& Schombert 1995, hereafter Paper II) and other recent work have shown that some order can be found in this confusion---for example, that elliptical-rich HCGs tend to have intragroup media while spiral-rich groups don't, and that HCG ellipticals are not as quiescent as they appear at first glance---in no sense is there a coherent picture of what Hickson compact groups are and how they evolve.  Observations alone can shed only a limited amount of light on the puzzle of Hickson compact groups.  While optical and x-ray data can give information about the matter that radiates in those wavebands as well as some indirect indication about quantities such as the total gravitating mass, they cannot provide direct and unambiguous information about invisible baryonic components such as hot gas whose emission is absorbed by Galactic neutral hydrogen or diffuse cold gas, much less non-baryonic matter. Furthermore, while optical signs of interaction in HCGs---such as shells around galaxies and diffuse light in the group potential---imply certain limits on the interaction history of a group, they do not provide many details about that history.  Simulations of galaxy formation and structure evolution can help answer some of these questions.  For example, recent simulations by Hernquist, Katz, \\& Weinberg (1995) have found that galaxies form in filaments, and thus many compact-appearing groups may be filaments of galaxies seen end on.  In such cases, even though the galaxies have projected separations of tens of kiloparsecs, they are physically separated by 2--10 Mpc along the filament.  While Hernquist et al.~did not propose any immediate tests to determine how many Hickson groups are probable projections (redshift-independent distance indicators would test their hypothesis, but such indicators are not sufficiently accurate to use for this purpose), their study illustrates how simulations can explain earlier, seemingly contradictory observations and can help guide future observations in order to test a theory.  One suggestion as to how to discriminate between projected groups and physically compact ones has been made by Ostriker, Lubin, \\& Hernquist (1995).  They claim that their $Q$ parameter should be a constant for all physically compact groupings, from compact groups up to rich clusters ($Q = L_X a_p^3 L_g^{-2} f(T)^{-1}$, where $a_p$ is the projected radius, $L_X$ the bolometric x-ray luminosity, $L_g$ the galaxy luminosity in the blue, and $f(T)$ gives the temperature dependence for the x-ray luminosity of a thermal plasma).  Using cluster data and the results that were given in Paper I, they conclude that HCGs are either elongated along the line of sight to the extent that they cannot be considered true groups or extremely gas-poor relative to rich clusters.  In order to interpret the observational results from Paper I and Paper II, we analyze some ``Hickson-like'' groups found in a simulation of galaxy formation in a universe dominated by cold dark (non-baryonic) matter.  Such an analysis improves our qualitative understanding of group formation and provides some level of quantitative detail on basic quantities related to their evolution and present structure.   Properties such as a group's baryon fraction, gas-to-stellar mass ratio, and temperature of its x-ray--luminous gas can be compared to observed quantities in HCGs.  In addition, the simulations provide the dynamical histories of the groups.  If the simulated groups have roughly the same properties as do the HCGs in our sample, then examining the time evolution of the simulated groups will provide insight into how compact groups form and evolve. ",
        "conclusions": "The fundamental result of this study is that compact-appearing groups in a simulation of a cold-dark-matter--dominated universe are similar in many ways to actual compact groups.  The simulated groups that we analyzed have lower gas-to-galaxy mass ratios than do simulated clusters, and their total and component masses are fairly similar to those measured for HCGs.  In addition, the x-ray properties of the simulated groups are in many cases similar to those of x-ray--bright HCGs, even when the simulated groups are not physically compact.   An important difference between these groups and HCGs is that the simulated groups are chosen by examining the dark matter distribution at a modest density contrast, then examining the galaxies within them.  Much larger simulations will be required before realistic selection (including isolation criteria) from ``optical'' catalogues will be available.  While the gross similarities between the simulated groups and the x-ray--luminous HCGs discussed in Paper I are striking, there are also some important differences.  While this simulation does produce gas-to-galaxy mass ratios nearly an order of magnitude smaller than those of clusters (both simulated and actual), the ratios are not as low as some observed in HCGs. Also, the emission-weighted temperatures for the hot baryons in the simulation are a factor of 2--5 less than the 0.9 keV temperatures measured in Hickson groups.  Both these discrepancies may be due to the lack of galactic winds and metallicity enhancement in the simulation, as well as systematic effects that arise in x-ray data reduction such as the choice of gas metallicity and background level (Paper I, \\cite{mul96}).  While incorporating star formation into the simulation did make some observable properties of the simulated groups (such as the spatial structure of the x-ray--luminous gas) more similar to those of HCGs, it increased the differences in other areas (e.g., the temperature of the intragroup medium).  However, in most respects, HCGs fall somewhere in between the extremes of the GO and SF runs.  The addition of the effects of of energy feedback into the intragroup medium is likely to reduce the remaining discrepancies considerably.  If g2 and g4 can be considered typical x-ray--bright compact groups, then perhaps the best description of how compact groups evolve is a combination of the bound group simulations of Barnes (1989) and the filamentary ``false groups'' of Hernquist et al.~(1995).  Until the final 2 Gyr of the simulation (when the groups become physically compact), the galaxy-like objects in each group are found mainly along a single filament. For a viewing angle down that filament, the ``galaxies'' would appear to be close together and to have concordant redshifts due to their infall to the intersection with a poorer filament, but actually be physically far apart along the filament.  At a lookback time of roughly 2 Gyr, the galaxies become physically close, and the group becomes truly compact and somewhat richer as additional galaxies infall from the poorer filament.  At this point, the evolution should proceed as in the simulations of Barnes (1989) and others (e.g., \\cite{gov91,bod93}).  Our results provide no clear method to discriminate between projected and physical groups.  The x-ray temperatures of the groups rise to a constant value very early in their evolution, due to shock heating of gas as it falls into galactic potential wells and filaments (\\cite{kat92}).  The behavior of the various baryonic mass components does not appear to be affected by the formation of a physically compact group, either.  The rise in the x-ray temperature in the final few Gyr of the simulation is the principal indication of the formation of a physically compact group, but this simulation finds gas temperatures considerably lower than those of HCGs (likely due to the absence of heating from galactic winds), making it difficult to determine what magnitude of a temperature increase should be searched for.  As discussed above, the $Q$ parameter of OLH is not appropriate for gas-poor objects such as compact groups.  While it may be difficult to determine whether an individual group is physically compact, this simulation indicates that, on the whole, many HCGs are likely to be true groups rather than simply projections.  Although the groups in this simulation spend more time as filaments than as physically compact groups, filaments would be identified as groups only from a small percentage of viewing angles ($<$20\\%).  Since the groups are compact for over 30\\% of their (x-ray--luminous) lifetimes, a significant proportion of x-ray--bright HCGs are probably physically compact groups.  A good criterion for distinguishing true groups from projected ones might be found from a greater ensemble of simulated groups.  Not only would trends seen with the groups analyzed here become more clear, but examining the evolution of more groups chosen at $z >$ 0 would illustrate the probable future state of HCGs.  Perhaps the final merger products would have some distinctive property that would permit a search for such objects.  The addition of supernova-driven galactic winds to the simulations would bring them closer to reflecting reality, and perhaps provide some indication of how one can distinguish projected groups from bound ones.  Such simulations would also predict the metallicity of the hot gas in compact groups, a measurement of which will be able to be made with the next generation of x-ray telescopes.  \\clearpage"
    },
    "9604/gr-qc9604044_arXiv.txt": {
        "abstract": "{\\if@twocolumn ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604003_arXiv.txt": {
        "abstract": "The evolution of the particle distribution functions inside a relativistic jet containing an \\emp pair plasma and of the resulting $\\gamma$-ray and X-ray spectra is investigated. The first phase of this evolution is governed by heavy radiative energy losses. For this phase, approximative expressions for the energy-loss rates due to inverse-Compton scattering, using the full Klein-Nishina cross section, are given as one-dimensional integrals for both cooling by inverse-Compton scattering of synchrotron photons (SSC) and of accretion disk photons (EIC).  We calculate instantaneous and time-integrated $\\gamma$-ray spectra emitted by such a jet for various sets of parameters, deducing some general implications on the observable broadband radiation. Finally, we present model fits to the broadband spectrum of Mrk~421. We find that the most plausible way to explain both the quiescent and the flaring state of Mrk~421 consists of a model where EIC and SSC dominate the observed spectrum in different frequency bands. In our model the flaring state is primarily related to an increase of the maximum Lorentz factor of the injected pairs. ",
        "introduction": "Accretion of matter onto a central black hole is the most relevant process to power active galactic nuclei (Lynden-Bell 1969, Salpeter 1969, Rees 1984). However, the details of the conversion processes of gravitational energy into observable electromagnetic radiation are still largely unknown. The discovery of many blazar-type AGNs (Hartman et al. 1992, Fichtel et al. 1993) as sources of high-energy \\gamr radiation dominating the apparent luminosity, has revealed that the formation of relativistic jets and the acceleration of energetic charged particles, which generate nonthermal radiation, are key processes to understand the energy conversion process. Emission from relativistically moving sources is required to overcome \\gamr transparency problems implied by the measured large luminosities and short time variabilities (for review see Dermer \\& Gehrels 1995).  Repeated \\gamr observations of AGN sources have indicated a typical duty cycle of \\gamr hard blazars of about 5 percent, supporting a ``2-phase\" model for the central regions of AGNs (Achatz et al. 1990, Schlickeiser \\& Achatz 1992). According to the 2-phase model the central powerhouse of AGNs undergoes two repeating phases: in a ``quiescent phase\" over most of the time ($\\sim $95 percent) relativistic charged particles are efficiently accelerated in the central plasma near the black hole, whereas in a short and violent ``flaring phase\" the accelerated particles are ejected in the form of plasma blobs along an existing jet structure.  We consider the acceleration of charged particles during the quiescent phase. The central object accretes the surrounding matter. Associated with the accretion flow is low-frequency magnetohydrodynamic turbulence which is generated by various processes as e.g.:  \\noindent (a) turbulence generated by the rotating accretion disk at large eddies and cascading to smaller scales (Galeev et al. 1979);  \\noindent (b) stellar winds from solar-type stars in the central star cluster deliver plasma waves to the accretion flow;  \\noindent (c) infalling neutral accretion matter becomes ionized by the ultraviolet and soft X-ray radiation of the disk. These pick-up ions in the accretion flow generate plasma waves by virtue of their streaming (Lee \\& Ip 1987);  \\noindent (d) if standing shocks form in the neighbourhood of the central object they amplify any incoming upstream turbulence in the downstream accretion shock magnetosheath (McKenzie \\& Westphal 1969, Campeanu \\& Schlickeiser 1992).  These low-frequency MHD plasma waves from the accretion flow are the source of free energy and lead to stochastic acceleration of charged particles out of the thermal accretion plasma.  The dynamics of energetic charged particles (cosmic rays) in cosmic plasmas is determined by their mutual interaction and interactions with ambient electromagnetic, photon and matter fields. Among these by far quickest is the particle-wave interaction with electromagnetic fields, which very often can be separated into a leading field structure $F_o$ and superposed fluctuating fields $\\delta F$. Theoretical descriptions of the transport and acceleration of cosmic rays in cosmic plasmas are usually based on transport equations which are derived from the Boltzmann-Vlasov equation into which the electromagnetic fields of the medium enter by the Lorentz force term. The quasilinear approach to wave-particle interaction is a second-order perturbation approach in the ratio $q_L\\equiv ({\\delta F/F_o})^2$ and requires smallness of this ratio with respect to unity. In most cosmic plasmas this is well satisfied as has been established either by direct in-situ electromagnetic turbulence measurements in interplanetary plasmas, or by saturation effects in the growth of fluctuating fields. Nonlinear wave-wave interaction rates and/or nonlinear Landau damping set in only at appreciable levels of $(\\delta F)^2$ and thus limit the value of $q_L\\le 1$. We assume the AGN plasma to have very high conductivity so that any large-scale steady electric fields are absent. We then consider the behaviour of energetic charged particles in a uniform magnetic field with superposed small-amplitude $(\\delta B)^2 \\ll B_o^2$ plasma turbulence ($\\delta \\vec{E}, \\delta \\vec{B}$) by calculating the quasilinear cosmic ray particle acceleration rates and transport parameters. This is by no means trivial since especially for the interaction of non-relativistic charged particles with ion- and electron-cyclotron waves thermal resonance broadening effects are particularly important (Schlickeiser \\& Achatz 1993, Schlickeiser 1994). The acceleration rates and spatial transport parameters are then used in the kinetic diffusion-convection equation for the isotropic part of the phase space density of charged particles $F(x,p,t)$ which for non-relativistic bulk speed $u \\ll c$ reads  $${{\\partial F}\\over {\\partial t}}-\\, S_o= {\\partial \\over {\\partial x}}[\\kappa {{\\partial F}\\over {\\partial x}}] $$ \\begin{equation} \\;\\;\\; - \\, u{{\\partial F}\\over {\\partial x}}+{p\\over 3}{{\\partial u}\\over {\\partial x}} {{\\partial F}\\over {\\partial p}}+\\,{1\\over p^2}{\\partial \\over {\\partial p}} [p^2A{ {\\partial F}\\over {\\partial p}} -p^2 \\dot{p}_{\\rm loss} F]. \\end{equation} Here $x$ denotes the spatial coordinate along the ordered magnetic field, $p$ the cosmic ray particle momentum, $\\kappa $ is the spatial diffusion coefficient, $A$ the momentum diffusion coefficient, and $S_o$ denotes the \"Stossterm\" describing the mutual interaction of the charged particles and their injection.  With respect to the generation of energetic charged particles, the basic transport equation (1) shows that stochastic acceleration of particles, characterized by the acceleration time scale $t_A=p^2/A$, competes with continuous energy loss processes $\\dot{p}_{\\rm loss}$, characterized by energy loss time scales $t_L=p/|\\dot{p}_{\\rm loss}|$. Dermer et al. (1996) have recently inspected the acceleration of energetic electrons and protons in the central AGN plasma by comparing the time scales for stochastic acceleration with the relevant energy loss time scales. At small proton momenta the Coulomb loss time scale is extremely sensitive to the background plasma density and temperature, and for slight changes in the values of these parameters cosmic ray protons may not be accelerated above the Coulomb barrier. Although at small particle momenta the plasma wave's dissipation and the interaction with the cyclotron waves become decisive and might modify the acceleration time significantly, the results of Dermer et al. (1996) demonstrate that reasonable central AGN plasma parameter values are possible where the low-frequency turbulence energizes protons to TeV and PeV energies where photo-pair and photo-pion production are effective in halting the acceleration (Sikora et al. 1987, Mannheim \\& Biermann 1992). According to the results of Dermer et al. (1996) it takes about $\\simeq 10M_8$ days for the protons to reach these energies, where $M_8$ is the mass of the central black hole in units of $10^8 \\, M_{\\odot}$. The corresponding analysis for cosmic ray electrons shows that the external compactness provided by the accretion disk photons (Becker \\& Kafatos 1995) leads to heavy inverse Compton losses which suppress the acceleration of low-energy electrons beyond Lorentz factors of $\\gamma \\approx 10-100$. It seems that due to their much smaller radiation loss rate cosmic ray protons are effectively accelerated during the quiescent phase in contrast to low energy electrons.  Now an important point has to be emphasized: {\\it once the accelerated protons reach the thresholds for photo-pair ($\\gamma_{p,th} = m_ec^2/<\\epsilon >=5\\cdot 10^4\\epsilon^{-1}_1$) and photo-pion production and the threshold for pion production in inelastic proton-matter collisions they will generate plenty of secondary electrons and positrons of ultrahigh energy} which are now injected at high energies ($E_s\\ge 25 {\\rm GeV} \\epsilon ^{-1}_1$) into this acceleration scheme. $<\\epsilon > = 10 \\epsilon_1 $ eV denotes the mean accretion disk photon energy. It is now of considerable interest to follow the evolution of these injected secondary particles.  Although many details of this evolution are poorly understood, it is evident that the further fate of the secondary particles depends strongly on whether they find themselves in a compact environment set up by the external accretion disk, or not. As has been pointed out by Dermer \\& Schlickeiser (1993b) as well as Becker \\& Kafatos (1995) the size of the \\gamr photosphere (where the compactness is greater unity so that any produced \\gamr photon is pair-absorbed) is strongly photon energy dependent. The \\gamr photosphere attains its largest size at photon energies $E_p\\simeq 50 \\epsilon_1^{-1}$ GeV. Secondary particles within the photosphere having energies $E_s\\le E_p$ will initiate a rapid electromagnetic cascade which has been studied by e.g. Mastichiadis \\& Kirk (1995), which might even lead to runaway pair production and associated strong X-ray flares (Kirk \\& Mastichiadis 1992), and/or due to the violent effect of a pair catastrophy (Henri \\& Pelletier 1993) ultimately lead to an explosive event and the emergence of a relativistically moving component filled with energetic electron-positron pairs.  In contrast, if the secondary particles are generated outside the \\gamr photosphere the secondary electrons and positrons will quickly cool by the strong inverse Compton losses generating plenty of \\gamr emission up to TeV energies. As solutions of the electron and positron transport equation (1) for this case demonstrate (Schlickeiser 1984, Pohl et al. 1992) a cooling particle distribution ($N(p) \\propto p^{-2}$) with a strong cutoff at low (but still relativistic) momentum $p_c$ develops, which grows with time as more and more protons hit the photo-pair and photo-pion thresholds. Such bump-on-tail particle distribution functions, which are inverted ($\\partial F/\\partial p>0$) below $p_c$, are collectively unstable with respect to the excitation of electromagnetic and electrostatic waves such as oblique longitudinal Langmuir waves (Lesch et al. 1989). As described in detail by Lesch \\& Schlickeiser (1987) and Achatz et al. (1990), depending on the local plasma parameters (mainly the electron temperature of the background gas and the density ratio $N_o/n_e$ of relativistic electrons and positrons to thermal electrons) these Langmuir waves either (1) lead to quasilinear plateauing of the inverted distribution function and rapid collective thermalization of the electrons and positrons, or (2) are first damped by nonlinear Landau damping, but ultimately heat the background plasma strongly via the modulation instability once a critical energy density in Langmuir waves $W_c = n_e(k_BT_e/m_ec^2)^2$ has been built up. Both relaxation mechanisms terminate the quiescent phase of the acceleration process. The almost instantaneous increase of the background gas entropy due to the rapid modulation instability heating again leads to an explosive outward motion of the plasma blob carrying the relativistic particles away from the central object.  As we have discussed, in both cases it is very likely that at the end of the quiescent phase an explosive event occurs that gives rise to the emergence of a new relativistically moving component filled with energetic electron-positron pairs. It marks the start of the flaring phase in \\gamr blazars. The initial starting height of the emerging blob $z_i$ entering the calculation of the \\gamr flux should be closely related to the size of the acceleration volume in the quiescent phase mainly determined by the maximum size of the \\gamr photosphere. This scenario is supported by the measurements of Babadzanhanyants \\& Belokon (1985) that in 3C 345 and other quasars optical bursts are in close time correlation with the generation of compact radio jets. A similar behaviour has also been observed during the recent simultaneous multiwavelength campaign on 3C~279 (Hartman et al. 1996). Further corroborative evidence for this scenario is provided by the recent discovery of superluminal motion components in the \\gamr blazars PKS~0528+134 (Pohl et al. 1995) and PKS~1633+382 (Barthel et al. 1995) that demonstrate a close physical connection between \\gamr flaring and the ejection of new superluminal jet components in blazars.  It is the purpose of the present investigation to follow the time evolution of the relativistic electrons and positrons as the emerging relativistic blob moves out. Because of the very short radiative energy loss time scales of the radiating electrons and positrons it is important to treat the spectral evolution of the radiating particles self-consistently. In earlier work Dermer \\& Schlickeiser (1993a) and Dermer et al. (1997) have studied the spectral time evolution from a modified Kardashev (1962) approach by injecting instantaneously a power law electron and positron energy spectrum at height $z_i$ at the beginning of the flare, and calculating its modification with height in the relativistically outflowing blob due to the operation of various continous energy loss processes as inverse Compton scattering, synchrotron radiation, nonthermal bremsstrahlung emission and Coulomb energy losses. Here we generalize their approach by accounting for Klein-Nishina effects and including as well external inverse-Compton scattering as synchrotron self-Compton scattering self-consistently.  The acceleration scenario described above leads us to the following assumptions on the distribution of pairs inside a new jet component at the time of their injection: If the pairs are created by photo-pair production, their minimum Lorentz factor is expected to be in the range of the threshold value of the protons' Lorentz factor for photo-pair production. We use the standard accretion disk model by Shakura \\& Sunyaev (1973), which we will describe in more detail in the next section, to fix this threshold, determined by the average disk photon energy $\\langle \\epsilon \\rangle$. The pair distribution above this cutoff basically reflects the acceleration spectrum of the protons, i. e. a power-law distribution with spectral index $2 \\ukl s \\ukl 3$ ($n(\\gamma) \\sim \\gamma^{-s}$) which extends up to $\\gamma_{2\\pm} \\sim 10^6$. This yields the initial pair distribution functions  \\begin{equation} $$ f_{\\pm} (\\gpm) = \\gpm^{-\\sigma} \\hskip 2cm \\gepm \\le \\gpm \\le \\gzpm $$ \\end{equation} with $4 \\ukl \\sigma \\ukl 5$. The differential number of particles in the energy intervall $[\\gpm, \\gpm + d\\gpm]$ per unit volume is then given by  \\begin{equation} d \\, n_{\\pm} = 4 \\, \\pi \\, N_{\\pm} \\, \\gpm^2 \\bpm \\, f_{\\pm} (\\gpm) d\\gpm \\end{equation} where $N_{\\pm}$ is a normalization factor related to the total particle density $n_{\\pm}$ through $N_{\\pm} \\approx n_{\\pm} {\\sigma - 3 \\over 4 \\, \\pi} \\gepm^{\\sigma - 3}$.  The detection of TeV $\\gamma$-rays from Mrk~421 suggests that such components must be produced/accelerated outisde the $\\gamma$-ray photosphere for photons of energy $\\sim 1$ TeV. The height of this photosphere due to the interaction of $\\gamma$-rays with accretion disk radiation will be determined self-consistently in section 4. Backscattering of accretion disk radiation by surrounding clouds is negligible in the case of BL Lac objects emitting TeV $\\gamma$-rays (B\\\"ottcher \\& Dermer, 1995).  We point out that most of our basic conclusions are also valid if the pairs inside the blob are accelerated by other mechanisms, e. g. by a relativistic shock propagating through the jet.  Beginning at the injection height (henceforth denoted as $z_i$), we follow the further evolution of the pair distribution and calculate the emerging photon spectra.   \\begin{figure} \\epsfxsize=6cm \\epsffile[20 150 450 650] {figure1.ps} \\caption[]{Model for the geometry of a relativistic AGN jet} \\end{figure}   The negligibility of pair absorption due to the interaction with the synchrotron and $\\gamma$-ray emission from the jet is checked self-consistently during our calculations.  Interactions of the jet pair plasma with dilute surrounding material will cause turbulent Alfv\\'en and Whistler waves. It has been shown by Achatz \\& Schlickeiser (1993) that a low-density, relativistic pair jet is rapidly disrupted as a consequence of the excitation of such waves. Thus, to insure stability of the beam over a sufficient length scale, we need that the density of pairs inside the jet exceeds the density of the surrounding material. In this case, pitch angle scattering on plasma wave turbulences leads to an efficient isotropization of the momenta of the pairs in the jet without destroying the jet structure. Thus, additional assumptions on our initial conditions are that the particle momenta are isotropically distributed in the rest frame of a new jet component (blob) and that the density of pairs in the jet $n_j \\gg n_b$ where $n_b$ is the density of the background material.  This study is devided into two papers. In the first (present) paper we investigate the details of electron/positron cooling due to inverse-Compton scattering and follow the pair distribution and photon spectra evolution during the first phase in which the system is dominated by heavy radiative losses. In the second paper we will consider the later phase of the evolution where collisional effects (possibly leading to thermalization) and reacceleration become important, and a plausible model for MeV blazars which follows directly from our treatment will be presented (B\\\"ottcher, Pohl \\& Schlickeiser, in prep.).  In section 2 of this first paper, we describe in detail how to calculate the energy-loss rates due to the various processes which we take into account and give useful approximative expressions for the inverse-Compton losses (as well scattering of accretion disk radiation as of synchrotron radiation), including all Klein-Nishina effects. In section 3, we discuss the relative importance of the various processes. The location of the $\\gamma$-ray photosphere for TeV $\\gamma$-rays is briefly outlined in section 4. The technique used to follow the evolution of the pair distributions is described in section 5. In section 6, we describe how to use the pair distributions resulting from our simulations in order to calculate the emanating $\\gamma$-ray spectra, and in section 7 we discuss general results of our simulations giving a prediction for GeV -- TeV emission from $\\gamma$-ray blazars which due to the lack of sensitivity of present-day instruments in this energy range could not be observed until now. Only two extragalactic objects have been detected as sources of TeV emission, namely Mrk~421 (Punch et al. 1992) and Mrk~501 (Quinn et al. 1995). In Section 8, we use our code to fit the observational results on the broadband emission during the TeV flare of Mrk~421 in May 1994 and on its quiescent flux. We summarize in section 9. ",
        "conclusions": "We gave a detailed discussion of the various radiative energy loss mechanisms acting in relativistic pair plasmas and concentrated on the application to jets from active galactic nuclei. In the first phase after the acceleration or injection of pairs radiative cooling due to inverse-Compton scattering is the dominant process. Here, Klein-Nishina effects are much more pronounced for inverse-Compton scattering of external radiation (from the central accretion disk) than for the SSC process since the synchrotron spectrum extends over a much broader energy distribution, allowing for efficient Thomson scattering at all particle energies.  We developed a computer code which allows to follow the evolution of the energy distributions of electrons/positrons self-consistently, accounting for all Klein-Nishina effects and for the time-dependence of the synchrotron, SSC and external radiation fields which contribute to the radiative cooling of the pairs and computing the instantaneous $\\gamma$-ray production from the resulting pair distributions. Since the intrinsic cooling timescales are much shorter than the time resultion of present-day instruments, we calculated time-integrated spectra and compared them to observations.  We found some general results from a series of different simulations:  (a) Due to the short cooling times involved in the SSC radiation mechanism, implying that we only see time-integrated SSC emission leading to an overprediction of hard X-ray flux from blazars, we favor EIC to be the dominant radiation process in $\\gamma$-ray blazars.  (b) Detailed simulations show that the EIC spectrum of a cooling ultrarelativistic pair distribution with maximum energies implying EIC scattering in the Klein-Nishina regime is harder in the case of pure Thomson scattering where we recover the classical results of Dermer \\& Schlickeiser (1993 a).  (c) The low-energy cutoff in the synchrotron spectra from dense pair jets is {\\it not} determined by synchrotron self-absorption, but by the Razin-Tsytovich effect. Only when the jet widens up, radio emission at $\\nu \\ukl 10^{11}$ Hz can escape the jet."
    },
    "9604/astro-ph9604151_arXiv.txt": {
        "abstract": "Using the near infrared fluxes of local galaxies derived from Cosmic Background Explorer (COBE)/Diffuse Infrared Background Experiment (DIRBE)\\footnote{COBE data sets were developed by the NASA Goddard Space Flight Center under the guidance of the COBE Science Working Group and were provided by the NSSDC.}  J(1.25 $\\um$) K (2.2 $\\um$) \\& L (3.5 $\\um$) band maps and published Cepheid distances, we construct Tully-Fisher diagrams for the nearby galaxies.  The measured dispersions in these luminosity-linewidth diagrams are remarkably small: $\\sigma_J = 0.09$ magnitudes, $\\sigma_K = 0.13$ magnitudes, and $\\sigma_L = 0.20$ magnitudes.  These dispersions include contributions from both the intrinsic Tully-Fisher relation scatter and the errors in estimated galaxy distances, fluxes, inclination angles, extinction corrections, and circular speeds.  For the J and K bands, Monte Carlo simulations give a 95\\% confidence interval upper limit on the true scatter in the Tully-Fisher diagram of $\\sigma_J \\le 0.35$ and $\\sigma_K \\le 0.45$.  We determine Milky Way's luminosity and place it in the Tully-Fisher diagram by fitting a bar plus exponential disk model of the Milky Way to the all-sky DIRBE maps.  For ``standard'' values of its size and circular speed (Sun-Galactic center distance $R_0 = 8.5 \\kpc$ and $\\Theta_0 =220 \\kms$), the Milky Way lies within $1.5 \\sigma$ of the TF relations.  We can use the TF relation and the Cepheid distances to nearby bright galaxies to constrain $R_0$ and $\\Theta_0$: $-\\log\\left(R_0 / 8.5 \\kpc\\right) +1.63\\log\\left(\\Theta_0 / 220 \\kms\\right) = 0.08 \\pm 0.03$. Alternatively, we can fix the parameters of the Galaxy to their standard values, ignore the Cepheid zero-point, and use the Tully-Fisher relation to determine the Hubble Constant directly: $H_0 = 66 \\pm 12$ km/s/Mpc.  We have also tested the Tully-Fisher relation at longer wavelengths, where the emission is dominated by dust.  We find no evidence for a Tully Fisher relation at wavelengths beyond 10$\\mu$m.  The tight correlation seen in L band suggests that stellar emission dominates over the 3.3 $\\mu$m PAH emission. ",
        "introduction": "The Tully-Fisher (TF) relation between the luminosity and linewidth of spiral galaxies (Tully \\& Fisher 1977) has been used extensively as a distance indicator and to map large scale flows of galaxies (cf., Strauss \\& Willick 1995 and Jacoby et al. 1992 for reviews). Its usefulness as a distance indicator is limited by the intrinsic scatter of the relation. This scatter is lowest in redder bands where dust extinction is low:  H band ($1.65 \\um$) (Aaronson, Huchra \\& Mould 1979, Aaronson, Mould \\& Huchra 1980, Aaronson et al. 1989, Freedman 1990, Pierce \\& Tully 1992) and I band ($0.90 \\um$) (Bernstein et al. 1994). In this paper, we extend the TF relation to longer wavelengths. The DIRBE experiment, with its excellent calibration and large beam width, is ideal for measuring the total flux of galaxies in the local galaxies. We describe our galaxies data set and present the results of our analysis in section 2.  The Milky Way has often been deemed unsuitable for zero point calibration of the TF relationship, mainly because of difficulties in estimating its total luminosity.  Such difficulties can be overcome at infrared wavelengths, where dust absorption is small.  In section 3, we use a three-dimensional model of the Milky Way based on the DIRBE J, K, and L band maps to obtain a measurement of the Galaxy's luminosity.  This luminosity can be used to place the Milky Way on the TF diagram with nearby galaxies, to constrain Galactic parameters, and to obtain an independent calibration of the Hubble constant. ",
        "conclusions": "We find that TF relation extends to longer wavelengths (2.2 $\\mu {\\rm m}$ and 3.5 $\\mu {\\rm m}$) than previously explored. The extinction corrections (cf. Mathis 1990) at 2.2 $\\mu {\\rm m}$ and $3.5\\mu {\\rm m}$ are about half and one-third as large as for the H-band at $1.65\\mu {\\rm m}$, so these bands may be usefully exploited for distance estimation. With the advent of imaging IR instruments and two major near-IR sky surveys (2MASS and DENIS), there is also potential to estimate the distances and hence the peculiar velocity flows for many more galaxies ( $\\sim 10^6$) in a greater part of the sky and nearer to the plane of the Milky Way.  Using Cepheid distances for galaxies in the Local Group (not including the Milky Way) and assuming standard Galactic parameters, we find that the Galaxy obeys the TF relation for nearby galaxies with Cepheid distances. This consistency is an independent check of the distance scale."
    },
    "9604/astro-ph9604008_arXiv.txt": {
        "abstract": "Much of the far-UV emission from elliptical galaxies is thought to arise from extreme horizontal branch stars and related objects. Only about 10\\% of the stellar population needs to evolve through this phase even in galaxies with the strongest UV upturn. However it is not yet clear if this population represents the extreme low-metallicity or high-metallicty tail of the distribution, or rather arises from the overall population through some metallicity-insensitive mechanism that causes increased mass loss in a small fraction of RGB stars. We investigate the utility of far-UV line strengths for deciding between these possiblities.  Complications include the fact that the line strengths reflect both the temperature distribution and the metallicity distribution of the stars, that there may be abundance anomalies introduced on the RGB, and that metals are likely to be redistributed by gravitational settling and radiative diffusion in the atmospheres of hot high-gravity stars. Line-strength measurements from Astro-2 HUT spectra are considered in this context. ",
        "introduction": "Giant elliptical galaxies show a large variation in the ratios of their far-UV to optical fluxes. Shortward of 2000{\\AA}, most ellipticals have spectra that rise in $f_\\lambda$ toward shorter wavelengths. This hot component has been known since the early days of space astronomy \\cite{CW79}, but it is only within the last two years that observations and theory seem to be converging on a consensus that the dominant component in UV bright galaxies is extreme horizontal branch (EHB) stars and their evolutionary progeny \\cite{FD93,DOR95,BCF94,BFD95}. This conclusion stems from the rather cool temperature (25000 K) derived for the dominant component in NGC1399 \\cite{Ferg91L}, and from computations that indicate that EHB stars can provide enough far-UV photons over their lifetimes to produce the elliptical galaxy fluxes, while other candidates such as PAGB stars cannot \\cite{GR90,DOR95}.  While it seems clear that EHB stars provide the far-UV flux, it is not at all clear how they got there. The general trend observed for globular clusters is that the horizontal branch (HB) becomes redder with increasing metallicity. Elliptical galaxies are even more metal rich than Galactic globular clusters; they must somehow be able to buck the trend. The HB morphology depends on age, metallicity, helium abundance, and the amount of mass loss on the red giant branch. The helium-burning core in HB stars has a mass (0.5 $M_\\odot$) that is nearly independent of these parameters. The position of stars along the HB thus depends on the envelope mass, which in turn depends on the main-sequence mass and the amount of mass lost during the RGB phase. EHB stars (those with $T_{eff} > 20000$~K) have envelope masses less than 0.05$M_\\odot$. Hence they must arise from stars that have lost nearly all the mass it was possible for them to lose and still ignite helium in their cores.  There are several plausible ways to produce a minority population of EHB stars in elliptical galaxies.  First, the giant elliptical galaxies may in general be {\\it older} than the galactic globular clusters. This argument is a natural extension of the interpretation of the second parameter effect in globular clusters as being due to variations in age \\cite{Lee94,PL95}. The EHB stars in this model represent the extreme metal-poor tail of the metallicity distribution, and are 2-4 Gyr older than the most-metal poor globular clusters. They show up in giant elliptical galaxies, which are on average metal rich, because these galaxies formed first, and hence have the oldest stars.  Second, elliptical galaxies may have high helium abundance $Y$ \\cite{GR90,BCF94,YADO95}. At fixed age and metallicity, the main-sequence lifetime decreases with increasing $Y$. Observations of nearby star-forming galaxies and the galactic bulge hint at a rather steep relation ($\\Delta Y / \\Delta Z > 2$) between helium abundance and metallicity \\cite{Pagel89p201,Renzini94}. If this is the case, then old metal-rich populations may have EHB stars, even with standard RGB mass-loss rates. The required ages ($>7$ Gyr) are not as extreme as in the Lee model. In these models, the EHB stars arise from the extreme high-metallicity tail of the abundance distribution.  Third, some other process may act to increase mass loss in a small fraction of the population. For example, EHB stars could arise from stars in close binary systems that have shed their envelopes during interactions with their companions, our they could arise only from stars with high rotation rates. Such mechanisms could produce EHB populations from anywhere in the abundance distribution, but might be enhanced in giant ellipticals through some secondary effect (for example binary fraction might somehow depend on galaxy metallicity or velocity dispersion). In this case, the EHB stars may come closer to reflecting the mean metallicity of the stellar population.  To distinguish between these possibilities, it is important to try to get some direct measure of the metallicities of the far-UV emitting population, as suggested for example by \\citeN{PL95} and \\citeN{YADO95}. It is now possible to attempt this with new observations from the Hopkins Ultraviolet Telescope (HUT), obtained March 1995 during the Astro-2 mission. In the rest of this contribution we summarize our preliminary attempts to do this, and outline some arguments why such efforts may in the end yield ambiguous results. ",
        "conclusions": ""
    },
    "9604/astro-ph9604134_arXiv.txt": {
        "abstract": "\\noindent  We analyse the ROSAT PSPC spectrum  of 19 X--ray selected Narrow Emission Line Galaxies (NELGs) discovered during the optical identification of  sources in the ROSAT UK Deep Survey. Their properties are compared to those of broad line Active Galactic Nuclei (AGN) in the same sample.  Counts in three spectral bands have been extracted for all the sources, and have been fitted with a power-law model assuming the Galactic value for $N_{H}$. The average slope of NELGs is $\\alpha$= 0.45 $\\pm$ 0.09 , whilst for the AGN it is $\\alpha $= 0.96 $\\pm$ 0.03. The power-law model is a good fit for $\\sim$ 90\\% of NELGs and $\\sim$ 75\\% of AGN.  Recent work shows that the fractional surface density of NELGs increases with respect to AGN at faint fluxes. Thus they are expected to be an important component of the residual soft ($< 2$ keV) X--ray background.  The slope of the X--ray background ($\\alpha \\sim $ 0.4, 1-10 keV) is harder than that of AGN ($\\alpha \\sim$ 1) but our results show that it is consistent with the summed spectrum of the NELGs  in the deep survey ($\\alpha \\sim$ 0.4). This may finally reconcile  the spectrum of the background with the properties of the sources that constitute it. ",
        "introduction": "The soft ($< 2$ keV) X--ray background is thought to arise from the integrated signal of individual unresolved sources (Fabian and Barcons, 1992, and references therein). Although broad line AGN are known to be the main contributors to the soft X--ray background at higher fluxes (Shanks \\etal 1991; Boyle \\etal 1994), their spectral shape (energy index, $\\alpha \\sim 1$, Maccacaro \\etal 1988) seems too steep to match that of the background ($\\alpha \\sim$ 0.4, 1-10~keV, Gendreau \\etal 1995, Chen \\etal 1996). Hence a new population of harder and fainter sources, with a steep Log N-Log S and which do not contribute significantly above fluxes $\\sim$ 10$^{-14}$ erg cm$^{\\rm -2}$ s$^{\\rm -1}$, has been sought in order to resolve this spectral paradox (Hasinger \\etal 1993). Narrow Emission Line Galaxies (NELGs) seem to be an attractive candidate for this new contributor to the X--ray background at fainter fluxes ( Jones \\etal 1995a, Boyle \\etal 1995).  NELGs are generally defined as galaxies which possess only emission lines with FWHM $< 1000$ km s$^{\\rm -1}$ in their optical spectra. They can be substantially brighter than normal galaxies in X--rays. Thus for example Fabbiano (1989) finds X-ray luminosities in the range $\\sim$ 10$^{38}$ - 10$^{42}$ erg s$^{-1}$ (0.2 - 3.5 keV) for normal galaxies, while our sample of NELGs yields $\\sim$ 10$^{41}$ - 2 x 10$^{43}$ erg s$^{-1}$ when extrapolated to the same energy band.  The number ratio of NELGs has recently been found to increase at faint X-ray fluxes (Jones \\etal 1995a, Boyle \\etal 1995), leading to the suggestion that they may be important contributors to the unresolved fraction of the X--ray background.  However, in order to solve the soft X--ray spectral paradox, NELGs would need to have a flatter spectrum than the X--ray background, to compensate for the steeper slope of AGN.  In this paper we examine X--ray data on faint NELGs to determine if their integrated spectrum can reproduce the spectral shape of the X--ray background. We present a total sample of 19 NELG from the UK Deep Survey, and compare their spectra to that of broad line AGN from the same survey, and to the spectrum of the soft X--ray background. In section 2, the deep survey sample is described. In section 3 we explain the data reduction and analysis process. We report and discuss our results in section 4, and in section 5 we present our conclusions. A preliminary report of this work was given in Romero-Colmenero \\etal (1996).  \\footnotesep=-1.5in \\begin{table*} \\caption {Spectral results on Deep Survey NELGs} \\begin{minipage}[t]{10in} \\begin{tabular}{c c c c c c r @{\\hspace{1mm}}c @{\\hspace{1mm}} l}\\hline  ID & extraction & z & $\\alpha$ & Hard Counts \\footnote[1]{Corrected for PSF and vignetting. The error is the uncertainty in the counts.  \\noindent Note: All errors quoted in the table are 1 $\\sigma$ deviation.} & Fx (0.5-2 keV) & \\multicolumn{3}{c}{Lx (0.5-2 keV)} \\\\ Number & radius (arcsec)& & & PHA channels 52-201 & 10$^{\\rm-14}$ erg cm$^{\\rm -2}$ s$^{\\rm -1}$ & \\multicolumn{3}{c}{10$^{\\rm40}$ erg s$^{\\rm -1}$} \\\\  \\hline 032 & 43 & 0.068 & -2.41 $\\pm$ 0.40 & 101.98 $\\pm$  13.03 & 1.96 $\\pm$ 0.43 &   34& $\\pm$ &  7 \\\\ 036 & 41 & 0.235 &  0.18 $\\pm$ 0.26 &  89.68 $\\pm$  11.49 & 1.31 $\\pm$ 0.20 &  328& $\\pm$ & 49 \\\\ 042 & 18 & 0.366 & -0.15 $\\pm$ 0.88 &  50.71 $\\pm$   9.78 & 0.76 $\\pm$ 0.22 &  430& $\\pm$ &116 \\\\ 043 & 25 & 0.382 &  0.81 $\\pm$ 0.25 &  69.76 $\\pm$  13.32 & 0.99 $\\pm$ 0.20 &  827& $\\pm$ &167 \\\\ 047 & 18 & 0.364 &  0.49 $\\pm$ 0.56 &  36.48 $\\pm$   8.79 & 0.52 $\\pm$ 0.15 &  357& $\\pm$ &101 \\\\ 051 & 54 & 0.062 &  0.78 $\\pm$ 0.19 &  64.77 $\\pm$  10.56 & 0.92 $\\pm$ 0.16 &   16& $\\pm$ &  3 \\\\ 060 & 26 & 0.580 &  0.51 $\\pm$ 0.37 &  38.47 $\\pm$   7.60 & 0.55 $\\pm$ 0.12 & 1062& $\\pm$ &237 \\\\ 067 & 54 & 0.554 & -1.85 $\\pm$ 1.77 &  44.71 $\\pm$   9.50 & 0.81 $\\pm$ 0.38 &  497& $\\pm$ &234 \\\\ 085 & 48 & 0.304 & -0.29 $\\pm$ 1.35 &  19.89 $\\pm$   7.50 & 0.30 $\\pm$ 0.17 &  113& $\\pm$ & 64 \\\\ 093 & 49 & 0.590 &  0.57 $\\pm$ 0.47 &  36.88 $\\pm$   8.81 & 0.53 $\\pm$ 0.14 & 1085& $\\pm$ &295 \\\\ 094 & 54 & 0.061 & -0.41 $\\pm$ 1.38 &  24.88 $\\pm$   8.48 & 0.39 $\\pm$ 0.21 &    6& $\\pm$ &  3 \\\\ 103 & 32 & 0.200 &  1.31 $\\pm$ 0.32 &  34.03 $\\pm$   9.81 & 0.48 $\\pm$ 0.14 &  100& $\\pm$ & 30 \\\\ 117 & 54 & 0.064 & -0.33 $\\pm$ 1.86 &  13.00 $\\pm$   7.67 & 0.22 $\\pm$ 0.18 &    4& $\\pm$ &  3 \\\\ 121 & 27 & 0.310 & -0.13 $\\pm$ 0.93 &  27.10 $\\pm$   7.06 & 0.41 $\\pm$ 0.15 &  166& $\\pm$ & 61 \\\\ 127 & 23 & 0.250 &  0.88 $\\pm$ 0.48 &  22.07 $\\pm$   6.40 & 0.31 $\\pm$ 0.10 &  103& $\\pm$ & 33 \\\\ 131 & 24 & 0.576 &  0.87 $\\pm$ 0.54 &  22.87 $\\pm$   7.21 & 0.32 $\\pm$ 0.11 &  720& $\\pm$ &246 \\\\ 132 & 28 & 0.223 &  1.63 $\\pm$ 0.26 &  19.18 $\\pm$   5.77 & 0.25 $\\pm$ 0.08 &   74& $\\pm$ & 24 \\\\ 134 & 49 & 0.250 &  0.38 $\\pm$ 0.83 &  26.33 $\\pm$   8.09 & 0.38 $\\pm$ 0.15 &  113& $\\pm$ & 43 \\\\ 135 & 32 & 0.520 &  0.93 $\\pm$ 0.58 &  17.15 $\\pm$   6.46 & 0.24 $\\pm$ 0.10 &  423& $\\pm$ &172 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "We have integrated the X--ray spectrum of 19 NELGs taken from the UK Deep Survey and compared this with the spectrum of 23 AGN of similar count rate in the same sample. We find that the NELG spectrum is harder than that of the AGN at more than 3~$\\sigma$ confidence.  Moreover, the mean spectral slope of the NELGs ($\\alpha=0.45 \\pm 0.09$) is consistent with the slope of the X--ray background between 1 and 10 kev ($\\alpha = 0.4\\pm0.1$; Gendreau \\etal 1995) whereas the slope of the AGN at similar count rates is not ($\\alpha \\sim 1.0$).  This work is important for understanding the origin of the soft ($< 2$ keV) X--ray background. Extrapolation of the source number counts suggests that NELGs are the dominant source population at fluxes four times fainter than the deep survey flux limit (i.e. at fluxes of $\\sim 5\\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ in the 0.5-2 keV band; Jones \\etal 1995a, McHardy \\etal 1996).   The increasing number of these sources at low fluxes and their spectral properties as shown in this paper, taken together, can reproduce both the flux and spectrum of the X--ray background.  This result adds considerable weight to the idea that NELGs are  the major contributor to the  residual unresolved soft X-ray background.   \\vskip 0.4cm"
    },
    "9604/astro-ph9604120_arXiv.txt": {
        "abstract": "We report the detection using data from the {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca) of a hard X-ray source in the vicinity of the radio pulsar PSR B1853+01, which is located within the supernova remnant (SNR) W44. PSR B1853+01, a 267~ms pulsar, has to date been detected only in the radio band.  Previous observations at soft X-ray energies (e.g., with  \\rosat HRI) have failed to detect any significant X-ray emission (pulsed or unpulsed) from the pulsar. In addition, no high energy emission (${_>\\atop^{\\sim}}$4~keV) has been detected previously from W44.  Over the 0.5--4.0~keV band, the \\asca data show soft thermal emission from W44, with a morphology very similar to that observed earlier by \\ein  and {\\it ROSAT}. In the high-energy band (4.0--9.5 keV), the SNR is, for the most part, invisible, although a source coincident with the position of PSR B1853+01 is evident. The observed \\asca spectra are consistent with a power-law origin (photon index $\\sim$2.3) for the X-ray emission from this source at a flux level (flux density $\\sim$0.5 $\\mu$Jy at 1 keV) consistent with previous upper limits. The maximum allowed size for the source is determined directly from the \\asca data ($<$5$^\\prime$), while the minimum size is derived from the nondetection of a point source in the \\rosat HRI data (${_>\\atop^{\\sim}}$30$^\\prime$$^\\prime$). Timing analysis of the hard X-ray source  failed to detect pulsations at the pulsar's period. Based on these lines of evidence, we conclude that the new hard source in W44 represents an X-ray synchrotron nebula associated with PSR B1853+01, rather than the beamed output of the pulsar itself.  This discovery adds W44 to the small group of previously known plerionic SNRs. This nebula lies at the low end of, but is consistent with, the correlation between X-ray luminosity and pulsar spin-down energy loss found for such objects, lending further support to our interpretation. ",
        "introduction": "The radio pulsar PSR B1853+01 was discovered by Wolszczan, Cordes, \\& Dewey (1991) inside the radio shell of the supernova remnant W44.  The estimated distances to the pulsar and the supernova remnant (SNR) are both on the order of 3~kpc, and the age deduced from the spin-down of the pulsar (20,000~yrs) is compatible with the dynamical age of the remnant inferred from X-ray observations (Harrus \\& Hughes 1994; Harrus et al. 1996). Previous work has failed to detect X-ray emission from PSR B1853+01 in any of the expected forms: pulsed or unpulsed emission from the pulsar itself, or nonthermal emission from an associated synchrotron nebula.  No high energy tail was seen in the \\ginga spectrum, leading to a 3\\,$\\sigma$ upper limit of $3.6\\times 10^{-12}$~ergs~cm$^{-2}$~s$^{-1}$ for the 2--10~keV flux of a Crablike power-law component ($dN/dE \\sim E^{-2.1}$) from W44 (Harrus \\& Hughes 1994; Harrus et al. 1996). Although there was a report of a high-energy component to the X-ray spectrum of W44 from \\exosat data (Rho et al.\\ 1994), an earlier analysis (Jones, Smith, \\& Angellini 1993) had already indicated that this component arose from contamination of the data by particle background events and that, in fact, the SNR was undetected above 5~keV. Moreover neither a point source nor pulsed emission was seen at the pulsar position in the \\rosat PSPC observation in the soft X-ray band (Rho et al.\\ 1994).  In this Letter, we report on the discovery of the anticipated nonthermal X-ray emission from the vicinity of PSR B1853+01 using spectral, imaging, and timing observations from {\\it the Advanced Satellite for Cosmology and Astrophysics} ({\\it ASCA}) (Tanaka, Inoue, \\& Holt 1994). ",
        "conclusions": "We have presented the results of morphological and spectral studies of the contribution from PSR B1853+01 to the X-ray emission of the SNR W44.  High-energy emission in both the GIS and the SIS is found at the position of the pulsar; the spectra from both instruments require a power-law component to model the observed emission.  A lower limit to the spatial extent of the emission region can be derived from the average \\rosat HRI surface brightness around the pulsar position, which, after subtraction of the nominal HRI background, is $\\sim$4$\\times$10$^{-3}$ counts~arcmin$^{-2}$~s$^{-1}$.  Based on our spectral analysis, we expect an HRI count rate from the power-law component of (2.0--3.5)$\\times$10$^{-3}$ counts~s$^{-1}$.  This translates into a minimum angular size of $\\sim$30$^{\\prime\\prime}$. Using the known distance to the SNR, we deduce an emission volume of $2.78\\,D_{\\rm 3\\,kpc}^3\\theta^3\\ {\\rm pc}^3$, where $D_{\\rm 3\\,kpc}$ is the distance to the SNR in units of 3~kpc, and $\\theta$ is the angular size of the nebula expressed in arcminutes.  An upper limit on the source size can be obtained from the \\asca data.  We have simulated images with Gaussian source profiles of various widths and convolved each with the PSF of the mirror and the detector, weighted according to the spectrum from the pulsar region.  We then cross-correlated the simulated images with the GIS image and deduced an upper limit to the source radius of $\\sim$5$^\\prime$.  The inferred unabsorbed flux from the power-law component extrapolated to the \\rosat energy band (0.4--2.0~keV) is 1.9$_{-1.4}^{+9.1}$$\\times$10$^{-12}$ ergs~cm$^{-2}$~s$^{-1}$; the value of 1.2$\\pm$0.3$\\times$10$^{-12}$ ergs~cm$^{-2}$~s$^{-1}$, measured for the 2--10 keV band flux, is in agreement with the upper limit of 3.6$\\times$10$^{-12}$ ergs~cm$^{-2}$~s$^{-1}$ obtained using the \\ginga spectrum (Harrus \\& Hughes 1994; Harrus et al. 1996). The unabsorbed X-ray luminosity in the \\ein band (0.2--4.0~keV) is (4$_{-3}^{+30}$)$\\times10^{33}\\,\\,D_{\\rm 3\\,kpc}^2$~ergs~s$^{-1}$, which should be compared to the luminosity predicted by the empirical \\.{E}/L$_{X}$ relation (Seward \\& Wang 1988) of 7.7$\\times$10$^{32}$ ergs~s$^{-1}$.  Frail et al.\\ (1996) present new radio data on the nebula surrounding PSR B1853+01 and, using the \\ginga\\ upper limit, construct a spectrum covering 10 decades in frequency from the radio to the X-ray bands. From these data, the equipartition value of the nebular magnetic field, ${\\rm B}_n$, and the energy in relativistic electrons, $E_e$, can be estimated by assuming that the break in the spectrum between the radio and X-ray regimes is due to synchrotron losses (Pacholczyk 1970, p~169). Our \\asca measurements of the X-ray flux and spectral index of the nebula allow better determination of the break frequency, that is, the frequency where the extrapolated X-ray and radio power-law spectra intersect. The best-fit spectral values indicate a value of $\\nu_{\\rm B} \\sim 2 \\times 10^{13}$ Hz. (Note that the break frequency of the Crab Nebula is $\\sim$$10^{13}$ Hz.)  We find for the nebular magnetic field ${\\rm B}_n \\simeq 70\\,\\mu{\\rm G}\\, (V_{\\rm R} / 3.1\\times10^{55}\\,{\\rm cm}^3)^{-2/7}$ (using the estimated volume of the radio nebula, $V_{\\rm R}$, and the luminosity integrated from $10^7$ Hz to $\\nu_{\\rm B}$ which is $2.3\\times 10^{34}$ ergs s$^{-1}$) and for the energy in electrons $E_e \\simeq 10^{46}\\,{\\rm ergs}\\, ({\\rm B}_n / 70\\,\\mu{\\rm G})^{-3/2}$, values that are consistent with other plerionic SNRs of comparable luminosity.  The lifetime of the electrons giving rise to the radio emission, $\\sim$20,000 yr, is of the same order as the age of the SNR and pulsar, while the electrons giving rise to the X-ray emission are quite short-lived, $\\sim$120 yr.  Consequently, the X-ray synchrotron nebula should be significantly smaller than the radio nebula and should be located close to the pulsar, as indicated by our \\asca\\ data. In particular, given the parameters derived above, we expect essentially no nonthermal X-ray emission from the bright diffuse radio-emitting nebula north of the pulsar, since the pulsar would have traversed this region over 2000 years ago, given the projected distance and the inferred transverse motion, 375 km s$^{-1}$, of PSR B1853+01 (Frail et al.\\ 1996).  A follow-up X-ray observation of the pulsar region with high spatial resolution should be able to confirm this conjecture and, in addition, provide a number of new constraints on the energetics of the synchrotron nebula.  In any event, the \\asca\\ data presented here provide strong evidence for the oldest synchrotron nebula yet detected, a result that should aid considerably in our attempts to understand the interaction of pulsar relativistic winds with their environments.  We thank Alex Wolszczan for providing an accurate position and a recent ephemeris for PSR B1853+01 and we thank E.~Churazov, M.~Gilfanov and A.~Finoguenov for allowing us to use their software for merging \\asca images.  Pat Slane, Olaf Vancura, Paul Callanan, Didier Barret, Rick Harnden, and Fred Seward are acknowledged for their helpful comments and discussions.  This research was partially supported by NASA Grant NAG 5-2605 and the Smithsonian Institution, through the Smithsonian Predoctoral Fellowship program. This is contribution 595 of the Columbia Astrophysics Laboratory."
    },
    "9604/hep-ph9604286_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmic rays of galactic and extra-galactic origin will interact in high energy collisions with the interstellar medium  of our galaxy and produce secondary particles \\cite{Domokos93,Berezinsky93,DePaolis95}. These are mainly mesons that decay and give rise to a flux of muons, neutrinos and photons. The very low density of the interstellar medium imply that the interaction lengths of the secondary particles is long compared to their decay length, such that the mesons will decay before loosing energy in secondary interactions. This is also the case for the muons which will decay giving neutrinos. This is the opposite of the situation for cosmic ray particles interacting in the Earth's atmosphere \\cite{GIT}, where meson typically loose energy in interactions before decaying. The fluxes of high energy neutrinos and photons from the interstellar medium could therefore be larger than those from the atmosphere, although the initial production rate of mesons is smaller.  A measurement of these fluxes could potentially give valuable information  about the distribution of matter and cosmic rays in the galaxy, which  could be of great importance in determining the origin of the cosmic rays.  Understanding the flux from the disc of the Milky Way could also be the  starting point in a search for baryonic dark matter in a spherical halo  around the Milky Way \\cite{Domokos93,DePaolis95}.  In addition, these fluxes from the interstellar medium constitute a background in searches for other, more spectacular cosmic sources and must therefore be known in order to extract the desired signal. For example, there is much interest in neutrinos from Active Galactic Nuclei (AGN) (see for example \\cite{Stecker95} and references therein).  In this paper we present a realistic calculation of the neutrino and photon fluxes to be expected from cosmic ray particle interactions in the interstellar medium. These fluxes are derived from complete events obtained by detailed Monte Carlo simulations based on state-of-the-art models for high energy particle collisions.  In section 2 we present the model including the cosmic ray energy  spectrum, the interstellar matter distribution in the Milky Way and the model for particle production in high energy elementary particle collisions. The resulting fluxes are presented in section 3, where also the attenuation of the photon flux due to interactions with the cosmic background radiation is demonstrated to give a significant effect. Section 4 considers the mentioned dark matter halo and the fluxes that would arise from it. We end, in section 5, with a discussion and comparison with estimated neutrino fluxes from active galactic nuclei. ",
        "conclusions": "\\label{sec:discussion}  The neutrino fluxes shown above in Fig.\\,\\ref{fig:nuflux}  are given for our basic `unit column'  of interstellar matter. Using these one may obtain the integrated  flux from any direction. In Fig.\\,\\ref{fig:compflux2} we show our resulting muon neutrino flux (full curves) from the direction towards  the center of the Milky Way (highest flux) and orthogonal to the galactic  plane (lowest flux).  In Fig.\\,\\ref{fig:compflux2}a  we compare with the original results of two earlier calculations by   Domokos \\etal\\ \\cite{Domokos93} and by Berezinsky \\etal\\ \\cite{Berezinsky93}. The differences between these curves are mainly related to the  different assumptions concerning the interstellar matter density profile and the normalisation of the cosmic ray spectrum. Thus, they illustrate the  uncertainty due to these inputs to the different calculations.  \\ffig{b}{comparison2.eps}{14cm} {\\it  The $E^3$-weighted fluxes of interstellar muon neutrinos in the direction  towards the center of the Milky Way (upper set of curves) and orthogonal to the galactic plane (lower set of curves).  Comparison of our results (full lines) with those by Domokos \\etal \\protect\\cite{Domokos93} (dashed lines) and by Berezinsky \\etal \\protect\\cite{Berezinsky93} (dotted lines).  In (a) the original results are used, whereas in (b) the results of \\protect\\cite{Domokos93} and \\protect\\cite{Berezinsky93} are modified to have the same galactic density profile and cosmic ray flux as in our  calculation. } {fig:compflux2}  In order show other differences between these different calculations we have  recalculated the results of \\cite{Domokos93} and \\cite{Berezinsky93}  using their formalisms but with our density and cosmic ray parametrisations. The results are shown Fig.\\,\\ref{fig:compflux2}b and demonstrate a close agreement for energies up to $10^5$--$10^6\\: GeV$. At higher energies there are, however, significant differences due to the treatment of the change in slope of the cosmic ray spectrum, {\\it cf.}\\ Eq.\\,(\\ref{eq:initflux}). In \\cite{Berezinsky93} energies above the `knee' are not considered, and the naive extrapolation to higher energies reults in an overestimated flux. About a factor three excess at the highest energies is expected based on  an estimate using the analytic method with spectrum-weighted $Z$-moments  \\cite{GIT,Gaisser90,Lipari}, in agreement with the effect seen in Fig.\\,\\ref{fig:compflux2}b. In \\cite{Domokos93} this change of slope is included, but its effect on secondary particles cannot be fully taken into account since their calculation is based on an analytic method. One must here make assumptions about from what average primary energy a given neutrino comes, and thereby specify how the `knee' from the change in the cosmic ray energy spectrum is transported into a `knee' in the final neutrino spectrum.  This problem does not occur in our Monte Carlo method, since the neutrino  spectra are here obtained through an event-by-event simulation taking the  fluctuations into account. Our results are therefore more reliable in this respect.  There is also a contribution to the interstellar flux from semileptonic decays of charmed and heavier hadrons.  We have estimated this based on the analytic method with spectrum-weighted  moments \\cite{GIT,Gaisser90,Lipari} using our previoulsy calculated $Z$-moments for charm particle production and decay \\cite{GIT}. Taking also the contribution from muon decay into account, we find a muon neutrino flux which is only contributing about $2\\cdot10^{-4}$ to the interstellar flux and a factor two higher contribution  for the electron neutrino flux (due to a lower interstellar electron neutrino flux). The smallness of this charm contribution justifies the neglect of it in our Monte Carlo treatment.  From Fig.\\,\\ref{fig:compflux2}, and these considerations, one can also conclude that the uncertainties originating from the particle physics input are not larger than other uncertainties. Our Monte Carlo calculation also confirmes that the approximation done with the analytic method are justfied except at the highest energies where special precautions must be taken.  \\ffig{bt}{comparison3.eps}{7cm}{\\it $E^3$-weighted fluxes of muon neutrinos at the Earth from cosmic ray interactions with the interstellar medium looking towards the galactic centre (upper solid curve) and orthogonal to the galactic plane (lower solid curve) as calculated in this study, and the flux from the Earth's atmosphere in the vertical (V) \\protect\\cite{GIT} and horizontal (H) \\protect\\cite{Volkova79} direction originating from conventional $\\pi ,K$-meson decays as well as from prompt (P) charm decays \\protect\\cite{GIT}. For comparison, the diffuse flux of muon neutrinos from  active galactic nuclei as predicted in \\protect\\cite{Protheroe} (dotted line) and in \\protect\\cite{Sikora} (dash-dotted line) are shown.}{fig:compflux3}  The fluxes from the interstellar medium should be compared with those from other sources in order to establish their observabelness and their significance  as a primary interest of study or as a potential background.  This is done in Fig.\\,\\ref{fig:compflux3} where our result on the muon  neutrino flux is compared with those from cosmic ray interactions in  the Earth's atmosphere and with two predictions of the diffuse flux from active galactic nuclei. The interstellar flux is seen to be considerably lower than the atmospheric  flux, except at the highest energies. It will therefore be very hard to observe with detectors on the Earth. At the highest energies  considered, the interstellar flux is of comparable magnitude as the  prompt atmospheric flux (from charm decays). The latter is almost direction independent up to $\\sim10^7\\,GeV$ such that the curve shown  for the vertical direction applies in essentially the whole energy range; the horizontal flux being slightly higher for the highest energies only.  Given the small absolute scale of the fluxes at these energies, the event  rate will be very low in the neutrino telescopes currently under construction. The interstellar flux becomes important compared to the vertical atmospheric flux only at energies above $10^4\\,GeV$. The corresponding event rate in a detector of $3\\cdot10^4\\,m^2$, \\eg {\\sc Amanda}, would be $\\sim0.5/year$ in a cone of opening angle $10^{\\circ}$ directed towards the centre of the Milky way.   Our calculated neutrino flux from the interstellar matter is significantly lower than the one from active galactic nuclei, as demonstrated in  Fig.\\,\\ref{fig:compflux3}, and should therefore not be a problematic  background for this extra-galactic source.  On the other hand, the atmospheric neutrino flux will be problematic in  this context, since it dominates in the lower energy range. At high energies where the atmospheric flux is lower, the absolute rate is very low.  Concerning the flux of electron neutrinos the situation is very similar to the case of muon neutrinos, except that the atmospheric background is significantly lower \\cite{GIT} and the interstellar flux is a factor $\\sim2$ lower than the corresponding muon neutrino flux (Fig.\\,\\ref{fig:nuflux}). Since the dominating production process in AGN is decays of pions followed by the decays of the muons, the electron neutrino flux from AGN's should also be a factor $\\sim2$ lower. Thus, the relative fluxes from these sources should be the same as for  muon neutrinos. There is, however, no good technique to detect high energy  electron neutrinos, so the prospects to search for such a signal is worse.  \\ffig{tb}{gyeild.eps}{7cm} {\\it Energy-weighted spectrum of photons from interactions of cosmic ray particles, having the specified energies, with the interstellar matter (\\ie the photon yield). } {fig:gyield}  The yield of photons in cosmic ray interactions with the interstellar matter has been considered before, \\eg  by Berezinsky \\etal\\ \\cite{Berezinsky93}. Our corresponding result obtained from the unattenuated photon flux in section 3.2 is  shown in Fig.\\,\\ref{fig:gyield}. It is about a factor two higher than that  of \\cite{Berezinsky93}, mainly due to the higher $\\pi^0$ multiplicity in our  complete event simulations.  The gamma flux can be used to set limits on the cosmic ray flux and the matter distribution in the galaxy.  The experimental situation is here, however, quite different from the case with the neutrino flux.  The gammas are detected either directly on satellite based experiments or  with ground based air shower arrays.  The signal rate will be higher than with neutrinos due to the larger  interaction cross section.  The flux of gammas is, however, very small ($\\sim 10^{-4}$) compared to the  inclusive cosmic ray flux dominated by protons. Differences in the air showers produced by gammas and hadrons may allow access to information on the gamma flux.  \\vspace{8mm} \\noindent {\\bf Acknowledgement:} We are grateful to Ph.~Jetzer for useful discussions."
    },
    "9604/astro-ph9604053_arXiv.txt": {
        "abstract": "COMPTEL imaging analysis revealed a patchy, asymmetric distribution of diffuse 1.8 \\MeV\\ emission along the Galactic plane, which is attributed to the decay of radioactive \\al26\\ in the ISM. If massive stars were the major source of Galactic \\al26, the 1.8 \\MeV\\ emission should be asymmetric and trace the spiral arms of the Galaxy, presumed site of massive star formation. Using model fits, we indeed find weak evidence in the COMPTEL data that the observed 1.8 \\MeV\\ emission is at least partly confined to spiral arms. We derive a total Galactic \\al26\\ mass of 2.5 \\Msol\\ from which at least 0.7 \\Msol\\ can be attributed to massive stars. ",
        "introduction": "Since the discovery of the 1.8 \\MeV\\ gamma-ray line emission from radioactive \\al26\\ by Mahoney \\etal\\ (1982), the questions of its origin and its distribution along the Galactic plane stimulated a wave of research (see review of \\cite{rf:pd95}). Core collapse supernovae (SNe), Wolf-Rayet (WR) stars, asymptotic giant-branch (AGB) stars, and O-Ne-Mg novae were suggested as possible sites of significant \\al26\\ creation. Early works assumed that the large mean lifetime of $\\tau_{26}\\sim10^6$ yr and the low \\al26\\ yield per source will lead to a smooth and symmetric distribution of 1.8 \\MeV\\ emission. Prantzos (1991) was the first who dropped the assumption of an axisymmetric source distribution in the Galactic plane if massive stars were the dominant \\al26\\ producers. He argued that star formation occurs predominantly inside the spiral arms, especially in the case of massive stars. Thus the 1.8 \\MeV\\ emission profile should reflect the structure of these arms which is thought to be asymmetric with respect to the Galactic centre. Previous analysis of COMPTEL data indeed revealed an asymmetry with more 1.8 \\MeV\\ emission from the southern (Galactic longitude $l$=180\\deg-360\\deg) than from the northern ($l$=0\\deg-180\\deg) Galaxy (\\cite{rf:diehl95}). Additionally, the 1.8 \\MeV\\ sky map shows lumpy emission and `hot spots'. Prantzos (1993) noted that some emission maxima coincide with the assumed tangential directions of Galactic spiral arms. Thus a detailed study of the spiral arm hypothesis is of interest. In this paper we will report on a comparison of COMPTEL phase I+II data (May 1991 - August 1993) to models of Galactic \\al26\\ distribution with special emphasis on spiral structure. ",
        "conclusions": "We have compared the COMPTEL 1.8 \\MeV\\ data from observation phases I+II to axisymmetric and spiral-arm models of Galactic \\al26\\ distribution. All models were detected at a significance level of $>20\\sigma$ above background. To first order, the observed 1.8 \\MeV\\ emission is well represented by axisymmetric models. However, details of the data like the north-south asymmetry and some regions with significant count excesses are better described by models which incorporate the Galactic spiral structure. Our best fit model holds a total Galactic \\al26\\ mass of $\\sim2.5$ \\Msol\\ from which at least $0.7$ \\Msol\\ are produced by massive stars. Thus, massive stars clearly contribute to the observed \\al26\\ in the Galaxy but we can certainly not exclude from this work that a large fraction of \\al26\\ is produced by low-mass AGB stars or novae."
    },
    "9604/astro-ph9604115_arXiv.txt": {
        "abstract": "Cooling of neutron stars with dipole magnetic fields is simulated using a realistic model of the anisotropic surface temperature distribution produced by magnetic fields. Suppression of the electron thermal conductivity of outer stellar layers across the field increases thermal isolation of these layers near the magnetic equator. Enhancement of the radiative and longitudinal electron thermal conductivities in quantizing magnetic fields reduces thermal isolation near the magnetic poles. The equatorial increase of the isolation is pronounced for $B \\ga 10^{10}$ G, while the polar decrease -- for $B \\ga 10^{12}$ G. The effects compensate partly each other, and the actual influence of the magnetic fields on the cooling is weaker than predicted by the traditional theories where the equatorial effects have been neglected. ",
        "introduction": "% Thermal X-ray radiation has been observed recently with $ROSAT$ from at least four cooling neutron stars: PSR 0833-45, PSR 1055-52, PSR 0656+14 and Geminga (see, e.g., \\\"Ogelman 1995). A soft ($T_{\\rm s} \\sim (5$ -- $10) \\times 10^5$~K) component of the observed spectra is most probably emitted from all the visible neutron star (NS) surface. It pulses with the pulsar period, and the pulsed fraction, $\\sim 10-30 \\%$, varies with photon energy. The spectra and the light curves of these middle-age ($10^4$ -- $10^6$ yr) NSs contain important information on NS parameters and on the properties of superdense matter in NS interiors.  The pulsations are most probably caused by the temperature variation over the NS surface. The natural cause of the variation is non-uniformity of the surface magnetic field and anisotropy of the heat transport in strongly magnetized subphotospheric layers of cooling NSs (e.g., Yakovlev \\& Kaminker 1994).  So far most of the cooling calculations of NSs (e.g., Nomoto \\& Tsuruta  1987, Van Riper 1991) have been performed under the traditional simplified assumption that the magnetic field is radial everywhere over the stellar surface. In this case the thermal energy is carried through the subphotospheric layers along the magnetic field lines, and the surface temperature $T_{\\rm s}$ is uniform and determined by the longitudinal (along the field $\\vect{B}$) thermal conductivity. The latter conductivity is mainly enhanced by the magnetic field, which increases $T_{\\rm s}$ at given internal temperature $T_{\\rm i}$. As a result, the traditional cooling theories predict (e.g., Van Riper 1991) that the strong magnetic fields $B \\ga 10^{12}$~G increase $T_{\\rm s}$ (and the photon luminosity) at the neutrino cooling stage. At this stage, a NS is not too old (age $t \\la 10^4$--$10^5$ yrs) and cools mainly via neutrino emission from the stellar interior. On the other hand, strong fields decrease $T_{\\rm s}$ and accelerate the cooling at the subsequent photon cooling stage when the neutrino emission becomes low and the star cools via the thermal surface radiation.  The main disadvantage of the traditional approach is that it neglects those parts of the NS surface where $\\vect{B}$ is essentially non-radial, and the transverse thermal conductivity is important. The aim of this paper is to use a realistic model of heat transport and the related surface temperature distribution valid for any magnetic field geometry in the NS surface layers. The model generalizes the results of Greenstein \\& Hartke (1983), Van Riper (1989), Schaaf (1990a), and Page (1995) who performed detailed studies of the NS surface temperature either for restricted magnetic field geometries or under specified assumptions on thermal conductivity of stellar matter (see Sect.~2, for details). Adopting this model we will carry out the cooling calculations (Sect.~3). Our results show (Sects.~4 and 5) that the effects of the magnetic fields on the NS cooling are more sophisticated than anticipated previously. ",
        "conclusions": "We have used a realistic model (\\ref{eq:Te}) of the anisotropic distribution of the effective temperature $T_{\\rm s}$ over the surface of a magnetized NS and have calculated NS cooling. We have considered both standard and rapid cooling of the NS with the dipole magnetic field. The results are noticeably different from the traditional results obtained for radial (monopolar) magnetic fields. We have shown (Figs.~2--7) that the dipole fields $10^{10} \\la B_{\\rm p} \\la 3 \\times 10^{12}$~G decrease the stellar photon luminosity at the neutrino cooling stage (by a factor of $\\sim 2$), and increase the luminosity at the photon cooling stage (by about one order of magnitude), as compared to a non-magnetized NS, contrary to the traditional theories developed for monopolar magnetic fields. The effect is produced by the growth of the thermal isolation of the equatorial surface layers due to the low thermal conductivity of degenerate electrons across the magnetic field. On the other hand, the dipole fields $B_{\\rm p} \\ga 3 \\times 10^{12}$~G increase the photon thermal luminosity at the neutrino cooling stage and decrease the luminosity at the photon cooling stage, in comparison with the luminosity at $B=0$. This conclusion agrees qualitatively with the results of the traditional theories but the actual magnetic field effect is much weaker than predicted by these theories due to the interference with the opposite effect in the equatorial regions. We have obtained also that the dipole field with $B_{\\rm p} \\approx 3 \\times 10^{12}$~G has almost no effect on the NS cooling.  Our results are based on a model of the surface temperature distribution applied to the dipole magnetic field. It would be easy to consider the NS cooling for other magnetic field geometries (quadrupole, shifted dipole, etc.), and the results are expected to be qualitatively the same as for the dipole field: the parts of the surface with essentially non-radial magnetic fields should noticeably increase the thermal isolation (Sect.~2).  For further studies of the magnetic field effects on the NS cooling, it would be highly desirable to perform detailed investigation of the heat transport in the outer NS layers using the best microscopic physics available (equation of state, thermal conductivities, etc), and obtain thus exact surface temperature distribution in magnetized NSs. The distribution (\\ref{eq:Te}) used above is expected to give an upper limit for the surface temperature variation produced by the magnetic fields (Sect.~2). It can be used for a test in more advanced theories. Even if (\\ref{eq:Te}) is not very accurate at large magnetic fields near the equatorial regions, it can yield quite accurate results since the cold equatorial surface parts do not contribute significantly into the photon luminosity (\\ref{eq:L}). \\begin{figure} \\begin{center} \\leavevmode \\epsfxsize=7.5cm \\epsffile[60 60 530 530]{fig7.ps}  \\end{center} \\caption[ ]{ Same as in Fig.~4 for the 1.44\\,$M_\\odot$ neutron star. } \\label{fig7} \\end{figure} The results of this work can be useful for theoretical interpretation of thermal X-ray radiation of NSs. All the objects from which this radiation has been detected (Sect.~1) are thought to possess magnetic fields of about several times of $10^{12}$~G. We can expect that such fields do not affect significantly the NS cooling but they can produce strongly anisotropic surface temperature variation and associated modulation of the surface thermal radiation.  The magnetic fields broaden the allowed ranges of the photon luminosities in the $L$--$t$ diagram for NSs of various masses, radii and equations of state with standard or enhanced neutrino energy losses (e.g., \\\"Ogelman 1994). According to the traditional theories, high magnetic fields {\\it increase} the upper boundary of the expected photon luminosity at the neutrino cooling stage (typically, at $10^2 \\la t \\la 3 \\times 10^5$~yrs) and decrease the lower boundary of $L$ at the photon stage. Our results (Figs.~4 and 7) show that the dipolar fields $B \\sim 10^{11}$~G {\\it  decrease} the lower boundary of $L$ at the neutrino stage and increase the upper boundary at the photon stage."
    },
    "9604/astro-ph9604005_arXiv.txt": {
        "abstract": "We analyze the anisotropy signature expected if the high energy (above $10^{19}$eV) cosmic ray (CR) sources are extra-Galactic and trace the distribution of luminous matter on large scales.  We investigate the dependence of the anisotropy on both the relative bias between the CR sources and the galaxy distribution and on the (unknown) intrinsic CR source density. We find that the expected anisotropy associated with the large scale structure (LSS) should be detected once the number of CR events observed above $10^{19}{\\rm eV}$ is increased by a factor of $\\sim10$. This would require $\\sim30$ observation-years with existing experiments, but less then $1$ year with the proposed $\\sim5000\\ {\\rm km}^2$ Auger detectors. We find that the recently reported concentration of the Haverah Park CR events towards the super-galactic plane is not consistent with the known LSS. If real, the Haverah Park result suggests that the CR sources are much more concentrated towards the super-galactic plane than the known LSS. Our results are not sensitive to the number density of CR sources. We show that once the number of detected events is increased by a factor of $\\sim10$, the number density would be strongly constrained by considering the probability for having repeating sources. ",
        "introduction": "Recent cosmic ray observations, reported by the Fly's Eye (\\cite{Fly}) and by the AGASA (\\cite{AGASA}) experiments, show two major features in the cosmic ray (CR) energy spectrum above $10^{17}{\\rm eV}$. First, a break in the shape of the spectrum is observed at $\\sim5\\times10^{18}{\\rm eV}$. Second, the CR composition changes from being predominantly heavy nuclei below the break to light nuclei above the break. Coupled with the lack of anisotropy, that would be expected if the CRs above $10^{19}{\\rm eV}$ were protons produced in the Galaxy, these features strongly suggest that the CR flux above $10^{19}{\\rm eV}$ is dominated by an extra-Galactic component of protons. This view is supported by the fact that the CR spectrum above $\\sim2\\times10^{19}{\\rm eV}$ is consistent with a cosmological distribution of sources, with a power law generation spectrum ${\\rm d}\\ln N/{\\rm d} \\ln E\\simeq-2$ (\\cite{Wb},c) (below $\\sim2\\times10^{19}{\\rm eV}$ a significant contribution from iron cosmic rays from Galactic sources is likely to be present; Bird {\\it et al.} 1994, Waxman 1995b).  If the particles observed are indeed protons of extra-Galactic origin, and if their sources are correlated with luminous matter, then the inhomogeneity of the large scale galaxy distribution, on scales $\\lesssim100$\\mpc, should be imprinted on the CR arrival directions. In this paper, we examine the expected anisotropy signature if the CR sources trace the large scale structure (LSS), and investigate its dependence on the relative bias between the CR sources and the galaxy distribution and on the (unknown) intrinsic CR source density. The galaxy distribution is derived from the \\iras 1.2 Jy redshift survey (Fisher {\\it et al.} 1995). We find that the expected anisotropy associated with the LSS should be detected once the number of CR events detected above $10^{19}{\\rm eV}$ is increased by a factor of $\\sim10$. \\cite{Stanev} have recently noted that the arrival directions of $E>4\\times10^{19}{\\rm eV}$ CR events detected by the Haverah Park experiment show a concentration in the direction of the Supergalactic Plane (SGP) which is inconsistent with the hypothesis that the CR sources are distributed isotropically.  We confirm this result, but also find that the Haverah Park CR distribution is unlikely to be explained by the hypothesis that the CR sources trace the known LSS. ",
        "conclusions": "We have shown that, if the distribution of CR sources trace the large scale structure, large exposure CR detectors should clearly reveal anisotropy in the arrival direction distribution of CRs above $4\\times10^{19}{\\rm eV}$. The exposure required for a northern hemisphere detector to discriminate between isotropic CR source distribution and an unbiased distribution that traces the LSS is approximately $10$ times the current Fly's Eye exposure. If the CR source distribution is strongly biased, the required exposure is $\\sim3$ times the current. Increasing the exposure by a factor of $10$ with existing experiments (AGASA and Fly's Eye), would require $\\sim30$ years of observation (taking into account the fact that the AGASA experiment triggering was recently improved). The required observation time would be reduced if new, larger, CR experiments become operative: $\\sim10$ observation-years would be required with the new High Resolution Fly's Eye experiment (\\cite{hires}), which is planned to become operative in two years; Less than $1$ year of observation would be required if the proposed $\\sim5000\\ {\\rm km}^2$ detectors of the Auger project are built (\\cite{huge1,huge2}).  An experiment with full sky coverage, such as the Auger experiment, would allow the CR event distribution to be analyzed conveniently in terms of a spherical multipole decomposition. In Fig. \\ref{fig6}, we show the probability distribution for dipole and quadrupole moments of the CR distribution for a full sky coverage detector with an exposure $\\sim 100$ times the current Fly's Eye exposure, comparable to the exposure of the Auger detectors after a few years of operation. The dipole moment has been defined as $D= \\langle \\cos \\theta \\rangle$ where $\\theta$ is the angle between the CR position and fixed reference direction, taken to be $(l=270\\arcdeg,\\ b=30\\arcdeg)$, while the quadrupole moment is defined as $Q= \\langle \\sin^2 b^G \\rangle - {{1}\\over{3}}$. From the figure, it is evident that even these low order statistics are sufficient to discriminate between the models at a high degree of confidence. Moreover, since the various multipole moments are independent, the confidence level can be increased by including higher order moments of the CR distribution.  The anisotropy signal is not sensitive to the currently unknown number density of CR sources (see Figs. \\ref{fig5} and \\ref{fig6}). We have shown, that a reliable lower limit to the source number density, $\\bar s_0$, may be obtained by considering the probability of observing repeating sources. As the source number density decreases, individual sources become brighter and, for a given number of detected events, the probability that one source would contribute more than one event increases. An absence of repeaters in the current Fly's Eye data above $5\\times10^{19}{\\rm eV}$, for example, would imply $\\bar s_0>10^{-5}\\ {\\rm Mpc}^{-3}$ with 90\\% confidence limit (see Fig. \\ref{fig2}). The number density would be strongly constrained by considering repeating sources once the exposure is increased by a factor of $\\sim10$.  \\cite{Stanev} have recently noted that the arrival directions of $E>4\\times10^{19}{\\rm eV}$ CR events detected by the Haverah Park experiment show a concentration in the direction of the SGP. In agreement with \\cite{Stanev}, we find that the probability to obtain the Haverah Park results assuming an isotropic CR source distribution is very low. However, we find that this probability is not significantly higher for models where the CR source distribution traces the LSS; thus, contrary to the claim by \\cite{Stanev}, we find that the concentration of the Haverah Park events towards the SGP is {\\it not} in strong support for the CR sources tracing the known LSS.  Our analysis addressed only the 2-dimensional (angular) LSS information contained in the distribution of CR arrival directions. However, the energy dependent distance cutoff of high energy CRs, shown in Fig. \\ref{fig1}, implies that the differential CR flux is dominated at different energies by sources which lie at different distances. Therefore, analyzing the angular distribution of CRs at different energy ranges would provide information on the LSS at different distances, and would therefore probe the 3-dimensional LSS. The exposure which would be required in order to extract 3-dimensional LSS information from the CR arrival distribution may be estimated as follows. Let us assume that we are interested in a LSS feature that lies within a distance range $d_1-d_2$ and occupies a solid angle $\\delta\\Omega$. Let's denote the fraction of the flux in the energy range $E_1-E_2$, that is contributed on average (i.e. for homogeneous source distribution) by sources within the distance range $d_1-d_2$, by $f(E_1,E_2,d_1,d_2)$. Our ``signal'' in the energy range $E_1-E_2$, i.e. the number of CRs in this energy range that are produce by sources that lie within the LSS feature of interest, is then given by $\\sim f(E_1,E_2,d_1,d_2) N(E_1,E_2)\\delta\\Omega/{\\rm d}\\Omega$, where $N(E_1,E_2)$ is the total number of events in the $E_1-E_2$ range and ${\\rm d}\\Omega$ is the experimental field of view. The \"noise\" from sources outside $d_1-d_2$ is $\\sim[(1-f)N\\delta\\Omega/{\\rm d}\\Omega]^{1/2}$, so that the signal to noise is $\\sigma(E_1,E_2,d_1,d_2)\\simeq f[N\\delta\\Omega/{\\rm d}\\Omega (1-f)]^{1/2}$. Figure \\ref{fig7} shows a contour map of the signal to noise $\\sigma(E_1,E_2)$ for a structure with radial extent of $50{\\rm Mpc}$ and angular extent of $10\\arcdeg\\times10\\arcdeg$ lying at a distance of $d=(d_1+d_2)/2=200,\\,300{\\rm Mpc}$. For this figure, the exposure was chosen to be the current Fly's Eye exposure. The map shows that, as indicated by Fig. \\ref{fig1}, the best signal to noise is obtained by choosing the energy range $5.5-6.5\\times10^{19}{\\rm eV}$ for $d=200{\\rm Mpc}$, and $4.5-5.5\\times10^{19}{\\rm eV}$ for $d=300{\\rm Mpc}$. Unfortunately, the exposure required to obtain a signal to noise of order $1$ is very large, $\\sim300$ times the current Fly's Eye exposure, which correspond to $\\gtrsim10$ observation years of the proposed Auger detectors. Thus, although it is possible in principle to probe the 3-dimensional LSS using UCHERs at different energy channels, it is not clear if it would be possible to do so in practice, even with the largest detectors planned today."
    },
    "9604/astro-ph9604143_arXiv.txt": {
        "abstract": "We present an empirical method that uses multicolor light curve shapes (MLCS) to estimate the luminosity, distance, and total line-of-sight extinction of Type Ia supernovae (SN Ia).  The empirical correlation between the MLCS and the luminosity is derived from a ``training set'' of nine SN Ia light curves with independent distance and reddening estimates.  We find that intrinsically dim SN Ia are redder and have faster light curves than the bright ones which are slow and blue.  By thirty-five days after maximum the intrinsic color variations become negligable.   A formal treatment of extinction employing Bayes' theorem is used to estimate the best value and its uncertainty.    Applying MLCS to both light curves and to color curves provides enough information to determine which supernovae are dim because they are distant, which are intrinsically dim, and which are dim because of extinction by dust. The precision of the MLCS distances is examined by constructing a Hubble diagram with an independent set of twenty SN Ia's.  The dispersion of 0.12 mag indicates a typical distance accuracy of $5\\%$ for a single object, and the intercept yields a Hubble constant on the Cepheid distance scale (Sandage et al 1994, 1996) of $H_0=65 \\pm$ 3 (statistical) km s$^{-1}$ Mpc$^{-1}$ ($ \\pm 6 $ total  error).  The slope of 0.2010 $\\pm$ 0.0035 mag over the distance interval 32.2 $< \\mu <$ 38.3 yields the most precise confirmation of the linearity of the Hubble law. ",
        "introduction": "Since the Curtis-Shapley debate of 1920 (Curtis 1921, Shapley 1921), the determination of supernova (SN) luminosities has been central to the discussion of extragalactic distances.  Shapley (1919) argued against the ``Island Universe'' hypothesis, because it required certain novae (such as S Andromedae of 1885) to reach the astonishing luminosity of $M=-16$.  He considered this to be ``out of the question''.   Curtis (1921) countered, concluding that ``the dispersion of the novae in spirals and in our galaxy may reach ten absolute magnitudes...a division into two classes is not impossible''.  This distinction between novae and supernovae, required by the extragalactic nature of the nebulae,  was later made explicit by Baade and Zwicky (1934).  They showed that in addition to their tremendous difference in absolute luminosity, the photometric and spectroscopic behavior of supernovae is distinct from novae.  Baade (1938) showed that supernovae were more uniform than novae, with a dispersion at peak of  1.1 magnitudes, making them suitable as extragalactic distance indicators.  The precision of supernova distance estimates has increased as the SN Ia class has been better understood and more narrowly defined.  The low dispersion in Baade's sample benefited from the fortuitous absence of Type II supernovae, which are significantly less luminous in the mean.  Beginning with SN 1940B, Type II supernovae were classified by the presence of hydrogen in their spectra (Minkowski 1941). The growing list of spectroscopically defined Type I supernovae had dispersions at peak of 0.8-0.6 magnitudes (Minkowski 1964, Kowal 1968, Kirshner et al. 1973, Oke and Searle 1974).  However, this sample included a number of ``peculiar'' SN Ia noted for their lack of silicon, which are now recognized to arise from massive stars that lose their envelope before core collapse (see Wheeler \\& Harkness 1990).  After removing these silicon-deficient objects, now classified as Ib's and Ic's (Doggett and Branch 1985, Uomoto and Kirshner 1985, Wheeler and Levreault 1985, Wheeler and Harkness 1986), the remaining Type Ia supernovae (SN Ia)  form a more homogeneous set which serve as even more precise indicators of astronomical distances.  Leibundgut (1989) devised a set of standard templates to describe the photometric behavior of SN Ia and to estimate the peak apparent magnitude.  Hubble diagrams constructed using the peak of photographic SN Ia light curves had observed dispersions ranging from $\\sigma_M$=0.65-0.36 magnitudes depending on which objects and color bands were used (Tammann and Leibundgut 1990, Branch and Miller 1993, Miller and Branch 1990, Della Valle and Panagia 1992, Rood 1994, Sandage and Tammann 1993, Sandage et al 1992,1993,1994). An ambitious program to calibrate nearby SN Ia through Cepheid variables observed with the Hubble Space Telescope has been undertaken by Sandage et al (1992, 1993, 1994, 1996).  Assuming SN Ia to be homogeneous yields a Hubble constant in the range 50-58 km  s$^{-1}$  Mpc$^{-1}$ (Sandage et al 1992,1993,1994, 1995, Schaefer 1994, 1995a,b, 1996,  Branch \\& Tammann 1992) with the most recent measurement giving 57 $\\pm \\ 4$ km  s$^{-1}$  Mpc$^{-1}$.  We show in \\S 6 and \\S 7 that the precision of this measurement is improved and the value of H$_0$ altered by including information contained in the light and color curve shapes.  The hypothesis that SN Ia are standard candles drew support not only from empirical studies, but also from the earliest theoretical models which suggested they arose from ignition of a carbon-oxygen white dwarf at the Chandrasekhar mass (Hoyle and Fowler 1960, Arnett 1969, Colgate \\& McKee 1969).  In these models a supersonic shock wave travels through the degenerate star, burning material into  $^{56}$Ni at a temperature of five billion degrees  (Khokhlov, Muller, and H\\\"{o}flich 1993, Mazurek \\& Wheeler 1980).  Because the detonation is supersonic, the pre-shock region cannot expand to decrease the pressure or the burning temperature.  Further, the Fermi pressure of the degenerate material in the post-shock region remains insensitive to temperature for longer than the burning timescale.  The result is a total incineration and the production of a pure mass of nickel.  Such a standard explosion of a uniform mass would lead to a homogeneous light curve and uniform luminosity.  Yet, these complete burning models of Ia's do not reproduce the intermediate mass elements which are seen in the spectra of SN Ia (Wheeler and Harkness, 1990).  A successful model (Nomoto, Thielemann, and Yokoi 1984) which matched the observational constraints was so persuasive that Arnett, Branch, and Wheeler (1985) and Branch (1992) suggested calibration of the Hubble constant based only on theoretical models of uniform nickel production.  However, a variety of models  (Livne 1990, Khokhlov, Muller, and  H\\\"{o}flich 1993,  Woosley \\& Weaver 1992, H\\\"{o}flich, Khokhlov, \\& Wheeler 1995) match the observed features of the spectra and produce a range of nickel masses, and a range of predicted luminosities.   These models employ subsonic deflagration fronts, pulsations, or off-center explosions to allow the surface layers to pre-expand and burn at low temperature.  The success of these models in reproducing the observed spectra opens a large range of theoretical possibilities.  Unlike the first monoenergetic models, these models suggest a wide range of luminosities might result from the ignition of a white dwarf.  Recently, precise observations of SN Ia made with CCD detectors show evidence for inhomogeneity in both luminosity and light curve shape.  One of the first SN Ia observed with a distinctly different light curve was 1986G (Phillips et al 1987, Cristiani et al 1992) which displayed a spectacularly rapid decline in its B and V light curve  and unique spectral characteristics including stronger-than-usual Si features.  Although SN 1986G was heavily reddened by dust, reddening cannot significantly alter the {\\it shape} of the light curve.  SN 1991bg in NGC 4374 is the most extreme SN Ia in an increasingly apparent photometric and spectral sequence.  Leibundgut et al (1993) (also Filippenko et al 1992) described a number of photometric abnormalities of 1991bg with respect to his templates.  These include the fastest post-maximum decline (2.05 and 1.42 mag decrease drop in B and V in the fifteeen days after maximum compared to 1.22 and 0.64 mag for the templates in B and V), a narrow luminosity peak, and an intrinsic red color near maximum. A simple and convincing argument that SN Ia have a large spread in luminosity is that SN 1957B, which occurred in the same galaxy, was 2.5 magnitudes brighter in B than SN 1991bg.  In addition, SN 1991bg was at least 2 magnitudes fainter than other SN Ia in the Virgo cluster, of which NGC 4374 is a member.  This extreme SN Ia seems to have a twin in the sub-luminous SN 1992K (Hamuy et al 1994), which strongly resembles 1991bg in photometric and spectral behavior.  At the opposite extreme of the SN Ia class, SN 1991T showed spectral and photometric peculiarities which were different from those seen in the rapid decliners.  Phillips et al (1992) found the light curves in B and V to rise and decline more slowly than the standard templates near maximum, and a month after peak, this shallower decline resulted in a light curve 0.2-0.3 mag brighter than the templates.  Although SN 1991T's host galaxy, NGC 4527, appears to lie south of the main Virgo cluster,  Phillips et al (1992, 1993) estimates the peak luminosity exceeded that of other SN Ia in Virgo by 0.3-0.5 mag.  From the narrow Na I D absorption line, Filippenko et al (1992) deduced that SN 1991T is dimmed by dust in NGC 4527 and concluded it may have been as much as $ \\sim 0.9 $ mag more luminous than a typical SN Ia.  More recently, the Cal\\'{a}n/CTIO supernova search yielded SN 1992bc and 1992bo (Maza et al 1994), two SN Ia with similar recession velocities of 6050 and 5660 km s$^{-1}$, but with peak apparent luminosities differing by 0.69 mag in B.  The large difference in apparent magnitude and the small difference in recession velocity imply that SN 1992bc was intrinsically brighter by $0.8\\pm0.2$ mag than 1992bo.  For a difference in distance to account for this difference in luminosity, the peculiar velocities of the two supernovae would have to differ by $\\sim 2,500$ km sec$^{-1}$, an unlikely alternative.  Interestingly, SN 1992bc declined more slowly than the average template while SN 1992bo's post-maximum fall was more rapid than the average template, a result in good accord with later analysis by Phillips (1993).  Even before these recent examples of inhomogeneity, less precise photographic measurements by Barbon et al (1973) suggested that there might exist two photometric classes; those with ``fast'' decline rates after maximum which were intrinsically brighter supernovae, and those with ``slow'' decline rates which were fainter.   With the poor quality of photographic and photoelectric photometry then available, such a real distinction was difficult to demonstrate convincingly.   Pskovskii (1977,1984) suggested a continuous photometric sequence of SN Ia light curves.  He introduced a parameter, $\\beta$: the slope of the B band post-maximum decline in magnitudes per 100 days.  Using 54 photographic SN Ia light curves, Psovskii found a weak correlation between $\\beta$ and the absolute B magnitude at maximum light which was {\\it opposite} to that of Barbon (1973).   Early difficulties in measuring a relation between SN Ia light curve shape and intrinsic luminosity resulted from noisy photographic data which were poorly sampled and calibrated.  These difficulties were compounded by the problem of measuring a decaying light curve on a bright galaxy with a non-linear photographic detector (Boisseau and Wheeler 1991).  With the advantage of better data measured with linear detectors, Phillips (1993) demonstrated conclusive evidence for a luminosity-light curve decline relation among SN Ia.  Using a set of well-sampled SN Ia light curves with precise optical photometry and accurate relative distances, Phillips found that the absolute luminosity in B,V, and I is correlated with the B band decline in the fifteen days following maximum light.  The sense of the correlation is that dimmer SN Ia fall more rapidly after maximum than bright SN Ia.  Application of this relation to his sample results in a significant reduction of the dispersion in B,V, and I luminosity from 0.8, 0.6, and 0.5 mag to 0.36, 0.28, and 0.38 respectively.  A more recent investigation by Tammann and Sandage (1995)  has examined the luminosity-light curve decline relation among ``normal'' SN Ia, in this case ``normal'' is defined as having a $(B-V)_{max}$ $\\leq 0.30$ mag.  They tentatively confirm a luminosity dependence on light curve decline, with a slope that is shallower than Phillips (1993) but consistent with Hamuy et al (1995) for a similar set of ``normal'' SN Ia.  In \\S 6 we show for a sample restricted to ``normal'' SN Ia, a significant correlation between light curve shape and luminosity exists.  This empirical relation between light curves and luminosity in SN Ia has been paced by an abundance of theoretical models which can account for the observed behavior (H\\\"{o}flich, Khokhlov, and Muller 1993, Woosley and Weaver 1994, Ruiz-Lapuente et al 1993, Livne and Arnett 1995, H\\\"{o}flich and Khokhlov 1995).  These new models include deflagration burning fronts, off-center detonations, surface helium burning, pulsed delayed detonations and sub-Chandrasekhar progenitors.  These models give plausible causes for the observed inhomogeneity of SN Ia and for the origin of the empirical relations between light curve shape and luminosity.  Most recently, Hamuy et al (1995) have employed a template-fitting approach and we (Riess, Press, and Kirshner 1995a, hereafter RPK 95a) have developed a linear estimation algorithm to use the distance-independent light curve shapes to improve the precision of distance measurements to SN Ia.  The techniques have much in common and both yield Hubble diagram dispersions of $\\sim$ 0.2 magnitudes for an overlapping set of objects.  The Light Curve Shape (LCS) method, described in RPK 95a, has the advantage of providing quantitative error estimates for the distance.  The present paper extends the LCS technique to use the shapes of B$-$V, V$-$R, and V$-$I color curves which provide enough information to determine the relation between absolute luminosity and intrinsic color.  Knowledge of an SN Ia's intrinsic color allows us to measure the extinction from the observed reddening.   For each well-observed SN Ia we measure the luminosity, extinction and the extinction-corrected distance.  The multicolor light curve shape (MLCS) method significantly increases the precision of distance estimates from SN Ia light curves as we show in \\S 6.  There are many potential applications for a bright distance indicator with $<$ $10 \\%$ precision.  Nearby $(0.01 \\leq z \\leq 0.1)$, it should be possible to measure the Hubble constant to an accuracy which is limited only by the underlying calibration of Cepheids.    It is important to compare the Hubble constant derived using the light curve shape-luminosity relation with determinations which have assumed a homogeneous luminosity for SN Ia (Sandage et al. 1992, 1993, 1994, 1996).  Using the velocity residuals from Hubble flow, we have measured the motion of the Local Group with respect to the rest frame of galaxies with supernovae (Riess, Press, and Kirshner 1995b).  At even greater distances $(0.3 \\leq z \\leq 0.6)$ MLCS distance measurements of all well-observed SN Ia could be used to determine the cosmological deceleration parameter, $q_o$ (Perlmutter et al 1995, Perlmutter et al 1996, Schmidt et al 1996, IAUC 6160).  Some have sought to improve the homogeneity of the observed SN Ia by restricting the sample to supernovae with ``normal'' spectra (Branch, Fisher, and Nugent 1993, Sandage et al 1994, 1996).  While spectroscopic information may prove useful in producing smaller dispersion in distance estimates,  the difficulty in obtaining good spectra of very distant SN Ia makes it hard to identify subtle spectral variations. A sample cut which cannot be applied with equal effectiveness for nearby and for distant SN Ia could lead to a bias in cosmological parameters determined by them.  We prefer to develop a method that can be applied to all SN Ia.  Given at least one light curve and one color curve, photoelectrically observed within ten days after maximum, our MLCS method can distinguish between the effects of distance, intrinsic dimness, and dust for all SN Ia.  In \\S 2-5 we develop the multicolor light curve shape method for measuring extinction-corrected distances.  In \\S 6, for an independent sample, we compare the extinction-corrected MLCS distances with distances determined by more limited assumptions.  In \\S 7, we estimate the Hubble constant and discuss sample membership, selection bias and the range of progenitors which may be responsible for the empirical range of SN Ia luminosities and colors ",
        "conclusions": ""
    },
    "9604/astro-ph9604157_arXiv.txt": {
        "abstract": "Deep {\\sl HST} WFPC2 images have revealed a population of very narrow blue galaxies which Cowie et al.\\ (1996) have interpreted as being a new morphological class of intrinsically {\\sl linear} star forming galaxies at $z=0.5-3$.  We show that the same population exists in large numbers at low redshifts ($z\\approx0.03$) and are actually the edge-on manifestation of low surface brightness disk galaxies. ",
        "introduction": "The launching of the Hubble Space Telescope has for the first time allowed astronomers to observe the morphology of galaxies at moderate redshifts.  Coupled with deep redshift surveys, {\\sl HST} imaging provides a powerful tool for studying the evolution of galaxies, in particular, the ubiquitous but poorly understood population of ``excess'' faint blue galaxies.  A recent paper by Cowie et al.\\ (1996; hereafter CHS) reports that many of the blue galaxies in deep F814W {\\sl HST} images of two fields from the Hawaii Redshift Survey appear to be very narrow, linear structures.  The galaxies tend to be very straight, and have sizes between 2-3\\arcsec\\ in the long dimension and are marginally resolved after deconvolution at 0.05-0.1\\arcsec\\ in the narrow dimension.  CHS argue that the extreme ellipticity of these galaxies and the lack of continuity between the linear ``chain'' galaxies with the rest of the galaxy population is incompatible with their being edge-on disk galaxies.  They therefore treat these galaxies as a distinct morphological class of {\\sl intrinsically} linear galaxies which they refer to as ``chain'' galaxies.  Two of the chain galaxies are emission line galaxies at redshifts of $z\\approx0.5$, another is argued to be at $z\\approx1.4$ based upon one strong emission line, and a fourth is suggested to be at $z=2.4$ based on interpreting a break in the continuum at $4000$\\AA\\ as the onset of intergalactic Ly-$\\alpha$ forest absorption.  On the basis of the emission lines and blue colors, coupled with the short lifetime expected for intrinsically linear structures, CHS argue that their new morphological class of chain galaxies are not only star-forming, but are in the process of formation (see also Ogorodnikov 1967).  The conclusions of CHS that chain galaxies represent a new morphological class of linear galaxies in the process of forming at $z=0.5-3$ is based on the following observations: 1) the extreme ellipticities 2) the lack of continuity between the chain galaxies and the rest of the galaxy population and 3) the lack of any low redshift counterparts.  In this paper we will first show that the extreme ellipticities are not incompatible with chain galaxies being intrinsically disky systems viewed edge-on.  We will then use data from an ongoing ground-based redshift survey of low surface brightness galaxies to show that there are large numbers of galaxies with similar morphology at very low redshifts ($z\\approx0.03$) and that they join seamlessly onto the population of edge-on normal galaxies. Furthermore, we will show that the low redshift galaxies with the ``chain'' morphology are consistent with their being the edge-on manifestation of face-on low surface brightness galaxies (LSBs) found in the same survey. ",
        "conclusions": "\\label{conclusions}  We have shown that the apparently unfamiliar population of very thin galaxies revealed in deep HST imaging do in fact have low redshift counterparts among the population of low surface brightness disk galaxies, and thus, that there is no immediate need to invoke any evolutionary scenario to explain their appearance at moderate redshifts.  We have shown examples of such very thin galaxies in existing catalogs of nearby galaxies, and shown the statistics of the ellipticity distributions of low-redshift LSB galaxies are consistent with the low redshift ``chain'' galaxies being edge-on manifestations of disks.  This suggests that, at low redshift, such thin systems are {\\sl not} necessarily intrinsically linear (as opposed to disky), bringing into question the assumption that equally thin systems at high redshift {\\sl cannot} be disks (CHS).  Our conclusion that the thin ``chain'' galaxies observed at moderate to high redshifts are analogs of low surface brightness disks is supported by both their morphologies and by their colors as well.  LSBs at low redshift are intrinsically asymmetric (Figure \\ref{inclinationplot}), often with one or two bright knots superimposed upon a diffuse low surface brightness disk (Figure \\ref{faceonfig}); such systems, when viewed edge-on, would be consistent with the lumpy structure of some of the chain galaxies.  The moderate redshift chain galaxy / low redshift edge-on LSB class may be viewed simply as a subset of the link between the faint blue galaxies and nearby LSBs posited by Ferguson \\& McGaugh (1995)."
    },
    "9604/astro-ph9604011_arXiv.txt": {
        "abstract": "The status of searches for gravitational microlensing events of the stars in our galaxy and in other galaxies of the Local Group, the interpretation of the results, some theory, and prospects for the future are reviewed. The searches have already unveiled $ \\sim 100 $ events, at least two of them caused by binaries, and have already proven to be useful for studies of the Galactic structure.  The events detected so far are probably attributable to the effects of ordinary stars, and possibly to sub-stellar brown dwarfs; however a firm conclusion cannot be reached yet because the analysis published to date is based on a total of only 16 events. The current searches, soon to be upgraded, will probably allow determination of the mass function of stars and brown dwarfs in the next few years; these efforts will also provide good statistical information about binary systems, in particular their mass ratios.  They may also reveal the nature of dark matter and allow us to detect planets and planetary mass objects. ",
        "introduction": "The topic of gravitational lensing has a long history, as described for example in the first book entirely devoted to the subject (Schneider et al. 1992).  The first known theoretical calculation of a light ray bending by massive objects was published by Soldner (1801), who used Newtonian mechanics, and determined that the deflection angle at the solar limb should be $ 0.''84 $, half the value calculated with the general theory of relativity (Einstein 1911, 1916).  The first observational detection of this effect came soon afterwards (Dyson et al. 1920).  Zwicky (1937) pointed out that distant galaxies may act as gravitational lenses.  Almost all essential formulae used today to analyze gravitational lensing were derived by Refsdal (1964). The first case of a double image created by gravitational lensing of a distant source, the quasar 0957+561, was discovered by Walsh et al. (1979).  Arc--like images of extended sources, the galaxies, were first reliably reported by Lynds \\& Petrosian (1989).  In these three cases the sun, a galaxy, and a cluster of galaxies, were acting as gravitational lenses.  There are many recent review articles (Blandford \\& Narayan 1992, Refsdal \\& Surdej 1994, Roulet \\& Mollerach 1996) and international conferences (Moran et al. 1989, Mellier et al. 1990, Kayser et al. 1992, Surdej et al. 1993, Kochanek \\& Hewitt 1996) on the subject of gravitational lensing.  A review of gravitational microlensing experiments has been published by Ansari (1995).  The effect of double imaging of a distant source by a point mass located close to the line of sight, and acting as a gravitational lens, has been proposed many times over.  Chang and Refsdal (1979) and Gott (1981) noted that even though a point mass in a halo of a distant galaxy would create an unresolvable double image of a background quasar, the time variation of the combined brightness of the two images could be observed. This way the effect of non-luminous matter in a form of brown dwarfs or Jupiters could be detected. The term ``microlensing'' was proposed by Paczy\\'nski (1986a) to describe gravitational lensing which can be detected by measuring the intensity variation of a macro--image made of any number of unresolved micro--images.  Paczy\\'nski (1986b) suggested that a massive search of light variability among millions of stars in the Large Magellanic Cloud could be used to detect dark matter in the galactic halo.  Luckily, the technology needed for such a search became available soon afterwards, and the 1986 paper is credited with triggering the current microlensing searches: EROS (Aubourg et al. 1993), MACHO (Alcock et al. 1993), OGLE (Udalski 1992), and DUO (Alard 1996b, Alard et al. 1995a).  Griest (1991) proposed that objects responsible for gravitational microlensing be called massive astrophysical compact halo objects (MACHO).  The name became very popular and it is commonly used to refer to all objects responsible for the observed microlensing events, no matter where they are located and what their mass may be.  It is not possible to discuss all theoretical and observational papers related to microlensing in the Local Group within the modest volume of this article. The selection of references at the end of this review is limited, and I apologize for all omissions, and for the way I made the selection.  Let us hope that somebody will write a more careful historical review before too long.  Fortunately, there is a bibliography of over one thousand papers that are related to gravitational lensing, and it is available electronically.  This bibliography has been compiled, and it is continuously updated, by J. Surdej and  by A. Pospieszalska. It can be found on the World Wide Web at:  \\centerline{http://www.stsci.edu/ftp/stsci/library/grav\\_ lens/grav\\_ lens.html}  The MACHO and the OGLE collaborations provide up-to-date information about their findings and a complete bibliography of their work on the World Wide Web and by anonymous ftp. The photometry of OGLE microlensing events, their finding charts, as well as a regularly updated OGLE status report, including more information about the ``early warning system'', can be found over Internet from the host:  \\centerline{sirius.astrouw.edu.pl   \\hskip 0.5cm (148.81.8.1) ,} \\noindent using ``anonymous ftp'' service (directory ``ogle'', files ``README'', ``ogle.status'', ``early.warning''). The file ``ogle.status'' contains the latest news and references to all OGLE related papers, and PostScript files of some publications. These OGLE results are also available over World Wide Web at:  \\centerline{http://www.astrouw.edu.pl \\hskip 1.0cm (Europe).} \\noindent A duplicate of this information is available at  \\centerline{http://www.astro.princeton.edu/\\~\\/ogle/ \\hskip 1.0cm (North America).} \\noindent A complete information about MACHO results is available at:  \\centerline{http://wwwmacho.mcmaster.ca/ \\hskip 1.0cm  (North America) } \\noindent with a duplicate at  \\centerline{http://wwwmacho.anu.edu.au/ \\hskip 1.0cm (Australia). } \\noindent The information about MACHO alerts is to be found at:  \\centerline{http://darkstar.astro.washington.edu/ }  No doubt other groups will provide similar information before too long. The main weakness of this electronic information distribution is the frequent change of address and mode of access. I shall do my best to keep a guide to this information as part of the OGLE home page on the WWW.  In the following section a simple model of lensing by an isolated point mass is presented -- this is all theory one needs to understand most individual microlensing events.  The third section presents a model for the space distribution and kinematics of lensing objects in order to provide some insight into problems in relating the observed time scales of microlensing events to the masses of lensing objects. The fourth section provides a glimpse of diversity of light curves due to lensing by double objects, like binary stars or planetary systems. Some of the special effects which make microlensing more complicated than originally envisioned are described in the fifth section. The most essential information about the current searches for microlensing events, and some of the results as well as problems with some results, are presented in section six.  The last section is a rather personal outline of the prospects for the future of microlensing searches. ",
        "conclusions": ""
    },
    "9604/astro-ph9604176_arXiv.txt": {
        "abstract": "The MACHO project has been monitoring about ten million stars in the Large Magellanic Cloud in the search for gravitational microlensing events caused by massive compact halo objects (Machos) in the halo of the Milky Way. In our standard analysis, we have searched this data set for well sampled, long duration microlensing lightcurves, detected several microlensing events consistent with Machos in the $0.1\\,\\msun \\simlt m \\simlt 1.0 \\,\\msun$ mass range, and set limits on the abundance of objects with masses $10^{-5}\\,\\msun \\simlt m \\simlt 10^{-1} \\,\\msun$. In this paper, we present a different type of analysis involving the search for very short time scale brightenings of stars which is used to set strong limits on the abundance of lower mass Machos. Our analysis of the first two years of data toward the LMC indicates that Machos with masses in the range $2.5\\ee{-7}\\,\\msun < m < 5.2\\ee{-4} \\,\\msun$ cannot make up the entire mass of a standard spherical dark halo. Combining these results with those from the standard analysis, we find that the halo dark matter may not be comprised of objects with masses $2.5\\ee{-7}\\,\\msun < m < 8.1\\ee{-2}\\,\\msun$. ",
        "introduction": "\\label{intro}  If a significant fraction of the dark halo of the Milky Way is made up of Machos (MAssive Compact Halo Objects), it should be possible to detect them by searching for gravitational microlensing \\cite{paczynski86,petrou}. As a Macho passes near the line of sight to a background star, the star appears to be magnified by a factor \\begin{equation} A = {{u^2 + 2}\\over{u\\sqrt{u^2 + 4}}} \\label{eqamp} \\end{equation} where $u=b/r_E$, $b$ is the distance from the Macho to the line of sight, and the Einstein ring radius $r_E$ is given by \\begin{equation} r_E = \\sqrt{{4GmLx(1-x)}\\over{c^2}} \\end{equation} where $m$ is the mass of the Macho, $L$ is the observer-star distance, and $x$ is the ratio of the observer-lens and observer-star distances.  Since Machos are in motion $(v\\sim v_\\odot = 220$km/sec) relative to the line of sight, this magnification is time dependent, with $A(t)=A(u(t))$, where \\begin{equation} u(t)=\\left[ u_{\\rm min}^2 + \\left({2 (t-t_0 ) \\over \\that}\\right)^2 \\right]^{1 \\over 2}. \\label{equt} \\end{equation} Here $t_0$ is the time of peak magnification, and $\\that$ is the event duration, which can be written \\begin{equation} \\that = 2r_E/v_{\\perp} \\sim 130\\sqrt{m/\\msun} \\, \\rm days, \\label{eqt_hat} \\end{equation} where $v_{\\perp}$ is the Macho velocity relative to the line-of-sight. For more detailed information, see \\citeN{paczynski86} and \\citeN{griest91a}.   If the halo consisted entirely of objects with masses under about $10^{-4} \\,\\msun$ the average duration of microlensing would be less than 1.5  days, and the events would last only about three hours if the halo were made of $10^{-6} \\,\\msun$ objects. In order to clearly see the shape of the microlensing curve for such low mass lenses, images of a lensed star must be taken in rapid succession during the event. Such an experiment was undertaken by the EROS collaboration \\cite{eros}, in which a total of about 82,000 stars were imaged up to 46 times per night for several months. No microlensing events were found, and it was reported that objects with masses $5\\times 10^{-8}\\,\\msun < m < 5\\times 10^{-4} \\,\\msun$ can not comprise the entire Galactic Halo at the 90\\% c.l. We have not followed this approach, but we have used a different technique also capable of setting limits on low mass Machos.  The MACHO collaboration has been monitoring the brightnesses of several million stars in the LMC, SMC, and Galactic Bulge since 1992 July using the 50' telescope at Mt. Stromlo, Australia. A dichroic beamsplitter and filters are used to provide simultaneous measurements in red and blue passbands \\cite{lmc1}. The observing strategy for the first two years of LMC data was designed to be sensitive to objects with masses $m > 10^{-3} \\,\\msun$, so a typical LMC field was generally imaged at most once or twice per clear night. Therefore microlensing events with durations under a few days will have very few magnified points on their light curves and would show up in our data as upward excursions of one, two or three consecutive measurements, occurring on stars which otherwise appeared completely normal. In this paper we search specifically for such short duration ``spike\" events. Clearly, if any spikes are detected, no conclusion could be drawn regarding their origin since there would be insufficient detail in the lightcurves. Therefore, the technique described here is most useful when few if any spikes are found, in which case useful upper limits can be placed on the prevalence of low mass Machos. After applying selection criteria described below, we do not find any such spikes and so are able to strongly constrain the existence of low mass objects in the halo of the Milky Way. ",
        "conclusions": "We have extended the sensitivity of the MACHO experiment to two orders of magnitude lower in mass using existing data and without changing observing strategy. Objects with masses $2.5\\ee{-7}\\,\\msun < m < 8.1\\ee{-2}\\,\\msun$ (roughly one Mars mass to 80 Jupiter masses) can not comprise the entire standard spherical halo mass, and Machos in the range $1.88\\ee{-6}\\,\\msun < m < 2.5\\ee{-2} \\,\\msun$ make up less than 20\\% of the halo. Independent of halo model, objects in the range of $3.2\\ee{-7}\\,\\msun < m < 1.87\\ee{-2}\\,\\msun$ can not make up the canonical halo mass inside $50$ kpc of $4.1\\ee{11}\\,\\msun$, and less than $10^{11} \\,\\msun$ of the halo is made from Machos with masses $1.85\\ee{-6}\\,\\msun < m < 6.5\\ee{-3}\\,\\msun$. These limits are the strongest published to date."
    },
    "9604/astro-ph9604030_arXiv.txt": {
        "abstract": "HS\\,2324+3944 is a peculiar PG\\,1159 star, with a high amount of H in its atmosphere (Dreizler et al. 1995). Its location in the $\\log T_{\\rm eff}$ -- $\\log g$ plane is well inside the GW Vir instability strip. In this paper I report the results of two photoelectric observations of HS\\,2324+39, which clearly show that the luminosity of this star presents periodical variations with a period of (2140 $\\pm$ 11) s. This photometric behaviour is most easily explained by the presence of non-radial oscillations. Therefore HS\\,2324+39 is a new member of the GW Vir group, and is characterized by the longest pulsation period found among these stars. Moreover HS\\,2324+39 appears to be the first pulsating PG\\,1159 star with a high H abundance in its atmosphere (H/He = 2 by number, Dreizler et al. 1995). Were the pulsation mechanism based on the C/O cyclic ionization (Starrfield et al. 1984) at work, the H abundance should drop to zero sharply in the driving regions. Such a strong decrease of hydrogen looks quite unlikely; for this reason the presence of pulsations in HS\\,2324+39 seems to be a very interesting phenomenon. ",
        "introduction": "HS\\,2324+3944 was recognized analyzing the Hamburg Schmidt Survey plates (Hagen et al. 1995, Wisotzki 1994). It has been classified as a lgEH\\,PG\\,1159 (Dreizler et al. 1995), following the scheme of Werner (1992). Dreizler et al. (1995, hereafter DWHE95) have made a detailed analysis of the spectral characteristics of HS\\,2324+39; here is a brief summary of what they have found. HS\\,2324+39 is a ``hybrid PG\\,1159 star'', showing H Balmer absorption features in its spectrum, and it is the only ``hydrogen PG\\,1159'' not surrounded by a planetary nebula. The high H abundance (H/He = 2.0 $^{+0.5}_{-0.6}$ by number) is about 10 times the upper limit found by Werner (1995) for PG\\,1159--035\\,! The abundances of He, C and N (C/He = 0.3, N/He $<$ 0.002) are ``normal'', whereas the O abundance (O/He = 0.006 $\\pm$ 0.004) is quite low respect to the ``standard'' PG\\,1159 stars (O/He = 0.13 for PG\\,1159--035). The effective temperature is equal to (130\\,000 $\\pm$ 10\\,000) K; the surface gravity, corresponding to log $g$ = 6.2 $\\pm$ 0.2, is one of the lowest for the PG\\,1159 stars without nebula. In fact, HS\\,2324+39 is located in a region of the $\\log T_{\\rm eff}$ -- $\\log g$ plane where most objects are central stars of planetary nebulae (CSPN) (DWHE95 Fig. 5).  With these values for temperature and gravity, HS\\,2324+39 results to be well within the GW Vir instability strip. Presently the group of the GW Vir stars is formed by 13 objects \\footnote{We refer here to the``pulsational definition'' of the GW Vir stars; from spectroscopy the GW Vir stars are only 8 (4 with nebula, 4 without), the remaining 5 pulsating CSPN are early WR stars.} \\hspace{-1.5mm} (9 CSPN, 4 not showing a nebula), whose luminosity has multi-frequency variations, due to the presence of non-radial g-mode pulsations (Bradley 1995, Silvotti et al. 1995 and references therein). The typical values of the pulsation periods are 10--35 minutes for the variable CSPN (also called PNNV), and 5--15 min for the pulsating stars without nebula (also called DOV).  The driving mechanism proposed to explain the pulsations in all these stars is the $\\kappa-\\gamma$ mechanism, based on the C/O cyclic ionization (Starrfield et al. 1984). It has been shown (Stanghellini et al. 1991) that even a small amount of H in the driving layers can be sufficient to inhibit the pulsation mechanism. Such a high sensitivity to the chemical composition could in principle explain why not all the PG\\,1159 stars do pulsate.  \\begin{figure*}[ht] \\vspace{88mm} \\special{psfile=hs1.eps voffset=-60 hoffset=0 vscale=45 hscale=85} \\caption[]{Light curves of HS\\,2324+3944 in October 19 and 20, 1995} \\end{figure*}  Now the high H abundance of HS\\,2324+39 would suggest that no pulsations are possible for this star. Therefore the presence or the absence of pulsations in HS\\,2324+39 seemed to be an interesting test for the GW Vir pulsation mechanism. For this reason I decided to investigate on the variability of HS\\,2324+39. The results of two photoelectric observations of this star are presented in the following section of this paper. A preliminary communication on the variability of HS\\,2324+39 has been already published (Silvotti 1995). ",
        "conclusions": "HS\\,2324+39 is a new H-rich peculiar PG\\,1159 pulsating star. It seems to be a very interesting object for several reasons: \\begin{enumerate} \\item{it is, so far, the only PG\\,1159 pulsating star with a high H abundance in its atmosphere (He/H=0.5 by number, DWHE95). This fact seems to be in contrast with the C/O pulsation mechanism proposed by Starrfield et al. (1984), unless we admit that the H abundance drops to zero sharply in the driving regions. The hydrogen could be maintained in a thin surface layer by the gravitational settling. This hypothesis would imply some limits to the strength of the stellar winds and to the extension of the surface convection zone. A weak stellar wind is also suggested by the shape of the blue part of the C\\,IV components near 4650 \\AA \\hspace{0.1mm} in the spectrum of HS\\,2324+39 (DWHE95 Fig. 4, Leuenhagen 1995).  The edges of the theoretical GW Vir instability strip obtained by Starrfield et al. (1984) and by Stanghellini et al. (1991) are calculated using the Los Alamos opacities. Starrfield et al. (1984) considered only carbon and oxygen in the surface, whereas Stanghellini et al. (1991) contemplated also the presence of helium. The presence of hydrogen was never considered. The most realistic models, with 50 $\\%$ carbon and 50 $\\%$ helium by mass (Stanghellini et al. 1991), are pulsationally unstable at effective temperatures much lower than the real GW Vir stars. More recently Saio (1996), using the OPAL opacities, has obtained overstable modes in models with a surface composition closer to the real GW Vir stars (Y=0.38, X$_{\\rm C}$=0.4, X$_{\\rm O}$=0.2, Z=0.02). A model sequence with a 3 $\\%$ (by mass) abundance of hydrogen was also computed; the stability of g-modes results to be hardly affected by the existence of hydrogen. The chemical abundances obtained by Werner (1995) for PG\\,1159--035 (X$<$0.015, Y=0.32, X$_{\\rm C}$=0.48, X$_{\\rm O}$=0.165, taking Z $\\simeq$ 0.02) are quite similar to those used by Saio (1996). But for HS\\,2324+39 (X=0.20, Y=0.405, X$_{\\rm C}$=0.365, X$_{\\rm O}$=0.01, DWHE95 considering Z=0.02) things are quite different: the H and O abundances, compared with PG\\,1159--035, are in practice exchanged. Nevertheless, the location of HS\\,2324+39 in the $\\log T_{\\rm eff}$ -- $\\log \\Pi$ diagram of Saio (1996, Fig. 5) is consistent with models having stellar masses between 0.58 and 0.60 $M_{\\odot}$, in good agreement with the value of 0.59 $M_{\\odot}$, obtained by DWHE95 comparing the location of HS\\,2324+39 in the $\\log T_{\\rm eff}$ -- $\\log g$ plane with evolutionary tracks.} \\item{The pulsation period of (2140 $\\pm$ 11) s is the longest ever observed in a PG\\,1159 star.} \\item{The duration of the pulsation period and the position of HS\\,2324+39 in the $\\log T_{\\rm eff}$ -- $\\log g$ plane (DWHE95) would suggest the presence of a planetary nebula (PN) around the star. Moreover all the other H-rich PG\\,1159 stars are CSPN. Presently it seems that HS\\,2324+39 does not have any PN remnant. A recent observation made at Calar Alto (Werner 1996) confirms this thesis.} \\end{enumerate}  For all these reasons, a detailed study of the HS\\,2324+39 pulsation should be undertaken. With precision asteroseismology some open questions regarding the structure and the mass of the external layers of HS\\,2324+39 could be solved. The best way to achieve this purpose would be to observe HS\\,2324+39 with the Whole Earth Telescope (Nather et al. 1990). The peculiar characteristics of HS\\,2324+39 make this star very interesting and give us a possibility to improve our knowledge on the GW Vir pulsation phenomenon."
    },
    "9604/astro-ph9604097_arXiv.txt": {
        "abstract": "We analyse the {\\it ROSAT} PSPC hardness ratio and the 0.5-2 keV to 2-10 keV flux ratio of 65 Active Galactic Nuclei (AGN) for which there are both {\\it ROSAT} archival observations available and 2-10 keV fluxes, mostly from the HEAO-1 MC-LASS survey.  We conclude that the simplest spectral model for the AGN that can accommodate the variety of X-ray colours obtained is a standard power law (with energy spectral index $\\alpha\\sim 0.9$) plus a $\\sim 0.1$ keV black body both partially absorbed.  In our sample, type 1 AGN require an absorbing column around $10^{22}\\, {\\rm cm}^{-2}$ with covering fractions between 20 and 100\\%, while type 2 AGN display larger columns and $\\sim 100\\%$ coverage. This simple model also provides a good link between soft and hard AGN X-ray luminosity functions and source counts. We also consider a warm absorber as an alternative model to partial covering and find that the the presence of gas in two phases (ionized and neutral) is required. ",
        "introduction": "Comparing the information obtained through the analyses of AGN data in different X-ray energy bands appears to be a difficult task which leads in some cases to striking results. For example, 2-10 keV X-ray source counts (which are dominated by AGN) obtained by the {\\it Ginga} fluctuation analyses (Warwick \\& Butcher 1992) are clearly above the number counts obtained from the {\\it Einstein Observatory Medium Sensitivity Survey} (EMSS) in the 0.3-3.5 keV band (Gioia et al. 1990) if a power law spectrum with energy spectral index $\\alpha\\approxgt 0.7$ and negligible photoelectric absorption are assumed.  Several explanations have been proposed to bring these results into consistency. Warwick \\& Butcher (1992) were able to fit the spectrum of the fluctuations in the 2-10 keV band (which should be close to the median X-ray spectrum of a source at the level where there is one source per {\\it Ginga} LAC beam $\\sim 5\\times 10^{-13}\\, \\ergpcmsqps$) to a power law with an energy spectral index $\\alpha \\approx 0.8$ with no evidence for photoelectric absorption (the derived upper limit is $N_{H}\\approxlt 3 \\times 10^{21} \\rm cm^{-2}$, Stewart 1992). In order to have a 2-10 keV to 0.3-3.5 keV flux ratio of about 2 (which is what is needed in order to reconcile source counts in both energy bands) a photoelectric absorption larger than the above upper limit is required. In this case, the contribution of clusters of galaxies, having a much softer spectrum will compensate the AGN photoelectric absorption in the spectrum of the fluctuations (Barcons 1993). If photoelectric absorption were ignored, an energy spectral index $\\alpha \\sim 0.4$, much flatter than the typical index for any class of source (and this indeed includes AGN) at that flux level, would be required. Therefore, absorption in AGN appears to be a necessity to solve the soft/hard X-ray source counts discrepancies.  On the other hand, soft X-ray selected AGN actually exhibit low-energy excesses over the average power law (Maccacaro et al. 1988, Turner and Pounds 1989, Hasinger 1992).  Franceschini et al. (1993) proposed a scenario to account for these facts: the existence of two different populations, the soft X--ray sources, with steep spectrum and high evolution rates and the hard X-ray sources, with a weak cosmological evolution and strong self-absorption. Recent models for the origin of the X-ray background (Madau, Ghisellini \\& Fabian 1994, Comastri et al 1995) based on the AGN unified scheme (Antonucci \\& Miller 1985, Antonucci 1993), rather suggest that there is a continuity between these two populations.  We analyse a sample of AGN for which there are 2-10 keV fluxes (mostly coming from the HEAO-1 MC-LASS survey, Wood et al 1984) and archival {\\it Rosat} PSPC observations. By analysing the PSPC hardness ratios versus 0.5-2 keV to 2-10 keV flux ratio, we conclude that the simplest spectral model that can accommodate the whole sample is a power law plus a blackbody both partially absorbed.  Furthermore the model we propose is able to bring into consistency the source counts as well as the AGN luminosity functions in both bands. Our comparison is relevant only to local ($z<0.2$) AGN and its extension to higher redshifts would require more data.  We also consider an alternative to this partial covering scenario which can account for the spectra of AGN as well as for the soft excess observed: a full obscuration of the X--ray continuum by partly ionized gas (see Netzer 1993 and references therein). However unless some neutral absorbing material is also present the spectra are invariably too soft.  In section 2 we describe the sample, present the broad-band hardness ratio versus flux ratio relation, and introduce the simplest model also able to accommodate the wide range of parameter space occupied by these sources. The warm absorber model is also introduced in this section as an alternative model to describe the spread in the observed parameters. In section 3 we show that a partial covering model is able to bring soft and hard X-ray AGN luminosity functions and source counts into agreement. We summarize the results and present some conclusions in section 4. ",
        "conclusions": "We have studied the hardness ratios for a sample of hard X-ray selected sources, all of them observed by ROSAT, and derived a source emission spectrum that would explain their X--ray colours. This spectrum has three components: a power law with an energy index $\\alpha\\sim 0.9$, a blackbody of $kT\\sim 0.1$ keV representing about 50\\% of the power law at 1~keV and a low-energy absorption by neutral gas with column density $N_{H}\\sim 10^{22}-10^{23}\\> \\rm cm^{-2}$ that partially covers the source ($f_{cov}\\sim 80-100\\%$). Type 2 AGN appear to be fully covered with large column densities while type 1 AGN have an average covering factor $\\sim 0.82$ and absorbing column density $\\approx 10^{22}\\, {\\rm cm}^{-2}$.  Qualitatively this model can account for the scatter observed in the hardness ratios as well as for the soft excess found in some sources.  A complementary result of our analysis is that starting from a single population of AGN with a distribution of covering factors, it is possible to describe the two distinct populations proposed by Franceschini et al (1993). Their soft X--ray class of active galaxies, showing steep power law spectra and being easily detected by soft X-ray band missions would correspond to those sources that in our model had lower covering factors and eventually exhibit soft excess emission. The hard population would be composed of those sources strongly self-absorbed by high covering factors. The selection of objects in the 2-10 keV band does not particularly favour high values of the covering factor, as it is demonstrated by the presence of a large fraction of partially covered type 1 AGN in the sample used in Section 2. However, for a soft X-ray selected sample, the average covering factor will certainly be smaller. Our model is more alike the one used by Comastri et al (1995) where a single AGN population is used to reproduce the spectrum of the X-ray background.  As far as the mismatch between the number counts in different X-ray bands is concerned, a typical spectrum with a 50\\% blackbody contribution at 1 keV, a power law energy index $\\alpha\\sim 0.9$, an absorbing column density between $N_{H}\\sim 10^{22}$ and $N_{H}\\sim 10^{23}$ and a covering factor $f\\approxgt 85\\%$ would result in a flux ratio $F(2-10\\,\\rm keV)/F(0.3-3.5\\, \\rm keV)\\sim 2$ (see Figure 5) which is the required value to solve the discrepancy.  Thus, the model presented here with the average parameters obtained in Section 2 also brings the soft and hard source counts into agreement.  We also tried to describe the dispersion on the X--ray colours through a warm absorber model and the conclusion we can extract from this analysis is that the presence of gas in two phases (neutral and partly ionized) is required to reproduce the X--ray colours of most of the sampled AGN.  Finally, we can derive an interesting conclusion from this analysis: the gas responsible for the absorption of the X--ray primary spectrum emitted by the active nucleus must have structure. This structure could be due to holes or to the coexistence of gas in two phases (neutral and ionized).  This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA-Goddard Space Flight Center. We acknowledge J.P.D.  Mittaz and F.J. Carrera at Mullard Space Science Laboratory for help and Ruth Carballo for careful reading. We also thank the referee for helpful comments and A. Fabian for interesting suggestions. Partial financial support for this research was provided by the Spanish DGICYT under project PB92-0501. MTC was supported by a fellowship from the Universidad de Cantabria."
    },
    "9604/astro-ph9604125_arXiv.txt": {
        "abstract": "Velocity dispersions $\\sigma$ and Mg absorption line-strengths Mg$_b$ have been measured for a sample of 16 ellipticals in 3 clusters at a redshift of 0.37. Like local cluster ellipticals, these objects show a correlation between Mg$_b$ and $\\sigma$.  However, at any given $\\sigma$, the mean Mg$_b$ of the ellipticals at $z=0.37$ is weaker than the mean Mg$_b$ of their local relatives in the Coma and Virgo clusters. The Mg$_b$ weakening is smallest for the most luminous ellipticals and larger for the fainter objects. This is unambiguous evidence for {\\it small but significant passive evolution} of the stellar populations of elliptical galaxies with redshift.  It requires that the bulk of the stars in cluster ellipticals has formed at $z>2$. The most luminous objects may even have formed at $z>4$. The Mg$_b-\\sigma$ test is a very reliable estimator for the evolution of old stellar populations because it is virtually independent from the stellar initial mass function (IMF) and from the metallicities of the galaxies. Furthermore, the influence of selection effects is minimal.  Consistent with the weakening of Mg$_b$ we find evidence for a B-band luminosity evolution of about $0.5\\pm 0.1$mag at $z=0.37$ from the Faber-Jackson relation.  The combined information about the evolution of Mg$_b$ and luminosity allows us to constrain both the slope of the IMF in ellipticals and the cosmological deceleration parameter $q_o$. Our present measurements are compatible with a standard Salpeter IMF and a $q_o$ of $0.5\\pm 0.5$. ",
        "introduction": "The redshift evolution of galaxies provides strong constraints on their ages and formation processes as well as on theoretical models of structure formation in the Universe. In the specific case of elliptical galaxies, there is the additional prospect that, because of their homogenous properties, it may eventually be possible to calibrate their evolution accurately enough to allow their use as cosmological standard candles or rods.   The luminosity and color evolution of elliptical galaxies or brightest cluster members is generally found to be weak.  For various samples of brightest cluster galaxies up to redshifts $z=1$, both optical and infrared colors were found to change only very slowly with $z$ indicating a high redshift of formation (e.g., Arag\\'on-Salamanca \\etal 1993, Stanford \\etal 1995). Similarly, Hamilton (1985) found only weak evolution in the strength of the 4000\\AA\\ break. The bright end of the galaxy luminosity function does not evolve significantly with redshift either, being compatible even with a no-evolution scenario (e.g., Glazebrook \\etal 1995, Lilly \\etal 1995, Ellis \\etal 1996). Observations of distant galaxies up to $z=1$ from the ground and with the Hubble Space Telescope revealed no significant deviations of the surface brightness evolution from the Tolman relation, again an indication of very small if any evolution (Sandage \\& Perelmuter 1991, Dickinson 1995, Pahre et al. 1996). At lower redshifts ($z<0.4$), Franx \\& van Dokkum (1996) showed that the mass-to-light ratios of ellipticals as obtained from the fundamental plane relation (Djorgovski \\& Davis 1987, Faber \\etal 1987) change only very slowly with redshift but in accordance with passive evolution of a very old population.  These results are all indicative of a virtually passive evolution of most cluster ellipticals up to $z\\approx 1$. This may be in part a selection effect because ellipticals that experience a merging or accretion event may be transformed into an E$+$A-type (Dressler \\& Gunn 1983) or even bluer object and therefore drop out of an ellipticals sample for some time. But in general, these events cannot have injected a significant fraction of young stars into brightest cluster ellipticals below redshifts of 1. Otherwise, the observed redshift evolution would be much stronger and, also, it would be difficult to explain the high uniformity of colors and absorption line indices in {\\it local} cluster ellipticals (Bower \\etal 1992, Bender, Faber \\& Burstein 1993).  This does not contradict the claim that many ellipticals were originally formed in {\\it major} mergers (e.g., Schweizer 1990, Bender 1990b).  It simply means that, in rich clusters, the {\\it major} merging and star formation epoch is to be found at higher redshifts (e.g., Bender, Burstein and Faber 1993, 1994). This picture has recently been shown to be even in agreement with cold dark matter models (Kauffmann 1995).  With the redshift evolution of elliptical galaxies being very weak, the challenge of measuring it accurately is correspondingly bigger. In this paper we will analyse the evolution of ellipticals with a method that is principally different from the methods applied previously. Our test is based on the tight relation between Mg-Index Mg$_b$ and velocity dispersion $\\sigma$ of elliptical galaxies and will be described in Section 2.  Section 3 briefly describes sample selection and observations, Section 4 the data analysis. Section 5 presents results and conclusions. ",
        "conclusions": "{\\it Figure 1} shows the \\mgs relation for cluster ellipticals at a redshift of $z = 0.37$ in comparison to local elliptical galaxies in the Coma and Virgo cluster. Two main conclusions can be drawn:  \\placefigure{fig1}  (a) The distant ellipticals also show a correlation between Mg$_b$ index and velocity dispersion $\\sigma$ as local ellipticals do. The slope of the \\mgs relation at higher redshift appears to be slightly steeper than today indicating that low luminosity ellipticals may be systematically younger than high luminosity ellipticals. This may support a recent claim by Faber \\etal (1995) according to which the mean ages of ellipticals should systematically decrease with decreasing luminosity. The distant objects with {\\it very} low Mg$_b$ could be genuinely younger or may only have experienced a more recent but smaller star formation event.   (b) There is clear evidence for evolution: at any given velocity dispersion, the mean Mg$_b$ of ellipticals at $z=0.37$ is significantly weaker than at $z=0$. The effect is relatively weak for big ellipticals and larger for small ellipticals.  On average, the Mg$_b$ absorption is reduced by about 0.3~\\AA. Note that this conclusion is not affected by selection biases since our color selection criterion (see Section 3) only cuts off ellipticals with Mg$_b < 3\\,$\\AA.  \\smallskip On the basis of Worthey's (1994) population synthesis models we can use this weaking of Mg$_b$ to estimate the ratio of the mean age of the $z=0.37$ ellipticals relative to the mean age of local ellipticals. We derive the following dependence of Mg$_b$ on age $t$ and metallicity $Z$: $\\log$~Mg$_b = 0.37 + 0.20~\\log t + 0.31~\\log Z/Z_o$ (valid for $t > 5\\,$Gyrs and $1/3 < Z/Z_\\odot < 3$). The proportionality factors of $\\log t$ and $\\log Z$ are very robust and only weakly dependent on IMF and on uncertainties in the population synthesis models (Bruzual 1996).  The small evolution we see in {\\it Figure 1} strongly supports the notion that ellipticals have evolved mostly passively between $z=0.37$ and $z=0$. Since then $Z \\approx const.$, the observed reduction of Mg$_b$ can be directly translated into a relative age difference. The three dashed lines in {\\it Figure 1} give the {\\it expected} \\mgs relation at $z=0.37$, if the mean formation redshifts of the stars in ellipticals were formed at $z_f = 4.5, 2, 1$, respectively. Independent from H$_o$ and only weakly dependent on q$_o$ we conclude that {\\bf the bulk of the stars in the {\\it luminous} cluster ellipticals must have formed at redshifts} $\\bf z>2$, the higher luminosity objects may even have formed at $\\bf z>4$. Because of the possible presence of some younger stars injected in minor accretion events and because we observe luminosity-weighted spectra, these age estimate is rather a lower limit.  \\placefigure{fig2}  \\smallskip The age-driven reduction of the Mg$_b$ absorption in ellipticals at higher redshifts should correspond to an increase in their luminosity. This should be detectable via the Faber-Jackson relation (Faber \\& Jackson 1976).  In order to derive the Faber-Jackson relation for our distant cluster ellipticals, their rest-frame $B$ luminosities were calculated as described in Section 4 and using $H_o= 50$~km/s/Mpc and $q_o=0.5$.  The Faber-Jackson relation for the $z=0.37$ cluster ellipticals and for the Coma and Virgo ellipticals is shown in the upper panel of {\\it Figure 2}. It is evident that the distant ellipticals are brighter than their local counter-parts by about 0.5$\\pm 0.1\\,$mag in the B-band.  Is this consistent with the weakening of Mg$_b$ in the Mg$-\\sigma$ relation? Using Worthey's models with a Salpeter IMF, we obtain that age variations lead to a correlated change of the B-band luminosity with the Mg$_b$ value according to $\\rm \\Delta  = 1.2 \\Delta\\, Mg_b/$\\AA. Thus, the evolution of the B \\lum as estimated from the weakening of Mg$_b$ should be $\\Delta B$($z\\approx0.4$) $\\approx -0.4$ mag, consistent with the actually observed brightening. This corresponds to an evolution in rest frame $(B-V)$ of $-0.04\\,$ mag and $(V-K)$ of $-0.12\\,$mag, in accordance with the measurements of Stanford et al. (1995).  We can also turn the above procedure around and correct the luminosity of each galaxy by its individually calculated evolution correction $\\rm \\Delta M_B = 1.2 \\Delta\\, Mg_b/$\\AA. The lower panel of {\\it Figure 2} shows the result: the $z=0.37$ ellipticals and the Coma and Virgo ellipticals now all fall on top of each other. Even the slope of the Faber-Jackson relation at $z=0.37$ is similar within the errors to the locally observed slope.   We conclude that the evolution of the stellar populations in ellipticals as derived from the \\mgs relation and from the Faber-Jackson relation are consistent. Within the limits of our errors, there is no evidence for a very unusual slope of the IMF, nor a very unusual value of the cosmological constant. Changing the slope of the IMF by $\\Delta x =1$ would cause a change of the luminosities at $z=0.37$ by ca. $\\pm 0.13$mag (Bruzual \\& Charlot 1996 models). Similarly, changing $q_o=0.5$ by $\\pm 0.5$ causes a change of the luminosities by ca. $\\pm 0.11$mag.  In a future paper, we will combine our spectroscopic data with HST photometry. This will allow us to investigate the evolution of the elliptical galaxies more accurately by employing the fundamental plane relations. In combination with the Tolman surface brightness test, one may then also be able to obtain better and, more importantly, independent constraints on the slope of the IMF and on the value of $q_o$."
    },
    "9604/astro-ph9604077_arXiv.txt": {
        "abstract": "The growth of the angular momentum $\\bfL$ of protogalaxies induced by tidal torques is reconsidered.  We adopt White's formalism and study the evolution of $\\bfL$ in Lagrangian coordinates; the motion of the fluid elements is described by the Zel'dovich approximation. We obtain a general expression for the ensemble expectation value of the square of $\\bfL$ in terms of the first and second invariant of the inertia tensor of the Lagrangian volume $\\Gamma$ enclosing the protoobject's collapsing mass. We then specialize the formalism to the particular case in which $\\Gamma$ is centred on a peak of the smoothed Gaussian density field and approximated by an isodensity ellipsoid. The result is the appropriate analytical estimate for the rms angular momentum of peaks to be compared against simulations that make use of the Hoffman-Ribak algorithm to set up a constrained density field that contains a peak with given shape. Extending the work of Heavens \\& Peacock, we calculate the {\\it joint} probability distribution function for several spin parameters and peak mass $M$ using the distribution of peak shapes, for different initial power spectra. The probability distribution for the rms final angular momentum $\\langle \\bfL_f^2\\rangle^{1/2}$ on the scales corresponding to common bright galaxies, $M\\approx 10^{11} M_{\\odot}$, is centred on a value of $\\approx 10^{67}\\, {\\rm kg}\\,{\\rm m}^2\\,{\\rm s}^{-1}$, for any cosmologically relevant power spectrum, in line with previous theoretical and observational estimates for $L_f$. Other astrophysical consequences are discussed. In particular, we find that typical values $\\lan \\lambda^2\\ran^{1/2}\\approx 0.1$ of the dimensionless spin parameter for peaks smoothed on galactic scales and of height $\\nu\\sim 1$, usually associated with late type galaxies, may be recovered in the framework of the Gaussian peak formalism. This partially relaxes the importance attributed to dissipative processes in generating such high values of centrifugal support for spiral galaxies. In addition, the values of the specific angular momentum versus mass -- as deduced from observations of rotational velocities and photometric radii of spiral galaxies -- are well fitted by our theoretical isoprobability contours. In contrast, the observed lower values for the specific angular momentum for ellipticals of the same mass cannot be accounted for within our linear regime investigation, highlighting the importance of strongly non-linear phenomena to explain the spin of such objects. ",
        "introduction": "It has been argued that tidal coupling between the inhomogeneities in the primordial matter distribution, in the context of a gravitational hierarchical scenario, may explain the acquisition of the angular momentum by a protogalaxy.  This idea, originally due to Hoyle (1949) and applied by Sciama (1955), has been first thoroughly examined by Peebles (1969), who demonstrated that the tidal spin growth of the matter contained in a spherical (Eulerian) volume is proportional to $t^{5/3}$ in an Einstein--de Sitter universe ($t$ is the standard cosmic time).  Specifically, Peebles' analysis is based on a {\\it second-order} perturbative description, since a spherical volume does not gain angular momentum from tidal torques in linear approximation, as pointed out by Zel'dovich and reported by Doroshkevich (1970).  An important result of Doroshkevich's paper is that a {\\it generic nonspherical} volume enclosing the protogalaxy acquires tidal angular momentum proportionally to the cosmic time $t$ during the linear regime.  This theoretical prediction has been confirmed by the $N$--body simulations of White (1984). In addition, White showed that the second-order growth described by Peebles is due to convective motion of matter across the surface of the initial volume $\\Gamma$ containing the protoobject.  An important point of Peebles and White's theoretical analyses is that they describe the tidal torques acting on a volume centred on a {\\it random} point in the smoothed density field, which does not necessarily enclose a bound protogalaxy. In contrast, galaxies are expected to form around (high) peaks on relevant scales in the density field.  This idea, in embryo in Doroshkevich (1970), developed into the biased galaxy formation scenario, where only density maxima above a given threshold (peaks) of the initial Gaussian density field can eventually form galaxies (Kaiser 1984; Politzer \\& Wise 1984; Peacock \\& Heavens 1985; Bardeen \\etal 1986).  The acquisition of angular momentum due to tidal torques by Gaussian high--density peaks has been recently analyzed, amongst others, by Hoffman (1986; 1988), Heavens \\& Peacock (1988), Ryden (1988), Quinn \\& Binney (1992) and Eisenstein \\& Loeb (1995). Comparisons with $N$--body simulations are displayed in Efstathiou \\& Jones (1979), Barnes \\& Efstathiou (1987) and Warren \\etal (1992).  In this paper, we re-examine the growth during the linear regime of the angular momentum $\\bfL$ of protogalaxies induced by tidal couplings with the surrounding matter inhomogeneities. The layout of this article is the following. In the next section we briefly review White's formalism describing the linear evolution of tidal galactic spin. The motion of the mass fluid elements is described by the Zel'dovich approximation: the invariance of the angular momentum with respect to the Eulerian and Lagrangian description is stressed.  Next, we derive a general (but approximate) expression for the ensemble expectation value of the square of $\\bfL$, $\\lan \\bfL^2\\ran$, in terms of the first and second invariant of the inertia tensor of the Lagrangian volume $\\Gamma$. We then specialize our formalism to the particular case in which the Lagrangian volume is centred on a peak of the underlying smoothed Gaussian density field and approximated by an isodensity profile ellipsoid: in this case, we obtain the correct constrained ensemble average for these objects with preselected inertia tensor. The result is the appropriate analytical estimate for the rms angular momentum of peaks to be compared against simulations that make use of the Hoffman-Ribak algorithm to set up a constrained density field that contains a peak with given shape (Hoffman \\& Ribak 1991; van de Weygaert \\& Bertschinger 1996).  Extending the work of Heavens \\& Peacock, we calculate the {\\it joint} probability distribution function for several spin parameters (angular momentum, specific angular momentum, angular momentum in units of $M^{5/3}$) and peak mass $M$ using the distribution of peak shapes (Bardeen et al. 1986), for different density power spectra.  Finally, we discuss astrophysical implications in the last section. Technical details are given in appendices. ",
        "conclusions": "In this paper we re-analyzed the problem of the acquisition of angular momentum by a protoobject, progenitor of a galaxy or a cluster, due to tidal interactions with the surrounding matter distribution. This process is of most interest for gravitational instability theories of galaxy and cluster formation. In section 2, we reviewed the dynamical description of the spin evolution in the version given by White (1984), hence the motion of the matter patch of fluid is followed by applying the Zel'dovich approximation. The expression for the linear tidal angular momentum $\\bfL$ of a protoobject obtained in this formalism and reported in equation~(7), contains a combination of the deformation tensor $\\calD_{\\al\\beta}$ and the inertia tensor $\\calJ_{\\al\\beta}$ of the matter inside the Lagrangian volume $\\Gamma$, which by definition encloses the collapsing protoobject. Such a combination of the elements of the tensors $\\calD$ and $\\calJ$ is zero if $\\Gamma$ is a spherical volume (or if the boundary of $\\Gamma$ is an equipotential of the gravitational potential). The temporal evolution of the spin is governed in the linear regime by the function $a(t)^2\\,\\dot{D}(t)$, where $a(t)$ is the scale factor and $D(t)$ is the growth factor of the linear density perturbations, which is proportional to the cosmic time $t$ in the Einstein-de Sitter universe (Doroshkevich 1970).  These results may be notably simplified by considering the ensemble expectation value of the square of the angular momentum $\\lan\\bfL^2\\ran$. The resulting general (albeit approximate) expression in equation~(16) is one of the main results of this paper. It highlights the very simple fact that, in linear regime, the ensemble average $\\lan\\bfL^2\\ran$ is a function of the first and second invariant of the inertia tensor $\\calJ$ alone and is in addition proportional to the mass variance $\\s_0(R)^2$ on the scale $R$. The fact that $\\lan\\bfL^2\\ran$ depends only on the invariants of the inertia tensor and not on the detailed shape of the surface boundary of $\\Gamma$ provides a considerable simplification of the calculation.  In section~2.3, we specialised the statistic $\\lan\\bfL^2\\ran$ to the case in which the volume $\\Gamma$ is centered on a peak of the underlying Gaussian density field. The ensemble average is therefore restricted to those realizations of the density field that are compatible with the preselected shape of the collapsing object. This constrained ensemble average corresponds to an unconstrained ensemble average over the off-diagonal shear (see the extensive discussion in Appendix~B). The analysis of the tidal torques acting on matter in the vicinity of local density maxima (peaks) during the linear regime has been first attempted by Hoffman (1986; 1988) and Heavens \\& Peacock (1988). In particular, Hoffman first analyzed the correlation between the height $\\nu$ of the peak and the amount of angular momentum it acquires by tidal interactions. Heavens \\& Peacock performed a more sophisticated investigation, taking into account also the asymmetric shape of the matter peak, described in terms of the eigenvalues $-\\lam_1, -\\lam_2, -\\lam_3$ of the mass tensor $\\p_\\al\\p_\\beta\\de$ (Bardeen et al. 1986). This represents an important improvement because the strength of the tidal interactions depends strongly on the shape of the object. Our analysis uses an approach complementary to the one of Heavens and Peacock and extends their work.  The resulting expression for $\\lan\\bfL^2\\ran$ appropriate for peaks has been recast in equation~(\\ref{eq:Ltor}) into two factors: $\\calL\\equiv\\sqrt{\\lan\\bfL^2\\ran}=\\ell\\,\\calL_{\\ast}$, where $\\ell=\\ell(\\nu,x,e,p;\\gamma)$ is dimensionless and contains the dependence on peak shape (height $\\nu$, sharpness $x$, ellipticity $e$ and prolateness $p$) and $\\calL_{\\ast} = \\calL_{\\ast}(t;R_\\ast)$ is the angular momentum unit which contains the growth rate and is fixed by the cut-off $R_\\ast$ of the power spectrum. This factorization should be compared with the one in Heavens \\& Peacock (1988) for the modulus $J$ of the angular momentum, their equation~(9), which corresponds to a single realisation of the shear field. We then proceed by computing the probability distribution of several spin variables (e.g. of $\\ell$, $\\ell/m$, $\\ell/m^{5/3}$ versus mass $m$ of the peak; $m$ is the dimensionless mass of the peak in units of the mass selected by $R_\\ast$) by using the probability distribution of peak shape parameters (Bardeen et al. 1986). This is an important extension of the work of Heavens \\& Peacock (1988), who analysed the rotational properties of galaxies directly in terms of spin parameters versus the peak height $\\nu$. The latter, however, does not allow to determine uniquely the mass-scale of the progenitor. As we argued at the end of section~2.3, we expect our distributions of angular momentum versus mass to be more reliable than the estimates for single objects.  In sections~3 and 4, we computed and discussed probability distributions for $\\ell$, $\\ell/m$ and $\\ell/m^{5/3}$ and the dimensionless spin parameter $\\lambda$, both in linear regime and at the time of maximum expansion. Our findings can be summarized as follows: \\begin{itemize} \\item linear values of the spin $\\ell$ have a strong dependence on $m$: $\\ell\\propto m^{5/3}$, see Fig.~\\ref{fig:l}  \\item for objects of a restricted range in $\\nu$, this dependence steepens to $\\ell\\propto m^{8/3}$, see Fig.~\\ref{fig:lm}  \\item final values (i.e., at maximum expansion time) $\\ell_f$ of $\\ell$ scale shallower with $m$, $\\ell_f\\propto m^{2/3}$, see Fig.~\\ref{fig:j53_final}  \\item median values of $\\ell_f$ decrease with $\\nu$ for high $\\gamma$ (i.e. steep spectra) but are rather insensitive to $\\nu$ for low $\\gamma$ (i.e. broad spectra)  \\item median values of $\\ell_f$ decrease with $\\gamma$ for given $\\nu$  \\item median values for $\\lambda$ are $\\overline \\lambda(\\nu=1)= 0.15$ and $\\overline\\lambda(\\nu=2)= 0.07$, comfortably close to the spin of spiral galaxies but marginally higher than the value typically quoted for ellipticals  \\item dimensional values of $\\calL_f/M$ on galactic scales are in good agreement with measured values for spirals but are too high to described ellipticals  \\item selecting objects on galactic scales with peak height $\\nu\\sim 1$ gives values of $\\calL_f/M$ typical of spiral galaxies. Selecting objects with $\\nu\\sim 2$ moves the position of the most probable $\\calL_f/M$ in the direction of the location of elliptical galaxies in the $(M,\\calL_f/M)$ plane but the shift falls short by a large margin to give good agreement between values of $\\calL_f/M$ for $\\nu=2$ peaks and those observed for ellipticals.  \\item typical values of the angular momentum of a spiral galaxy are predicted to be $\\calL_f\\approx 1.8\\times 10^{67}\\,{\\rm kg~m}^2{\\rm ~s}^{-1}$ with a corresponding circular velocity $v_{c,{\\rm spiral}}\\approx 140\\,{\\rm km~s}^{-1}$. For a rich cluster, we find a typical value $v_{c,{\\rm clus.}}\\approx 5\\,{\\rm km~s}^{-1}$ while for a supercluster $v_{c,{\\rm sup. clus.}}\\le 10\\,{\\rm km~s}^{-1}.$ \\end{itemize} The reliability of our estimate for the ensemble averaged value $\\Lambda_f\\equiv\\sqrt{\\langle\\lambda^2_f\\rangle}$ depends on the extent to which our approximations capture the main features of the dominant processes. We have made the following assumptions:~(1) linear perturbation theory, extrapolated to the mildly non-linear regime (Zel'dovich approximation), can be used to describe the growth of $\\ell$ until the maximum expansion time ;~(2) tidal torques spin up the matter until the maximum expansion time of the protoobject and are negligible thereafter;~(3) the mass of the object can be identified with the mass inside an isodensity ellipsoidal surface around the peak;~(4) the spherical model can be used to calculate the clump's binding energy $E$ -- needed to compute $\\lambda$;~(5) Gaussian peaks formalism (Bardeen et al. 1986) can be used to compute probability distributions for $\\ell_\\beta$.  We briefly comment on the appropriateness of these approximations. The use of the Zel'dovich approximation, very powerful in describing the mildly non-linear evolution of matter before shell-crossing, may give only a partial description of the evolution of the tidal angular momentum of the clumps themselves. Certainly, it cannot predict the very final stages of evolution when clumps merge and interact non-linearly, which leads to the present-day galactic configurations. In fact, as shown by the numerical simulations of e.g. Barnes \\& Efstathiou (1987), Frenk (1987), Zurek, Quinn \\& Salmon (1988), Navarro \\& Benz (1991) and Navarro \\& White (1994), the tidally acquired spin may change drastically as objects merge or as angular momentum gets transported between core and halo. The merging history of the surrounding protohalo is a key factor in the determination of the morphology of a galaxy and the merging processes of dense clumps are associated with substantial $loss$ of angular momentum to the halo. $\\calL_f$ -- as obtained from the extrapolation of the linear theory -- is typically a factor of $\\sim$ 3 larger than the final spin of the non-linear object (Barnes \\& Efstathiou 1987; Frenk 1987). At present, we see little hope of computing theoretically such a reduction factor, which is due to non-linear interactions as well as late infall and dissipation processes. Consequently, these processes constrain the applicability of perturbation theory to the period before the maximum expansion time (Catelan \\& Theuns 1996).  The determination of the surface boundary of the Lagrangian volume $\\Gamma$ is tricky. This boundary determines the mass of the object. We assumed that it can be described by an isodensity contour, which is ellipsoidal when the object is centered on a high peak. However, a non-negligible part of the luminous matter might have been captured from the very outer regions of the proto-galaxies (Ryden 1988; Quinn \\& Binney 1992), where the ellipsoidal model is presumably a poor description. In addition, the formalism applied here is also based on Taylor expanding the density and the velocity fields around the centre of the peak. This expansion breaks down far from the peak's centre and consequently is unable to accommodate late infall (see the discussion in Hoffman 1988), a process definitely important during the late stages of galaxy formation (e.g. Navarro \\& White 1994).  In assuming the spherical model to estimate the clump's binding energy $E$, we have lost the explicit dependence of $E$ on the properties of the underlying linear gravitational potential field $\\psi$. Consequently, our calculation could be improved by taking this dependence into account (see Hoffman 1988, section II), but this would complicate considerably the computation of the ensemble average and is beyond the scope of the present paper.  Finally, in comparing our predictions against measurements of luminous parts of galaxies, one has to bear in mind that dissipative processes, not taken into account here, surely have had a major influence in shaping those objects."
    },
    "9604/astro-ph9604056_arXiv.txt": {
        "abstract": "The theory of diffusive acceleration of energetic particles at shock fronts assumes charged particles undergo spatial diffusion in a uniform magnetic field. If, however, the magnetic field is not uniform, but has a stochastic or braided structure, the transport of charged particles across the average direction of the field is more complicated. Assuming quasi-linear behaviour of the field lines, the particles undergo sub-diffusion on short time scales. We derive the propagator for such motion, which differs from the Gaussian form relevant for diffusion, and apply it to a configuration with a plane shock front whose normal is perpendicular to the average field direction. Expressions are given for the acceleration time as a function of the diffusion coefficient of the wandering magnetic field lines and the spatial diffusion coefficient of the charged particles parallel to the local field. In addition we calculate the spatial dependence of the particle density in both the upstream and downstream plasmas. In contrast to the diffusive case, the density of particles at the shock front is lower than it is far downstream. This is a consequence of the partial trapping of particles by structures in the magnetic field. As a result, the spectrum of accelerated particles is a power-law in momentum which is steeper than in the diffusive case. For a phase-space density $f\\propto p^{-s}$, we find $s=\\sdiff[1+1/(2\\rcomp)]$, where $\\rcomp$ is the compression ratio of the shock front and $\\sdiff$ is the standard result of diffusive acceleration: $\\sdiff=3\\rcomp/(\\rcomp-1)$. A strong shock in a monatomic ideal gas yields a spectrum of $s=4.5$. In the case of electrons, this corresponds to a radio synchrotron spectral index of $\\alpha=0.75$. ",
        "introduction": "\\label{transport} Our approach to the transport process involves separating it into two parts (Chuvilgin \\& Ptuskin~\\cite{chuvilginptuskin93}). The first of these consists of macroscopic magnetic fluctuations of length scale large compared to the gyro radius of the particle concerned, which we characterise by a relative amplitude $b\\equiv\\left<|\\delta B|\\right>/\\left<B\\right>$. As described in Duffy et al.~(\\cite{duffyetal95}) we  assume these fluctuations lead to a quasi-linear type diffusion or wandering of the field lines in the plane perpendicular to the direction of the average field. This is described by a magnetic diffusivity $\\bdiff$, which, in terms of cartesian coordinates with the $z$-axis along the average field, is defined by: \\eqb {\\left<\\Delta x^2\\right>\\over 2 s} &=& \\bdiff \\enspace, \\eqe where $\\Delta x$ is the change in the $x$ co-ordinate upon travelling a distance $s$ along the field line. If the turbulence is characterised by correlation lengths $\\lpar$ and $\\lperp$ along and across the average field, then assuming quasi-linear behaviour and adopting the normalisation of Kadomtsev \\& Pogutse~(\\cite{kadomtsevpogutse79}), we have \\eqb \\bdiff&=&{b^2\\lpar\\over 4} \\enspace. \\eqe (Note that different normalisations are used by Achterberg \\& Ball~(\\cite{achterbergball94}) and Isichenko~(\\cite{isichenko91a})).  The second component of the transport concerns length scales comparable to that of the gyro radius. We assume that any anisotropy in the distribution function results in the rapid growth of magnetic fluctuations (Alfv\\'en waves) which, in their turn, scatter the particles so as to remove the anisotropy. This process we model in terms of two spatial diffusion coefficients $\\kpar$ and $\\kperp$ responsible for transport along and across the local field direction, respectively. In keeping with other treatments (e.g., Achterberg \\& Ball~\\cite{achterbergball94}, Jokipii~\\cite{jokipii87}) we parameterise these coefficients in terms of the gyro-Bohm diffusion coefficient $\\kbohm\\equiv\\gamma v^2 mc/(3eB)$, for a particle of mass $m$, charge $e$ moving at speed $v$, with Lorentz factor $\\gamma=(1-v^2/c^2)^{-1/2}$ in a magnetic field $B$: \\eqb \\kpar &=& {\\kbohm\\over\\epsilon} \\nonumber\\\\ \\kperp &=& {{\\epsilon\\kbohm}\\over{(1+\\epsilon)}} \\label{eqkperp}\\enspace. \\eqe Here $\\epsilon\\le1$ is the ratio of the energy density in microscopic fluctuations to that in the average magnetic field. In the picture in which the scattering is modelled by the \\lq\\lq$\\tau$\\rq\\rq\\ operator (Chuvilgin \\& Ptuskin~\\cite{chuvilginptuskin93}), $\\epsilon$ is the ratio of the collision rate to the gyro-frequency.  Thus we have on the one hand microscopic fluctuations of relative amplitude $\\sqrt{\\epsilon}$ which are responsible for the diffusion of particles along and across the local field and, on the other, macroscopic fluctuations of relative amplitude $b$ which are responsible for diffusion of the field lines. We shall assume $\\epsilon\\ll1$, which means that on the microscopic scale, particles are closely tied to field lines along which they diffuse. Transport across the local field direction is therefore severely restricted on the microscopic scale.  The important physical process in such a picture has been described by Rechester and Rosenbluth~(\\cite{rechesterrosenbluth78}): because of the exponential divergence of neighbouring field lines, a particle's orbit will take it out of regions of correlated magnetic field. This can happen either because the particle diffuses a small distance across the field by scattering off the microscopic fluctuations, or, in the case of very small $\\epsilon$, because the finite size of its gyro orbit causes it to encounter uncorrelated field lines (Isichenko~\\cite{isichenko91a}). The latter case we term \\lq\\lq ballistic\\rq\\rq\\ propagation. It occurs if the particle decorrelates from the field before having a chance to diffuse microscopically, which translates into the condition \\eqb {b^2\\over\\epsilon}>{\\lperp^2\\over\\lpar\\rgyro}\\log\\left({\\lperp\\over\\rgyro}\\right) \\eqe (Duffy et al.~\\cite{duffyetal95}).  Ballistic propagation has been considered by Achterberg \\& Ball~(\\cite{achterbergball94}) in connection with the radio emission of supernovae. The particles propagate diffusively, with an effective spatial diffusion coefficient across the field given by \\eqb \\kperp^{\\rm ballistic}&=& v\\bdiff \\enspace, \\eqe and the standard theory of diffusive acceleration at shocks applies. However, in this paper we will assume the self-excited microscopic turbulence is sufficiently strong to enable particles to scatter before they decorrelate from the magnetic field. In this case, one can define a dimensionless parameter \\eqb \\Lambda&=&{b^2\\lpar\\over\\sqrt{2}\\epsilon\\lperp} \\eqe such that for $\\Lambda\\lesim1$ the macroscopic braiding of the field is irrelevant, whereas it dominates cross-field transport for $\\Lambda\\gg1$. If braiding is irrelevant, there is no anomalous transport regime, and no modification of the diffusive acceleration picture, provided the distribution can still be considered isotropic (cf.~Achterberg \\& Ball~\\cite{achterbergball94}).  The most interesting parameter regime is that in which both braiding and microscopic scattering are important, which occurs for the parameter range \\eqb {\\lperp\\over\\sqrt{2}\\rgyro}\\log\\left({\\lperp\\over\\rgyro}\\right) > \\Lambda\\gg1 \\enspace. \\eqe Two kinds of propagation are then possible: for times less than that needed to decorrelate from the field $\\tdecorrel$, the particles undergo sub-diffusion, which is a combination of diffusion along a fixed field line, which itself diffuses. The defining characteristic of sub-diffusion is that the mean square deviation of a particle in the $x$ direction from its position at time $t=0$ is not proportional to $t$ as in ordinary diffusion, but rather to $\\sqrt{t}$.  The propagator $P_{\\rm sub}(x,t)$, which is the probability of finding a particle in the interval ($x$,$x+\\diff x$) at time $t$, given that it was at the origin $x=0$ at $t=0$, is found by simply folding the two gaussian propagators appropriate to diffusion of a particle along the field (i.e., in $s$) and of the field in $x$: \\eqb P_{\\rm sub}(x,t)&=& {1\\over 4\\pi\\sqrt{\\bdiff\\kpar t}} \\nonumber\\\\ &&\\int_{-\\infty}^{+\\infty} \\diff s {1\\over\\sqrt{|s|}} \\exp\\left(-{x^2\\over4\\bdiff |s|}-{s^2\\over4\\kpar t}\\right) \\label{subdiffprop} \\enspace. \\eqe (Rax \\& White~\\cite{raxwhite92}, Duffy et al.~\\cite{duffyetal95}). It is straightforward to confirm that the mean square value of $x$ increases as $t^{1/2}$ with this propagator. Equation~(\\ref{subdiffprop}) describes the motion of an injected particle for times shorter than a decorrelation time, i.e., \\eqb t<\\tdecorrel&=&{(\\lperp\\log\\Lambda)^2\\over2\\kpar} \\label{decorrtime} \\enspace. \\eqe Subsequently, the particle decorrelates and undergoes \\lq\\lq compound diffusion\\rq\\rq, which is the collisional transport regime discussed by Rechester \\& Rosenbluth~(\\cite{rechesterrosenbluth78}). Here the propagation is diffusive in character, with an effective diffusion coefficient given by a combination of macro- and microscopic effects: \\eqb \\kappa_{\\rm comp} &\\approx&\\kperp\\left(1+{\\Lambda^2\\over\\log\\Lambda}\\right) \\label{kcompound} \\enspace \\eqe The associated gaussian propagator is \\eqb P_{\\rm comp}(x,t)&=& {1\\over \\sqrt{4\\pi\\kappa_{\\rm comp}t}} \\exp\\left[-{x^2\\over 4\\kappa_{\\rm comp}t}\\right] \\label{compprop} \\enspace. \\eqe  Although we have derived the sub-diffusive propagator on the basis of quasi-linear macroscopic turbulence of the magnetic field, it is a phenomenon which probably has a much wider importance. This is demonstrated,  for example, by the numerical work of Rax \\& White~(\\cite{raxwhite92}), where sub-diffusion is found using a simple Taylor-Chirikov mapping to model the macroscopic field turbulence. Of course, the formulae we quote for quantities such as the decorrelation time (Eq.~\\ref{decorrtime}) are not generally applicable to other models of turbulence. Furthermore, they apply only in the absence of significant particle drifts and only if the magnetic field can be considered static. Both drifts and time-dependent fields may in fact be responsible for decorrelating the particle from a field line (Isichenko~\\cite{isichenko91b}) before this is achieved by either diffusion across the field or by the finite size of the gyro orbit. Nevertheless, even though it does not appear possible to model the macroscopic turbulence in an astrophysical source in such detail, we may nevertheless expect the qualitative picture of a sub-diffusive regime on short time scales followed by a diffusive one on longer time scales to be a generic feature. ",
        "conclusions": "\\label{discussion}   The main conclusion of this paper is that stochastic test particle acceleration at a strong shock which is predominantly perpendicular does not necessarily produce the canonical $n(p)\\propto p^{-2}$ spectrum expected from diffusive particle acceleration. Although much depends on the length scales and relative strengths of both the macroscopic inhomogeneities and the microscopic scattering, which, taken together, determine the decorrelation time $\\tdecorrel$, Eqs.~(\\ref{subindex}) and (\\ref{generals}) show that the spectrum produced depends in general on the properties of the turbulence, through the parameter $\\beta$ where $\\left<\\Delta x^2\\right>\\propto t^{\\beta}$.  We have made several simplifying assumptions in order to arrive at an analytic treatment of this problem. As in the case of diffusive acceleration, we have assumed that the presence of the shock front does not have an important effect on the spatial transport of the particles, i.e., that the solution of the particle transport problem at a shock front can be obtained from a consideration of the case in which an imaginary boundary moves through a uniform stationary medium. This is certainly an idealisation. In a realistic situation the plasma upstream and downstream of a shock front probably supports turbulence of different character and amplitude. Furthermore, the plasma velocity and the magnetic field strength are discontinuous. Thus, we should allow for some change in the parameters of the propagator on crossing the shock, and perhaps also for a change of its functional form. Nevertheless, we feel that the essential physical feature introduced by the anomalous transport process -- the steepening of the spectrum -- is captured in our simple approach. This is because the nature of the transport in the upstream medium has very little effect on the particle spectrum, provided all particles entering this region are returned to the shock. The slope of the spectrum is determined by competition between energy gain on crossing the shock and escape downstream. In the case of sub-diffusion, a particle is always tied to the same field line. If we consider a particle at the shock front moving into the downstream medium, it is evident that its escape probability is determined by the average distance it can travel along its field line before encountering the shock again and does not depend on the nature of that portion of the field line which lies upstream. The situation is different if we are interested  in the spatial dependence of the particle density upstream, or in the time taken to perform a cycle of crossing and recrossing, since then the nature of propagation in the upstream field is important.  Another of our assumptions -- that the magnetic field lines themselves diffuse -- is not readily lifted. Our calculation requires a specific model of the magnetic field, and we have too little knowledge of the properties of turbulence in the vicinity of shocks to do anything more realistic than simply assume diffusive behaviour. This corresponds to the case of sub-diffusion, $\\beta=1/2$. However, any situation in which the magnetic field plays a role in inhibiting particle transport is likely to resemble that of sub-diffusion, or, more generally, the case $\\beta<1$. In this connection, it is interesting to note that steep spectra have been found in numerical simulations of acceleration in random magnetic fields by Ballard \\& Heavens~(\\cite{ballardheavens92}). Although there is no well-defined average field direction in their computations, so that the shock cannot be described as quasi-perpendicular, the fact that particles diffuse freely in only one dimension, which does not always correspond with the direction of the shock normal, leads one to suspect that here too, anomalous transport may be an important factor. However, as well as the effects described above, their results may also be influenced by the relativistic flow speeds they assumed.  If particles are accelerated over several decades of momentum, as is thought to be the case in supernova remnants, it may be that different types of transport dominate in different ranges of $p$. One would expect that low energy particles which have relatively small gyro-radii compared to the correlation length of the magnetic field might be more closely tied to the magnetic field lines than high energy ones, and so suffer more from sub-diffusive type effects. We have not presented a detailed discussion of this effect here, which may, nevertheless, be important in the problem of the acceleration of cosmic rays (Duffy et al.~\\cite{duffyetal95}).  Finally, the relative steepening of the spectral index for $\\beta=1/2$ also has important implications for the acceleration of electrons in astrophysical sources of synchrotron emission. In particular, a correlation should arise between the obliquity of the shock and the spectral index of the radiation. At perpendicular shocks, sub-diffusion can lead to particle spectra given by Eq.~(\\ref{subindex}), which, even for uncooled electrons, would produce a relatively steep synchrotron spectrum of $F_{\\nu}\\propto \\nu^{-0.75}$.  \\noindent{\\bf Acknowledgements:} We are grateful to A.R.~Bell, R.O. Dendy and L.O'C.~Drury for helpful and stimulating discussions. This research was supported in part by the Commission of the European Communities under Contract ERBCHRXCT940604."
    },
    "9604/astro-ph9604038_arXiv.txt": {
        "abstract": "We present here the first results from two recently completed, fully sampled redshift surveys comprising 3703 IRAS Faint Source Survey (FSS) galaxies.  An unbiased counts-in-cells analysis finds a clustering strength in broad agreement with other recent redshift surveys and at odds with the standard cold dark matter model.  We combine our data with those from the QDOT and 1.2~Jy surveys, producing a single estimate of the IRAS galaxy clustering strength.  We compare the data with the power spectrum derived from a mixed dark matter universe. Direct comparison of the clustering strength seen in the IRAS samples with that seen in the APM-Stromlo survey suggests $b_O/b_I=1.20\\pm0.05$ assuming a linear, scale independent biasing. We also perform a cell by cell comparison of our FSS-$z$ sample with galaxies from the first CfA slice, testing the viability of a linear-biasing scheme linking the two.  We are able to rule out models in which the FSS-$z$ galaxies identically trace the CfA galaxies on scales 5-20$h^{-1}$Mpc.  On scales of 5 and 10$h^{-1}$Mpc no linear-biasing model can be found relating the two samples.  We argue that this result is expected since the CfA sample includes more elliptical galaxies which have different clustering properties from spirals.  On scales of 20$h^{-1}$Mpc no linear-biasing model with $b_O/b_I<1.70$ is acceptable. When comparing the FSS-$z$ galaxies to the CfA {\\em spirals}, however, the two populations trace the same structures within our uncertainties. ",
        "introduction": "Following the success of IRAS Point Source Catalog (PSC) galaxy redshift surveys in probing the large scale structure of the local Universe, which presented serious problems for the standard cold dark matter (CDM) theory \\cite{e90,snature,k91,moore,2jy_lss,fish_pow} two new fainter redshift surveys have been undertaken to confirm these results at significantly greater depths.  The first of these surveys (FSS-$z$~I\\footnote{These surveys have been referred to previously as `QCCOD'}) was based on samples drawn from the IRAS Faint Source Data Base [FSDB, see Moshir \\et~\\shortcite{fss}].  On the basis of deep IRAS coverage and freedom from cirrus contamination, three regions in the Northern Galactic Hemisphere were selected (these areas are defined in Table~\\ref{areas}). Within these areas we believe the FSDB to be $\\sim 99$ per cent complete for $S_{60}\\ge 0.2$~Jy (Lonsdale~\\et~in preparation). We thus constructed a catalogue of all FSDB sources within these areas having good to moderate $60\\mu$m fluxes greater than 0.2~Jy.  (In area E an additional sample was selected with $0.15\\le S_{60}<0.2$ but this is not considered further in this paper.)  The majority of such sources are galaxies [e.g. Rowan-Robinson~\\et~ \\shortcite{rr86}, Lawrence~\\et~\\shortcite{al86}].  No formal colour-cuts were employed to exclude either stars or cirrus sources.  The former were excluded upon inspection of the POSS plates or APM cartoons together with a case by case examination of the colours.  In the combined FSS samples only five `stars' lay close to the galaxy colour locus and they had the same colours as the $60\\mu$m excess stars. In addition 13 galaxies were excluded because a bright star lay in the field making identification and acquisition impossible, these are assumed to be random line of sight coincidences with little effect on our analysis. Cirrus sources were rare because of our careful selection of areas (four sources have been excluded upon examination of maps made from the raw IRAS data).  This survey was particularly designed to clarify the evolutionary behaviour of the IRAS population; see contrasting conclusions of Saunders \\et \\shortcite{s90} and Fisher \\et~\\shortcite{fish_evol}.    The second survey (FSS-$z$~II) was constructed from the Faint Source Catalog Version 2 \\cite{fss}. The area selected traversed the North Galactic Pole connecting the FSS-$z$~I areas. In other respects the FSS-$z$~II samples were selected in the same manner as FSS-$z$~I. Designed principally for large scale structure studies, the flux completeness in areas P and (to a lesser extent) X is not as good as other areas and is estimated to be 90 per cent at 0.25 Jy.  The combined surveys cover an area of 1310 deg$^2$ (0.4 sr) and contain 3728 sources, more than 3600 of which are galaxies.   \\begin{table}\\centering \\caption{Definition of FSS-$z$ areas.  Areas N, A and E constitute FSS-$z$~I while P and X constitute FSS-$z$~II.} \\label{areas} \\begin{tabular}{cccr} &           \\multicolumn{2}{c}{Definition}            & Area/\\\\ &                        &                            &sq deg\\\\ N  & $150\\deg\\le l\\le 210\\deg, $ & $ 50\\deg\\le b\\le 70\\deg$  & 594\\\\ A  & $ 30\\deg\\le l\\le  90\\deg, $ & $ 60\\deg\\le b\\le 68\\deg$  & 178\\\\ E  & $ 70\\deg\\le l\\le  90\\deg, $ & $ 50\\deg\\le b\\le 55.5\\deg$&  67\\\\ P  & $26.5\\deg\\le\\delta\\le44.5\\deg,$ & $ b\\ge 70\\deg$ if $\\alpha\\le 12^h$ & 476\\\\ && or $ b\\ge 68\\deg$ if $\\alpha>12^h$&\\\\ X  & $32.5\\deg\\le\\delta\\le38.5\\deg,$ & $b<50\\deg$ \\& $\\alpha\\ge8^h$         & 94\\\\  \\end{tabular} \\end{table}  Redshifts were obtained from the literature for 872 sources; a number of redshifts were kindly provided in advance of publication by: John Huchra, Ray Wolstencroft, Quentin Parker \\& Roger Clowes and Marc Davis \\& Michael Strauss. Major observation programmes were instigated to obtain the redshifts for the remaining sources using the FOS and FOS2 instruments on the INT and WHT facilities. Using automatic, optimal extraction techniques and line fitting procedures we obtained an average redshift accuracy of $\\sim 190\\kms$ \\cite[ and Oliver~\\et in preparation]{phd}, c.f. an accuracy of $\\sim 250 \\kms$ for the QDOT survey (Lawrence \\et, in preparation).  Of all 1931 FSS-$z$~I galaxies we currently have 1769 redshifts giving an overall redshift completeness of 91.6 per cent; while for the FSS-$z$~II project we have redshifts for 80.4 per cent of the galaxies. A mask has been constructed which, as well as defining the survey boundaries, excludes a number of sectors that have not been exhaustively followed up and those sources lying close to the IRAS coverage gap. Upon application of this mask the completeness statistics improve to 92.5 per cent and 86.8 per cent respectively. Above 0.25Jy (where FSS-$z$~II suffers less from flux incompleteness) we have 95 per cent of the FSS-$z$~I redshifts and 90 per cent for FSS-$z$~II. Detailed numbers are listed in Table~\\ref{comp}.  Of the failures the most interesting will be those with either very faint or no optical counterparts.  VLA maps have been obtained for all sources in the FSS-$z$~I that are either blank within the IRAS error ellipse (to the limit of the sky survey plates) or have very faint identifications.  $R$-band CCD images have also been obtained for many of these sources. These have yielded a number of further identifications and redshifts have been obtained for several of these. Work is continuing to obtain the redshifts for the remainder.    \\begin{table}\\centering \\caption{Redshift completeness statistics. The subsample has sources lying within our `mask' excluded.  Figures in parenthesis indicate the number of galaxies for which we have redshifts.} \\label{comp} \\begin{tabular}{crrrrrrr} & Total &\\multicolumn{2}{c}{galaxies} &\\multicolumn{2}{c}{subsample} &\\multicolumn{2}{c}{$S_{60}>0.25$Jy}\\\\ \\\\ N    & 1410 & 1369 & (1253) & 1369 & (1253) &  924 &  (874)\\\\ A    &  474 &  457 &  (413) &  383 &  (361) &  263 &  (251)\\\\ E    &  107 &  105 &  (103) &  105 &  (103) &   71 &   (70)\\\\ P    & 1416 & 1391 & (1133) & 1234 & (1062) &  921 &  (826)\\\\ X    &  321 &  312 &  (236) &  224 &  (203) &  146 &  (134)\\\\ \\hline & 3728 & 3634 & (3138) & 3315 & (2982) & 2325 & (2155)\\\\ \\end{tabular} \\end{table}   This paper quantifies the large scale clustering seen in both surveys and directly compares these samples with optical surveys.  These surveys are also being used for many other studies.  It was during the first of these projects that the unique object F10214+4724 was discovered \\cite{f10214}.  The completeness and reliability of the FSS-$z$~I survey will be detailed in future papers (Lonsdale \\et \\& McMahon \\et, in preparation).  Simple tests reveal significant evolution within this sample \\cite{hx94}.  This evolution will be elaborated in greater detail by Broadhurst \\et (in preparation). The appearance of the Bo\\\"{o}tes void which overlaps with areas A \\& E will be discussed by Oliver \\et (in preparation). The data for both surveys will be presented by Oliver~\\et~(in preparation). ",
        "conclusions": "We have completed two new large, deep IRAS redshift surveys.  The large scale clustering properties of these surveys have been analysed and found to be consistent with previous IRAS redshift surveys, strengthening previous conclusions that the IRAS galaxy distribution is inconsistent with the standard CDM predictions.  We noted two particular shortcomings in the counts-in-cells analysis. Firstly, the variances determined using different grid patterns on the same data show a scatter comparable with the quoted errors.  Secondly the method is not completely independent of the selection function and we demonstrate one way of overcoming this bias.  We are able to combine all the IRAS samples together to obtain a single estimate of the variance on large scales with errors that will not be matched until the completion of the full PSC-$z$ survey \\cite{pscz1,pscz2}.  Converting these variances into the form of a dimensionless power-spectrum (correcting for redshift-space distortions), we compare them with a mixed dark matter (MDM) model.  The slope of this model spectrum appears slightly steeper than that observed but the assumptions involved in transforming theory to data are sufficiently uncertain that we can not read too much into this.  A comparison of our variance estimates with those from the APM-Stromlo redshift survey suggests $b_O/b_I=1.20\\pm0.05$, assuming a linear-biasing model with ratio independent of scale. This is higher than that found by the APM-Stromlo team themselves \\cite{lday_cincell} because we find smaller IRAS variances than E90.  Using our FSS-$z$ sample and the first CfA strip we directly compare the distributions of optical and IRAS galaxies (using a physically meaningful definition of the two classes).  This is done on a point-by-point basis so we are able to test the hypothesis that the two populations are related by a linear-bias ratio ($b_O/b_I$).  We are able to reject the hypothesis that $b_O/b_I=1$ with a very high degree of significance on the scales 5, 10 and 20 $h^{-1}$ Mpc.  This hypothesis is, however, acceptable if we exclude ellipticals and S0 galaxies from the optical sample, suggesting that the discrepancies are due to the morphology density relation and the under-representation of ellipticals in IRAS samples.  On scales of 30 $h^{-1}$ Mpc our conclusions would be reversed but we suspect that our statistics are too poor at this scale for this to be very meaningful.  Allowing $b_O/b_I$ to vary away from 1 we find that only on the largest scales are {\\em any} optical/IRAS linear-biasing models acceptable, if elliptical galaxies are included in the optical samples.  Even on scales of 20 $h^{-1}$ Mpc, linear-biasing with $b_O/b_1<1.7$ is ruled out with more than 95 per cent confidence.  It should of course be stressed that in this particular volume the optical galaxy distribution is dominated by the Coma cluster and so the distorting effect of the elliptical galaxies has to some extent been amplified.  This non-linear clustering relation between ellipticals and IRAS galaxies may present complications for the interpretation of any galaxy surveys that include both elliptical and spiral galaxies.  As yet we have been unable to determine any differences in the clustering properties of IRAS galaxies and spiral galaxies and so surveys composed of either of these can be meaningfully compared.  \\vspace{1cm}  {\\bf ACKNOWLEDGMENTS}  SJO and ANT acknowledge PPARC support.  WS was supported by a PPARC Advanced Fellowship.  We thank John Huchra, Ray Wolstencroft, Quentin Parker \\& Roger Clowes and Marc Davis \\& Michael Strauss for providing redshifts in advance of publication.  We acknowledge use of the ING telescopes and thank the staff for their invaluable assistance.  Much use has been made of the STARLINK resources."
    },
    "9604/astro-ph9604104_arXiv.txt": {
        "abstract": "Four magnetic white dwarfs have been found in the course of the Hamburg/ESO Survey for bright QSOs. The objects have  been selected as QSO candidate on the basis of its blue continuum and the apparent absence of strong hydrogen or helium lines. One star, HE\\,1211-1707,  shows a rather fast spectral variability: both the strength and the position of the shallow absorption features  change on a time scale of 20\\,minutes. We  interpret this variability as being   due to a magnetic field on the surface of a rotating white dwarf, having a relatively uniform magnetic field on one hemisphere and a much larger spread of field strengths visible during other phases of the rotational period. All attempts to determine the magnetic field structure in detail with the help of synthetic spectra have failed so far because the star must have  a rather complicated field geometry. However, both the optical and the UV spectra indicate that a significant  part of the surface is dominated by  a magnetic field strength of  about 80\\,MG.  The spectrum of HE\\,0127-3110 is also rotationally modulated. This star (approximate range of magnetic fields: 85-345\\,MG) as well as HE\\,2201-2250 (a spectroscopic  twin of HE\\,0127-3110)  and HE\\,0000-3430 (43-118\\,MG) could be reasonably well reproduced with the help of theoretical spectra calculated assuming magnetic dipoles which are offset by 0.1 and 0.2 stellar radii along the magnetic axis. This result is in agreement with the assumption that  Ap stars, also showing significant deviations from a centered dipole, are the progenitors of magnetic white dwarfs. ",
        "introduction": "Degenerate stars with strong magnetic fields have been studied intensively in recent years (see Chanmugam 1992 for a review), ever since Greenstein et al. (1985) could show that in Grw$ +70^{\\circ}8247$, whose peculiar absorption bands (Minkowski, 1938) resisted for decades a successful identification, these bands are due to hydrogen in magnetic fields of several hundred MG. In the meantime, there has been considerable progress in numerical quantum mechanical calculations of the hydrogen levels in strong magnetic fields (R\\\"osner et al. 1984; Forster et al. 1984; Henry \\& O`Connell, 1984; Wunner et al., 1985), in applying these results to stellar atmospheres and line spectrum synthesis calculations (Achilleos and Wickramasinghe 1989, Putney and Jordan 1995 and references therein) and in discovering more single degenerate stars field strengths of several hundred Megagauss, cf. Table\\,2 in Schmidt \\&\\ Smith  (1995). Several of these stars like PG\\,1031+234 (Schmidt et al. 1986, Latter et. al. 1987) or PG\\,1015+014 (Wickramasinghe \\&\\ Cropper 1988) in addition exhibit strong rotational modulation of their spectral lines with periods as short as 1.65 hours. As relics of stellar cores, the study of magnetic fields in white dwarfs should shed important light on the role such fields play in stellar formation and main sequence evolution. In this paper we present spectra of four newly discovered magnetic white dwarfs. Three of the objects have been successfully modelled with offset dipole configurations. We also present time resolved spectra of HE\\,1211-1707, probably one of the fastest rotating white dwarfs with a rather complicated magnetic field configuration.  \\newpage  \\begin{figure}[htbp] \\fbox{{\\centering \\epsfxsize=8.5cm \\epsffile{he1211fc.eps}}}  \\vspace{1.2mm} \\fbox{{\\centering \\epsfxsize=8.5cm \\epsffile{he0127fc.eps}}} \\caption[]{Finding charts for HE\\,1211-1707 and HE\\,0127-3110 produced with the help of the Digitized Sky Survey } \\label{ffc1} \\end{figure}  \\begin{figure}[htbp] \\fbox{{\\centering \\epsfxsize=8.5cm \\epsffile{he2201fc.eps}}}  \\vspace{1.2mm} \\fbox{{\\centering \\epsfxsize=8.5cm \\epsffile{he0000fc.eps}}} \\caption[]{Finding charts for HE\\,2201-2250 and HE\\,0000-3430 } \\label{ffc2} \\end{figure} ",
        "conclusions": "Beside quasars, the Hamburg/ESO Survey has turned out to be a rich source of interesting blue stars. Together with HE1045-0908 the survey has enlarged the number of magnetic white dwarfs with field strengths above 10\\,MG by 25\\% (Schmidt \\&\\ Smith 1995).  With HE\\,1211-1707  we have discovered one of the  hottest magnetic white dwarfs showing a rather fast rotational modulation of the spectrum. Most likely, the star has a rather complicated structure of the magnetic field. The hemisphere that is seen during the rotational phase where the absorption features are strongest is probably dominated by a magnetic field of about 80\\,MG, while much higher field strengths are present on the rest of the star. Without better time resolved spectra (which are also necessary to determine the rotational period) and high S/N UV spectroscopy an unambiguous identification of the absorption features remains impossible at the moment.  Recently, Muslimov et al. (1995) performed theoretical calculations demonstrating that non-dipole components can survive much longer on  the white dwarf cooling sequence than previously believed. This is confirmed by our result that all four stars have a spread of magnetic field strengths larger than in the case of a centered dipole-field.  The spectra of HE\\,0127-3110, HE\\,2201-2250, and HE\\,0000-3430 have been successfully modelled with the help of synthetic spectra in the framework of offset dipoles geometries. Putney \\&\\ Jordan (1995) have also considered  dipole-quadrupole combinations as  a more physical model. However, for a detailed analysis which can distinguish between both these geometries it would be necessary to obtain measurements of the polarization first. The flux spectrum alone can be reproduced by  several slightly different offset-dipole models or dipole-quadrupole combinations.  The fact that the spectra of many magnetic white dwarf show evidence for non-centered magnetic dipoles strengthens the connection to Ap stars as their precursors since these objects also exhibit large deviations from the centered dipole structure.  Time resolved observations  of the rotating white dwarfs HE\\,1211-1707 and possibly HE\\,0127-3110 will allow to determine the field geometry over a large portion of the stellar surface. This will provide a  rather important test for  theoretical calculations of the magnetic field structure."
    },
    "9604/astro-ph9604042_arXiv.txt": {
        "abstract": "We present high resolution ($\\approx 8$ \\kms) spectra of the QSO Q0201+365 obtained with HIRES, the echelle spectrograph on the 10m W.M. Keck Telescope. Although we identify over $80\\%$ of the absorption features and analyze several of the more complex metal-line systems, we focus our analysis on the damped \\Lya system at $z=2.462$. Ionization simulations suggest the hydrogen in this system is significantly neutral and all of the observed metals are predominantly singly ionized. We measure accurate abundances for Fe, Cr, Si, Ni and place a lower limit on the abundance of Zn: [Fe/H] = $-0.830 \\pm 0.051$, [Cr/H] = $-0.902 \\pm 0.064$, [Si/H] = $-0.376 \\pm 0.052$, [Ni/H] = $-1.002 \\pm 0.054$ and [Zn/H] $> -0.562 \\pm 0.064$. We give evidence suggesting the actual Zn abundance is [Zn/H] $\\approx -0.262$, implying the highest metallicity observed at a redshift $z \\geq 2$.  The relative abundances of these elements remains constant over essentially the entire system ($\\approx 150$ \\kms in velocity space), suggesting it is well mixed. Furthermore, we use the lack of abundance variations to infer properties of the dust responsible for element depletion. Finally, we discuss the kinematic characteristics of this damped \\Lya system, comparing and contrasting it with other systems.  The low-ion line profiles span $\\approx 200$ \\kms in velocity space and have an asymmetric shape with the strongest feature on the red edge.  These kinematic characteristics are consistent with a rotating disk model. ",
        "introduction": "This paper is the second in a series devoted to studying the metal content of high-redshift galaxies and their progenitors. Our primary objectives are  \\indent (1) to record the emergence of metals in galaxies,  \\indent (2) to trace the mean cosmic metallicity from $z$ $\\approx$ 4.5 to the present,  \\indent (3) to determine the kinematic state of galaxies from $z$ $\\approx$ 4.5 to the present.  \\noindent We are implementing this study using HIRES, the echelle spectrograph on the Keck 10m telescope (\\cite{vgt92}), to obtain high-resolution spectra of QSOs with foreground damped {\\Lya} systems. The damped {\\Lya} systems are a population of neutral gas layers exhibiting properties indicating they are either galaxy progenitors, or well formed galaxies detected during an early evolutionary phase. Recent studies indicate the comoving density of neutral gas in damped systems at $z \\approx 3.3$ is comparable to the density of visible stars in current galaxies.  At lower redshifts, the comoving density of neutral gas decreases with time in a manner consistent with gas consumption by star formation (\\cite{wol95}). Therefore, studies of the metal content of the damped \\Lya systems enable one to trace the chemical evolution of representative galaxies from a presumably metal-poor gaseous progenitor phase to metal rich epochs when most of the baryons are in stars. As a result, the age-metallicity relation, kinematic conditions, etc., deduced from the damped \\Lya systems should tell us more about the history of galaxies at large redshifts than analogous relations deduced from old stars found in the solar neighborhood (\\cite{evd93}).  In a previous paper we presented echelle spectra of the $z$ = 2.309 damped system toward PHL 957 at a spectral resolution of $\\approx$ 8 {\\kms} (FWHM) and signal-to-noise ratio of $\\approx$ 35:1 (\\cite{wol94}). By fitting multiple Voigt velocity components to low-ion transitions such as Zn II 2026, Ni II 1741, and Cr II 2062 we obtained accurate abundances for Zn, Ni, and Cr in the neutral gas.  The Zn and Cr abundances were accurate because we resolved Zn II 2062.664 from Cr II 2062.234  for the first time in a QSO absorption system. We found that the abundances relative to solar were low: [Zn/H] = $-$1.55$\\pm$0.11, [Cr/H] = $-$1.79$\\pm$0.10, and [Ni/H] = $-$2.13$\\pm$0.08. The Zn abundance is especially significant because Zn is relatively undepleted by grains in the ISM of the Galaxy (\\cite{sem95}) and is presumably unaffected by dust which may be present in damped \\Lya systems (\\cite{fal93}). We also found the line profiles to be asymmetric in the sense that the low column density gas was found in absorption only at velocities higher than the high column-density gas. The kinematics can be explained by the passage of the line of sight through a rotating disk in which the density of clouds decreases with radius and with perpendicular distance from midplane.   The purpose of this  paper is to present HIRES spectra for Q0201+365, a $V$ = 17.5 QSO with emission redshift $z_{em}$ = 2.49. While we identify more than $80\\%$ of the absorption features and find a total of 13 metal-line redshift complexes, the focus of this paper is on the damped \\Lya system at $z$ = 2.462 (\\cite{lzwt91,lwt93}).  \\cite{lwt93} studied this system at a resolution of $\\approx$ 50 {\\kms}. They fitted a Voigt damping  profile to the \\Lya absorption trough and found $\\N{HI} = 2.4 \\sci{20} \\cm2$.  They also identified the metal transitions Si II 1190, Si II 1193, Si III 1206, Si II 1260, and Fe II 1144. Because these transitions are (a) saturated and (b) in the \\Lya forest, neither the abundances nor kinematics were accurately measured. The present study represents a major improvement over the previous work for the following reasons. First, we obtain spectra at a resolution of $\\approx 8$ {\\kms} and a typical signal-to-noise ratio of 33:1. Second, in contrast to the previous work, we focus on metal lines redward of \\Lya emission where confusion with \\Lya forest lines is absent. Third, because of the higher accuracy of the data we focus on weak unsaturated transitions of ions expected to dominate the ionization state of gas in neutral clouds.  In fact we establish accurate element abundances for Fe, Si, Ni, and Cr, and a lower limit on the abundance of Zn. Furthermore, we use computer simulations to investigate the ionization of this system. We also analyze the relative metal abundances and comment on the characteristics of dust grains in this system.  In addition, we examine a \\Lya absorption system at $z$ = 1.955 and use kinematic and abundance arguments to suggest it may be a damped \\Lya system.  Finally, we discuss the kinematics of the $z$=2.462 damped absorption system, contrasting its features with several other systems toward Q0201+365, as well as other damped \\Lya systems measured with HIRES (\\cite{wol94}).  The paper is organized as follows.  In $\\S$ 2 we describe the data acquisition and reduction techniques, present the spectra (Figure~\\ref{sptra}) and give a nearly complete absorption line list in Table~\\ref{orders}. We detail the analytic methods utilized throughout the paper in $\\S$ 3.   In $\\S$ 4 we present velocity profiles of the most significant metal line systems along with the VPFIT package solutions where applicable.  $\\S$ 5 presents the ionic column densities of the two systems associated with the damped \\Lya profile at $z$=2.46. In $\\S$ 6 we argue that the degree of photoionization of the damped \\Lya system is low. $\\S$ 7 gives the results of the abundance measurements and discusses the possible depletion of the gas-phase metals by dust grains.  We also describe the kinematics of the \\Lya system. Finally, $\\S$ 8 summarizes the results and gives concluding remarks. ",
        "conclusions": "This paper presented HIRES spectra obtained with the Keck 10m telescope of absorbing gas toward toward Q0201+365.  We identified over $80\\%$ of the absorption features and have analyzed several of the more interesting metal-line systems. We have focused on the damped \\Lya system at $z$=2.462 as part of an ongoing program to investigate the chemical content and kinematics of damped systems within the redshift interval $z \\approx 2 - 4$. We summarize our results as follows.  (1)  Based on the analysis of ionization simulations, we predict the damped \\Lya system to be significantly neutral.  Although it is possible the system is partially ionized, our analysis predicts the metals are all essentially in the singly ionized state, and that the total hydrogen column density is well within a factor of two of the adopted value from the $\\N{H^0}$ measurement.  (2) A hidden component analysis of the Ni II 1741, 1751 transitions did not reveal any significant hidden saturated components.  We expect this to hold true for the other low-ion transitions.  (3)  With the VPFIT least squares line profile fitting package we have measured ionic column densities for the damped \\Lya system at $z$=2.462 (as well as several other systems). We performed similar measurements with the apparent optical depth method and found the two results to be in agreement, further eliminating  the possibility of hidden line saturation.  (4)  We measured the following abundances of Si, Fe, Cr, and Ni for the damped \\Lya system: [Si/H] = $-0.376 \\pm 0.052$, [Fe/H] = $-0.830 \\pm 0.051$, [Cr/H] = $-0.902 \\pm 0.064$, and [Ni/H] = $-1.002 \\pm 0.054$. We placed limiting values on the abundances of the s-process elements Pb (3$\\sigma$ detection) and Ge (upper limit), [Pb/H] = $2.233 \\pm 0.121$ and [Ge/H] $<$ 0.664,  and a lower limit value on the abundance of Zn, [Zn/H] $> -0.562$.  Based on the VPFIT solution of the low-ions, we expect the metallicity is [Z/H] $\\approx -0.262$. This damped \\Lya system has the highest metallicity measured to date at $z \\geq 2.0$.  (5)  Comparing individual features of the damped \\Lya system, we find the relative abundance between Si, Fe, Cr, Ni and Zn remains nearly constant throughout our system (Figure~\\ref{depl}).  This suggests a well mixed system with an age large compared to the internal dynamical time scale at the epoch of detection.  (6)   We have used the relatively minor variations observed in the Si, Cr, and Fe abundances relative to Zn to place limits on the expected variation in the Hydrogen volume density throughout the damped \\Lya system, having assumed the presence of dust grains and the Jenkins relation (\\cite{jen87}).  Our measurements of [Cr/Zn] place a maximum variation of $n_H$ at $\\approx 1$ dex. The lack of $n_H$ (and [X/Zn]) variations could be evidence of weaker supernova input in the past, but we believe they are more likely due to the absence of grains with the properties of dust found in the ISM.  (7)  Plotting the measured abundances versus condensation temperature (Figure~\\ref{Tcond}), we do find evidence for a depletion pattern, but the overall depletion level of Si, Fe, Cr and Ni with respect to Zn is indicative of a relatively dust free ISM cloud.  Although gas with a lower dust-to-gas ratio than evident in the ISM can account for the pattern, one can also explain the pattern in terms of nucleosynthetic yields from type II supernovae (\\cite{lu95b}).  Both explanations are problematic and are under continued debate.  Determining the proper explanation is particularly important as they predict different metallicities for damped \\lya systems, which will significantly affect the investigation of galactic chemical evolution.  (8) The low-ion profiles of the damped system exhibit an edge leading asymmetry as predicted by a simple model of rotation.  The shape is similar to the other damped system observed with HIRES (PHL 957; \\cite{wol94}), though the velocity interval is significantly greater."
    },
    "9604/gr-qc9604029_arXiv.txt": {
        "abstract": "Several applications of spectral methods to problems related to the relativistic astrophysics of compact objects are presented. Based on a proper definition of the analytical properties of regular tensorial functions we have developed a spectral method in a general sphericallike coordinate system. The applications include the investigation of spherically symmetric neutron star collapse as well as the solution of the coupled 2D\\slsh Einstein\\slsh Maxwell equations for magnetized, rapidly rotating neutron stars. In both cases the resulting codes are efficient and give results typically several orders of magnitude more accurate than equivalent codes based on finite difference schemes. We further report the current status of a 3D\\slsh code aiming at the simulation of non\\slsh axisymmetric neutron star collapse where we have chosen a tensor based numerical scheme. ",
        "introduction": "Compact objects in astrophysics such as neutron stars and black holes are subjected to the strong field regime of gravitation and have hence to be treated within the framework of general relativity. The growing interest in the numerical solution of the Einstein equations for astrophysically relevant systems has given rise to a new branch of computational physics --- {\\em numerical relativity\\/} \\cite{SM79,HB89,IN92}. This development is due to the increasingly powerful computational resources which make these problems accessible to a numerical investigation. It is further stimulated by the prospects of gravitational wave astronomy which will turn into an observational science toward the end of this decade thanks to gravitational wave observatories like LIGO, VIRGO and GEO600 that are now under construction \\cite{AB92,BR90,HD94}.\\nspc We use the (3+1)\\slsh formalism of general relativity \\cite{SY78-2} which consists in foliating spacetime into a sequence of spacelike hypersurfaces which represent curved three\\slsh space at a fixed coordinate time $t$. The fabric of spacetime is then determined by the three\\slsh metric $h_{ij}$ and four additional quantities, the {\\em lapse function} $N$ and the {\\em shift vector} $N^i$ which fix the propagation of the spacelike hypersurfaces in time and the change of the spatial coordinate system between adjacent hypersurfaces. This Hamilton type approach to general relativity results in a temporal first order evolution scheme for the dynamical variables which is completed by some constraint equations which ensure the consistency of gravitational and matter fields. Furthermore $N$ and $N^i$ have to be determined by the choice of appropiate gauge conditions which typically lead to elliptic equations that have to be solved at each time step. For stationary configurations all time derivatives vanish and one obtains a system of coupled elliptic equations for the gravitational fields. The efficient solution of elliptic equations is hence of central interest for us.\\nspc Let us consider a covariant Poisson equation $N^{|i}{}_{|i}\\!=\\!S$ in a conformally flat axisymmetric space where the line element reads \\bgeq dl^2\\!=\\!A^4(r,\\theta)\\,(dr^2\\!+\\!r^2\\,d\\theta^2\\!+\\!r^2\\sin^2\\!\\theta\\, d\\phi^2).\\edeq The former equa\\-tion can be rewritten to yield a Poisson\\slsh like equation for $N$ where we have isolated the flat space Laplacian $\\Delta_f$ and contributed the curvature terms to the source. Here $\\alpha$ denotes $\\ln A$. \\bgeq \\quad\\;\\Delta_f N=\\tilde{S}\\;\\;\\mbox{with}\\;\\;\\tilde{S}=A^4 S-2\\, (\\partial_r\\alpha\\partial_r N+\\frac{1}{r^2}\\,\\partial_\\theta\\alpha \\partial_\\theta N). \\edeq This equation has to be solved by iteration. The solution of $\\Delta_f N\\!=\\! \\tilde{S}$ at each iteration has hence to be accomplished sufficiently fast in order to keep the total computation cost at a reasonable level.\\nspc After outlining the basic features of our spectral method \\cite{BM86,BM90}, we will proceed in a first step to the investigation of black hole formation due to spherically symmetric neutron star collapse which has proved the high aptitude of spectral methods in this field \\cite{GO91,GH93,GHG95}. The second part is devoted to the study of axisymmetric stationary rotating bodies which has been applied to model rapidly rotating neutron stars \\cite{BGSM93,SBGH94}. This work has been extended recently to include strong magnetic fields for the first time into neutron star models \\cite{BBGN95}. Special emphasis in all cases has been put on the extensive use of external and intrinsic tests \\cite{BO73,BG94,GB94} of the self\\slsh consistency and the attained accuracy of the numerical results. The resulting neutron star models provide us with the required initial value models for the investigation of 3D\\slsh gravitational collapse of neutron stars which will reveal the whole range of gravitational wave emission associated with this phenomenon. We give an overview about the inset of spectral methods in this project which is currently in work. Here a new method for the efficient inversion of a generalized 3D\\slsh vector Poisson equation is a first major result. ",
        "conclusions": "We have presented the application of spectral methods to several problems of numerical relativity. In each case they proved to be a highly valuable tool which lead to results typically several orders of magnitude more accurate than corresponding codes based on finite difference schemes. Especially in sphericallike coordinates the advantages of a spectral method which allows a rigorous treatment of the associated regularity conditions, while improving the efficiency of the code at the same time, are remarkable. Particularly important properties for our problems are the negligible numerical viscosity in temporal evolution schemes which enabled us to capture subtle details in the time\\slsh dependence of evolved variables as observed for equilibrium configurations of neutron stars in Sec.~\\ref{SUBSEC:NSR}, as well as the very natural treatment of boundary conditions and the efficient solution of elliptic equations which is a frequently encountered task in our investigations. Our so far very positive experiences with spectral methods give us confidence to dispose of the appropriate numerical tool to tackle the exciting problem of black hole formation by 3D\\slsh gravitational collapse of neutron stars. \\\\[7.5ex] {\\em Acknowledgements.} J. Frieben gratefully acknowledges financial support by the {\\sc Gottlieb Daimler\\slsh und Karl Benz\\slsh Stiftung}."
    },
    "9604/hep-ph9604229_arXiv.txt": {
        "abstract": "\\noindent The cosmological baryon asymmetry can be explained by the nonperturbative electroweak reprocessing of a lepton asymmetry generated in the out-of-equilibrium decay of heavy right-handed \\mbox{Majorana} neutrinos. We analyze this mechanism in detail in the framework of a SO$(10)$-subgroup. We take three right-handed neutrinos into account and discuss physical neutrino mass matrices. ",
        "introduction": "One of the most striking features of the observable universe is the baryon asymmetry, which is usually expressed as the ratio of the baryon density $n_B$ to the entropy density $s$. From measurements of the abundances of the light elements one finds\\footnote{For a review and references, see \\cite{kt1}}: \\beq Y_B={n_B\\over s}=(0.6-1)\\times10^{-10}\\;. \\eeq With the appearance of grand unified theories (GUTs) it became possible to explain this asymmetry by the baryon number $(B)$ violating decays of Higgs or gauge bosons at the GUT scale. However, these models of baryogenesis are not easily reconciled with inflation. Indeed, a baryon asymmetry present before inflation would be diluted by a huge factor, while the reheating temperature after the inflationary phase is in general too low for these baryogenesis mechanisms to work.  Then, it was realized that anomalous baryon number violating processes are unsuppressed at high temperatures \\cite{sphal2}. These so called sphaleron transitions violate $(B+L)$ and conserve $(B-L)$, where $L$ is the lepton number. As sphalerons are in thermal equilibrium for temperatures between $\\sim\\!\\!10^{12}\\,$GeV and $\\sim\\!\\!10^2\\,$GeV, they will strongly modify any primordial $(B+L)$ asymmetry. The connection between the baryon asymmetry and a primordial $(B-L)$ asymmetry is given by \\cite{sphal} \\beq Y_B=\\left({8N_f+4N_H\\over22N_f+13N_H}\\right)Y_{B-L}\\;, \\eeq where $N_f$ is the number of fermion families and $N_H$ is the number of Higgs doublets.  The needed primordial $(B-L)$ asymmetry can be realized as a lepton asymmetry generated by the out-of-equilibrium decay of heavy right-handed Majorana neutrinos, as suggested by Fukugita and Yanagida \\cite{fy2}.  $L$ is violated by Majorana masses, while the necessary $CP$ violation comes about through phases in the Dirac mass matrix of the neutrinos. In a detailed quantitative analysis Luty showed that the scenario works for a wide range of parameters \\cite{luty}.  In order to generate a lepton asymmetry of the correct order of magnitude, the right-handed neutrinos have to be numerous before decaying. This is only possible if they are in thermal equilibrium at high temperatures. In the original model \\cite{fy2} the right-handed neutrinos were only interacting through Yukawa couplings, which are far too weak to bring the neutrinos into equilibrium at high temperatures.  Hence one had to assume an equilibrium distribution for the neutrinos as initial condition in previous analyses.  An appealing way to solve this problem is to study this baryogenesis mechanism in the framework of an extended gauge symmetry, since right-handed neutrinos appear naturally in unified theories based on the gauge groups SO$(10)$ or E$_6$. As we shall see, the gauge interactions in which the right-handed neutrinos take part, are strong enough to bring them into thermal equilibrium at high temperatures. Of course, the neutrinos have to be out of equilibrium when decaying, i.e.\\ the reaction rates for the gauge interactions have to fall fast enough, so that they cannot significantly reduce the number density of the neutrinos before they decay.  In this paper we investigate this mechanism in the framework of an SO$(10)$ subgroup. After a short discussion of the relevant Boltzmann equations in the next section, we present our model and calculate the needed reaction rates in section \\ref{model}. In section \\ref{solutions} we solve the Boltzmann equations, first for one and then several heavy neutrino families. We explicitly show that the lepton asymmetry is mainly determined by the lightest of the right-handed neutrinos. Finally we will look at physical mass matrices for the neutrinos coming from an additional abelian gauged family symmetry. These mass matrices give a lepton asymmetry in the right order of magnitude and predict light neutrino masses and mixings of the magnitude needed to explain the solar neutrino deficit and a $\\t$-neutrino mass of a few eV which is needed in the cold-plus-hot dark matter models. ",
        "conclusions": "We have seen that the cosmic baryon asymmetry can be explained by the lepton number violating decays of heavy Majorana neutrinos combined with the anomalous electroweak $(B+L)$ violation. The lepton asymmetry generated in our model is independent of the initial conditions on the heavy neutrino density, which was not the case in a previous analysis of this mechanism \\cite{luty}. This is a consequence of the new gauge interaction which we have introduced, and which is related to the spontaneous breaking of lepton number.  We have also considered the decay of more than one heavy neutrino and have seen that the generated asymmetry is determined by the properties of the lightest heavy neutrino, if the right-handed neutrinos have a pronounced mass hierarchy.  By performing explicit calculations we could show that the generation of a lepton asymmetry is possible at every temperature between the intermediate breaking scale and the electroweak breaking scale. Furthermore we have checked that neutrino mass matrices may explain the generation of the asymmetry and low energy neutrino phenomenology at the same time.  In supersymmetric models of inflation the reheating temperature has to be lower than $10^5$ to $10^8\\,$GeV to solve the gravitino problem \\cite{sarkar,camp}. Therefore in these models baryogenesis has to take place at relatively low temperatures. Since this is possible in our model, a supersymmetric generalization should be viable. \\vspace{1cm}\\mbox{ }\\\\ \\noindent \\setlength{\\parskip}{1ex} {\\bf\\Large Acknowledgments}  \\mbox{ }\\\\\\noindent I would like to thank W.~Buchm\\\"uller, who suggested this investigation, for continuous support and encouragement."
    },
    "9604/astro-ph9604167_arXiv.txt": {
        "abstract": "The Las Campanas Redshift Survey (LCRS) consists of 26418 redshifts of galaxies selected from a CCD-based catalog obtained in the $R$ band.  The survey covers over 700 square degrees in 6 strips, each 1.5$\\arcdeg$ x 80$\\arcdeg$, three each in the North and South galactic caps.  The median redshift in the survey is about 30 000 km~s$^{-1}$.  Essential features of the galaxy selection and redshift measurement methods are described and tabulated here.  These details are important for subsequent analysis of the LCRS data.  Two dimensional representations of the redshift distributions reveal many repetitions of voids, on the scale of about 5000 km~s$^{-1}$, sharply bounded by large walls of galaxies as seen in nearby surveys.  Statistical investigations of the mean galaxy properties and of clustering on the large scale are reported elsewhere.  These include studies of the luminosity function, power spectrum in two and three dimensions, correlation function, pairwise velocity distribution, identification of large scale structures, and a group catalog. The LCRS redshift catalog will be made available to interested investigators at an internet web site and in archival form as an Astrophysical Journal CD-ROM. ",
        "introduction": "Redshift surveys reveal surprising structures in the large-scale distribution of galaxies, which provide clues to physical properties of the Universe (see \\cite{gio91} and \\cite{strwil95} for reviews.) Early investigations (such as \\cite{KOSSI}, or \\cite{huchra83}) suggested that the nearby Universe might be strongly inhomogeneous with nearly empty voids and thin, high contrast regions of galaxy overdensity. Convincing demonstrations that this pattern is a general feature of the Universe have come from extensive surveys which cover large angles on the sky, (\\cite{deL86}, \\cite{gel89}:summarizing the CfA surveys, and \\cite{dacosta94a}: the SSRS2 survey.) Galaxies appear to lie on networks of filaments or sheets extending over $100 \\ h^{-1}$~Mpc that encompass sharply bounded voids (Hubble constant $H_0 = 100 \\ h$~km~s$^{-1}$~Mpc$^{-1}$). Despite the power of the CfA2 (\\cite{huchra95}) and SSRS2 surveys in conveying an impression of the galaxy distribution on large scales, the largest features they reveal are very near the upper limit set by the depth of the survey at about 12000~km~s$^{-1}$.  Are these the largest features in the Universe?   There are hints of possible structure on larger scales from apparent periodicities in redshifts (\\cite{broadhurst}), from the flows suggested by Lauer \\& Postman (1994), or from possible variations in the luminosity density (\\cite{dacosta94a}). Only a deep extensive survey can provide evidence on the reality of these suggestions.  Similarly, the statistical measures of galaxy clustering derived from redshift catalogs (\\cite{Efstathiou}, \\cite{park})\\ which can be used to constrain models for the formation of structure in the Universe, have only modest precision at the large scales which are most interesting. To explore structure in the local Universe on a scale of 30000~km~s$^{-1}$ and to improve our measures of galaxy statistics, we have been working since 1987 on the Las Campanas Redshift Survey (LCRS).  We have compiled a CCD-based galaxy catalog, selected the galaxies, developed the fiber-optic equipment for mass producing galaxy spectra, and observed from 1988 to 1994.  We have measured more than 26000 galaxy redshifts averaging $z = 0.1$ over an area of 700 square degrees arranged in six long thin strips, three in the North Galactic Cap, and three in the South.  The Las Campanas Redshift Survey provides a reconaissance of present-day structure on the largest scales mapped to date.  The Las Campanas Redshift Survey is a direct descendent of earlier surveys aimed at measuring the average properties of galaxies: the luminosity function and the space density of galaxies (\\cite{KOSI}; \\cite{KOSII}; \\cite{KOSSII})  Annoying variations in the average properties, such as the luminosity density, signalled the presence of strong inhomogeneities in the galaxy distribution. Most startling was a large void of diameter $60 \\ h^{-1}$~Mpc in the galaxy distribution beyond the constellation Bo\\\"{o}tes (\\cite{KOSSI}; \\cite{KOSSIII}; \\cite{KOSSIV}). Because similar structures are now frequently seen in wider-angle surveys and because the structures are nearly as large as the survey dimensions, a much larger and deeper survey seems needed in order to encompass a fair sample of the universe, as well as to search for structure on larger scales. The LCRS began with a goal of obtaining a galaxy sample of sufficient size, sky coverage, and depth to permit a reliable characterization of the average properties of galaxies and of their distribution.  While this work was in progress, some important large wide-angle redshift surveys have been carried out. Recent large-scale structure analyses include (see \\cite{strwil95} for a more complete list): (1) nearly all-sky samples of objects selected from the IRAS Point Source Catalog, specifically the IRAS 1.2 Jy survey (5321 galaxies with 60~$\\mu$m flux $f_{60} > 1.2$~Jy; \\cite{fisher95}) and the 1-in-6 sparse-sampled IRAS QDOT survey (2184 galaxies with $f_{60} > 0.6$~Jy; \\cite{law95}); (2) the Stromlo-APM survey, selected from the optical APM galaxy catalog, covering $4300$ square degrees of the Southern sky to a median survey depth $cz = 15000$~km~s$^{-1}$, but 1 in 20 sparse sampled (1787 galaxies, $b_J \\leq 17.15$; \\cite{1:lov92}; \\cite{1:lov95}; \\cite{lov96}); and (3) the combined CfA2 and Southern Sky (SSRS2) redshift surveys, selected from the Zwicky catalog and from plate scans, respectively, covering one-third of the sky over the North and South galactic caps out to a mean depth of about $7500$~km~s$^{-1}$ (14383 galaxies with $m_{B(0)} \\leq 15.5$; \\cite{huchra95}, \\cite{dacosta94a}, 1994b). The CfA2+SSRS2 sample represents the state of the art for a wide-angle survey, with a large sample size and dense sky coverage.  Sparse sampling provides an efficient way to measure a particular statistic, while dense sampling reveals structures that require the coherent arrangement of many galaxies and which are not easily characterized by low-order statistics.  Our survey has some of the desirable properties of each: the samples are strips on the sky 1.5$\\arcdeg$ wide and  80$\\arcdeg$ long which are separated by 3$\\arcdeg$. Within each strip, the galaxies are sampled randomly from a magnitude-limited catalog, but the sampling is quite dense, averaging about 70\\% of the magnitude-limited list.  The distinctive properties of our survey are its depth (out to 60000 km~s$^{-1}$) and the number of redshifts (26000).  The Las Campanas Redshift Survey provides improvements in sample size, volume, and depth in a limited region of the sky. Unlike most previously completed surveys, where spectra were taken of individual galaxies, one at a time, the LCRS employs a fiber-optic spectrograph system on the Las Campanas 2.5 meter DuPont telescope to observe over 100 objects at once. The greater depths explored by the LCRS mean that we observe galaxies which have sufficient surface density on the sky to make efficient use of a multi-fiber system.  Some details of the galaxy selection and observing procedure which add to the complexity of the analysis were shaped by the properties of the fiber system, but the overall gain for exploring large-scale structure is very large. In this sense the LCRS represents part of a fundamental technological change in the way that large-scale redshift surveys are carried out.  Other surveys in progress, such as the Century Survey (\\cite{century}) and the ESO-based survey (\\cite{eso}), have a comparable depth but cover smaller areas. Each has selection criteria that differ from ours in interesting ways which should make the comparison of our results with theirs a fruitful enterprise. The next generation of surveys, such as the 2 Degree Field project (\\cite{ellis93}) and the the Sloan Digital Sky Survey (\\cite{gunwein95}), will use the multiplex advantage of fiber-fed spectrographs to obtain redshifts of up to a million galaxies, to about the same depth ($z \\approx 0.1$) explored by the LCRS.  It is our hope to illuminate some of the most interesting subjects for more thorough investigation by these ambitious projects.  In addition to providing a map of the galaxy distribution on the largest scales, the LCRS can help to constrain the physical history of the Universe and the properties of its constituents. Current theories of structure formation start with a spectrum of fluctuations produced in the early universe, follow the growth of structure due to gravitation, and strive to match both the subtle structure observed at recombination and the high-contrast distribution of the galaxies today mapped by redshift data. Most models presume that the Universe is dominated by non-baryonic dark matter (\\cite{blum84}) whose properties are to be inferred from the match to the data. The luminous galaxies trace the underlying mass, but may be biased:  galaxies may form preferentially in peaks of the matter distribution (\\cite{bard86}, \\cite{white87}). Such models, with the aid of numerical N-body simulations, predict the clustering properties of galaxies from small nonlinear scales of less than $10 \\ h^{-1}$~Mpc, all the way up past the largest scales sampled by existing redshift surveys. The scales sampled with the LCRS, up to about $300 \\ h^{-1}$~Mpc, correspond to the horizon size at interesting stages in the early universe and are closer to the scales probed by microwave backgound observations than can be measured by smaller surveys. By observing galaxy fluctuations on these scales with LCRS, we can learn about primordial density fluctuations and the processes by which they grow.  Combining evidence from large-scale structures and from microwave background observations holds the potential to answer fundamental questions about the dark matter that makes up most of the universe (\\cite{scott95}).  This paper gives details about the construction of the LCRS which are essential to any analysis of the survey data.  Some preliminary results and descriptions of the survey have already been published (\\cite {she92}, \\cite {oem93}, \\cite {mapscale}, \\cite {moriond}, \\cite {she95}, \\cite {tuc95}.) Tucker's (1994) Ph.D. thesis at Yale discusses the construction of the survey's photometric catalog and statistical analyses of the early 50-fiber sample.  Tucker describes the galaxy correlation function, the properties of galaxy groups, and galaxy clustering as a function of galaxy color.  A more complete analysis, using the entire LCRS data set to investigate these topics, will be carried out by Tucker (Tucker et al.\\ 1996c,a,b, respectively). Huan Lin's Ph.D. thesis at Harvard (\\cite {thesis}) is the basis for the present paper and for several forthcoming papers on the LCRS. Lin et al. (1996a) details the derivation of the luminosity function and space density of LCRS galaxies. Lin et al. (1996b) deals with the power spectrum, a basic statistic which characterizes the clustering properties of galaxies.  Lin et al. (1996c) examines redshift-space distortions in the clustering of LCRS galaxies, derives the velocity dispersion of galaxy pairs, and estimates the value of the cosmological density parameter $\\Omega$. The analysis of Landy et al. (1996) uses the tools of 2D Fourier analysis to search for coherent structures on very large scales in the LCRS data.  Doroshkevich et al. (1996) examines the typical scales and types of large scale structure in the LCRS sample. Our intention is to make results from this survey and the survey data itself available to interested investigators.  We have established a home page for the LCRS at ``http://manaslu.astro.utoronto.ca/\\~{ }lin/lcrs.html'' which will provide rapid access to the data.  An archival table of the redshift data will accompany this paper as an Astrophysical Journal CD-ROM. ",
        "conclusions": ""
    },
    "9604/astro-ph9604173_arXiv.txt": {
        "abstract": "We have observed near--IR H$_2$ line emission on large scales in the Galactic center. Paper~I discussed our 400~pc long strip map and 50~pc map of the H$_2$ $v=1\\rightarrow 0\\ S(1)$ line. In this paper, we present observations of the higher vibrational lines (H$_2$ $v=2\\rightarrow 1\\ S(1)$ and $v=3\\rightarrow 2\\ S(3)$) at selected positions and conclude that strong far--UV radiations excites the H$_2$. We compare the H$_2$ $v=1\\rightarrow 0\\ S(1)$ emission to far--IR continuum emission and show that the ratio of these two quantities in the Galactic center equals the ratio seen in the starburst galaxies, M82 and NGC~253, and in ultraluminous infrared bright galaxies. ",
        "introduction": "The central kpc regions in starburst galaxies and ultraluminous IR bright galaxies are powerful emitters of near-IR H$_2$ emission (Puxley, Hawarden, \\& Mountain 1990; Goldader et al. 1995). Ro--vibrational lines of H$_2$ can trace both photon--dominated regions (PDRs), where far--UV photons excite the H$_2$, and shocked regions, where the H$_2$ is thermally excited. Vigorous star formation in these galaxies produces large numbers of UV photons which fluorescently excite H$_2$, while subsequent supernovae shock--excite the H$_2$.  We have used the University of Texas near--IR Fabry--Perot Spectrometer, to survey giant molecular clouds (GMCs) on $1-10$~pc scales (Luhman et at. 1994; Luhman \\& Jaffe 1996; Luhman et al. 1996). In Orion~A, for example, the H$_2$ $v=1\\rightarrow 0\\ S(1)$ line emission extends up to 8~pc (1\\deg) from the central UV source, $\\theta^1$~Ori~C. The detection of higher vibrational state H$_2$ lines, e.g., $v=6 \\rightarrow 4\\ Q(1)$ and $v=2 \\rightarrow 1\\ S(1)$, showed that far--UV photons excite the H$_2$. Although the shock--excited H$_2$ emission is intense in the Orion $BN-KL$ region, the emission region is relatively compact ($\\sim 1\\arcmin$). The total H$_2$ luminosity in the $BN-KL$ region is only $\\sim 1 \\%$ of the Orion PDR H$_2$ luminosity. Similarly, UV--excited H$_2$ dominates the large--scale H$_2$ emission from other GMCs.  We have observed the H$_2$ emission in the inner $\\sim 400$ pc ($\\sim 3\\deg$) of our Galaxy in order to investigate H$_2$ emission on a more global scale and to compare the Galactic center with central $\\sim 1$ kpc regions in external galaxies. The physical conditions in the interstellar medium of the Galactic center are significantly different from those in the solar neighborhood. The thin disk (diameter of 450~pc, height of 40~pc) of dense interstellar material in the Galactic center contains M(H$_2$) $>$ $2 \\times 10^7 M_{\\sun}$ (G\\\"{u}sten 1989; Hasegawa et al. 1996). The molecular clouds in the Galactic center have higher density, higher metallicity, and higher internal velocity dispersion than the clouds in the solar neighborhood (Blitz et al. 1993). There is strong radio continuum radiation from giant H~{\\small II} regions (Sgr~A, Sgr~B, Sgr~C, and Sgr~D) and extended low--density (ELD) ionized gas. The spectral index in the areas away from the discrete H~{\\small II} regions shows that thermal bremsstrahlung from ionized gas can account for about half of the emission from the extended gas (Sofue 1985). Another indicator of the intense UV radiation in the central 400~pc is strong far--IR continuum emission (Odenwald \\& Fazio 1984). About 90\\% of the far--UV energy is absorbed by dust and reradiated in the far--IR. From the far--IR intensity, we estimate that the far--UV radiation field is $\\sim 10^3$ times the value in the solar neighborhood ($I_\\circ = 4 \\times 10^{-4}$ ergs~s$^{-1}$~cm~$^{-2}$~sr$^{-1}$, Draine 1978). The energetic conditions in the Galactic center mean that the center can provide a unique view of the interaction between stellar UV radiation and molecular clouds, and serve as a nearby model for the nuclei of galaxies.  \\begin{figure}[t] \\vbox to 7.9cm { \\plotfiddle{fig1.ps} {7.9cm} {0} {40} {40} {-160} {0} } \\caption{ \\label{fig-1} Observed intensity distribution of H$_2$ $v=1\\rightarrow 0\\ S(1)$ ($\\lambda = 2.121\\ \\micron$) along the Galactic plane at $b = -0\\fdg05$. The open circles were taken at the McDonald 0.9 m telescope with a 3\\farcm3 beam (Paper I) and the filled circles at the CTIO 1.5 m telescope with a 1\\farcm35 beam. The intensities have not been corrected for interstellar extinction. The error bars represent $1\\sigma$ measurement uncertainties. } \\end{figure}  In paper I (Pak, Jaffe, \\& Keller 1996) we showed the distribution of H$_2$ $v=1\\rightarrow 0\\ S(1)$ emission along a 400 pc--long strip and in the inner 50~pc of the Galactic center. We detected H$_2$ emission throughout the surveyed region. The typical dereddened ($A_K = 2.5$ mag) H$_2$ $v=1\\rightarrow 0\\ S(1)$ intensity, $\\sim 3 \\times 10^{-5}$ ergs~s$^{-1}$~sr$^{-1}$, is similar to the surface brightness in Galactic PDRs (Luhman \\& Jaffe 1996). In this Paper, we present observations of several H$_2$ lines, discuss the excitation mechanism, and compare the Galactic center observations to observations of other galaxies. ",
        "conclusions": ""
    },
    "9604/astro-ph9604035_arXiv.txt": {
        "abstract": "s{This paper gives an overview of the properties of all possible equilibrium sequences of compact strange-matter stars with nuclear crusts, which range from strange stars to strange dwarfs. In contrast to their non-strange counterparts, --neutron stars and white dwarfs--, their properties are determined by two (rather than one) parameters, the central star density and the density at the base of the nuclear crust. This leads to stellar strange-matter configurations whose properties are much more complex than those of the conventional sequence.  As an example, two generically different categories of stable strange dwarfs are found, which could be the observed white dwarfs.  Furthermore we find very low-mass strange stellar objects, with masses as small as those of Jupiter or even lighter planets.  Such objects, if abundant enough in our Galaxy, should be seen by the presently performed gravitational microlensing searches. Further aspects studied in this paper concern the limiting rotational periods and the cooling behavior of neutron stars and their strange counterparts.} ",
        "introduction": "The theoretical possibility that strange quark matter may be absolutely stable with respect to iron (energy per baryon below 930 MeV) has been pointed out by Bodmer\\,\\cite{bodmer71:a}, Witten\\,\\cite{witten84:a}, and Terazawa\\,\\cite{terazawa89:a}.  This so-called strange matter hypothesis constitutes one of the most startling possibilities of the behavior of superdense nuclear matter, which, if true, would have implications of greatest importance for cosmology, the early universe, its \\begin{figure}[tb] \\begin{center} \\leavevmode \\psfig{figure=qlcomp.bb,width=4.5in,height=3.1in} \\caption[Relative densities of quarks and leptons in cold, beta-stable, electrically charge neutral quark matter versus density.]{Relative densities of quarks and leptons in absolutely stable strange-quark-star matter versus density. $n_i$ and $n$ denote partial and total densities, respectively.} \\label{fig:1.5} \\end{center} \\end{figure} evolution to the present day, astrophysical compact objects, and laboratory physics\\,\\cite{aarhus91:proc}.  Unfortunately it seems unlikely that QCD calculations will be accurate enough in the foreseeable future to give a definitive prediction on the absolute stability of strange matter, such that one is left with experiments and astrophysical tests, as performed here, to either confirm or reject the hypothesis.  One striking implication of the hypothesis would be that pulsars, which are conventionally interpreted as rotating neutron stars, almost certainly would be rotating strange stars (strange pulsars). Part of this paper deals with an investigation of the properties of such objects.  In \\begin{figure}[tb] \\begin{center} \\leavevmode \\psfig{figure=S95_1-.ps.bb,width=10.0cm,height=9.0cm,angle=90} \\caption[Neutron star--white dwarf sequences and their strange counterparts]{Neutron star (NS)--white dwarf (wd) sequence (solid line). The dotted and dashed curves refer to strange star (SS)--strange dwarf (sd) sequences with inner crust densities as indicated (in $\\gcmt$).  Vertical bars mark minimum mass stars, crosses mark the termination points of the strange star sequences.} \\label{fig:sequence} \\end{center} \\end{figure} addition to this, we develop the complete sequence of strange stars with nuclear crusts, which ranges from the compact members, with properties similar to those of neutron stars, to white-dwarf-like objects (strange dwarfs) and discuss their stability against acoustical vibrations\\,\\cite{weber93:b}.  The properties with respect to which strange-matter stars differ from their non-strange counterparts are discussed.  \\goodbreak ",
        "conclusions": "This work deals with an investigation of the properties of the complete sequences of strange-matter stars that carry nuclear crusts. The following items are particularly noteworthy:  \\begin{enumerate} \\item The complete sequence of compact strange stars can sustain extremely rapid rotation and not just those close to the mass peak, as is the case for neutron stars!  \\item If the \\smh is correct, the observed white dwarfs and planets could contain strange-matter cores in their centers. The baryon numbers of their cores are smaller than $\\lsim 2 \\times 10^{55}$!  \\item The strange stellar configurations would populate a vast region in the mass-radius plane of collapsed stars that is entirely void of stars if strange quark matter is not the absolute ground state!  \\item If the new classes of stars mentioned in (2) and (3) exist abundantly enough in our Galaxy, the presently performed gravitational microlensing experiments could see them all!  \\item Due to the uncertainties in the behavior of superdense nuclear as well as strange matter, no definitive conclusions about the true nature (strange or conventional) of observed pulsar can be drawn from cooling simulations yet.  As of yet they could be made of strange quark matter as well as of conventional nuclear matter. \\end{enumerate}  Of course, there remain various interesting aspects of strange pulsars, strange dwarfs and strange planets, that need to be worked out in detail. From their analysis one may hope to arrive at definitive conclusion about the behavior of superdense nuclear matter and, specifically, the true ground state of strongly interacting matter."
    },
    "9604/astro-ph9604099_arXiv.txt": {
        "abstract": "Recent analysis of relativistically expanding shells of cosmological $\\gamma$-ray bursts % standard and not peak luminosity ($L_0$).  Assuming a flat Friedmann cosmology ($q_o = 1/2$, $\\Lambda = 0$) and constant rate density ($\\rho_0$) of bursting sources, we fit a standard candle energy to a uniformly selected log~$N$-log~$S$ in the BATSE 3B catalog correcting for fluence efficiency and averaging over 48 observed spectral shapes. We find the data consistent with $E_0 = 7.3^{+0.7}_{-1.0} \\times 10^{51}$ ergs and discuss implications of this energy for cosmological models of $\\gamma$-ray bursts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604102_arXiv.txt": {
        "abstract": "The production of \\li and \\be during the explosive hydrogen--burning that occurs in nova explosions is computed by means of a hydrodynamic code able to treat both the accretion and the explosion stages. Large overproduction factors with respect to solar abundances are obtained, the exact value depending mainly on the chemical composition of the envelope. Although the final ejected masses are small, these results indicate that novae can contribute to the \\li enrichment of the interstellar medium. Furthermore, since \\be decays emitting a gamma--ray (478 KeV), with a half--life of 53.3 days, the synthesis of \\li could be tested during the INTEGRAL mission. ",
        "introduction": "The origin of lithium and other light elements is still an unsolved problem in astrophysics. It is widely accepted that \\li isotopes are produced  during the Big Bang and by \"spallation\" reactions in the interstellar medium by galactic cosmic rays or in flares (see \\cite{Ree93} for a recent review). Standard Big Bang nucleosynthesis underproduces \\li with respect to solar by more than an order of magnitude (see however the recent paper by \\cite{Del95}), whereas spallation reactions by galactic cosmic rays, produce \\li and \\lil simultaneously, as well as $^{9}$Be, $^{10}$B and $^{11}$B. These two mechanisms are unable to account alone for the present \\li abundance (\\li/H $\\approx$ 2\\power{-9}). Furthermore, they are unable to produce neither the high isotopic ratio \\li/\\lil observed in the solar system (\\li/\\lil = 12.5 $\\pm$ 0.2) nor the $^{11}$B/$^{10}$B one ($^{11}$B/$^{10}$B$\\approx$ 4). Recent measurements of the lithium isotopic ratio in the interstellar medium (\\cite{Lem93}, \\cite{Mey93}) yield values similar to those found in the solar system, indicating that it has remained nearly constant or even decreased during the last 4.5--5 Gyr. The contribution of a low energy component of the galactic cosmic rays, confined at the source or by stellar flares (\\cite{Men71}, \\cite{Can75}, \\cite{Pra93}) can account for the boron isotopes but \\li is still underproduced. Therefore, an extra stellar source able to produce this  \\li without generating \\lil has to be invoked. The interplay of these sources in the galactic evolution of lithium has been extensively studied (\\cite{DAM91}, \\cite{Abi95}).  The synthesis of \\li by a stellar source requires the formation of \\be\\, which transforms into \\li by an electron capture, being \\be half--life 53.3 days. As \\li is very easily destroyed, \\be has to be transported to zones cooler than those where it was formed with a time scale shorter than its decay time. This {\\em beryllium transport} mechanism, as first suggested by \\cite{Cam55}, requires a dynamic situation, like that encountered in asymptotic giant branch (AGB) stars and novae. Another possibility is the production of lithium and boron isotopes by neutrino induced synthesis, during gravitational supernova explosions (\\cite{Woo90} and \\cite{Woo95}). The importance of such a mechanism is still a matter of debate (\\cite{Mat95}).  The production of lithium in AGB stars has been extensively studied and these stars represent the unique observational evidence of an autogenic stellar origin, since it has been observed in them (\\cite{Abi91}, \\cite{Abi93}). The huge abundances of lithium found in some AGB stars are a clear proof that these stars are currently injecting important quantities of lithium to the interstellar medium. However it is hard to estimate their total contribution, since it depends on the estimated number of such stars, that are buried by their own wind (\\cite{Abi93}).  The production of \\li in explosive hydrogen burning and, in particular, in accreting white dwarfs exploding as classical novae, was first studied with a parametrized one--zone model by \\cite{Arn75}. Later on \\cite{Sta78} computed the \\li yields by means of a hydrodynamic code. This code simulated the explosive stage of novae, without considering the accretion phase, i.e., with an initial envelope already in place. The conclusion of this work was that, depending on the initial abundance of \\he and on the treatment of convection, \\li could be formed in substantial amounts during explosive hydrogen--burning in novae. This problem has been revisited by \\cite{Bof93}. On the basis of an extended nuclear reaction network and of updated nuclear reaction rates, but adopting again a parametrized one-zone model, they showed that \\li could only be produced in significant amounts at peak densities lower than 10$^{3}$g.cm$^{-3}$, which are lower than those predicted by hydrodynamic simulations. They argued that the reason of the discrepancy was the neglect of the $^{8}$B(p,$\\gamma)^{9}$C reaction in the calculations of \\cite{Sta78}. However, large overproductions of \\be were found by \\cite{novae95} (using the semi--analytical model of \\cite{Mac83} to obtain the temperature and density profiles and a complete reaction network that included $^{8}$B(p,$\\gamma)^{9}$C) showing that the origin of the different results was not that reaction. In fact, \\cite{Bof93} also showed by means of a two--zone approximation that the efficiency of mixing by convection is a very critical parameter, and they stressed the need of a detailed hydrodynamic model to study \\li production more accurately.  The purpose of this letter is to compute the synthesis of \\li in both carbon-- oxygen (CO) and oxygen--neon--magnesium (ONeMg) novae by means of an implicit hydrodynamic code, that includes a full reaction network, able to treat both the hydrostatic accretion phase and the explosion stage. An estimation of the contribution of novae to galactic enrichment is made, on the basis of the overproductions and ejected masses obtained. The importance of the initial chemical composition of the envelope is analyzed.  The detection of \\li in novae would confirm our theoretical predicition. Furthermore, the detection of gamma--ray emission at 478 KeV, corresponding to the decay of \\be to \\li (half--life 53.3 days) in the early phases of novae by the future mission INTEGRAL would also confirm the thermonuclear runaway model for novae and the nucleosynthesis related to it. ",
        "conclusions": "Our results confirm that nova explosions can produce significant amounts of \\li. Overproduction factors as large as 2000 are obtained. Our results are quite different from those obtained from one--zone models (\\cite{Bof93}), as the most important contribution to \\li enrichment comes from the external shells, where this element has been transported by convection from the burning shell.  Comparison with the results of \\cite{novae95} and \\cite{Pol95} shows that the behavior of \\be abundances can only be correctly predicted if the evolution of \\he during the accretion phase is accurately followed. It is also necessary to stress that the final results strongly depend on the chemical composition at the onset of the explosion. If the underlying white dwarf is a CO one, the \\li abundances are about one order of magnitude larger than if the white dwarf is an ONe one.  Since the decay of \\be to \\li emits a photon with energy 478 KeV, during a phase in which the envelope is very transparent, this transition could be detected by the future INTEGRAL mission (with a sensitivity around 6\\power{-6} at this energy). The flux of the \\be decay line is:  $\\rm{F}(\\rm{counts.s}^{-1}.cm^{-2})= 2.2\\times 10^{-6} \\frac{\\rm {X}(^{7}Be)} {10^{-6}} \\frac{M_{ej}}{10^{-4} M_{\\odot}} \\frac{1}{\\rm{D}^{2}(\\rm{Kpc})} e^{-t/76d}$  \\noindent For an ejected mass of 10$^{-5}$ \\msun, with an abundance by mass of X($^{7}$Be)=8\\power{-6}, the \\be decay line would be detectable just after the outburst only for a nova closer than 0.5 Kpc. But for an ejected mass of 10$^{-4}$ or 10$^{-3}$ \\msun, more in accordance with observations, the lower limit for the distance would be 1.7 or 5.4 Kpc, respectively. This detection would provide a confirmation of the theoretical models of novae and also ensure that \\li is produced in these scenarios, encouraging a deep search of this element in novae."
    },
    "9604/astro-ph9604178_arXiv.txt": {
        "abstract": "We present nine color CCD intermediate-band spectrophotometry of a two square degree field centered on the old open cluster M67, from 3890$\\rm \\AA$ to nearly 1$\\mu$.  These observations are taken as a part of the BATC (Beijing-Arizona-Taipei-Connecticut) Color Survey of the Sky, for both scientific and calibration reasons.  The BATC program uses a dedicated 60/90 cm Schmidt telescope with a 2048 $\\times$ 2048 CCD and 15 specially-designed intermediate-band filters to be applied to both galactic and extragalactic studies.  With these data we show that the BATC survey can reach its goal of obtaining spectrophotometry to a zero point accuracy of 0.01 mag, and down to V =  21 with 0.3 mag random error. Nine-band spectrophotometry of 6558 stars is presented. Systematic issues studied include the effect of image undersampling, astrometric accuracy and transformation from BATC photometric system to broad-band systems.  We fit the color-magnitude diagrams (CMDs) with Worthey's theoretical models. The net result is the excellent fit of the 4.0 Gyr, [Fe/H] = $-0.10$ model to our data, including a good fit to the main sequence (MS) turn-off.  This fit predicts $\\rm (m-M)_0 = 9.47 \\pm 0.16$ and E(B--V) between 0.015 and 0.052. We show that 16\\% of stars in M67 are binaries with mass-ratio larger than 0.7.  Our data are consistent with a toy model with 50\\% of the stars in M67 being binaries and a random distribution of binary mass-ratios, although other models with different mass-ratio distributions cannot be ruled out.  The spatial distribution and mass function (MF) of stars in M67 show marked effects of dynamical evolution and evaporation of stars from the cluster. Blue stragglers and binary stars are the most condensed within the cluster, with degree of condensation depending on mass. The inner part of M67 is missing most of the lower mass MS stars.  We find M67 to have an elongated shape, oriented at an angle of $15^{\\circ}$ relative to the galactic plane.  Within its tidal radius, the observed MF of M67 between 1.2 $\\rm M_\\odot$ and $\\rm 0.8 M_\\odot$ has a Salpeter slope $\\rm \\eta = -1.93 \\pm 0.66$.  For stars of mass below 0.8 $\\rm M_\\odot$, $\\rm \\eta \\sim 0$. It is plausible that the leveling-off of the MF at lower masses is a result of evaporation of lower mass stars in this mass range at a rate of one every $\\sim 10^7$ years. If so, it is plausible that the IMF of M67 has the canonical field value of $\\rm \\eta = -2.0$.  Overall, we find the stellar distribution as a function of mass within M67, and the observed MF, are in good agreement with theoretical predictions of dynamical evolution and evaporation of an old galactic cluster.  Moreover, the fraction of binary stars and inferred IMF for higher mass main sequence stars for this old galactic cluster are consistent with known field star values. This implies a similarity of IMF that persists for at least 4 Gyr in the disk of our Galaxy. ",
        "introduction": "M67 is one of the most-studied old open clusters, because it is reasonably close, populous ($\\sim 1000$ stars) and of similar age as the Sun (cf. Janes \\& Phelps 1994).  As such, photometry of its stars ranges from photographic (e.g., Johnson \\& Sandage  1955, Racine  1971) to CCD (e.g., Gilliand et al. 1991, Montgomery, Marschall \\& Janes 1993; hereafter MMJ), and from UBVRI (Gilliand et al.; MMJ) to uvby (Nissen, Twarog \\& Crawford 1987) to DDO (Janes \\& Smith 1984).  The goals of the photometry have been wide ranging.  They include establishment of photometric standard stars in the cluster (e.g., Schild 1983, Chevalier \\& Ilovaisky 1991); determination of cluster age and metallicity (e.g., Gilliland et al.; MMJ; Burstein, Faber \\& Gonzalez 1986); and determination of the physical properties of the cluster, such as luminosity function (LF), spatial distribution, binary fraction, blue stragglers, and dynamical state of the cluster, etc. (cf. Racine 1971; Francic 1989; MMJ).  In this paper we report a new photometric investigation of M67 that studies this cluster over out to a large radius, $\\sim 1^\\circ$, and down to faint apparent magnitudes (equivalent of V $\\sim$ 20).  The photometric system we use is new, part of a 15-filter intermediate-band system designed to cover the wavelength range $\\rm 3300 \\AA - 1 \\mu$, which avoids known bright or variable sky lines (cf. Thuan \\& Gunn  1976), and is equally applicable to galactic and extragalactic studies.  The M67 observations reported here were obtained both for scientific reasons and as part of the calibrations done for the Beijing-Arizona-Taipei-Connecticut (BATC) Color Survey of the Sky.  The BATC survey is a cooperative long term program by the institutions represented by the authors of this paper. The observational goal of the BATC survey is to obtain accurate ($\\approx 1\\%$) spectrophotometry for all stellar and diffuse objects in 500 one deg$^2$ areas of the sky centered on nearby spiral galaxies, active galaxies, QSO's and various calibration fields for Galactic and extragalactic objects, as well as random fields at high galactic latitudes. M67 was chosen as one of our first targets specifically because it has been well-studied in the past, and because its large angular size provides data for a number of tests of the BATC Survey CCD+filter+telescope system.  The observations and reduction of the M67 data are described in \\S 2. Special aspects of how BATC data are obtained and reduced are highlighted. In \\S 3, the combination of large field of view with accurate intermediate-band photometry permit us to comprehensively investigate both the photometric and the astrometric accuracy we can attain with BATC survey data. Various color-magnitude diagrams we can form with these data are presented in \\S 4, where they are used in conjunction with theoretical models to derive estimates of the cluster age, metallicity, reddening and distance modulus. The binary fraction of stars on the MS, and the relationship of binaries to intrinsic scatter on the MS are discussed in \\S 5. In \\S 6 we discuss the spatial distribution of stars within M67 as a function of mass, as well as the overall shape of the cluster.  The luminosity function (LF) and MF of M67 are discussed in \\S 7 in the context of predictions from dynamical evolution of stars within a cluster that is over 100 relaxation times old.  Our main conclusions regarding calibration of the BATC survey in general, and our scientific results on M67 in particular, are summarized in \\S 8. ",
        "conclusions": "These M67 observations are but the first of many such observations of galactic clusters and galactic fields pertaining to studies of Galactic structure, which we will obtain with the BATC survey.  By using this open cluster as a calibration source, we are able to show the accuracy of the data we can obtain with the 0.6m/0.9m BAO Schmidt telescope at its Xinglong, China site.  It is evident we can obtain spectrophotometry from the ultraviolet to $\\sim 1\\mu$ down to an intrinsic accuracy of better than 0.02 mag for all objects in the nearly 1 deg$^2$ field of the CCD, using Oke--Gunn primary standard stars (Oke \\& Gunn 1983).  Similarly, we can define the positions of these objects to an accuracy of 0.15$''$ for bright stars, using the Guide Star Catalog (Jenkner et al. 1990). To reach these accuracies, and to be able to observe through as many as five filters in a given night, we have used a diffuser plate in combination with our dome flat fields to obtain flat fields at least as accurate as we could obtain using sky observations.  Our observations of M67 span an area $1.92^\\circ \\times 1.92^\\circ$ centered on the cluster in nine BATC intermediate-band filters.  The color magnitude diagrams formed from these data show morphologies consistent with previous CMDs of M67, as well as defining better than most the gap in the MS near the MS turn-off.  Isochrones of a range of appropriate stellar populations has kindly been made available by Dr. G. Worthey.  These models, which convolve our filter senstivity cures with his theoretical spectral energy distribuitions (based on the theoretical stellar evolutionary tracks of VandenBerg (1985) and the synthetic stellar spectra of Kurucz (1992)), have been fitted to our data.  We show that the combined Worthey-VandenBerg-Kurucz model fits the $\\rm (m_{3890} - m_{6075})$ vs. $\\rm m_{3890}$ CMD very well for an age of 4 Gyr and [Fe/H] = --0.10, yielding a reddening of E(B--V) between 0.015 and 0.052 mag and a distance modulus $\\rm (m-M)_0 = 9.47 \\pm 0.16$ mag. The uncertainties in the derived distance modulus and cluster reddening are dominated by uncertainties in the predicted Kurucz fluxes below 4000$\\rm \\AA$.  As such, we use the range of known reddenings for M67 to place limits on what are these uncertainties.  It is well known that M67 has a parallel binary main sequence. We find that 16\\% of stars in M67 are binaries with mass-ratio larger than 0.7, in agreement with previous studies. In a departure from previous work, here we model the binary star population of M67 in terms of a random distribution of secondary/primary mass-ratios from 0.0 to 1.0.  We find that even if the mass-ratio distribution of binaries in the cluster is random, there will still be an apparent offset, parallel MS in the CMD.  Although we cannot accurately determine the fraction of low mass-ratio binaries in the cluster, our data are consistent with 50\\% of the observed stars in M67 being binaries.  As our survey of M67 combines deep images, accurate photometry and a wider field coverage than previous surveys, we are able to explore dynamical evolutionary issues pertaining to this old galactic cluster.  We find much evidence of substantial dynamical evolution of M67.  The spatial distribution of stars is clearly dependent on their masses.  Blue stragglers are the most centrally--condensed among the stars, consistent with previous observations (cf. Mathieu \\& Latham 1986). The assumption that blue stragglers are binary stars of nearly equal mass ratio is consistent with the steady decrease in central concentration observed from blue stragglers to lower mass binaries with secondary/primary ratios $>$ 0.7. A similar trend is seen among the ``single'' stars --- giant star to lower mass main sequence stars --- where the word single is in quotes owing to a likely contamination by low-mass-ratio binary stars.  Overall, the binary stars as a group are more centrally condensed within M67 than are the ``single'' stars.  With these data we can also investigate the two-dimensional shape of M67, which we find to be elongated along an angle of $15^{\\circ} \\pm 45^{\\circ}$ relative to the Galactic plane.  Similarly, we find that the luminosity function (LF) of M67 is dependent on the volume samples. If we observe only the inner core of the cluster, we find results similar to those previously obtained, in which there are more stars on the upper MS than on the lower MS.  However, if we include the cluster stars out to a radius of 33.33$'$ (within the tidal radius estimated by Francic (1989)), we find a LF that rises from the MS turnoff and then flattens out at fainter absolute magnitudes.  We use the isochrone fit of the Worthey-VandenBerg-Kurucz model, combined with our model of binary star distribution, to derive a mass function (MF) for M67 at its present age.  Using the Salpeter (1955) definition of a MF power law, we find the slope $\\eta$ (as in $\\rm M^\\eta$) to be near the canonical field star value of --2.0 for MS stars with masses betwen 1.2 and 0.8 $\\rm M_\\odot$, but which levels off for lower mass stars to the limit of our observations (at 0.5 $\\rm M_\\odot$).  If we attribute this leveling off as due to evaporation of stars through dynamical evolution of this old cluster, we estimate that M67 has lost one star of mass between $\\rm 0.8 M_\\odot$ and $\\rm 0.5 M_\\odot$ every $\\sim 10^7$ years.  Hence, with these data we have been able to observe both direct and implied evidence of substantial dynamical evolution of M67, consistent with theoretical expectations.  We note that our data show an old open cluster whose most likely binary fraction and IMF are remarkably consistent with that of the field.  Files in standard ADC format, giving the full BATC data for all 6558 stars can be accessed either via anonymous ftp from samuri.la.asu.edu (IP 129.219.144.156), in the subdirectory pub/m67, or from the Astronomical Data Center (ADC)."
    },
    "9604/astro-ph9604159_arXiv.txt": {
        "abstract": "We have analysed new R-band photometry of globular clusters in NGC 6166, the cD galaxy in the cooling flow cluster A2199.  In combination with the earlier B photometry of Pritchet \\& Harris (1990), we obtain B$-$R colours for $\\sim$ 40 globular clusters in NGC 6166.  The mean B$-$R is 1.26 $\\pm$ 0.11, corresponding to a mean [Fe/H] = $-$1 $\\pm$ 0.4.  Given that NGC 6166 is one of the most luminous cD galaxies studied to date, our result implies significant scatter in the relationship between mean cluster [Fe/H] and parent galaxy luminosity.  We obtain a globular cluster specific frequency of S$_N$ $\\sim$ 9, with a possible range between 5 and 18.  This value is inconsistent with the value of S$_N$ $\\leq$ 4 determined earlier by Pritchet \\& Harris (1990) from B-band photometry, and we discuss possible reasons for the discrepancy.   Finally, we reassess whether or not cooling flows are an important mechanism for forming globular clusters in gE/cD galaxies. ",
        "introduction": "Broadband colours provide a relatively easy way to estimate metallicities of extragalactic globular clusters. The mean cluster metallicity (we will use colour and metallicity interchangeably in this paper, implicitly assuming that we are dealing with old stellar populations where age effects are not important) gives the overall level of chemical enrichment in the globular cluster system (GCS).  There appears to be a relationship between mean cluster [Fe/H] and parent galaxy luminosity (e.g. Brodie 1993) in the sense that more luminous galaxies have, on average, more metal-rich globular clusters.  However, recent work (e.g. Zepf, Ashman, \\& Geisler 1995a; Secker et al 1995) has shown that there is considerable scatter about this relationship.  The cluster metallicity {\\it distribution} (MD) provides far more information than the mean metallicity alone. The width of the MD is presumably indicative of the inhomogeneity of the protogalactic gas from which clusters formed and/or their subsequent chemical enrichment.  In addition, multimodality in cluster MDs has been detected in several ellipticals (e.g. M87: Lee \\& Geisler 1993, Whitmore et al 1995; NGC 3311: Secker et al 1995; NGC 1399: Ostrov, Geisler, \\& Forte 1993; NGC 3923: Zepf, Ashman, \\& Geisler 1995a), and is most naturally explained by distinct epochs of cluster formation.  The existence of spatial metallicity gradients is also interesting, since such gradients are predicted by many cluster formation scenarios, including classical dissipative collapse (e.g. Eggen, Lynden-Bell, \\& Sandage 1962), galaxy mergers (Ashman \\& Zepf 1992; Zepf \\& Ashman 1993), and possibly cooling flows if clusters have been forming out of the (subsonic) flow over several Gyr (e.g. Fabian 1994).  The number of early-type galaxies with measured globular cluster colours is still small ($<$ 20).  NGC 6166, the central cD in Abell 2199, is an interesting addition, given that it is one of the most luminous galaxies known (M$_V$ $\\simeq$ $-$23.6 for H$_0$=75).  It is also very X-ray luminous, together with a substantial cooling flow of 100--150 M$_{\\odot}$/yr (Edge, Stewart, \\& Fabian 1992). NGC 6166 has long been regarded as the classic multiple-nucleus galaxy, though two of the `nuclei' are in fact not bound to the brightest one (Tonry 1984; Lachieze-Rey, Vigroux, \\& Souviron 1985; Lauer 1986). Lachieze-Rey et al. and Peletier (private communication) both find a colour gradient in NGC 6166, with $\\delta$(B$-$R) $\\simeq$ $-$0.4 from the centre out to $\\sim$ 1 arcmin from the galaxy centre (see Figure 5). Cardiel, Gorgas, \\& Aragon-Salamanca (1995) also find stellar metallicity gradients from long-slit spectroscopy, and there is some evidence  from Einstein and Ginga data of a metallicity gradient in the X-ray gas.  Two of us (Pritchet \\& Harris 1990; hereafter PH) previously detected a GCS in NGC 6166, using B-band images taken at the CFHT.  The determination of the globular cluster specific frequency S$_N$ was hampered by the lack of a background frame, but PH estimated that S$_N$ $\\leq$ 4. In this respect then, NGC 6166 is clearly unlike M87 and some other central cD/gE galaxies where S$_N$ $\\simeq$ 15--20 (e.g. Harris 1991).  PH used this result, together with the large cooling flow, to argue that most globular clusters in central cD/gE galaxies do not form from cooling flows.  We will return to this point in our Discussion (Section 4). The motivation for the present R-band study is two-fold: first, to confirm or not the specific frequency found by PH, and second to obtain B$-$R colours and thus metallicities for the brightest globular clusters in NGC 6166. ",
        "conclusions": "We have obtained photometry to a limiting magnitude of R $\\simeq$ 25 for globular clusters in NGC 6166.  Our main conclusions are:  \\bigskip  $\\bullet$~~The {\\it local} globular cluster specific frequency is S$_N$ = 9, with a possible range between 5 and 18, which is inconsistent with the earlier result of S$_N$ $\\leq$ 4 found by Pritchet \\& Harris (1990) from their B-band images, yet is consistent with a local S$_N$ in a similar region of M87.  $\\bullet$~~The mean B$-$R colour for 37 globular clusters is 1.26 $\\pm$ 0.11, which corresponds to a mean [Fe/H] = $-$1 $\\pm$ 0.4 using a calibration based on Galactic globular clusters.  This mean cluster metallicity is inconsistent with the [Fe/H] $\\sim$ $-$0.4 found for the X-ray gas by White et al. (1994). Our result confirms and extends recent findings of significant scatter in any possible relation between mean cluster metallicity and parent galaxy luminosity at the high-luminosity end.  $\\bullet$~~There is no apparent trend between R mag and B$-$R, and no apparent globular cluster colour gradient.  $\\bullet$~~The available data and theoretical/numerical models allow us to rule out cooling flows as a mechanism for forming significant numbers of globular clusters in gE/cD galaxies, and cooling flows appear {\\it not} to be responsible for the high S$_N$ phenomenon. However, some small fraction of globular clusters may form in cooling flows, and more and better data are needed of globular cluster systems in cooling flow and other environments.  There also remain several interesting theoretical issues to explore (the metallicity evolution of ICM X-ray gas and the expected spread in metallicity, and the role of SN; numerical modelling of star/cluster formation in cooling flows, particularly the role of magnetic fields; a better understanding of cluster formation in merging galaxies)."
    },
    "9604/astro-ph9604184_arXiv.txt": {
        "abstract": " ",
        "introduction": "As I write this in spring 1996 \\cite{Snowmass95}, there is still concern about a crisis in cosmology.  The first article \\cite{Freedman94} using HST observations of Cepheid variable stars to determine a distance to a relatively distant galaxy, the beautiful face-on spiral M100, was published about a year ago.  The distance obtained was $17.1 \\pm 1.8$ Mpc.  With the additional assumptions that M100 lies in the core of the Virgo cluster and that the recession velocity of Virgo corrected for infall is about 1400 km s$^{-1}$, the value obtained for the Hubble parameter is at the high end of recent estimates: $H_0=80 \\pm 17 \\kmsMpc$. Using $h=0.8$ gives, for $\\Omega=1$ and a vanishing cosmological constant $\\Lambda=0$, a very short age for the universe $t_0=8.15$ Gyr, almost certainly younger than the ages of Milky Way globular clusters and even some nearby white dwarfs.  Even with $\\Omega_0=0.3$, about as low as permitted by observations, and with $\\Omega_\\Lambda \\equiv \\Lambda/(3H_0^2) = 0.7$, perhaps even higher than current data allow, $t_0= 11.8$ Gyr for $h=0.8$, which is also uncomfortably short.  Is this a crisis?  Does it undermine the strong evidence for the standard Big Bang?  I don't think so.  Given the considerable uncertainties reflected in the large quoted error on $H_0$, I think even $\\Omega=1$ models are not excluded.  But this Cepheid measurement of the distance to M100 bodes well for the success of the HST Key Project on the Extragalactic Distance Scale, which seeks to measure $H_0$ to 10\\% within a few years.  There has also been recent progress in using Type Ia supernovae as distance indicators, for measuring both $H_0$ and the deceleration parameter $q_0=\\Omega_0/2-\\Omega_\\Lambda$.  The expectation that accurate measurements of the key cosmological parameters will soon be available is great news for theorists trying to construct a fundamental theory of cosmology, and helps motivate the present summary.  In addition to the Hubble parameter $H_0 \\equiv 100 h$ km s$^{-1}$ Mpc$^{-1}$, I will discuss the age of the universe $t_0$, the average density $\\Omega_0$, and the cosmological constant $\\Lambda$.  But there are several additional cosmological parameters whose values are critical for modern theories: the densities of ordinary matter $\\Omega_b$, cold dark matter $\\Omega_c$, and hot dark matter $\\Omega_\\nu$, and, for primordial fluctuation spectra $P(k)=Ak^{n_p}$, the index $n_p$ and the amplitude $A$, or equivalently (for a given model) the bias parameter $b \\equiv 1/\\sigma_8$, where $\\sigma_8 \\equiv (\\delta M/M)_{rms}$ on a scale of $8 \\hMpc$.  A full treatment of these parameters would take a much longer article than this one, so to focus the discussion I will concentrate on the issue of the value of the density $\\Omega_0$ in currently popular cosmological models in which most of the dark matter is cold.  Although much of the following discussion will be quite general, it will be helpful to focus on two specific cosmological models which are perhaps the most popular today of the potentially realistic models: low-$\\Omega$ Cold Dark Matter with a Cosmological Constant ($\\Lambda$CDM, discussed as an alternative to $\\Omega=1$ CDM since the beginning of CDM \\cite{BFPR84,Peeb84}, and worked out in detail in \\cite{CenOstGn}), and $\\Omega=1$ Cold + Hot Dark Matter (CHDM, proposed in 1984 \\cite{CHDM84}, and first worked out in detail in 1992-3 \\cite{DSS92,KHPR}).  I will begin by summarizing the rationale for these models. ",
        "conclusions": "The main issue that I have tried to address is the value of the cosmological density parameter $\\Omega$.  Strong arguments can be made for $\\Omega_0 \\approx 0.3$ (and models such as $\\Lambda$CDM) or for $\\Omega=1$ (for which the best class of models that I know about is CHDM), but it is too early to tell for sure which is right.  The evidence would favor a small $\\Omega_0 \\approx 0.3$ if (1) the Hubble parameter actually has the high value $H_0 \\approx 75$ favored by many observers, and the age of the universe $t_0 \\geq 13$ Gyr; or (2) the baryonic fraction $f_b=M_b/M_{tot}$ in clusters is actually $\\sim 15$\\%, about 3 times larger than expected for standard Big Bang Nucleosynthesis in an $\\Omega=1$ universe. This assumes that standard BBN is actually right in predicting that the density of ordinary matter $\\Omega_b$ lies in the range $0.009 \\leq \\Omega_b h^2 \\leq 0.02$.  High-resolution high-redshift spectra are now providing important new data on primordial abundances of the light isotopes that should clarify the reliability of the BBN limits on $\\Omega_b$.  If the systematic errors in the $^4$He data are larger than currently estimated, then it may be wiser to use the deuterium upper limit $\\Omega_b h^2 \\leq 0.03$, which is also consistent with the value $\\Omega_b h^2\\approx 0.024$ indicated by the only clear deuterium detection at high redshift, with the same D/H$\\approx 2.4\\times 10^{-5}$ observed in two different low-metallicity quasar absorption systems \\cite{Tytler}; this considerably lessens the discrepancy between $f_b$ and $\\Omega_b$.  Another important constraint on $\\Omega_b$ will come from the new data on small angle CMB anisotropies --- in particular, the height of the first Doppler peak \\cite{DGS95}, with the latest data consistent with low $h\\approx 0.5$ and high $\\Omega_b \\approx 0.1$.  The evidence would favor $\\Omega=1$ if (1) the POTENT analysis of galaxy peculiar velocity data is right, in particular regarding outflows from voids or the inability to obtain the present-epoch non-Gaussian density distribution from Gaussian initial fluctuations in a low-$\\Omega$ universe; or (2) the preliminary report from LSND indicating a neutrino mass $\\geq 2.4$ eV is right, since that would be too much hot dark matter to allow significant structure formation in a low-$\\Omega$ $\\Lambda$CDM model.  The statistics of gravitational lensing of quasars is incompatible with large cosmological constant $\\Lambda$ and low cosmological density $\\Omega_0$.  Discrimination between models may improve as additional examples of lensed quasars are searched for in large surveys such as the Sloan Digital Sky Survey.  It now appears to be possible to measure the deceleration parameter $q_0=\\Omega_0/2-\\Omega_\\Lambda$ on very large scales using the objects that may be the best bright standard candles: high-redshift Type Ia supernovae. It is very encouraging that the Perlmutter group \\cite{Perlmutter96} now has discovered $\\sim25$ high-redshift Type Ia supernovae, that other groups are also succeeding in finding such supernovae, and that a theoretical understanding of the empirical correlation between SN Ia light curve shape and maximum luminosity may be emerging.  If the high value $q_0 \\sim 0.5$ in the preliminary report \\cite{Perlmutter96} is right, $\\Omega_\\Lambda$ is probably small and $\\Omega_0\\sim 1$.  The era of structure formation is another important discriminant between these alternatives, low $\\Omega$ favoring earlier structure formation, and $\\Omega=1$ favoring later formation with many clusters and larger-scale structures still forming today.  A particularly critical test for models like CHDM is the evolution as a function of redshift of $\\Omega_{gas}$ in damped Ly$\\alpha$ systems.  Reliable data on all of these issues is becoming available so rapidly today that there is reason to hope that a clear decision between these alternatives will be possible within the next few years.  What if the data ends up supporting what appear to be contradictory possibilities, e.g. large $\\Omega_0$ {\\it and} large $H_0$?  Exotic initial conditions (e.g. ``designer'' primordial fluctuation spectra) or exotic dark matter particles beyond the simple ``cold'' vs. ``hot'' alternatives (e.g. decaying intermediate mass neutrinos) could increase the space of possible inflationary theories somewhat.  But it may ultimately be necessary to go outside the framework of inflationary cosmological models and consider models with large scale spatial curvature, with a fairly large $\\Lambda$ as well as large $\\Omega_0$.  This seems particularly unattractive, since in addition to implying that the universe is now entering a final inflationary period, it means that inflation did not happen at the beginning of the universe, when it would solve the flatness, horizon, monopole, and structure generation problems.  Therefore, along with most cosmologists, I am rooting for the success of inflation-inspired cosmologies, with $\\Omega_0 + \\Omega_\\Lambda = 1$.  With the new upper limits on $\\Lambda$ from gravitational lensing of quasars, number counts of elliptical galaxies, and high-redshift Type Ia supernovae, this means that the cosmological constant is probably too small to lengthen the age of the universe significantly.  So I am hoping that when the dust finally settles, $H_0$ and $t_0$ will both turn out to be low enough to be consistent with General Relativistic cosmology. But of course the universe is under no obligation to live up to our expectations.   \\bigskip \\noindent ACKNOWLEDGEMENTS. I have benefited from conversations or correspondence with S.  Bonometto, S.  Borgani, E. Branchini, D. Caldwell, L. Da Costa, M. Davis, A. Dekel, S. Faber, G. Fuller, K. Gorski, K. Griest, J. Holtzman, A.  Klypin, C. Kochanek, K. Lanzetta, P. Lilje, R. Mushotzky, R.  Nolthenius, J. Ostriker, P.J.E. Peebles, S. Perlmutter, D.  Richstone, A. Sandage, R. Schild, J. Silk, L. Storrie-Lombardi, R.B. Tully, E.L. Turner, M. Turner, A. Wolfe, E. Wright, D.  York, and D. Zaritsky, and fruitful interactions with UCSC graduate students R. Dav\\'e, M. Gross, and R. Somerville.  This research was supported by NASA, NSF, and UC research grants at UCSC.  \\def\\MNRAS{Mon.\\ Not.\\ R.\\ Astron.\\ Soc.} \\def\\ApJ{Astrophys.\\ J.} \\def\\AJ{Astron.\\ J.} \\def\\AA{Astron.\\ Astroph.}"
    },
    "9604/astro-ph9604071_arXiv.txt": {
        "abstract": "The evolution of the two-point correlation function, $\\xi(r,z)$, and the pairwise velocity dispersion, $\\sigma(r,z)$, for both the matter, \\xirho, and halo population, \\xihh, in three different cosmological models: (\\ome,\\lam)=(1,0), (0.2,0) and (0.2,0.8) are described. If the evolution of $\\xi$ is parameterized by $\\xi(r,z)=(1+z)^{-(3+\\eps)}\\xi(r,0)$, where $\\xi(r,0)=(r/r_0)^{-\\gamma}$, then $\\epsrho$ ranges from $1.04 \\pm 0.09$ for (1,0) and $0.18 \\pm 0.12$ for (0.2,0), as measured  by the the evolution of \\xirho\\ at 1 Mpc (from $z \\sim 5$ to the present epoch). For halos, \\eps\\ depends indeed on their mean overdensity. Halos with a mean overdensity of about 2000 were used to compute the halo two-point correlation function, \\xihh, tested with two different group finding algorithms: the {\\it friends of friends} and the spherical overdensity algorithm. It is certainly believed that the rate of growth of this \\xihh\\ will give a good estimate of the evolution of the galaxy two-point correlation function, at least from  $z \\sim 1$ to the present epoch. The values we get for \\epshh\\ range from 1.54 for (1,0) to -0.36 for (0.2,0), as measured by the evolution of \\xihh\\ from $z \\sim 1.0$ to the present epoch. These values could be used to constrain the cosmological scenario.  The evolution of the pairwise velocity dispersion for the mass and halo distribution is measured and compared with the evolution predicted by the Cosmic Virial Theorem (CVT). According to the CVT, $\\sigma(r,z)^2 \\sim G Q \\rho(z) r^2 \\xi(r,z)$ or $\\sigma \\propto (1+z)^{-\\eps/2}$. The values of $\\eps$ measured from our simulated velocities differ from those given by the evolution of $\\xi$ and the CVT, keeping $\\gamma$ and $Q$ constant: $\\eps = 1.78 \\pm 0.13$ for (1,0) or $\\eps = 1.40 \\pm 0.28$ for (0.2,0). ",
        "introduction": "The large scale structure of the Universe that we see today is believed to have developed from the growth of small perturbations in the matter density driven by gravitational instability. The evolution of the clustering of the mass density field depends on the initial conditions via the density power spectrum and the mean density of the universe and is therefore a powerful constraint on theories of structure formation. The evolution of the galaxy clustering, however, need not necessarily follow that of the collisionless component of the mass density field. Galaxies have been subject to external phenomena such as tidal interactions, satellite accretion, mergers, etc. or internal phenomena such as galactic winds, that it would be unlikely to see the galaxy clustering evolution being the same as the clustering evolution of the dark matter.  Two straightforward statistical tools which describe the clustering properties of galaxies, positions and velocities, are the two-point correlation function, $\\xi(x,t)$, and the pairwise velocity dispersion, $\\sigma(x,t)$, hereafter comoving coordinates are denoted by $x$ while proper coordinates are denoted by $r$. In a flat universe where initial conditions were generated by a power-law spectrum with spectral index $n$ the two-point correlation function should scale as \\begin{equation} \\xi(x,t)= \\xi(s), \\end{equation} where $s= x/t^\\alpha$ and $\\alpha = 4/[3(3+n)]$ (\\cite{peeb}). Furthermore, it has been shown that even in the case where the hypothesis of scaling is broken, for instance, by a scale-dependent power spectrum, relation (1) has proved to be a very good approximation (\\eg\\ \\cite{ham}; \\cite{paddy}).  In the regime where the density perturbations grow linearly \\begin{equation} \\xi(r,t) = b(t)^2 \\xi(r,t_i), \\end{equation} where $b(t)$ is the growing mode of the density perturbations and $\\xi(r,t_i)$ is the initial correlation function ($b(t_i) = 1$). In the particular case of \\omeone, where $b$ is just the expansion factor of the universe, $a = (1+z)^{-1}$, and $P(k) \\propto k^n$ : $\\xi(x,z) \\propto (1+z)^{-2} x^{-(n+3)}$ or $\\xi(r,z) \\propto (1+z)^{-(n+5)} r^{-(n+3)}$. On the other hand, in the highly non-linear regime where the hypothesis of stable clustering is supposed to work, $\\xi(r,z)= \\xi(r,0) (1+z)^{-3}$. A convenient form of parameterizing the evolution of $\\xi\\ $ is \\begin{equation} \\xi(r,z) = (1+z)^{-(\\eps+3)} \\xi(r,0), \\end{equation} as it removes the universal expansion (\\cite{gpeeb}). If the hypothesis of stable clustering is satisfied $\\eps = 0$ while $\\eps = 2+n$ in the linear regime (\\omeone\\ and $P(k) \\propto k^n$).  The present observed galaxy two-point correlation function  $\\xi(r,0)$ is to a good approximation a power-law $\\xi_0 = (r/r_0)^{-\\gamma}$. \\cite{dp} from the CFA survey find $\\gamma = 1.77 \\pm 0.04$ and $r_0 = 5.4 \\pm 0.3 h^{-1} \\mpc$. \\cite{lmep} from the Stromlo--APM survey find $\\gamma = 1.71 \\pm 0.05$ and $r_0 = 5.1 \\pm 0.2 h^{-1} \\mpc$. Values for $\\gamma$ consistent with those found locally have been measured at moderate redshifts (Shepherd \\et\\ 1996; \\cite{cfrs}).  It is seen from the studies by \\cite{shep_cnoc} and Le F\\`{e}vre \\et\\ (1996) that the correlation length, $r_0$, has evolved a great deal. Le F\\`{e}vre \\et\\ find that $r_0$ has decreased by a factor of 10, assuming $\\ome = 1$, from the present epoch to $z \\sim 0.6$. An $\\eps_{g g} \\sim 1 \\pm 1$ is derived from these two studies.  There exists an extensive literature both observational and theoretical that at least mention the name of $\\xi$. Here we are interested in those works that deal directly with the time dependency of the two-point correlation function for both the density field and the halo population. In this sense, the paper by Davis \\et\\ (1985) is pioneer. They studied, among other things, the evolution of \\xirho\\ and `\\xigg'\\ in different cosmological models, their numerical simulations consisted of $32^3$ particles in a $64^3$ grid. They found no way of reproducing simultaneously the observed amplitude and slope of the correlation function with \\xirho\\ in their (\\ome,\\lam)=(1.0,0.0) model (a model hereafter is represented by a pair of coordinates, where the first coordinate is \\ome\\ and the second coordinate \\lam). A better match was obtained with their low-density models (flat and open). A biased galaxy formation scenario was invoked to save the (1.0,0.0) model and a \\xigg\\ was computed. Their \\xigg\\ was always above of their corresponding \\xirho. Much of the clustering of these `halos' was due to the pattern imposed by the initial conditions. This work was not intended to study the evolution of $\\xi$ although it certainly showed it. On the other hand, the evolution they found for \\xigg\\ is certainly a rough approximation because we know halos do not necessarly arise from high peaks and high-peak particles do not necessarly end up in halos. The evolution of $\\xi$ has also been showed and studied by other authors (\\eg\\ \\cite{carl91}; \\cite{bv1}; \\cite{bv2}). Carlberg used fof to compute the evolution of \\xihh, and found that while the correlation length of the density field continues to grow in comoving coordinates, \\xihh\\ did not change. Brainerd \\& Villumsen (1994), hereafter BV, with the analysis of simulations of $128^3$ particles, in a standard CDM scenario, were deeper in redshift (they started their analysis at $z=5$ as opposed to $z=2.15$ by Carlberg) and found a non-monotonic growth of \\xihh. Halos found by fof initially produce a biased \\xihh\\ (they follow closely the large-scale pattern of filaments and sheets just as particles located in high-density peaks did) and then decreases because of mergers. The shape of \\xihh\\ will be mainly set by four competing phenomena: (1) mergers, (2) formation, (3) dynamics, and (4) disruption. By the present epoch, because mergers dominates the scene, one would expect to have a \\xihh\\ that lies below \\xirho, how high will this bias be? is a question whose answer will depend on how much are these halos affected by merging and disruption (this will be of course an environmental effect).  In neither of the studies mentioned above was a value for $\\eps$ computed. Recently, two authors, at least, have computed the rate of evolution of \\xirho\\ using equation (3). Jain (1996) computed \\eps\\ as a function of $r/r_0(a)$ for a standard CDM scenario. He founds values for \\eps\\ that range from $\\sim -0.4$ to $\\sim 2.0$. The rate of growth could be much faster that those commonly cited numbers 0 or $\\sim 1$. Peacock (1996) found, on the other hand, that the rate of growth in an open universe could solve the apparent contradiction that: (1) the slope of \\xigg\\ seems not to change at all up to $z \\simeq 1$ and (2) the rate of evolution is relatively rapid with an \\eps\\ value close to 1.  In dissipationless N-body simulations a great effort has been made to find collapsed objects which subsequently  could be asociated with real galaxy halos ( \\cite{defw}; \\cite{cc}; \\cite{bgelb}; \\cite{wqsz}; \\cite{lc}; \\cite{sde}; \\cite{van} ; \\cite{knp}). In particular, a still often used group finding algorithm is {\\it friends of friends}, fof (\\cite{defw}). This algorithm find groups of particles, halos, that are connected more closely than a specified link length, $l$ (\\ie\\ particles that are in an overdensity region in excess of $\\delta_{min} \\sim [\\frac{4\\pi}{3}(l/2)^3]^{-1} \\sim 2/l^3$). It is well known that it suffers from the defect of joining, once in a while, two halos that are physically distinct. To avoid this problem others group finding algorithms were developed. One of these is the spherical overdensity algorithm described by Lacey \\& Cole (1994). Although the identification of these collapsed objects with galaxy halos seems at first not a bad approximation, it was soon realized that they suffered from the defect of {\\it overmerging} (\\eg\\ \\cite{fwde}). Most of these halos identified in an early epoch will be destroyed by the time they reach the present epoch and this is a function of the mean overdensity of the halo and mass resolution of the simulation (halos with the same mass are more tight in simulations with better mass resolution). Some very nice algorithms have been invented to solve this problem and here we cite again the paper by Summers \\et\\ where a complete discussion about this problem can be found (see also \\cite{van}). However, they are very model dependent. In this paper we try to avoid the problem of overmerging by considering only halos whose mean overdensity is rather high ($\\sim 2000$). We still expect the clustering of halos be less than the clustering of galaxies at recent epochs, specially at small scales where very massive halos composed of only one particle should contain many galaxies. Therefore, the analogy of halos with real galaxy halos should not be taken beyond its scope. The statistics of halos or, other galaxy tracers, has been extensively studied at the present epoch and compared with observations (\\eg\\ \\cite{wfde}; \\cite{cc}; \\cite{gelbb}).  An estimate of the galaxy pairwise velocity dispersion is given by the Cosmic Virial Theorem (CVT, \\cite{peeb}) \\begin{equation}\\sigma(r,z)^2 = \\frac{3 J(\\gamma) H(z)^2 \\ome(z)Q r_0(z)^\\gamma r^{2-\\gamma}}{4 (\\gamma-1)(2-\\gamma)(4-\\gamma)} \\end{equation} where $Q$ is the three-point parameter ($J(1.7)=4.14$). By assuming that $Q$ and $\\gamma$ do not vary with $z$ and that $\\xi \\propto (1+z)^{-(3+\\eps)}$, the behavior of $\\sigma(z)$ is required to be: $\\sigma \\propto (1+z)^{-\\eps / 2}$. An estimate of $\\sigg (h r = 1 \\mpc)$ from the CFA survey by \\cite{dp} gives $300 \\pm 40 \\kms$, which in turn produces a value for $\\ome = \\rho_0 / \\rho_c$ ($\\rho_c =  3 H_0^2 / 8 \\pi G = 1.879 h^2 \\times 10^{-29} \\gcm$, $h$ is the Hubble constant in units of 100 \\kms \\mpc$^{-1}$); however, the value of $\\sigg$ can be as high as $\\sim 1000\\ \\kms$ if Coma is included (\\cite{mjb}). Because galaxies may not dynamically represent the background mass density field, the accuracy of the CVT as an estimator of \\ome\\ depends on how well galaxies follow the background dark matter.  This paper is focused to three main goals: (1) do a more systematic analysis of the evolution of $\\xi$ by computing it for both the density field and a halo population (tested with two group finding algorithms), with a high mean overdensity, in three cosmologicals models: (1.0,0), (0.2,0), and (0.2,0.8). (2) Include the evolution of the pairwise velocity dispersion. (3) Put results in a convenient manner so as an observer can use them as they are in the paper. The outline of the paper is as follows. In \\S2 the characteristics of the simulations are described. In \\S3 the evolution of \\xirho\\ is discussed and values for the parameter $\\eps$ are given. In \\S4 the evolution of \\xihh\\ is discussed. In \\S5 the evolution of the first and second moment of the mass and halo velocity field is presented. And finally in \\S6 a summary is presented. ",
        "conclusions": "We have measured the evolution of the two-point correlation function and the pairwise velocity dispersion of the mass density field and halo population. The evolution is parameterized mostly by the \\eps\\ parameter. Our \\eps\\ values depend on the scale and the time period where evolution is measured, and for halos, they also depend on their specified mean overdensity. Results were quoted just for a mean overdensity of about 2000 because we believe halos of this overdensity suffer less from the overmerging problem. The \\eps\\ values for \\xirho\\ range from 0.4 (for (\\ome,\\lam)=(0.2,0.0)), when evolution is measured at 0.2 Mpc, to 1.5, when evolution is measured at 1 Mpc (for (\\ome,\\lam)=(1.0,0.0), both covering a period of time from $z \\sim 1$ to present epoch. The range of \\epshh\\ values covered by halos is:  $-1.0 \\simless \\epshh \\simless 1.5$ for (1.0,0.0), $-1.4 \\simless \\epshh \\simless -0.3$  for (0.2,0.0), and $-2.1 \\simless \\epshh \\simless -0.6$ for (0.2,0.8). The degree to which these results constrain the mean density of the Universe depends on how well the evolution of the galaxy clustering is traced by the evolution of the mass density field or halo population. More and better observations of $\\xigg$ at diferent redshifts along with better numerical determinations of its evolution are needed to constrain more the cosmological parameter space.  The correlation length, $r_0(z)$, was computed in two ways: (1) using the ``standard'' definition $\\xi(r_0,z) = 1$ and (2) fitting $\\xi$ to a power-law function at each epoch; \\ie, assuming that $\\xi(r,z) = (r/r_0(z))^{-\\gamma}$, the shape of $\\xi$ is a power-law and does not change with time. The range of data used for the fit was from the softening length to the correlation length (as found by the first method), to correspond to observations. We like to point out that the analytic evolution does not really do a good job of predicting the details of the nonlinear correlation function, errors as large as 50\\% are obtained, so these simple fits are considerable valuable.  The evolution of $\\sigma(1 \\mpc)$ was assumed to be a power-law with an exponent $-\\eps_v / 2$ (justified by the Cosmic Virial Theorem, CVT). All values found for $\\eps_v$ (halo and mass density field) are systematically higher than those predicted by the CVT; \\ie we see more evolution in the velocities than that predicted by the evolution of $\\xi$ and the CVT for constant $Q$ and $\\gamma$."
    },
    "9604/astro-ph9604161_arXiv.txt": {
        "abstract": "We present a catalog of morphological and color data for galaxies with $21 < I < 25$ mag in the {\\em Hubble Deep Field} (Williams et al. 1996). Galaxies have been inspected and (when possible) independently visually classified on the MDS and DDO systems. Measurements of central concentration and asymmetry are also included in the catalog.  The fraction of interacting and merging objects is seen to be significantly higher in the {\\em Hubble Deep Field} than it is among nearby galaxies. Barred spirals are essentially absent from the deep sample.  The fraction of early-type galaxies in the Hubble Deep Field is similar to the fraction of early-types in the Shapley-Ames Catalog, but the fraction of galaxies resembling archetypal grand-design late-type spiral galaxies is dramatically lower in the distant HDF sample. ",
        "introduction": "Recently published results from the {\\em Hubble Deep Field} (HDF) survey (Abraham et al. 1996a), in conjunction with earlier data from the {\\em Medium Deep Survey} (MDS) (Griffiths et al. 1994, Glazebrook et al. 1995, Driver et al. 1995, Abraham et al. 1996b), and samples of local galaxies such as that given in the Shapley-Ames Catalog (SAC) (Shapley \\& Ames 1932), allow one to study the morphological evolution of galaxy populations as a function of look-back time.  On the basis of bulk measurements of central concentration, $C$, and asymmetry, $A$, for galaxies in the HDF and MDS (calibrated using an artificially redshifted sample of local galaxies with Hubble types earlier than Scd), Abraham et al. (1996a) conclude that by $I=25$ mag the fraction of ``peculiar'' objects has risen to at least 30\\% of the galaxy population. The exact nature of these peculiar systems remains enigmatic: they may be luminous very-late-type spirals or irregulars seen in the rest-frame ultraviolet, mergers, or else systems with no local counterpart. The UV-optical colors of these peculiar systems suggests that a substantial fraction of faint peculiars might be at very high redshifts ($z>3$).  Parameters such as $C$ and $A$ have the important benefit of being objective {\\em measurements}, but these simple parameters describe only a subset of the morphological information contained in the HDF images, and are not designed to detect relatively subtle features (eg. bars and tidal tails), which at present still require subjective visual inspection (and human expertise) in order to be detected.  In the present paper we present a morphological catalog of $21<I<25$ mag galaxies in the {\\em Hubble Deep Field}, including both visual classifications on the MDS (Glazebrook et al. 1995; Abraham et al. 1996b) and DDO (van den Bergh 1960a, b) systems, as well as measurements of the morphological parameters $C$ and $A$ studied in Abraham et al. (1996a).  The plan of this paper follows.  In \\S2 we describe the format of the HDF catalog, and outline the steps that have been taken in order to allow meaningful comparisons to be made between the visual classifications of objects in the HDF and visual classifications in the MDS and SAC.  In \\S3 we show specific examples of different classes of peculiar objects, in order to clarify the terminology used in the catalog.  In Section~4 we discuss the relative fractions of various morphological types in the HDF, MDS, and SAC catalogs. A number of model-independent statements can be made directly from this comparison, which both support and expand upon the conclusions given in Abraham et al.  (1996a). A major new conclusion is that both barred and ``grand-design'' spiral structure is rare in HDF galaxies.  In Section~5 we attempt to understand more about the nature of peculiar systems in the {\\em Hubble Deep Field}, by determining the fraction of peculiar galaxies that show evidence for tidal interactions.  New evidence is presented suggesting that a substantial fraction of the peculiar systems in the HDF and MDS are tidally distorted. \\S6 summarizes our conclusions. ",
        "conclusions": ""
    },
    "9604/astro-ph9604027_arXiv.txt": {
        "abstract": "Three fields located close to the center of the globular cluster $\\omega$~Cen were surveyed in a search for variable stars. We present V-band light curves for 32 periodic variables. This sample includes 10 SX~Phe stars, 19 eclipsing binaries, and three likely spotted variables (FK~Com or RS~CVn type stars). Only 4 of these variables were previously known (including two objects reported in Paper I). All SX~Phe stars and 8 eclipsing binaries from our sample belong to blue stragglers. Five binaries occupy positions on the upper main-sequence of the cluster and another four are located at the base of the subgiant branch. Of particular interest is detection of a detached binary system with a period $P=4.64$ day. Further study of this star can provide direct information about properties of turnoff stars in $\\omega $~Cen. A detached or semi-detached system located slightly above the horizontal branch of $\\omega $~Cen was also discovered. All SX~Phe stars and most of eclipsing binaries discovered in our survey are likely cluster members. We present $V$ vs. $V-I$ color-magnitude diagrams for the monitored fields. A populous group of likely hot subdwarfs as well as numerous candidates for blue stragglers can be noted in these diagrams. Our data indicate that $\\omega $~Cen possesses the largest number of blue stragglers among all galactic globulars. ",
        "introduction": "The Optical Gravitational Lensing Experiment (OGLE) is a long term project with the main goal of searching for dark matter in our Galaxy by identifying microlensing events toward the galactic bulge (Udalski et al. 1992, 1994). At times  the Bulge is unobservable we conduct other long-term photometric programs. A complete list of side-projects attempted by the OGLE team can be found in Paczy{\\'n}ski et al. (1995). In particular, during last three observing seasons we monitored globular clusters NGC~104 (=47~Tuc) and  NGC~5139 (=$\\omega$~Cen) in a search for variable stars of various types. Of primary interest was detection of detached eclipsing binaries. In Paper I (Kaluzny et al. 1996) we presented results for three fields in $\\omega$~Cen. Here we report on variables discovered in another three fields of that cluster. ",
        "conclusions": "This section summarizes results presented above as well as those described in Paper~I. We surveyed a total of six fields covering most of the central part of $\\omega $~Cen and identified 70 periodic variables. This sample includes 5 previously known variables. RR~Lyr stars are excluded from these statistics -- a separate paper devoted to them is in preparation. The most interesting result is identification of 3 detached eclipsing systems which are located at the base of the subgiant branch of the cluster.  These are OGLEGC15, OGLEGC17 and OGLEGC51. Further observations of these systems can lead to direct determination of masses and radii of turnoff stars in $\\omega $~Cen. An accurate determination of the cluster distance would be also obtained from such data. Another clearly detached binaries discovered in our survey are OGLEGC14 and OGLEGC52. The former one is a  blue straggler (assuming it is a cluster member) while the latter is located slightly above the horizontal branch of the cluster. Determination of parameters for components of OGLEGC52 would provide information about masses of horizontal branch stars in $\\omega$~Cen.  We obtained light curves for 34 SX~Phe stars. All of them are located among blue stragglers and are likely members of the cluster. 20 out of 34 of identified SX~Phe stars shows variability with a full range not exceeding 0.1 mag in the V-band. Detected SX~Phe stars constitute about 11\\% of the blue straggler population of $\\omega $~Cen.  Several contact or nearly-contact binaries were identified. Most of them are located among blue stragglers but some systems are present among upper main-sequence and turnoff stars. Most of identified contact binaries are likely members of the cluster.  Three low-amplitude variables with periods from the range 19-33 days were identified on the subgiant branch (OGLEGC22, OGLEGC35 and OGLEGC69). These stars belong most likely to the so called \"spotted\" variables. Another likely variable of this type is OGLEGC30 whose position on the CMD is unknown due to lack of information about its color."
    },
    "9604/astro-ph9604049_arXiv.txt": {
        "abstract": "We study the combined effects of relaxation, tidal heating and binary heating on globular cluster evolution, exploring the physical consequences of external effects and examining evolutionary trends in the Milky Way population.  Our analysis demonstrates that heating on circular and low-eccentricity orbits can dominate cluster evolution. The results also predict rapid evolution on eccentric orbits either due to strong relaxation caused by the high densities needed for tidal limitation or due to efficient bulge shocking of low density clusters.  The combination of effects leads to strong evolution of the population as a whole.  For example, within the solar circle, tidally-limited $10^5 M_{\\odot}$ clusters lose at least 40\\% of their mass in 10 Gyr.  At high eccentricity most of these clusters evaporate completely.  Bulge shocking disrupts clusters within 40 kpc which have less than 80\\% of their mass within their pericentric inner Lagrange point.  Our results are consistent with suggestions that the shape of the cluster luminosity function results from evaporation and disruption of low mass clusters; they further predict that the net velocity dispersion of the cluster system in the inner Galaxy has decreased with time.  Preliminary constraints on formation models are also discussed.  We conclude that the observed cluster system has largely been shaped by dynamical selection. ",
        "introduction": "\\label{sec:intro}  Many studies of globular cluster evolution have focused on internal mechanisms which drive evolution.  This work has produced a clear picture in which initial stellar evolution causes significant mass loss from a nascent cluster (e.g. Chernoff \\& Weinberg 1990); two-body relaxation leads to mass segregation (e.g. Inagaki \\& Saslaw 1985) and core collapse in surviving clusters (e.g. Cohn 1980); binary heating halts collapse (e.g. Lee et al. 1991); and the cluster continuously loses mass due to the escape of stars, eventually undergoing complete evaporation (e.g. Lee \\& Goodman 1995).  It is also recognized that the Galaxy influences cluster evolution.  The time-dependent tidal field heats clusters and tidal limitation aids in the removal of escaping stars.  Previous investigations have considered disk shocking, bulge shocking and tidal limitation, concluding that each will play a role, particularly in the inner Galaxy (e.g. Spitzer \\& Chevalier 1973; Chernoff \\& Shapiro 1987; Aguilar, Hut \\& Ostriker 1988; Weinberg 1994).  In addition, recent observational studies showing correlations of cluster properties with Galactocentric position indicate the measurable influence of the Galaxy (e.g Chernoff \\& Djorgovski 1989; Djorgovski et al 1993; Djorgovski \\& Meylan 1994).  The principal tool used in studies of cluster evolution over the last decade-and-a-half has been direct solution of the Fokker-Planck equation (Cohn 1979).  However, most of these calculations have excluded external effects.  Recently, using time-dependent perturbation theory to investigate disk shocking, Weinberg (1994) demonstrated that resonant forcing can perturb internal stellar trajectories beyond the limit expected from adiabatic invariance.  This indicates that the Galaxy plays a greater role in cluster evolution than previously thought and motivates new studies of cluster evolution which combine internal and external effects.  The importance of external heating requires us to re-examine the current picture of internally-driven evolution.  In particular, external effects will influence the collapse rates, evaporation times and general physical properties derived in previous calculations.  The present work compares this behavior with and without heating over a wide range of cluster properties to present a revised view.  This study also examines the survival and disruption characteristics of clusters on a range of Galactic orbits to shed light on the initial conditions of the cluster system.  The results demonstrate that evolution does indeed depend strongly on position and orbit, further implying that observed cluster properties have been largely determined through dynamics.  Our study rests on a linear theory of external heating-- based on Weinberg's (1994) treatment of disk shocking-- which we include in numerical solutions of the Fokker-Planck equation.  Nearly all previous work has emphasized impulsive shock heating due to a single passage through the disk or bulge.  The work presented here describes resonant heating due to the time-varying tidal field encountered on periodic orbits of the cluster in the Galaxy- an effect we refer to as {\\it orbit\\rm} heating.  In this context, shock heating is seen to result from the broad resonances caused by an impulsively applied external force.  Although our treatment of external heating can include the influence of any component of the Galactic potential, here we consider only the spheroid in order to allow precise definition of the physical behavior and preliminary description of the evolutionary trends.  The present study includes heating on cluster orbits in the isothermal sphere and is used to study cluster evolution from initial King model states to the point of complete evaporation on a range of orbits in the Galaxy.  Our conclusions, therefore, place only lower limits on the overall rate of cluster evolution but are significant nonetheless.  The plan of the paper is as follows.  We derive the linear theory of external heating in \\S\\ref{sec:eh} and discuss its physical interpretation in \\S\\ref{sec:sumep}.  In \\S\\ref{sec:calc}, the numerical implementation is described.  In \\S\\ref{sec:results} we present the results of our study of cluster evolution under the combined influence of internal and external effects.  Finally, in \\S\\ref{sec:disc}, we discuss the implications of the results for the Milky Way globulars.  Readers concerned primarily with the effects of heating and its evolutionary consequences may skip \\S\\ref{sec:eh} without loss of continuity. ",
        "conclusions": "\\label{sec:sumep}  A cluster orbiting in the Galaxy feels a time-dependent tidal field.  A typical orbit is periodic and introduces a periodic external force on orbits of cluster stars.  As described in \\S\\ref{sec:eh}, resonant heating occurs when the periods of stellar orbit and external force coincide, leading to repeated acceleration and increase in the energy of individual orbits.  Integrated over many periods, the energy gained by the orbit increases linearly with time (c.f. eq. \\ref{eq:res}).  Energy absorption eventually leads to the evolution of individual orbits. (see Appendix \\ref{sec:imp} for discussion and numerical implementation of finite duration resonances).  This in turn drives the secular evolution of the cluster potential.  Orbits can either gain or lose energy to the tidal field depending on the particular resonance.  For example, in disk galaxies with flat rotation curves it is well-known that the inner Lindblad resonance loses energy to a perturbation while an outer Lindblad resonance gains energy.  However, for isotropic distribution functions with ${df/dE}>0$, the perturbation always heats the system on average though some localized regions of phase space may lose energy.  Non-resonant interaction has no net effect on an orbit. Successive maxima in the external force tend to accelerate and decelerate the star equally, leading to asymptotic cancellation as long as the initial transients remain linear (i.e. do not alter the intrinsic frequency of the star with an initial jolt).  Over short times, non-resonant heating does occur because the time duration is insufficient for complete cancellation to occur.  Non-linear transient or {\\it impulsive} heating leads to rapid change in orbital energies as a rapidly applied force `kicks' a star regardless of its orbital frequency.  However, the standard impulse approximation, when used to describe a periodic perturbation, ignores the long-term decay of transient energy in the linear limit as well as the linear growth in energy at the resonances.  For most cases of interest, heating rates are in the linear limit, implying that external influence depends primarily on secular transfer of energy through orbital resonances.  To illustrate the behavior of transients and transient decay, Figure \\ref{fig:k5.in.3.comp} compares the exact time-dependent energy input given by equation (\\ref{eq:de}) with the energy input defined by the asymptotic heating rate equation (\\ref{eq:res}).  Transients decay rapidly at low energy and more slowly at high energy.  Empirically, we find that two to three Galactic orbital periods are required before the asymptotic limit is effectively reached.  This treatment therefore adequately describes all but the outermost halo clusters for which initial transients may still be important.  The comparisons of perturbation theory with N-body simulation shown in Appendix \\ref{sec:compsim} demonstrate the validity of the approach.  \\begin{figure} \\epsfxsize=20pc \\epsfbox[12 138 600 726]{k5.in.3.comp.fig} \\caption{The mean change in energy as a function of internal orbital energy in a $W_0=5$ King model due to heating on an eccentric $\\kappa=0.3$ $(e\\approx0.7)$ orbit after one orbital period. Comparison of simulation (histogram), exact time-dependent perturbation theory (equation 6, solid) and heat input calculated from asymptotic heating rate (equation 8, dotted) shows that initial transients decay strongly at low energy while impulsive energy change persists at high energy.  Horizontal dotted line indicates the level of accuracy in the simulation.} \\label{fig:k5.in.3.comp} \\end{figure}  The magnitude of the heating rate is determined by the cluster profile, density and orbit.  The profile and density define the distribution of internal orbital frequencies and the cluster orbit defines the external forcing frequencies and amplitude.  For a cluster of fixed profile and mass, the density is determined by the tidal radius.  Individual clusters may not be tidally limited due to initial conditions or heating-driven expansion.  Therefore we use the function $M(x_p)$ to parameterize the fraction of the total cluster mass enclosed within the instantaneous pericentric inner Lagrange point, $x_p$.  This function depends on the profile and the ratio of cluster mean density to the mean density required by tidal limitation.  A tidally-limited cluster has limiting radius, $R_c=x_p$, and therefore $M(x_p)=1$, while a tidally-unlimited has $R_c>x_p$ and therefore $M(x_p)<1$.  Heating rates for a given orbit increase as $M(x_p)$ decreases.  The perturbing potential in the logarithmic sphere, equation (\\ref{eq:perpot}), heats clusters on orbits of any eccentricity.  The tide transfers energy and angular momentum to the cluster through the resonances, which unbinds stars.  On circular orbits, the tidal field creates a triaxial perturbation of constant amplitude proportional to $\\Omega_0^2$ rotating with fixed pattern speed $\\Omega_0$.  On eccentric orbits, conservation of center-of-mass angular momentum introduces time-dependent amplitude and rotation rate.  This produces more resonances.  Tidal torquing can also induce a net spin.  The rate of external heating is also influenced by our choice of equilibrium phase space distributions.  For example, according to Jeans' theorem, one can define equilibria in the rotating frame of a circular cluster orbit using the Jacobi constant, $E_J$ (e.g. Heggie \\& Ramamani 1994).  By transforming to the frame in which the perturbation is time-independent, we remove the resonances from the problem.  We can therefore choose a bound distribution of orbits in $E_J$ using the limiting zero-velocity surface, so no stars are lost and the cluster experiences no net tidal heating, although inertial energies and angular momenta are not conserved.  Using $f(E)$ instead of $f(E_J)$ leads to heating in the analogous case because we cannot choose orbits which are strictly bound.  In any case, a real cluster cannot reach true equilibrium because it is bound and therefore undergoes relaxation.  In fact, as is shown below, it is typically a competition between external heating and relaxation due to strong resonances with diffused core stars that strongly influences cluster evolution.   \\label{sec:sum}  The key conclusions are as follows:  \\begin{enumerate}  \\item Time-dependent heating on low-eccentricity orbits accelerates evolution and sharply reduces evaporation times.  \\item Tidally limited clusters on high-eccentricity orbits have high internal density, leading to short evaporation times even though heating rates are negligible.  \\item Bulge shocking on high-eccentricity orbits can rapidly disrupt clusters over a wide range in mass and apogalactic radius when their densities are roughly a factor of 10 below the mean density required for tidal limitation.  \\item Evaporation and disruption have shaped the mass, orbit and density distribution of clusters.  In particular, clusters at the peak of the luminosity function had at least $\\sim 30\\%$ more mass depending on orbit. Evaporation on high-eccentricity orbits has decreased the velocity dispersion in the cluster kinematic distribution.  \\end{enumerate}  Secondary results are as follows:  \\begin{enumerate}  \\item Evaporation times do not strongly depend on concentration in most cases.  However, heating can lead to rapid disruption in massive clusters with low concentration because of the low binding energy and long relaxation time.  \\item Clusters disrupting due to heating may still show signs of mass segregation due to continued relaxation.  \\item Heating accelerates evolution over a range of mass spectral index and reduces the dependence of evaporation time on different initial mass spectra.  \\item The development of anisotropy through relaxation in the core will increase evolutionary rates found in the isotropic distributions investigated here.  \\end{enumerate}"
    },
    "9604/astro-ph9604175_arXiv.txt": {
        "abstract": "Using cosmological simulations that incorporate gas dynamics and gravitational forces, we investigate the influence of photoionization by an ultraviolet radiation background on the formation of galaxies. In our highest resolution simulations, we find that photoionization has essentially no effect on the baryonic mass function of galaxies at $z=2$, down to our resolution limit of $\\sim 5 \\times 10^9 M_\\odot$. We do, however, find a strong interplay between the mass resolution of a simulation and the microphysics included in the computation of heating and cooling rates. At low resolution, a photoionizing background can appear to suppress the formation of even relatively massive galaxies. However, when the same initial conditions are evolved with a factor of eight improvement in mass resolution, this effect disappears.  Our results demonstrate the need for care in interpreting the results of cosmological simulations that incorporate hydrodynamics and radiation physics. For example, we conclude that a simulation with limited resolution may yield more realistic results if it ignores some relevant physical processes, such as photoionization. At higher resolution, the simulated population of massive galaxies is insensitive to the treatment of photoionization and star formation, but it does depend significantly on the amplitude of the initial density fluctuations.  By $z=2$, an $\\Omega=1$ cold dark matter model normalized to produce the observed masses of present-day clusters has already formed galaxies with baryon masses exceeding $10^{11}M_\\odot$. ",
        "introduction": "Hierarchical models of structure formation, in which small scale perturbations collapse gravitationally and merge into progressively larger objects, offer an attractive theoretical setting for galaxy formation.  As first emphasized by White \\& Rees (1978), the dissipation of the gas component inside dark matter halos can explain the gap between the masses of large galaxies and the masses of rich clusters, and the combination of cooling arguments with the formation and merger history of the dark halos allows a prediction of the baryonic mass function of galaxies.  The White \\& Rees picture of galaxy formation fits naturally into the more encompassing inflation plus cold dark matter (CDM) theory (Peebles 1982; Blumenthal \\etal\\ 1984) and its variants involving a cosmological constant, space curvature, or an admixture of hot dark matter.  However, attempts to predict the galaxy luminosity function in such scenarios, using analytic and semi-analytic extensions of the Press-Schechter (1974) formalism, have revealed a persistent difficulty: these hierarchical models predict an abundance of faint galaxies that far exceeds that in most estimates of the galaxy luminosity function (e.g., White \\& Frenk 1991; Kauffmann, White, \\& Guideroni 1993; Cole \\etal\\ 1994; Heyl \\etal\\ 1995). This defect is of great interest, since it could undermine not just specific cosmological models but the entire class of hierarchical theories of galaxy formation.  Various ideas have been proposed to overcome this ``faint galaxy excess'' problem, including feedback from star formation (e.g., Dekel \\& Silk 1986; Cole 1991; Lacey \\& Silk 1991; White \\& Frenk 1991) pre-heating of the intergalactic medium (e.g., Blanchard, Valls-Gabaud \\& Mamon 1992; Tegmark, Silk \\& Evrard 1993), and the possibility that faint galaxies are present in the real universe but are undercounted in most surveys because of their low surface brightness (e.g., Ferguson \\& McGaugh 1995). Efstathiou (1992; see also Ikeuchi 1986; Rees 1986; Babul \\& Rees 1992) suggested that the formation of low mass galaxies might instead be suppressed by photoionization from an ultraviolet (UV) radiation background, which can heat diffuse gas to temperatures of a few$\\,\\times\\, 10^4\\,$K and eliminate the dominant sources of atomic cooling at temperatures below $5\\times 10^5\\,$K. The photoionization solution is an attractive one, for the observed quasars alone should produce a rather strong UV background at redshifts $z \\sim 2-4$, and massive star formation in young galaxies would only increase the effect.  In an isolated, coherent collapse, one expects the influence of photoionization to be small for objects with virial temperatures above $\\sim 5\\times 10^5\\,$K, since at these temperatures primordial composition gas is highly ionized by collisions alone.  However, in a hierarchical scenario galaxies form by mergers of smaller subunits, and processes that affect these subunits may be able to percolate their influence to larger mass scales.  In this paper, we examine the effects of photoionization on galaxy formation using 3-dimensional simulations that follow the evolution of gas and dark matter in an expanding universe. Our simulations use \\tsph\\ (Hernquist \\& Katz 1989; Katz, Weinberg, \\& Hernquist 1996, hereafter KWH), a code that combines smoothed particle hydrodynamics (SPH; see, e.g., Lucy 1977; Gingold \\& Monaghan 1977; Monaghan 1992) with a hierarchical tree method for computing gravitational forces (Barnes \\& Hut 1986; Hernquist 1987) in a periodic volume (Bouchet \\& Hernquist 1988; Hernquist, Bouchet \\& Suto 1991). This numerical treatment avoids the physical idealizations required in Efstathiou's (1992) semi-analytic treatment, which made his quantitative conclusions rather uncertain. Our simulation volume is large enough to contain many galaxies embedded in an appropriate large scale environment, so our analysis complements the numerical studies of Steinmetz (1995), Quinn, Katz, \\& Efstathiou (1996; hereafter QKE) and Thoul \\& Weinberg (1996; hereafter TW), which examine collapses of individual objects with higher resolution (see discussion in \\S 5).  The middle phrase of our tripartite title may seem out of place: photoionization is a {\\it physical} process that might plausibly influence galaxy formation, but numerical resolution is not. However, numerical resolution can have a critical effect on simulations of galaxy formation.  Furthermore, the interaction between resolution and assumptions about gas microphysics (photoionization in particular) can be quite subtle. To demonstrate these points, it is most revealing to describe our results in the order in which we obtained them. After summarizing the physical effects of photoionization in \\S 2, we first present results from a series of low resolution simulations in \\S 3, then move to high resolution simulations in \\S 4. In \\S 5 we discuss the implications of our results for numerical simulations and for the ``faint galaxy'' problem. ",
        "conclusions": " photoionization suppresses the formation of systems with $v_c \\simlt 35$ \\kms, reduces the amount of cooled gas by about 50\\% for $v_c \\approx 50$ \\kms, and has progressively less effect towards higher circular velocities.  The agreement between the two studies indicates that TW's results are not sensitive to their idealized geometry and that QKE's results are not sensitive to their choice of UV background parameters or to their finite resolution.  The simulations in this paper, which represent a much larger comoving volume, cannot resolve systems of such low circular velocities, but they do show that no surprising new photoionization effects come into play during the hierarchical assembly of more massive galaxies.  Taken together, these studies suggest that photoionization alone will not solve the problem of excessive numbers of faint galaxies in the CDM model.  Kauffmann, Guideroni \\& White (1994), for example, find that star formation must be suppressed in halos with circular velocities up to 100 \\kms\\ in order to achieve acceptable agreement with observations.  As an alternative to photoionization by the ambient UV background, one can appeal to local feedback within star-forming systems (Dekel \\& Silk 1986; White \\& Frenk 1991; Cole \\etal\\ 1994; Heyl \\etal\\ 1995).  This mechanism can produce acceptable luminosity functions, but only if feedback suppresses the cooling of infalling gas very efficiently, perhaps unrealistically so. A third possibility, hinted at by a variety of recent studies (e.g., Marzke, Huchra \\& Geller 1994; De Propris \\etal\\ 1995; Ferguson \\& McGaugh 1995) is that large numbers of faint galaxies are present in the real universe but are undercounted in conventional estimates of the luminosity function. Finally, there is the possibility that our current theoretical pictures of galaxy and structure formation are missing an important basic ingredient.  \\begin{figure} \\epsfxsize=4.5truein \\centerline{\\epsfbox[85 260 550 765]{figGalPhys.ps}} \\caption{\\protect \\label{figGalPhys} Galaxy distributions at $z=2$ from {\\it (a)} the high resolution, photoionized simulation illustrated in Figures~\\ref{figGalRes}--\\ref{figMatchGal}, {\\it (b)} a simulation with the same initial conditions and UV background but including star formation and feedback, {\\it (c)} a simulation with the same initial conditions and star formation but a UV background spectrum taken from Haardt \\& Madau (1996), and {\\it (d)} a simulation with star formation, the Haardt--Madau background, and initial conditions normlized to $\\sigma_8=1.2$.  Galaxy baryon masses (represented by circle areas on the same scale as in Figure~\\ref{figGalZ2}) include the contributions from stars and cold, dense gas. The first three simulations have the same initial conditions and yield similar galaxy populations despite their different treatments of microphysics.  The fourth simulation has a higher fluctuation amplitude and produces more massive galaxies. } \\end{figure}  For $v_c > 100$ \\kms, Figures~\\ref{figGalRes}--\\ref{figGalMass} show that simulations with sufficient numerical resolution yield similar galaxy populations with and without a UV background. In Figure~\\ref{figGalPhys}, we show that this robustness to microphysical assumptions holds more broadly. Panel (a), repeated from Figure~\\ref{figGalRes}, displays the galaxy distribution at $z=2$ from the high resolution simulation with $J_0=10^{-22}$ and $\\alpha=1$.  Panel (b) shows a simulation that has the same initial conditions and UV background and includes star formation, using the algorithm described in KWH. Galaxies are identified by applying SKID to the combined distribution of star particles and cold, dense gas particles ($\\rhob/{\\bar\\rho}_{\\rm b}>1000$, $T<30,000\\;$K); the area of the circle representing a galaxy indicates its mass of stars plus cold, dense gas. The galaxy populations are nearly identical even though the first simulation treats the baryon component hydrodynamically throughout the calculation and the second simulation steadily converts cold, dense gas into collisionless particles and injects thermal energy from supernova feedback into the surrounding medium. Figure 5 of KWH demonstrates similar insensitivity to assumptions about star formation in low resolution simulations with no UV background evolved to $z=0$.  The simulations illustrated in Figures~\\ref{figGalPhys}a and~\\ref{figGalPhys}b both assume a simple $\\nu^{-1}$ spectrum for the UV background.  Figure~\\ref{figGalPhys}c shows a simulation with star formation and the UV background spectrum of Haardt \\& Madau (1996), who compute the ambient radiation field that would be produced by the observed population of quasars after absorption and re-emission by the Ly$\\alpha$ forest. We use their spectra for $q_0=0.5$ (kindly provided by P.\\ Madau) to compute photoionization parameters and photoionization heating parameters as a function of redshift (see \\S 3 of KWH for formal definitions of these quantities).  We reduce the amplitude of the spectrum by a factor of two so that the simulation roughly reproduces the observed mean Ly$\\alpha$ opacity towards high-$z$ quasars (Press, Rybicki \\& Schneider 1993) for our adopted baryon density of $\\Omega_b=0.05$. Despite the changes in the spectral shape, intensity, and evolution of the UV background, the galaxy population in Figure~\\ref{figGalPhys}c is almost indistinguishable from that in~\\ref{figGalPhys}b, except for the absence of a few of the smallest systems. The assumed UV background does influence the amount of Ly$\\alpha$ absorption produced by intergalactic hydrogen, but the effect enters mainly through a single parameter, the background intensity weighted by photoionization cross-section ($\\Gamma_{\\gamma {\\rm H}_0}$ in the notation of KWH).  Figure~\\ref{figGalPhys}d shows the galaxy population from a similar CDM model, but with the power spectrum normalized to $\\sigma_8=1.2$, close to the level implied by COBE.  We again incorporate star formation and the Haardt \\& Madau (1996) ionizing background spectrum. While the overall structure of the galaxy distribution is similar in this simulation (the initial fluctuations were increased in amplitude but otherwise unaltered), the largest galaxies are considerably more massive.  There are also a number of low mass galaxies present in this model that have not formed by this redshift in the lower normalization simulation. The galaxy mass functions for the three $\\sigma_8=0.7$ simulations are nearly identical, but the mass function for the $\\sigma_8=1.2$ simulation is shifted systematically towards higher masses.  Our results show that numerical resolution is important for studying galaxy formation in simulations like these, and that the required resolution can depend on the microphysical treatment in a complicated way. They also suggest that there are critical regimes where a small change in resolution can produce a major, qualitative change in results, by shifting a simulation from a regime where gas in a typical halo can cool to a regime where it cannot. Figure~\\ref{figGalPhys} adds an encouraging coda to this disconcerting story: when the numerical resolution is adequate, the simulated galaxy population is insensitive to uncertain microphysical assumptions, at least within the range that we have examined here.  The one model that stands out in Figure~\\ref{figGalPhys} is the one with a different primordial fluctuation amplitude, and it is easily distinguished from the others. Simulations that model galaxy formation in a large volume to $z=0$ will be computationally demanding, but we can expect them to provide good constraints on theories of the origin of cosmic structure."
    },
    "9604/hep-ph9604288_arXiv.txt": {
        "abstract": " ",
        "introduction": "High energy cosmic ray particles, mainly nucleons, interacts with matter to produce secondary particles that are mainly mesons and some baryons. This applies in particular to cosmic rays entering the atmosphere of the Sun. The produced particles propagate through the Sun until they either decay or make secondary interactions producing new particles that contribute further to develop a cascade. Decays of particles in such cascades will  produce neutrinos and other leptons, \\eg  muons that in turn decay into  neutrinos.  This scenario is similar to the cascades induced by cosmic rays  in the Earth's atmosphere, which we have studied extensively \\cite{GIT}.  However, the solar atmosphere is less dense at the typical interaction heights  and, therefore, a larger fraction of the mesons will decay instead of interacting. This leads to relatively more neutrinos produced in the Sun as compared to the Earth.  It has been proposed \\cite{Moskalenko} that the Sun might be used as a `standard candle' for neutrino telescopes, which is only possible if the flux is significantly higher than the Earth's atmospheric flux.  A large such neutrino flux from the Sun would, on the other hand, be a severe background for searches of neutrinos from neutralino annihilation  in the Sun \\cite{Kamionkowski}.  The hypothetical neutralinos appear in theories based on  supersymmetry (SUSY) \\cite{SUSY} and are of fundamental interest in particle  physics as well as in cosmology since they could contribute to the dark matter in the Universe.  In this paper we study the production of muon and electron neutrinos in  cosmic ray interactions in the Sun, as well as their propagation through the Sun and to the Earth where they could be detected in neutrino telescopes,  such as {\\sc Amanda} \\cite{Amanda}, {\\sc Baikal} \\cite{Baikal}, {\\sc Dumand}  \\cite{Dumand} and {\\sc Nestor} \\cite{Nestor}. The cascade interactions in the  Sun are treated in detail using Monte Carlo methods to simulate the high  energy particle interactions. In particular, the Lund model \\cite{Lund} and Monte Carlo programs \\cite{Pythia} are invoked.  The resulting neutrino fluxes at the Earth are compared with the fluxes from  cosmic ray interactions in the Earth's atmosphere \\cite{GIT}.  This solar neutrino flux is discussed in terms of the above `standard candle'  idea and neutralino search. We also investigate the possibility of neutrino oscillations taking place between the source at the Sun and the detector at the Earth. ",
        "conclusions": "\\label{sec:Summary}  We have calculated the high energy muon and electron neutrino fluxes arising from the interactions of cosmic ray particles with the solar matter. Our resulting muon neutrino flux agrees with that obtained by Seckel \\etal\\  \\cite{SSG}, but our result extends a few orders of magnitude higher in energy. The muon neutrino flux in ref.~\\cite{Moskalenko} is in disagreement with  both these results, due to an oversimplified model where secondary interactions in the Sun are not taken into account.  These solar neutrino fluxes are one to two orders of magnitude larger than  those from cosmic ray interactions in the Earth's atmosphere, when integrated  over the solid angle of the Sun as seen from the Earth. This opens a  possibility to use the solar neutrino flux as a `standard candle' for  neutrino telescopes. However, one must here also consider the angular spread introduced by the charged current interaction producing the detectable muon relative to the incoming neutrino direction.  This deflection is typically $\\sim10^{\\circ}\\sqrt{10\\,GeV/E_{\\mu}}$. In addition, the experimental measurement of the muon direction also has a limited resolution. Taking these two effects together, the solar disc will typically  cover less than $\\sim10\\%$ of the solid angle that has to be integrated over. The solar muon neutrino flux, which stays the same, would then have to be compared with a factor ten, or more, increased atmospheric flux such that the two  would be of the same order of magnitude. Under these conditions, one would  only have a factor two increase towards the Sun and the use of  the Sun as `standard neutrino candle' does not look so promising. One should here also be aware of the very low absolute rate of events in  a neutrino telescope of the size now under consideration.  For example, we estimate the rate of neutrino events with $E_\\nu >100\\: GeV$ to be one per year in a detector covering $6\\cdot 10^4\\: m^2$.  We have also investigated the potential for observing neutrino oscillations  during the passage from the source in the Sun and a detector at the Earth. In principle, one could access an interesting region in the parameter plane  of $sin^22\\theta$ and $\\delta m^2$, but with the very low absolute rates it is beyond present neutrino telescopes.  The positive consequence of these small solar neutrino fluxes is that they will cause less of a background problem in attempts to detect neutrinos from other sources. Of particular interest here is the search for neutrinos from neutralino annihilation in the Sun, where the predicted rate can be up to an order of magnitude larger depending on the supersymmetric parameters \\cite{Edsjo96}. A clear observation of this phenomenon would both demonstrate supersymmetry, \\ie physics beyond the standard model in particle physics, and the presence of non-baryonic dark matter in the Universe."
    },
    "9604/astro-ph9604012_arXiv.txt": {
        "abstract": "We present UV ($\\sim 2300$ \\AA) images, obtained with the {\\it Hubble Space Telescope} ({\\it HST}) Faint Object Camera, of the central $20''$ of five galaxies containing circumnuclear star-forming rings. The five galaxies are from a well-defined sample of 103 normal, nearby galaxies we have observed with {\\it HST}. At the {\\it HST} resolution ($0.05''$), the rings  break up into discrete star-forming clumps. Each clump is composed of many luminous ($L_{\\lambda} (2300{\\rm  \\AA})\\approx 10^{35}-10^{37}$ erg s$^{-1}$  \\AA$^{-1}$) and compact ($R\\ltorder5$ pc) star clusters. These objects are similar to those that have been recently reported in colliding and starburst galaxies, and in several other circumnuclear rings. A large fraction, 15\\%--50\\%, of the UV emission originates in these compact clusters. Compact clusters therefore may be the preferred mode of star formation in starburst environments.  For one galaxy, NGC 2997, we measure the UV-optical colors of the individual clusters using an archival {\\it HST} WFPC2 image at $\\sim 6000$ \\AA. Comparing the colors and luminosities to starburst population synthesis models, we show that the clusters are less than 100 Myr old and have masses of at least a few $10^3 M_{\\odot}$, with some as high as $10^5 M_{\\odot}$. The UV extinction to those clusters that are detected in the UV is at most a factor of 10. In NGC 2997, the limits on the masses and the ages of the young clusters indicate that these objects will remain bound and evolve into globular clusters. However, data in additional wavebands are needed to critically test this hypothesis.  The luminosity function of the clusters in the rings is similar in shape to those measured for super star clusters in other star-forming galaxies, and extends to luminosities lower by several orders of magnitude. All five of the UV-detected circumnuclear rings occur in barred or weakly barred spiral galaxies of type Sc or earlier. None of the five rings has an active nucleus at its center, arguing against a direct correspondence between circumnuclear star formation and nuclear activity. ",
        "introduction": "One of the most intriguing results of galaxy imaging with the high angular resolution of the {\\it Hubble Space Telecope} {\\it (HST)} has been the discovery of extremely luminous, compact, young star clusters in a variety of starburst environments. These include the cooling-flow/merger-remnant galaxy NGC 1275 (Holtzman et al.\\ 1992), merging galaxies (Whitmore et al.\\ 1993; Conti \\& Vacca 1994; Whitmore \\& Schweizer 1995), ``amorphous'' peculiar galaxies (Hunter, O'Connell, \\& Gallagher 1994; O'Connell, Gallagher, \\& Hunter 1994; O'Connell et al. 1995), two  barred galaxies with circumnuclear rings (Barth et al. 1995), and nine starburst galaxies with a range of morphologies and luminosities (Meurer et al. 1995). Prior to {\\it HST\\/}, only a few objects of this type were known to exist (e.g., Arp \\& Sandage 1985; Melnick, Moles, \\& Terlevich 1985); the severe crowding in most starbursts prevents their resolution into individual clusters in ground-based images. Similar objects have been detected recently in high-resolution infrared images of the starburst galaxy NGC 1808 (Tacconi-Garman, Sternberg, \\& Eckart 1996).  Such ``super star clusters,'' having diameters not exceeding a few tens of parsecs and luminosities greater than $M_V = -10$ mag, result from extreme episodes of rapid star formation which are virtually absent in ordinary galactic disks.  The small radii, high luminosities, and presumably high masses of these clusters have led to suggestions that, as opposed to standard OB associations, they may remain bound systems. If so, after $10^{10}$ years they would have magnitudes comparable to those of old Galactic globular clusters, and therefore could be present-day versions of young globular clusters. Van den Bergh (1995) has argued that the sizes and luminosity functions of the young clusters more closely resemble those of Galactic open clusters than those of globular clusters. Meurer (1995b) counters that the sizes of the young clusters are often poorly constrained in the {\\it HST} data, and that the luminosity functions of the globular clusters and the super star clusters should not be directly compared because the super star clusters in a given galaxy are not all necessarily of the same age and so are viewed at different epochs (and luminosities) in the clusters' evolution. Note that luminous, compact clusters are also known in our Galaxy and the Large Magellanic Cloud (LMC). For example, recent {\\it HST} and near-infrared observations of the R136 cluster in 30 Doradus in the LMC indicate a mass of $\\sim 10^4 M_{\\odot}$ within a radius of 2.5 pc (Brandl et al. 1996). The dynamical crossing time for such a cluster is $3\\times 10^5$ yr, while the measurements constrain the stellar ages to 3--5 Myr. This system is therefore bound, and could also evolve into a low-mass globular cluster. The limits on luminosity used to define super star clusters and distinguish them from the more common and slightly less-luminous compact clusters are arbitrary; physically, these are all similar objects. Clearly, learning more about these young clusters can shed light on rapid star formation, globular-cluster formation, and the possible connections between them.  The bulges of some nearby spiral galaxies display intense star formation in a ring of optical ``hot-spots'' (S\\'ersic \\& Pastoriza 1967). Such galaxies offer us an opportunity to study starburst properties at a level of detail that would be impossible for more distant or more heavily obscured systems. Kiloparsec-sized nuclear rings can form as a result of bar-driven inflow to an inner Lindblad resonance (e.g., Tubbs 1982; Athanassoula 1992; Piner et al.\\ 1995), yielding a large concentration of dense gas in a small region surrounding the nucleus (Kenney et al.\\ 1992). Several nuclear rings have been studied extensively from the ground (e.g., Hummel, van der Hulst, \\& Keel 1987), but the small sizes of the rings (typically $\\sim10$\\arcsec) have precluded detailed studies of their structure.  {\\it HST\\/} imaging has opened a new vista on these hot-spot galaxies.  WFPC, WFPC2, and FOC images have revealed large populations of super star clusters in nuclear rings in four nearby galaxies: NGC 4314 (Benedict et al. 1993), NGC 1097 and NGC 6951 (Barth et al. 1995), and NGC 3310 (Meurer et al. 1995).  The process of cluster formation in these starbursts appears to be physically distinct from normal star formation in spiral galaxies. Even the most luminous disk \\ion{H}{2} regions in late-type galaxies seldom contain compact clusters that might be potential young globulars (Kennicutt \\& Chu 1988). The occurrence of super star clusters may be unique to starburst environments. Hot-spot \\ion{H}{2} regions are also spectroscopically distinct from disk \\ion{H}{2} regions; the hot-spot \\ion{H}{2} regions have higher heavy-element abundances and possess lower emission-line equivalent widths than disk \\ion{H}{2} regions of the same luminosity, suggesting that continuous, intense star formation occurs within the hot-spots (Kennicutt, Keel, \\& Blaha 1989; Mayya 1995).  Circumnuclear hot-spots contain over 50\\% of the massive star formation in some galaxies (Phillips 1993); thus, determining the physical parameters of hot-spot \\ion{H}{2} regions is important for understanding the global star formation properties of spirals.  We have carried out a UV imaging survey with the {\\it HST} Faint Object Camera (FOC) of the central regions of 103 large, nearby galaxies, selected randomly from a complete sample of 240 galaxies. An atlas of the survey is presented by Maoz et al. (1996), where full details of the sample selection, observations, and data reduction are given. One of the main purposes of the survey is to study star formation in the central regions of the galaxies. By observing at $\\sim2300$ \\AA, the survey images detect predominantly the youngest existing stellar populations and thus provide a clean probe of the most recent sites of active star formation, uncontaminated by light from more evolved stars.  In this paper, we present results for five out of the 103 galaxies, whose UV images display circumnuclear star-forming rings. In a forthcoming paper (Ho et al. 1996c) we will analyze star formation based on the UV images of the other sample galaxies. In \\S 2, below, we describe the observations and data analysis. Notes on the individual galaxies are given in \\S 3. We discuss the physical implications of the data in \\S 4 and summarize our results in \\S 5. ",
        "conclusions": "We have presented {\\it HST} data for five nearby galaxies displaying circumnuclear star-forming rings in the UV. The rings are composed of numerous compact ($<$ few pc) and luminous objects, similar in all respects to those that have been recently seen in other starburst environments and presumed to be young star clusters. One of the galaxies, NGC 2997, has an {\\it HST} image in visual light, providing color information for these objects at the {\\it HST} angular resolution.  We summarize our main conclusions as follows.\\\\ 1. A substantial fraction (10\\%--40\\%) of the UV light in these galactic circumnuclear regions is emitted by the compact sources. Considering our detection limits, the true fraction may be larger, of order 30\\%--50\\%. Compact clusters are therefore a common, and possibly dominant mode of star formation in circumnuclear starbursts.\\\\ 2. In NGC 2997, the typical compact source has the color of a late B to early A star and the luminosity of $\\sim 10^4-10^5$ such stars, or the luminosity of tens of OB stars. The object must therefore be either a $\\sim 30$ Myr-old unreddened cluster, or a young reddened cluster. Either way, the mass of the typical object is $\\gtorder 10^4 M_{\\odot}$.\\\\ 3. The blue color of the NGC 2997 clusters which are visible in both optical and UV bands means that these clusters are not highly reddened, and are absorbed in the UV by less than a factor of 10. From the present data we cannot say whether the clusters that are detected only in the visual band (75\\% of the total) are fainter than the UV detection limit, more reddened, or older than the others.\\\\ 4. Most of the clusters with measurable size have radius $< 5$ pc. When combined with the age or mass estimates above, this implies they are bound, and they will evolve into objects similar to globular clusters.\\\\ 5. UV-selected circumnuclear rings occur preferentially in barred galaxies of Hubble type S0 to Sc. There is no clear correspondence between their occurrence and that of a Seyfert or LINER-type AGN.  Further observations in additional wavebands at the {\\it HST} resolution can significantly constrain the nature of these objects, and disentangle the effects of age, reddening, and star-forming initial conditions and history."
    },
    "9604/astro-ph9604154_arXiv.txt": {
        "abstract": "In the course of an extensive campaign to measure radial velocities of galaxies in a selected sample of compact groups photometrically studied by de Carvalho et al. (1994), we report the discovery of a system very rich in starburst galaxies and AGNs. This is the system HCG 16 of Hickson's (1982) catalog of CGs. The 7 brightest galaxies form a kinematical group with bi-weighted estimate mean velocity of V$_{BI} = 3959 \\pm 66$ km s$^{-1}$, dispersion $\\sigma_{BI} = 86 \\pm 55$ km s$^{-1}$, a median radius $\\langle {\\rm R} \\rangle  = 0.197$ Mpc, a mean density of $\\langle {\\rm D} \\rangle = 217$ gal Mpc$^{-3}$ and a total absolute magnitude of M$_{B} = -22.1$. From their spectral characteristics, we have identified one Seyfert 2 galaxy, two LINERs and three starburst galaxies. Thus, HCG 16 appears to be a dense concentration of active galaxies. In our sample of 17 Hickson groups, HCG 16 is unique in this regard, suggesting that it is an uncommon structure in the nearby universe. ",
        "introduction": "Compact groups (CGs) of galaxies may represent some of the densest concentrations of galaxies known in the Universe and so may provide ideal laboratories for studying the effects of strong interactions on the morphology and stellar content of galaxies. This concept has motivated several recent observational programs aimed at establishing the dynamical reality of these structures and signs of galaxy interactions. Whereas most of these works (Rubin et al. 1991, Pildis et al. 1995) have mainly considered objects in  which clear morphological signs of interactions are evident,  studies made with larger samples have surprisingly shown that the frequency of mergers in CGs is significantly less than that predicted by simple dynamical arguments (Zepf  1993). Confirming this result, in a search for tidal-tail induced dwarf galaxies in 42 Hickson's (1982) CGs, Hunsberger et al. (1995) have shown that only 7 of them exhibit clear signs of such objects.  Another question that may be addressed by studying CGs refers to the environmental origin of the nuclear activity of galaxies (AGNs) represented by the presence of nuclear emission lines that cannot be explained in terms of normal stellar population. This longstanding question has been debated in the literature with no clear answer.  For instance, whereas the studies  of Kennicut \\& Keel (1984) and Keel et al. (1985) have shown that the AGN phenomenom occurs more often in binary or interacting systems, other studies found no relevant correlations between nuclear activity and interaction parameters (Dahari \\& de Robertis 1988; Laurikanen \\& Salo 1995). However, it seems clear from studies of optically selected samples that nearby AGNs avoid systems which are strongly interacting. Activity, if present, seems rather due to intense starburst formation induced by the interaction itself (Bushouse 1986). This may not be true for the ultraluminous infrared galaxies which are mostly interacting galaxies and show an increasing probability of being Seyferts with increasing infrared luminosity (Veilleux et al., 1995). One may speculate that if AGNs do really prefer interacting systems but avoid those that are strongly interacting, then an ideal place to find them would be the CGs, or at least a subclass of them which, although apparently very dense as deduced from their radial velocities, show no signs of violent interactions.  In this {\\sl Letter} we spectroscopically revisit one compact group which has been previously noticed as presenting morphological signs of interaction among its galaxies. This is the system HCG 16 of Hickson's (1982) catalog of CGs. We have found this group to be very rich in starburst galaxies and AGNs, being possibly a very rare case of such a high concentration of active galaxies in such a dense environment.  \\section {Data and Kinematic Properties}   HCG 16  and its neighbouring galaxies were spectroscopically observed in the course of an extensive campaign aimed to measure radial velocities of galaxies of a selected sample of CGs photometrically studied by de Carvalho et al. (1994, hereafter dCAZ94). The spectra were taken at the 4m CTIO telescope, using the ARGUS fiber feed spectrograph. The details of the instrumental setup and data reduction are discussed  by de Carvalho et al. (1996). Table 1 lists the galaxy number as given by dCAZ94 (column 1), positions R.A. and Dec. (columns 2 and 3), magnitudes in B (dCAZ94, column 3), and heliocentric velocities and errors (columns 4 and 5). We list here only the seven galaxies defining the group. A complete list of all the galaxies measured in the field is presented by de Carvalho et al. (1996).  We have used the ROSTAT statistical package (Beers et al. 1990) in order to analyse the velocity distribution of our sample of galaxies around HCG16 (0.5$^{\\circ}\\times$0.5$^{\\circ}$ around the center), which is complete down to B = 18.6$^{m}$. The analysis revealed the presence of a kinematical group consisting of the 7 brightest galaxies in the field, with mean velocity (as given by the bi-weighted estimate) of V$_{BI} = 3959\\pm 66$ km s$^{-1}$ and dispersion $\\sigma_{BI} = 86\\pm 55$ km s$^{-1}$ (90\\% confidence errors). A more detailed study is presented in Ribeiro et al.(1996a)  The compact group HCG 16 is  a larger group than originally noted by Hickson, being composed of 7 galaxies. In Figure 1, we present the distribution of galaxies in HCG 16. The original group described by Hickson is composed of 4 galaxies: 1, 2, 4 and  5. To these galaxies, we added 3 others, 3, 6 and 10. Galaxies 3 and 6 are quite luminous, which suggests that this new addition is of great importance for the dynamical structure of the group. The median of the projected separations of the galaxies in the group is $\\langle {\\rm R}\\rangle = 0.197$ Mpc (H$_{\\circ} = 75~{\\rm km~ s^{-1} Mpc}^{-1}$), and the mean density is $\\langle{\\rm D}\\rangle = 217$ gal Mpc$^{-3}$. The density is therefore  smaller than the density $\\langle {\\rm D}\\rangle = {10^4}$ gal Mpc$^{-3}$ originally determined by Hickson. While this density is lower than that in the central part of rich clusters of galaxies, HCG 16 has a density roughly 30 times higher than those found for loose groups of galaxies, as determined by Maia et al. (1989). The total absolute magnitude of the system is M$_{B} = -22.1$ (as compared to M$_{B} = -21.5$ determined by Hickson).   \\section {Spectral Properties}  In Figure 2, we present the spectra for the 6 emission--line galaxies out of the 7 which constitute HCG 16. The seventh member of the group, galaxy 10, does not show any emission lines. Because the spectra are not flux calibrated, we divided the number counts by their mean values, to compare the relative intensity of the emission lines in the different galaxies. The galaxy 4 shows the most intense emission lines. Except for the unusually high ratio of [NII]/H$\\alpha$ and the presence of a faint [OI]$\\lambda$6300 line, this spectrum is very similar to those of disk spiral HII regions. This galaxy is clearly experiencing an intense starburst phase in or near its nucleus. Based on the similarity of the spectra, the galaxies 3, 5 and 6 are also starburst galaxies, although at relatively lower intensities than 4. Galaxies 1 and 2 show a different type of spectrum, both with the high [NII]/H$\\alpha$ ratio typical of AGNs.   The classification of the kind of activity encountered in the galaxies of HCG 16 is based on the different line ratios shown by galaxies of different activity classes. The criteria that we have used for our spectral characterization are explained in detail in Ribeiro et al. (1996b).  The presence or absence of a wide Balmer emission line component allows us to distinguish between a Seyfert 1 and a Seyfert 2.  We distinguished between Seyfert 2 and LINERs based on the ratio [OIII]$\\lambda$5007/H$\\beta > 2.5$ (Coziol 1996). We adopted this definition because in many cases the [OII]$\\lambda$3727 line, used by Heckman (1980) to characterize the LINER type, was not available. For the starburst galaxies, we distinguished also between HII galaxies and Starburst Nucleus Galaxies (SBNGs). This distinction is based on a correlation between spectroscopic characteristics and morphologies (Coziol et al. 1994).  In general, the HII galaxies are high--excitation ([OIII]$\\lambda$5007/H$\\beta > 2.5$) small metal poor galaxies, while the SBNGs are low--excitation massive and metal rich galaxies. Usually, the spectra of SBNGs indicate a mean excess of 0.2 dex in the [NII]/H$\\alpha$ ratio as compared to normal HII regions (Coziol et al. 1996).  In Figure 3, we present the diagnostic diagram of [OIII]$\\lambda$5007/H$\\beta$ vs. [NII]/H$\\alpha$ for all the emission--line galaxies in HCG 16. In this diagram, the dotted line represents our criterion to distinguish between high and low--excitation galaxies. The solid line is the empirical separation established by Veilleux \\& Osterbrock (1987) between galaxies ionized by an AGN and galaxies ionized by stars. The uncertainties on the line ratios are determined based on Poisson statistics. The high uncertainties in [OIII]$\\lambda$5007/H$\\beta$ for galaxies 1 and 2 reflect the weakness of the emission lines. In those galaxies for H$\\beta$ the stellar absorption dominates over the nebular emission. Following our classification, HCG 16 contains 3 AGNs (2 LINERs and 1 Seyfert 2) and 3  starbursts. Based on the unusually high intensity of the lines [SII]$\\lambda\\lambda$6716,6734 (see Fig. 2), galaxy 2 looks more like a LINER than a Seyfert 2 (Rubin et al. 1991). The starburst nature of galaxy 4 was already suggested by its high  emission in infrared (Sparks et al. 1986). The equivalent widths EW(H$\\alpha$+[NII]) for the galaxies 3, 4, 5 and 6 are 44, 146, 50 and 10 \\AA respectively. Except for galaxy 6, those values are comparable to those found in typical starburst galaxies (Kennicutt 1992). ",
        "conclusions": ""
    },
    "9604/astro-ph9604140_arXiv.txt": {
        "abstract": "We report on the initial results of a spectroscopic investigation of galaxies in the Hubble Deep Field which exhibit spectral discontinuities between the F450W and F300W passbands, indicative of the presence of the Lyman continuum break in the redshift range $2.4 \\simlt z \\simlt 3.4$.  We have employed color selection criteria similar to those we have used for selecting high redshift galaxy candidates from ground--based images. We find that, as for the ground--based color selection, the criteria are very successful in selecting high redshift objects. Of the 8 galaxies observed (selected from a list of 23 candidates with magnitudes equivalent to ${\\cal R}\\le 25.3$), 5 have confirmed redshifts in the range $2.59 \\le z \\le 3.22$, with the remaining 3 being indeterminate because of contamination from nearby brighter objects. As expected, the HST filter system is sensitive to a somewhat broader range of redshifts than our ground--based $U_n G {\\cal R}$ filter system,  and therefore the surveyed volume per unit area on the sky is correspondingly larger. The distribution of candidates on the plane of the sky is clearly non--uniform, consistent with the available ground--based data on the high redshift galaxies. Most Lyman break objects in the Hubble Deep Field exhibit a similar range of morphological properties to the $z>3$ galaxies we have previously identified in other fields, characterized by very compact cores (some with multiple components) with half--light radii of $0.2-0.3$ arc seconds, often surrounded by more diffuse and asymmetric ``halos''. A few of the brighter HDF Lyman break galaxies, however, have particularly unusual morphologies. ",
        "introduction": "The ``Hubble Deep Field'' (hereinafter ``HDF'') (Williams \\et 1996) presents an opportunity to assess the colors and morphologies of galaxies down to unprecedentedly faint magnitude levels. Given that the field was observed across a wide color baseline, including the UV F300W filter, it is a natural place to extend our ongoing studies of the galaxy populations at $z\\sim 3$ (Steidel \\et 1995, 1996; Giavalisco \\et 1996; hereinafter Papers III, IV, V) that have been based on the Lyman continuum break entering the UV passband at substantial redshifts, giving the objects colors which distinguish them from the rest of the faint field galaxies . Our ground--based photometric system and selection criteria (see Steidel \\et 1993, 1995 for details) have been shown to be sensitive to galaxies in the redshift range $3.0 \\le z \\le 3.5$, based on direct spectroscopic follow--up of the $z>3$ candidates using the W. M. Keck telescope (Paper IV). It is possible to use very similar selection criteria-- essentially ``flat'' rest--UV spectra across the observed--frame optical passbands, with a dramatic ``drop--out'' in the ultraviolet passband-- in the HST/HDF photometric system to pick out the candidate high redshift galaxies in the Hubble Deep Field. In this paper, we report the first spectroscopic observations of these ``Lyman break'' galaxies in the HDF. We have also performed a morphological analysis of these bright candidates following the lines of our previous work presented in Paper V.  \\section {SAMPLE SELECTION}  Given the very small amount of time (a matter of only a few days) between the availablility of the Hubble Deep Field data and our observing run on the W. M. Keck telescope, galaxies were selected as Lyman break candidates in a manner that was not necessarily optimized for the HST filter system. We simply adopted criteria very similar to those used in our ground--based survey. All of the magnitudes on the HST system were converted to ``AB'' magnitudes (such that a galaxy with equal magnitudes in each passband has a spectrum that is ``flat'' , i.e. $f_{\\nu} \\propto \\nu^0$). We then approximated our ground--based ${\\cal R}$ passband, which has an effective wavelength of 6930 \\AA\\ for a flat--spectrum source, by averaging the F606W and F814W AB magnitudes. For the colors of the objects of interest, this is likely to be very close to a ${\\cal R}$ magnitude in both normalization and effective wavelength. (We will call this magnitude ${\\cal R}$ hereinafter). We then formed the colors $F300W - F450W$ and $F450W - {\\cal R}$, and applied a ``spectral curvature'' criterion, $$F300W - F450W > 1.2 + (F450W-{\\cal R}), $$ which is equivalent to requiring that the break across the two bluer passbands is more than 3 times the break across the two redder ones. In addition, we expect that the unabsorbed continua of high redshift galaxies will be very blue (even with the blanketing effect of the Lyman $\\alpha$ forest, up to redshifts of $z \\sim 3.5$), so that we require that $F450W -{\\cal R} < 1.2$. In practice, essentially all of the objects which satisfied the ``spectral curvature'' criterion also satisfied the ``blueness'' criterion. All of the colors were measured using the isophotal aperture defined in the Version 1 catalog produced by the HDF team; such apertures are conservatively large for objects as bright as those considered here, and are certainly not optimal for obtaining colors with the highest S/N.  We then took the entire catalog of objects having ${\\cal R} \\le 25.3$ (this has been shown to be the practical limit for follow--up spectroscopy with Keck, and also ensures that the limits in F300W for Lyman break objects will be extremely robust) and applied the color selection criteria. Figure 1 shows the two--color diagram for all of the HDF objects with ${\\cal R} \\le 25.3$, together with the region of this plane satisfying the above color criteria. The filtering resulted in a sample of 25 resolved objects culled from the three detectors of the WFPC2 camera, and one stellar object. Two of the resolved objects were rejected as spurious after close inspection of the HDF images. The colors of the stellar object may have been slightly affected by saturation; a long slit spectrum obtained on the same night as the slit mask observations described below show it to be a subdwarf star. All of the candidates for follow--up spectroscopy, with accompanying photometry and positions, are listed in Table 1\\footnote{In the original generation of the list of candidates for spectroscopy, an error was made in the conversion of the F606W magnitude to the AB scale, so that the $F450W - {\\cal R}$ color was measured to be 0.2 magnitudes bluer, and the ${\\cal R}$ magnitude 0.2 magnitudes fainter, than the values that are now tabulated. In addition, when the catalog F300W magnitudes did not exceed 1 $\\sigma$ above sky in the isophotal aperture, we assigned lower limits on the F300W magnitude equivalent to $+1\\sigma$ above sky--in some cases this made the $F300W - F450W$ break less pronounced than in the original list. As a result, there are objects in Table 1 which no longer satisfy the adopted color criteria completely. We have retained all of the original candidates for spectroscopy in Table 1, however.}. A mosaic of all 23 of the galaxies is shown in Figure 2.  We expected {\\it a priori} that the color selection applied to the HST data set would be sensitive to somewhat lower redshifts than our ground--based criteria, since the Lyman limit of a galaxy is well within the F300W passband for any redshift $z \\simgt 2$. To better quantify the expected redshift range encompassed by our selection criteria, we convolved the throughput curves of the HST passbands with the same model galaxy template (Bruzual and Charlot 1993; see Paper III) we have used in our ground--based system, accounting for blanketing in the Lyman alpha forest following the prescriptions of Madau (1995) and by assuming that photons shortward of the Lyman limit in the galaxy rest frame are completely absorbed.   This assumption is probably valid, since recent {\\it Hopkins Unltraviolet Telescope} observations of nearby star forming galaxies have shown no detectable flux emerging shortward of the Lyman limit [Leitherer \\et 1995], and because of the known high opacity of intergalactic H~I for redshifts $z \\simgt 2.5$ (cf. Paper III, Madau 1995). The model galaxy colors, which should be considered rough estimates, satisfy our criteria for candidate selection in the redshift range $2.4 \\le z \\le 3.4$, i.e. a much broader range of redshifts than in our ground--based high redshift galaxy selection ($3.0 \\le z \\le 3.5$). This then explains (in part) the surprisingly large number of Lyman break candidates; whereas we would have expected $\\sim 2$ Lyman break galaxies in the 4.7 square arc minute HDF field of view to ${\\cal R}=25$ based on our ground--based statistic of $0.40$ galaxies per arc square minute (see Paper IV) we see 10 such candidates in the HDF. After accounting for the probed volume, which is larger by a factor of $\\sim 2$ than the equivalent volume probed by the $U_nG{\\cal R}$ selection criteria (because of the larger redshift range), and the fact that one probes a few tenths of a magnitude fainter in the star--forming galaxy luminosity function at the average redshift probed, these numbers are probably consistent with one another within the errors.  A detailed analysis of the luminosity function of Lyman break galaxies in the HDF is postponed to a separate paper.   \\section {OBSERVATIONS}  The spectroscopic observations were obtained on the night of 22 January 1996 (UT) on the W. M. Keck telescope with the Low Resolution Imaging Spectrograph (Oke \\et 1995).  Two slit masks were constructed, each containing slits for 8 of the Lyman break candidates in the HDF, plus 15--20 additional slits for galaxies in the HDF ``flanking fields'', which were much shorter (1--2 orbit) HST exposures. The candidate objects for spectroscopy were assigned slits based on a weighting scheme that was largely subjective, but was loosely based on apparent magnitude (brighter objects being given larger weight), and efficiency of how many slits could be placed on candidates within the HDF as a function of position angle, etc. The first mask, which is the one actually used for observations, was optimized to have a slit falling roughly along the major axis of one particularly interesting galaxy, C4-06, which also happens to be the brightest candidate in the HDF, and to allow another slit to be placed on C4-09, another interesting candidate with a peculiar morphology.  The weather was far from optimal on the night the observations were made, with variable cirrus and mediocre seeing ($\\sim 1.3-1.5$ arc seconds FWHM). We obtained a total of 7200s of integration through the first slit mask, in six separate exposures of 1200s each. (This integration represents $\\sim 90$\\% of the workable time we experienced in the entire two night observing run.) The telescope was moved small amounts ($\\sim 2$\\arcs) parallel to the slit direction  between each exposure, so that the spectrum of each object fell at 3 independent spatial positions on the Tektronix 2048 $\\times$ 2048 detector. We note that this integration time is only about half that obtained for the $z > 3$ galaxy spectra which were presented in Paper IV. The spectra were obtained with a 300 line/mm grating blazed at 5000 \\AA\\ in first order, resulting in a spectral resolution of $\\sim 12$ \\AA\\ through the 1\\secpoint 4 slits.  Two of the slits assigned to HDF Lyman break candidates (C4-02 and C2-01 in Table 1) were contaminated by light from nearby mask alignment star holes, so that no useful data were obtained for them. A third galaxy, C3-01, is contaminated by two nearby galaxies at smaller redshifts (see Figure 2) due to the poor seeing and the unfavorable alignment of the slit on the plane of the sky. We have obtained redshifts for the remaining 5 candidate $z>2.4$ galaxies, and indeed they all fall within the expected range of redshifts.  The spectra of the five confirmed high redshift galaxies in the HDF are presented in Figure 3. ",
        "conclusions": "We have obtained confirming spectra of 5 out of 23 galaxies from the 3 Wide Field Camera chips of the Hubble Deep Field, selected to have ${\\cal R} \\le 25.3$ and to satisfy color criteria essentially identical to those we have used successfully in our ground--based survey for very high redshift galaxies. The {\\it HST} data are, as expected, sensitive to a larger range of redshifts than the ground--based $U_nG{\\cal R}$ photometric system; the 5 redshifts all fall within the expected range $2.4 \\le z \\le 3.4$ based on the simple model used (see Madau \\et 1996 for a color prescription optimized for the {\\it HST} filter system). Thus, the Lyman break technique for isolating high redshift galaxies is again shown to be extremely efficient.  It is perhaps significant that in the sample of 5 redshifts available thus far, 3 are within $\\pm 2600$ \\kms of one another, near redshift $z \\approx 2.80$. We also point out the distribution of the candidates among the three WFC chips, with 12 on WF2, 9 on WF4, and only 2 on WF3. These two findings indicate the possible presence of large--scale clustering in the distribution of actively star--forming galaxies at high redshifts. Such inhomogeneities and signals of very large scale clustering are also present in our larger--field ground--based data as well, as we will discuss in future work.  We reiterate a conclusion which was reached in Paper V-- that the high redshift galaxies are in general very compact, with scales comparable to the cores of present--day luminous galaxies. The bulk of the star formation at high $z$ is occuring in very compact regions, of very high surface brightness, and which generally exhibit a relatively high degree of azimuthal symmetry.  We view the spectra presented in this paper as further evidence of the effectiveness of using the Lyman continuum break criterion for isolating well-defined populations (essentially volume limited) of high redshift galaxies. There is a great deal of science that may be done, from studies of the luminosity function of galaxies at extreme redshifts, to studies of large--scale structure, using the high resolution images and extremely accurate colors afforded by the Hubble Deep Field data alone. These avenues are all being explored at the moment."
    },
    "9604/astro-ph9604006_arXiv.txt": {
        "abstract": "An interesting way to calibrate the absolute magnitudes of remote Type Ia supernovae (SNe Ia) that are well out in the Hubble flow$^1$, and thus determine the value of the Hubble constant, $\\bf H_0$, has been introduced by van den Bergh$^2$.  His approach relies on calculations$^3$ of the peak absolute magnitudes and broad--band colors for SN Ia explosion models. It does not require any corrections for extinction by interstellar dust, and no SNe Ia are excluded on grounds of peculiarity.  Within the last few years distances have been determined to the parent galaxies of six SNe Ia by means of Cepheid variables$^{4-10}$. Cepheid--based distances also have become available for three other SNe Ia if one is willing to use the distance to a galaxy in the same group in lieu of the distance to the parent galaxy itself.  Here we determine the value of $\\bf H_0$ in a way that is analogous to that of van den Bergh, but now using Cepheid--based distances instead of calculated light curves.  We obtain $\\bf H_0 = 55 \\pm 5\\ \\rm \\bf km\\ s^{-1}\\ Mpc^{-1}$.  This value, with $\\bf \\Lambda=0$ and $\\bf \\Omega=1$, corresponds to a cosmic expansion time of $\\bf 12 \\pm 1$ Gyr, which is consistent with several recent determinations of the ages of globular clusters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604076_arXiv.txt": {
        "abstract": "We present a consistent model for the UV and supersoft X-ray emission from the symbiotic nova SMC3 (=~RX~J0048.4--7332). Following the present picture of symbiotic stars, the model consists of radiation from a hot star and an emission nebula excited by that star. The observations were compared to  theoretical models in which the hot star's emission is calculated with the help of hydrostatic and Wolf-Rayet-type non-LTE  model atmospheres. Our analysis clearly shows evidence for mass loss rates of several $10^{-6}$~\\Msolar/yr% . The minimum effective temperature compatible with both the observed UV and X-ray flux is about $260\\,000$~K, which is higher than in any other star analyzed with sophisticated NLTE model atmospheres. Since the hydrostatic surface is hidden by the stellar wind no upper limit for the temperature can be determined. However, we were able to determine the total luminosity of a symbiotic nova with reasonable accuracy ($L_{\\mbox{SMC3}}=10^{4.05\\pm 0.05}$\\,L$_\\odot$). This value is well below the Eddington limit ($\\approx50\\,000$~L$_\\odot$). In order to reproduce the observed energy distribution a carbon-to-helium ratio $>2\\cdot 10^{-4}$ --- leading to an absorption edge at 0.39\\,keV ---  is necessary. ",
        "introduction": "ROSAT observations established  the class of `supersoft X-ray sources' (`SSS'; e.g.\\ Hasinger 1994) with almost no flux at $h\\nu\\ge 0.5$~keV; normally, close-by low-luminosity objects like single white dwarfs or  cool coronal stars are excluded from this definition. The nature of many SSS is still a matter of controversial discussion. RX~J0048.4--7332 is one of ten  SSS confirmed as members of the Magellanic Clouds (Kahabka \\& Pietsch 1993; Pietsch \\& Kahabka 1993). In this paper we investigate the hypothesis  that the X-ray flux of this object is due to photospheric emission of a very hot white dwarf. Given a sufficiently high temperature the photospheric  emission of white dwarfs  is measurable in the ROSAT window (e.g. Jordan et al. 1994b, Wolff et al. 1996). A hot white dwarf is hardly the correct explanation for all supersoft X-ray sources. However, at least for some sources, the presence of a very hot star is also in agreement with other evidence: For instance, in symbiotic systems there are high excitation nebular emission lines combined with a hot continuum in the UV. The supersoft appearance is also expected from the theoretical side; in particular, for symbiotic novae Sion \\& Starrfield (1994) predict phases in which a very hot white dwarf can be detected.  Indeed, the supersoft source RX~J0048.4--7332 coincides with the symbiotic nova SMC3 that has been in outburst since 1981 (Morgan 1992). Thus, for this system it appears natural to propose that a hot white dwarf is  the source of the X-ray flux. Symbiotic stars are interacting binaries consisting of a red giant and a very hot white dwarf (typically $T\\tief{eff}\\sim100\\,000$~K), whose  radiation ionizes the emission nebula. In the optical spectral range we observe a composite spectrum from the cool star and the nebula. The UV spectrum is dominated by the nebular emission with only a small contribution from the hot star that becomes dominant at the shortest wavelengths. The spectral range best suited for a direct observation of the hot component of a symbiotic system is the soft X-ray and EUV part of the electromagnetic spectrum, where these stars emit the bulk of their energy (e.g. RR~Tel, see Jordan et al.\\ 1994a, hereafter JMW).  Vogel \\& Morgan (1994; Paper I) obtained an IUE spectrum of SMC3 which looks typical for a symbiotic star of moderate excitation. On the other hand, the optical spectrum (Morgan 1992; M\\\"urset et al.\\ 1995 = Paper III) reveals nebular ionization that is unusually high for symbiotics, up to Fe$^{+9}$. Even more peculiar, however, are the X-ray observations presented by Kahabka et al.\\ (1994): Despite of the large distance, the ROSAT count rate of SMC3 is about the same as of  RR~Tel,  hitherto the  X-ray brightest of all known galactic symbiotic novae (Jordan et al.\\ 1994a; hereafter JMW). Thus, the ratio of the optical to the X-ray flux of SMC3 is  a challenge for  models of symbiotic stars. On the other hand, the known distance of this extragalactic system provides a unique opportunity to investigate the energetics of interactions and outbursts in symbiotic systems.  Kahabka et al.\\ (1994) have presented black body fits to the ROSAT data, which resulted in an extremely unrealistic luminosity of up to  $10^{13}$~L$_\\odot$. However, we will show that a reasonable value is obtained if the data are compared to appropriate model atmospheres, and if we consistently take into account the constraints from both the  X-ray and UV regions of the electromagnetic spectrum.  In Sect.~\\ref{data} we list the data used in our analysis, in Sections \\ref{atmo} and \\ref{anal} we describe the methods, the results are discussed in Sect.~\\ref{resu}, and Sect.~\\ref{conclu} presents some concluding remarks. ",
        "conclusions": "}  SMC3 is extraordinary, even among symbiotic stars. It is remarkably X-ray bright and shows unusually high ionization in the nebula. We succeeded to model these properties and the UV flux with a very hot WR-type model atmosphere. Our result clearly confirms that at least some of the supersoft X-ray sources contain very hot photospheres of degenerate stars. It shows in particular, that such hot stars can be encountered in symbiotic binary systems. Reasonable fits are achieved  at temperatures above $260\\,000$~K and mass-loss rates of $\\Mdot\\gappr1\\cdot10^{-6}$~\\Msolar/yr, (assuming $v_\\infty=1000$~km/sec). Due to the properties of the wind models (the hot star itself is hidden by the stellar wind) no  upper limit for the effective temperature could be specified. However, this uncertainty does not effect the total luminosity: $L_{\\mbox{SMC3}}=10^{4.05\\pm 0.05}$\\,L$_\\odot$.  Even the lower limit for the effective  temperature is extraordinarily high (see the temperature distribution for hot components of symbiotic systems in Fig.~7 of Paper III). To our knowledge SMC3 has the hottest stellar atmosphere ever derived in a sophisticated NLTE analysis. Until now, the hottest known \\hbox{(pre-)} white dwarfs are the DA central star of the planetary nebula WDHS1 with \\Teff=160\\dots 200\\,kK (Liebert et al. 1994; Napiwotzki 1995) and the PG1159 stars H1504+65 (Werner 1991) and RXJ2117+3412 (Werner et al. 1995b) with \\Teff=170\\,kK. Note, however, that the temperature of SMC3 refers to the hydrostatic radius $R_*$ and not to $\\tau_{R}=2/3$.  On the basis of LTE model atmospheres Heise et al. (1994) have shown that the ROSAT PSPC observations of the  super-soft X-ray source 1E0056.8-7154, associated with the planetary nebula N67 (Wang 1991; Cowley et al. 1995) can be explained assuming a hot pre-white dwarf with an effective temperature as high as 450\\,kK. However, at such high temperatures NLTE effects are expected to be strong so that this result has to be checked using NLTE model atmospheres.  Strong mass loss is indispensible to explain the spectrum of SMC3. Direct or indirect signatures of strong mass loss are also encountered in other symbiotic novae (see Sion et al.\\ 1993, JMW, Nussbaumer et al.\\ 1995, Nussbaumer \\& Vogel 1995, and the discussion in M\\\"urset et al.\\ 1995a). If SMC3 has been shedding its wind constantly since the beginning of the outburst, it must have lost at least a mass of $\\Delta M\\sim10^{-4}$~M$_\\odot$. This is much larger than the mass burnt during the same time, which means that the duration of the outburst of SMC3 will be limited rather by the mass loss than by the nuclear processes. $\\Delta M$ is also a lower limit to the mass of the layer accreted by the white dwarf prior to outburst.  Kenyon et al.\\ (1993), M\\\"urset \\& Nussbaumer (1994), and JMW found for the known galactic symbiotic novae luminosities clearly below the Eddington limit. This result is now confirmed with a result that, due to the well known distance, is probably more reliable than the galactic ones. Here we encounter a basic difference between symbiotic and classical novae which seem to exceed the Eddington limit (Starrfield et al.\\ 1993; Shore, personal communication). If a core-mass -- luminosity  relation holds for the symbiotic novae we can derive a mass $M\\approx0.8$~\\Msolar for the hot star (with the relation from Joss et al.\\ 1987).  As a final comment we would like to stress that a photospheric X-ray source like SMC3 can only be reasonably investigated with i) multi-frequency data and ii) sophisticated \\underline{non}-LTE atmosphere models. Kahabka et al.\\ (1994) impressively demonstrated the failure of black body fits to the ROSAT PHD by deriving a luminosity of $L\\sim10^{13}$~L$_\\odot$ for SMC~3. The reasons for the failure are obvious: Fitting the ROSAT data alone is a dangerous extrapolation, because they represent only an extreme tail of the stellar spectrum. Even worse, this extreme part is exactly the spectral region where hot atmospheres differ very strongly from Planck curves."
    },
    "9604/astro-ph9604183_arXiv.txt": {
        "abstract": "We consider accretion disks consisting of counter-rotating gaseous components with an intervening shear layer. Configurations of this type may arise from the accretion of newly supplied counter-rotating gas onto an existing co-rotating gas disk. For simplicity we consider the case where the gas well above the disk midplane is rotating with angular rate $+\\Omega$ and that well below has the same properties but is rotating with rate $-\\Omega$. Using the Shakura-Sunyaev alpha turbulence model, we find self-similar solutions where a thin (relative to the full disk thickness) equatorial layer accretes very rapidly, essentially at free-fall speed. As a result the accretion speed is much larger than it would be for an alpha disk rotating in one direction. Counter-rotating accretion disks may be a transient stage in the formation of counter-rotating galaxies and in the accretion of matter onto compact objects. ",
        "introduction": "The widely considered models of accretion disks have gas rotating in one direction with a turbulent viscosity acting to transport angular momentum outward (\\cite{SS}). However, recent observations point to more complicated disk structures in both active galactic nuclei and on a galactic scale. Warped and tilted nuclear accretion disks have been detected for example in NGC 4753 (\\cite{TWIST}). Recent high spectral resolution studies of normal galaxies has revealed counter rotating gas (\\cite{CIRI}) and/or stars in many galaxies of all morphological types - ellipticals, spirals, and irregulars (see reviews by Rubin 1994 and Galletta 1996). NGC 4826 (\\cite{NGC4826}) and IC 1459 (\\cite{IC1459}) are examples of galaxies with central counter-rotating gas and stellar disks, respectively. In elliptical galaxies, the counter-rotating component is usually in the nuclear core and may result from merging of galaxies with opposite angular momentum. Newly supplied gas with misaligned angular momentum in the nuclear region of a galaxy may have important consequences for nuclear activity if there is a rotating massive black hole at the galaxy's center (\\cite{SCH}). In contrast, in a number of spirals and S0 galaxies, counter-rotating disks of stars and/or gas have been found to co-exist with the primary disk out to large distances ($10 - 20$ kpc), with the first example, NGC 4550, discovered by Rubin, Graham, and Kenney (1992). It is not likely that the large scale counter-rotating disks result from mergers of flat galaxies with opposite angular momenta because of the large vertical thickening observed in simulation studies of such mergers (\\cite{BARNES}). Thakar and Ryden (1996) discuss different possibilities, (a) that the counter-rotating matter comes from the merger of an oppositely rotating gas rich dwarf galaxy with an existing spiral, and (b) that the accretion of gas occurs over the lifetime of the galaxy with the more recently accreted gas counter-rotating. Subsequent star formation in the counter-rotating gas may then lead to counter-rotating stars. The two-stream instability between counter-rotating gas and co-rotating stars may enhance the rate of gas accretion (\\cite{RVEL}).   An important open problem is how counter-rotating gas disks form and what their structures are on galactic scales and on the scale of disks in active galactic nuclei. Here, we investigate accretion disks consisting of counter-rotating gaseous components with gas at large $z$ rotating with angular rate $+\\Omega(r)$ and gas at large negative $z$ rotating at rate $-\\Omega(r)$. The interface between the components at $z\\!\\sim\\! 0$ constitutes a supersonic shear layer and is shown in Figure \\ref{FIG1}.  \\begin{figure}[htb] \\begin{center} \\leavevmode \\epsfysize=1.9in \\epsfbox{Fig1.eps} \\end{center} \\caption{Structure of two apposed, counter-rotating accretion disks and the midplane boundary layer. The inset shows the three dimensional view of the velocity field for the $n=1/2$ case shown in Figure 2. The velocity variation is analogous to that in the Ekman layer of a rotating fluid (such as the ocean) where the Coriolis force balances the viscous force (\\cite{GKB}).} \\label{FIG1} \\end{figure}  \\noindent Configurations of this type may exist in astrophysical settings such as halo-disk interactions and the accretion of newly supplied counter-rotating gas onto an existing co-rotating disk. It might at first be supposed that powerful Kelvin-Helmholtz instabilities heat the gas to escape speed and rapidly destroy the assumed configuration. However, supersonic shear layers exist and exhibit gross stability in stellar and extra-galactic jets (\\cite{HARDEE}). In the counter-rotating disk, matter approaching the equatorial plane from above and below has angular momenta of opposite signs with the result that there is angular momentum annihilation at $z=0$, the matter loses its centrifugal support and accretes at essentially free-fall speed. On the other hand, accretion disks rotating in one direction are modeled assuming a turbulent viscosity which is crucial for the outward transport of angular momentum (\\cite{SS}). The counter-rotating disks can also be expected to be turbulent owing in part to the Kelvin-Helmholtz instability, and turbulent viscosity can transport angular momentum outward in the large $|z|$ regions of the disk. ",
        "conclusions": "Counter-rotating gas supplied to the outer part of an existing co-rotating gas disk will increase the mass accretion rate. Comparing accretion rates of a standard $\\alpha-$disk ($\\dot{M}_{SS}$) with that of a counter-rotating Keplerian disk ($\\dot{M}_{CR}$) of the same $\\Sigma$ and $M$, we find from equation (\\ref{MDOT}) $\\dot{M}_{CR}/ \\dot{M}_{SS} \\simeq 1 + 0.58(\\delta /\\epsilon^{2})$. For the values of Figure 2 this ratio is $ \\simeq 580$.  Thermal and/or dynamical instabilities may destroy the counter-rotating accretion flows described above. The relative importance of thermal and dynamical instabilities can be estimated by comparing the thermal dissipative time scale, $\\tau_{Q} \\approx \\Sigma c_{s}^2/D_{CR}(r) \\simeq (r/h)(H/V_{c})(c_{s}/V_{c})^2$, with dynamical and viscous time scales, $\\tau_{z}\\equiv h/c_{s},\\,\\tau_{\\perp}\\equiv r/V_{c}$, and $\\tau_{\\nu}\\equiv r^2/\\nu_{\\!\\perp}$. We find $\\tau_{Q}\\sim \\epsilon^2\\delta^{-3}\\tau_{z} \\sim \\epsilon^2 \\delta^{-5/2}\\tau_{\\!\\perp}\\sim \\epsilon^{4} \\delta \\tau_{\\nu}$. For $\\epsilon = 0.01$ and $\\delta = 0.1$, the thermal dissipation time $\\tau_{Q} \\ll \\tau_{z} <\\tau_{\\!\\perp} \\ll \\tau_{\\nu}$. In contrast to the standard $\\alpha-$disk, thermal equilibrium is maintained on time scales of possible flow instabilities of the inner shear layer.  The Kelvin-Helmholtz instabilities are likely to be the most important (\\cite{RAY} and \\cite{CL}).  If the counter-rotating disk is treated as a vortex sheet, then a local stability analysis indicates unstable warping for wave numbers $|k_\\phi/k_r|< \\sqrt2(c_s/\\Omega r) \\ll 1$.  Accretion of counter rotating gas by an existing co-rotating gas disk may be a transient stage in the formation of counter-rotating galaxies and in the accretion of matter onto rotating black holes in active galactic nuclei. We find that newly supplied counter-rotating gas drags inward the old co-rotating gas with an equal mass of old and new gas accreting rapidly. Thus the old co-rotating gas may be entirely ``used up'' (dragged to the center of a galaxy or into a black hole) if the mass of newly supplied gas exceeds that of the old gas disk. Accretion onto the faces of an existing thin disk may not have the symmetry shown in Figures \\ref{FIG1} and \\ref{PLOT}.  \\begin{figure}[htb] \\begin{center} \\leavevmode \\epsfysize=.82in \\epsfbox{Fig3.eps} \\end{center} \\caption{Schematic drawing of accretion of newly supplied counter-rotating gas ($\\odot$) induced by viscous interaction with an existing disk of co-rotating gas ($\\oplus$). } \\label{MERGER} \\end{figure}  \\noindent There may instead be two layers of rapid radial inflow bounding the existing gas disk near the midplane as sketched in Figure \\ref{MERGER}. Solutions for this configuration can be composed from those of Figure \\ref{PLOT} if $h \\ll H$.    \\vspace{5mm}"
    },
    "9604/hep-ph9604231_arXiv.txt": {
        "abstract": "In an Abelian gauge symmetry, spontaneously broken at a first order phase transition, we investigate the evolution of two and three bubbles of the broken symmetry phase. The full field equations are evolved and we concentrate in particular on gauge invariant quantities, such as the magnetic field and the integral around a closed loop of the phase gradient. An intriguing feature emerges, namely, the geodesic rule, commonly used in numerical simulations to determine the density of defects formed is shown not to hold in a number of circumstances. It appears to be a function of the initial separation of the bubbles, and the coupling strength of the gauge field. The reason for the breakdown is that in the collision region the radial mode {\\it can} be excited and {\\it often} oscillates about its symmetry  restoring value rather than settling to its broken symmetry value. This can lead to extra windings being induced in these regions, hence extra defects (anti-defects) being formed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604043_arXiv.txt": {
        "abstract": "In this paper we show that perturbations of the accretion flow within the central engines of  some active galactic nuclei (AGN) are likely to form shock waves in the accreting plasma. Such shocks, which may be either collisional or collisionless, can contribute to the observed high energy temporal and spectral variability. Our rationale is the following: Observations show that the continuum emission probably originates in an optically thin, hot plasma in the AGN central engine. The flux and spectrum from this hot plasma varies significantly over light-crossing timescales. Several authors have suggested that macroscopic perturbations contained within this plasma are the sources of this variability. In order to produce the observed emission the perturbations must be radiatively coupled with the optically thin hot matter and must also move with high velocities. We suggest that shocks, which can be very effective in randomizing the bulk motion of the perturbations, are responsible for this coupling. Shocks {\\it should} form in the central engine, because the temperatures and magnetic fields are probably reduced below their virial values by radiative dissipation. Perturbations moving at Keplerian speeds, or strong nonlinear excitations, result in  supersonic and superAlf\\'venic velocities leading to shock waves within the hot plasma.  We show that even a perturbation smaller than the emitting region can form a shock which significantly modifies the continuum emission in an AGN, and that the spectral and temporal variability from such a shock generally resembles those of radio quiet AGN. As an example, the shock inducing  perturbation in our model is a small main sequence star, the capturing and eventual accretion of which are known to be a plausible process. We argue that shocks in the central engine may also provide a natural triggering mechanism for the `cold' component of Guilbert \\& Rees two-phase medium and an efficient mechanism for angular momentum transfer. Current and future missions, such as Asca, XTE, XMM, AXAF and ASTRO-E may determine the importance of shock related emission from the central engines of AGN. ",
        "introduction": "Recent observations indicate that the continuum emission from AGN originates within a hot plasma (e.g., Mushotzky, Done \\& Pounds 1993, hereafter MDP93, and references therein). Variability arguments suggest that the plasma is contained within a region of size $R_X \\sim t_d c$, where $t_d$ is the flux doubling timescale. Within the commonly accepted accretion model for AGN, the hot plasma is probably located in the vicinity of a central supermassive black hole (Rees 1984 and references therein; Blandford 1990, hereafter B90). In this model, most of the continuum emission originates from within about twenty gravitational radii of the black hole in the region generally referred to as the AGN central engine.  An extensively considered accretion model of AGN, the simple optically thick geometrically thin disk model which does not include hot plasma, does not adequately explain the continuum UV to $\\gamma$ ray spectrum from AGN. Several modifications to the thin disk geometry have been proposed: (i) the geometrically thick accretion disk supported by radiation pressure (Abramowicz, Calvani \\& Nobili 1980) with a geometry which is susceptible to global perturbations (Papaloizou and pringle 1984), (ii) the geometrically thick, optically thin disk model with two temperature relativistic plasmas, first suggested for an X-ray binary (Shapiro Lightman \\& Eardley 1976) but which has been extended to AGN counterparts (e.g., Rees et al. 1982; White \\& Lightman 1989, hereafter  WL89; Tritz and Tsuruta 1989; and Narayan \\& Li 1994), and (iii) a disk-corona model (Liang \\& Price 1977; Haardt \\& Maraschi 1991; \\.Zycki, Collin-Soufrin \\& Czerny 1995, hereafter ZCC95; and Tsuruta and Kellen 1995, hereafter TK95). For the two-temperature torus model Narayan \\& Li (1994) propose that at lower accretion rates advection lowers the efficiency of the accretion process and stabilizes the otherwise unstable hot electrons, which suggests that this model may be more appropriate for sub-Eddington sources (see also Artemova et. al. 1996). A common feature of all models that attempt to reconstruct the higher energy ($\\sim 1$ keV -- $1$MeV) power-law spectrum is the introduction of an optically thin hot plasma (with the electron temperature $T_e> 10^8$K) which may dominate the emission from the central engine.  Large amplitude temporal variability of the hot plasma emission is present in at least 50\\% of radio quiet AGN, and significant spectral variability of similar timescales is also very common (MDP93 and references therein, Green 1993, McHardy 1988). The temporal and spectral variability has been generally explained in the context of perturbations on the thin, cold disk model, but rarely in the context of the geometrical models mentioned above. Models for the variability of radio quiet AGN include, for example, the hot spots model (Abramowicz 1991; Wiita 1993, hereafter W93), where  a multitude of randomly distributed flares or `hot spots' embedded within a thin disk are responsible for the total variability. Other models also assume randomly distributed perturbations on the thin disk (Pudritz \\& Fahlman 1982, DeVries \\& Kuipers 1989). The temperature of the thin disk in the central engine is found to be of order $T \\sim 10^6$K giving a sound speed of order $c_s \\sim 10^{-2}c$ or less (Frank, King \\& Raine 1992, hereafter FKR92). The total volume affected by {\\it {a single}} perturbation is therefore too small to easily account for the doubling of the entire continuum luminosity, making these models dependent on the existence of {\\it {many}} perturbations of characteristic size, time and brightness occurring on the disk surface. In contrast, within the hot plasma region, a disturbance may propagate at a much greater speeds, affecting a larger volume of the continuum emitting material. A single perturbation effectively coupled to the hot plasma may transfer enough energy to the radiating region within typical flux doubling timescales to account for the observed continuum variability.  Shocks have been investigated in the context of radio loud AGN with shocks propagating down a jet and interacting with irregularities in the jet material (Qian et al. 1991), or other forms of shocks (W93). In this paper we investigate the possibility that a strong perturbation, moving at supersonic or superAlf\\'venic speeds in the hot plasma of the AGN central engine, forms a shock which modifies the continuum emission (see also Sivron \\& Tsuruta 1994). Such shocks may effectively couple macroscopic perturbations with the optically thin hot matter, modifying the continuum spectrum. Shocks are also effective in randomizing the bulk motion of the perturbations which may have important consequences for the outward transfer of angular momentum.  In our work, we utilize recent studies of the effect of pairs on the central engine structure (Guilbert \\& Rees 1988, hereafter GR88; Lightman \\& White 1988, hereafter LW88; Coppi \\& Blandford 1990, hereafter CB90; Sivron \\& Tsuruta 1993, hereafter ST93; and Ghisellini \\& Haardt 1994, hereafter GH94). These authors suggest that in compact central engines, dampening processes readily cool a significant fraction of the hot matter, radiating the excess energy away. The cooled matter then forms a cold `phase' component. We show that in central engines with near-Eddington luminosity the post-shock matter is, in effect, an `increased compactness' region in which the shocked matter is cooled. The shocks are thus a natural mechanism by which radiative power is increased, and cold phase matter is created. The cooling of post shock matter results in radiation from the shock front which may be responsible for the `soft' flares, whereas cold phase matter behind the shock front may be responsible for the relatively hard `dips' in the light-curve (see section 4).  For mathematical simplicity we use a model in which the source of the perturbations is local, for example, a small main sequence star. Such a star can be captured, and eventually accreted into the central engine (Syer, Clarke \\& Rees 1990, hereafter referred to as SCR90). For a wide range of accretion parameters, such an event naturally results in the creation of shocks. Such shocks can significantly modify the light-curve from the X-ray emitting region of AGN.  In section 2 we show that  macroscopic inhomogeneities, which we assume travel at Keplarian speeds, naturally posses supersonic or superAlf\\'venic speeds when present in the hot plasma of the AGN central engine for the three models discussed above. In section 3 we show that the shocks effectively couple the perturbations with the hot plasma, and large enough perturbations can therefore result in the modification of the total emission. In section 4 we calculate the effects of a large shock on the light-curve of a model AGN. The main purpose of this exercise is to demonstrate the type of spectral variability expected when shocks modify the emission from the central continuum source. We discuss our results in section 5, and present our conclusions in section 6. ",
        "conclusions": "We have shown that shocks in the central engines of AGN form for a wide range of accretion parameters. We have also shown that perturbations in the central engine will effectivly radiate their energy away through shocks, and that for a wide range of accretion parameters that radiation will be observable. Given the generality of our results we believe that shocks should be fairly common and that their effects on observations should be significant.  Acknowledgments:  We thank the anonymous referee for some very helpful remarks. We thank Dr. M.J. Rees, Dr. N. Iwamoto and Mrs. N.Sivron for useful discussions and comments. We also thank Mr. W.McHargue and Mr. M.J.Kellen for their help with supporting software. This work was supported in part by NASA grants NAGW-2208 and NAG8-230 and NSF grant RII-8921978.  \\appendix"
    },
    "9604/astro-ph9604105_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the medium term, the cosmic microwave background (CMB) shall undoubtedly prove an extremely powerful tool in constraining cosmological parameters. Already, the {\\it COBE} observations of large-angle anisotropies provide the most accurate and unambiguous constraint on the spectra of perturbations in the universe. However, at the present time the best route to constraining cosmological parameters is not through the CMB alone, but from the combination of microwave data with a large number of measures of the power spectrum from large-scale structure observations.  In recent work, my collaborators and I have sought to test a wide parameter space of large-scale structure models, using linear and quasi-linear theory. We have written three papers on this topic, covering cold dark matter (CDM) models in open universes$^{1)}$ and in flat universes with a cosmological constant$^{3)}$, and the case of a critical density universe$^{2)}$ where we also allow a fraction of the dark matter to be hot. The key ingredient of our work is to take the inflationary hypothesis seriously, and to take advantage of the extra parameters that slow-roll inflation lends to large-scale structure modelling.  There isn't space here to give the full details of this work, so instead I'll concentrate on a couple of aspects and illustrate the outcome by showing some results from our investigation of CDM models in flat universes$^{3)}$. ",
        "conclusions": "Cosmologists are beginning to take seriously the possibility that one can determine the whole range of cosmological parameters. Within that context, one appreciates that it is possible to also include information from inflation, and attempt to fit for inflationary parameters at the same time as the cosmological parameters. The most popular inflationary paradigm, the slow-roll approximation, only introduces two extra parameters ($n$ and $R$) that one didn't have to consider anyway, and there is good reason to be optimistic that one can constrain these.  However, once one takes the extra inflationary input into account, it is clear that present observational data fall some way short of providing any telling constraints. We have found that it is possible to get an adequate fit to present data within almost any context. There are viable regions of parameter space for \\begin{itemize} \\item {\\bf CDM models$^{2)}$:} Requires some or all of low $h$, high $\\Omega_{{\\rm B}}$ or tilt to $n <1$. Gravitational waves don't help much, but they are not very strongly constrained. It is however very hard to fit the data for $h \\geq 0.50$. Adding extra massless species or decaying dark matter will also work though we haven't investigated them in detail ourselves. \\item {\\bf CHDM models$^{2)}$:} The same general picture as CDM models, but allows a higher value of $h$, at least up to 0.6, provided the amount of hot dark matter is chosen wisely. \\item {\\bf Low density CDM$^{1,3)}$:} Can be made to work either in the open case or in the flat case with a cosmological constant. Observationally, no strong preference between the open and flat cases. \\end{itemize} This situation should not remain for long. We stand at a tantalizing time, where observations are just good enough to exclude the more extreme inflationary models. We can look forward in the near future to a time when inflationary and cosmological parameters are extremely well determined, at which point we can expect most inflation models to be ruled out. Or maybe even all!  \\vspace*{12pt} \\noindent {\\bf Acknowledgments:} I thank my collaborators on work described herein, namely David Lyth, Dave Roberts, Bob Schaefer, Qaisar Shafi, Pedro Viana and Martin White. I also thank Pedro (again!) for producing the figure. I am supported by the Royal Society and acknowledge use of the Starlink computer system at the University of Sussex.  \\frenchspacing"
    },
    "9604/astro-ph9604111_arXiv.txt": {
        "abstract": "The discovery of HDF~J123652+621227, a candidate gravitational lens in the HDF, is reported.  This lens may be multiply imaging several optical sources at different redshifts.  If follow-up spectroscopy of the lens and the brightest image confirms this hypothesis, observations of this system alone can be used to obtain an estimate of the redshift distribution at extremely faint flux levels. ",
        "introduction": "Cosmologically distant galaxies ought to act as multiply imaging gravitational lenses for a fraction $\\sim0.002-0.005$ of background sources (Turner, Ostriker \\& Gott 1984; Blandford \\& Narayan 1992; Schneider, Ehlers \\& Falco, 1992). This prediction is being borne out by surveys of flat spectrum radio sources (Patnaik et al 1992; Myers et al 1995) and optical surveys (Maoz et al 1992; Glazebrook et al 1994; Ratnatunga et al 1995).  The incidence and character of strong gravitational lenses provide an important constraint on the source redshift distribution and world model at magnitudes too faint for direct spectroscopy (Kneib et al 1994; Kochanek 1992).  Recent Hubble Space Telescope (HST) observations of the Hubble Deep Field (HDF; Williams et al 1995) permit the optical lensing rate to be estimated in a uniform manner, using a single observation.  The HDF images are the deepest images in the visible ever taken: U, B, V and I-band images with point-source detection limits near 27, 29.5, 29.5 and 28.5~mag, respectively.  Approximately 2500 faint ``galaxies'' can be identified over 4~square arcmin, so the total number on the whole sky amounts to $\\sim9\\times10^{10}$~sources, a number roughly thirty times the product of the local bright galaxy density and the volume of an Einstein-de~Sitter Universe out to $z\\sim3$.  Possible explanations of this excess include fading (Babul \\& Rees 1992) or merging (Guiderdone \\& Rocca-Volmerange, 1991) of galaxies, counting multiple sub-galactic star formation sites within a common potential well as individual galaxies (Katz 1992; Colley et al 1996), or extreme cosmological models with large amounts of comoving volume per unit luminosity distance (Fukugita et al 1990).  An important means of distinguishing between explanations is to determine the redshift distribution of these sources.  In the HDF, $\\sim3-10$ cases of multiple imaging are expected (e.g., Turner et al 1984; Miralda-Escud\\'e \\& Lehar 1992); the actual number constrains the redshift distribution of the very faint sources observed in the HDF relative to brighter populations (such as quasars) for which both the lensing rate and the redshift distribution are better known.  This measurement can be used to constrain the redshift distribution of sources at much fainter levels than the limits of current spectroscopic surveys.  We have begun a partially automated search for multiply imaged sources derived from that used in the CLASS survey (Myers et al 1995).  After inspection of two dozen candidate lens systems, the most probable case was found to be \\HDF. This system consists of 16 components, all within 5~arcsec of a red, $F814W=23.3$~mag elliptical galaxy (component~0, hereafter c.~0) at RA\\,$12^h\\,36^m\\,52^s\\!\\!.01$, Dec\\,$+62^{\\circ}\\,12'\\,27''\\!\\!.3$ (J2000), identified as the lens candidate (Figure~1). The most striking of these companion images are a thin, $F606W=26.1$~mag tangential arc (c.~3) on one side and a $F606W=27.4$~mag counterimage (c.~1) on the other.  This configuration is seen in other gravitational lenses, e.g., FSC10214+4724 (Eisenhardt et al 1996).  It occurs when a source is located close to a cusp of the lens mapping in the source plane (Blandford \\& Narayan 1992; Schneider et al 1992). In addition to the arc and counterimage, there are a number of other faint sources surrounding c.~0 which may be multiply imaged.  If this hypothesis is confirmed by follow-up observations, this lens system alone will provide significant constraints on faint source redshift distributions, because the redshifts of all the lensed and unlensed components nearby c.~0 can be estimated or at least constrained with lens models. ",
        "conclusions": ""
    },
    "9604/astro-ph9604057_arXiv.txt": {
        "abstract": "We report the negative results of our searches in COMPTEL data for 1.809 \\MeV\\ gamma-ray line emission from four localized regions which contain nearby supernova remnants (SNRs). The upper flux limits (2$\\sigma$) are found to be in the range of $1.4 \\times 10^{-5}$ to $2.4 \\times 10^{-5}$ photons s$^{-1}$ cm$^{-2}$. These upper limits do not severely constrain the theoretical \\al26\\ yields from individual core collapse supernovae due to large uncertainties in the SNR distances and the nature of the progenitor stars. ",
        "introduction": "One of the outstanding achievements of the Compton Imaging Telescope (COMPTEL) aboard the {\\em Compton Gamma Ray Observatory} has been the first sky map in the light of the 1.809 \\MeV\\ \\gray\\ line which is attributed to radioactive decay of \\al26\\ ($\\tau=1.04\\times10^6$ yr). The observed emission is clearly of Galactic origin since it is concentrated on the Galactic plane. Its distribution along the plane is strikingly lumpy with extended emission features and `hot-spots' (\\cite{rf:drea95a}). One of these features, situated in the Vela region, is of particular interest (Diehl \\etal\\ 1995b). A recent survey of candidate \\al26\\ sources (core collapse supernovae (SNe), Wolf-Rayet (WR) stars, asymptotic giant-branch (AGB) stars, and O-Ne-Mg novae) in this region of the sky by Oberlack \\etal\\ (1994) identified the Vela supernova remnant (SNR) as most likely source of the emission if its progenitor star was massive ($\\sim35$ \\Msol) and its distance is $\\le350$ pc. The distance constraint is in line with recent X-ray based distance estimates of 350, 400-600, and 125-160 pc (Aschenbach 1993, Aschenbach et al. 1995, Becker 1995), respectively, all pointing towards distances below the canonical 500 pc. These findings raise the question if there are other SNRs detectable by COMPTEL as \\al26\\ sources. {}From the 1.8 \\MeV\\ \\gray\\ line sensitivity of $\\sim10^{-5}$ \\funit\\ (for an observation of $10^6$ seconds) and an optimistic \\al26\\ supernova yield of $3\\times10^{-4}$ \\Msol\\ for type II SNe (Hoffman \\etal\\ 1995) we estimate that SNe could be detectable by COMPTEL up to distances of $\\sim600$ pc. In this paper we report on our attempt to identify probable candidate SNRs and on the search for their 1.809 \\MeV\\ \\gray\\ line emission. ",
        "conclusions": "Observations of four nearby supernova remnants with COMPTEL do not provide evidence for 1.8 \\MeV\\ line emission from radioactive decay of \\al26. We obtain upper flux limits for the Cygnus Loop, HB 21, the Monoceros Nebula, and the Lupus Loop. Our limits are the most stringent to date, yet the large uncertainties in SNR distances and the nature of the progenitor stars do not allow to put severe constraints on \\al26\\ yields from individual supernovae. Based on current nucleosynthesis models, we found lower distance limits which are consistent with other distance estimates for the investigated SNRs."
    },
    "9604/astro-ph9604147_arXiv.txt": {
        "abstract": "Several models which have been constructed to explain the faint galaxy excess in observed number counts are used to predict the intensity of the extragalactic background light (EBL). Special attention is given to irregular and dwarf galaxies, which seem to be more common in the universe than once thought, and to low surface brightness galaxies (LSB), which can in principle be altogether missed from galaxy counts. The nature of the latter objects is still unclear, but some plausible models predict that LSB galaxies can increase the intensity of the EBL by a factor of up to 5 from a standard, no-evolution model in the optical and near infrared and by an order of magnitude in the UV.  If the faint excess population consists of low-luminosity dwarfs, whose luminosity function has a steep faint end, the EBL can well increase by a factor of 3 to 5, while still being consistent with current number count data. The resulting values of the EBL are not far from the observed upper limits. In the future the overall level of the EBL and its spectral distribution could be used to differentiate between galaxy population models. ",
        "introduction": "The importance of the extragalactic background light (EBL) for cosmology has long been recognized.  This integrated diffuse background radiation in the optical, ultraviolet and infrared wavebands contains information about otherwise difficult-to-observe or completely unobservable periods of the universe's past, particularly the era of galaxy formation.  The EBL may also be useful in discriminating between cosmological models.  For a review of the history of the subject, see Harrison (1990); and for both cosmology and galaxy evolutionary effects see Partridge \\& Peebles (1967) and the many papers by Tinsley (\\eg 1973, 1977).  In observational cosmology the nature of a background brightness measurement has in principle an advantage over the number count observations.  When counting galaxies, whether in magnitude or redshift bins, one needs to consider many kinds of selection effects which affect the completeness of the sample.  Measurements of the EBL are not plagued by this particular problem.  However, so far the EBL has not had much success as a cosmological probe or as a tool to investigate the evolution and origin of galaxies.  This is because the accurate elimination of the foreground components of the sky brightness has proved to be difficult and we lack a generally accepted measured value of the EBL (for a review of observational status see Mattila \\ea 1991).  And even if we had such a measurement, it would still not be easy to disentangle the roles of cosmological parameters, galaxy evolution and luminosity functions of galaxies (\\eg Tinsley 1973).  Although this work concentrates on optical wavelengths the treatment is essentially the same in the IR and UV (Franceschini \\ea 1991; Lonsdale 1995; Jakobsen 1995) which have recently been more active research areas than than the optical EBL. In the IR (see Franceschini \\ea 1991; Hauser 1995) analysis is currently being carried out on data from the Diffuse Infrared Background Experiment (DIRBE) on board COBE.  In the near future there will be additional EBL measurements  in the IR by Infrared Space Observatory (ISO) and later by the Space Infrared Telescope Facility (SIRTF). In the UV, recent observations and arguments by Sasseen \\ea (1995) indicate that the component formerly interpreted as extragalactic is in fact produced by galactic cirrus.  Their method utilized the power spectrum of background light; the result implies that the  {\\em optical} background fluctuations detected by Shectman (1973, 1974), using the same method, were also galactic in origin.   During the past years, much of observational cosmology has focused on deep galaxy-counts reaching ever fainter limits.  This in turn has produced new models for the galaxy population. The EBL has not generally been used as a further constraint on these models, because of the difficulties mentioned above.  However, it continues to be  a vital part of observational cosmology, especially in anticipation of near-future IR space observations.  The well-known apparent excess in the number counts of galaxies at faint magnitudes, most notably in the $B$-band, has led to numerous investigations as to the nature of this effect (\\eg Ferguson \\& McGaugh 1995, FMG95, and references therein). Over the years the discrepancy has been between the counts and the predictions made by using standard cosmology and no, or very modest, galaxy evolution.  The suggestions for solving the puzzle have included altering either the cosmology (\\eg introducing a non-zero cosmological constant) or the galaxy population (\\eg introducing new galaxy populations or altering the properties of giant galaxies via density or luminosity evolution.)  In the past year observations of faint galaxies have led to advances in the understanding to the faint excess and galaxy evolution. The HST Medium Deep Survey (MDS) provided evidence which indicates that the blue excess in number counts results from an excess population (relative to standard Hubble class -mixes) of late-type/irregular galaxies (Glazebrook \\ea 1995; Driver \\ea 1995a, 1995b; Casertano \\ea 1995).  Early results from the Hubble Deep Field also support the same conclusion (Abraham \\ea 1996). The Canada-France redshift survey (Lilly \\ea 1995a, 1995b) has also provided new data: results show a nearly unevolving early type population and a brightening LF of bluer galaxies. The mild evolution of elliptical galaxies was also found in an analysis of HST data (Im \\ea 1996). Finally, Cowie \\ea (1995) announced evidence of massive galaxies forming in the redshift range $z=1-2$ and Steidel \\ea (1996) at redshifts $z>3$.  In recent years there has also been cumulative evidence for a significant population of galaxies with very low surface brightnesses (\\eg Schombert \\ea 1992; de Blok \\ea 1995; Davies \\ea 1988; for a detailed rewiew of the field see especially McGaugh 1995 -- hereafter MG95); i.e.\\ surface brightnesses comparable to or fainter than the level of the night sky. It has been argued that the LSB population could actually {\\em be}, at least partially, the local counterpart of the faint blue population (McGaugh 1994). In this line of thought, a population of intrinsically LSB galaxies would have gone undetected in the local galaxy surveys due to selection effects; at the same time they would be more easily detected at larger distances in deep counts (which have much lower isophotal limits). There has been some work recently on quantifying the effect of observational selection criteria on the properties of observed galaxy populations; see Davies 1990; Yoshii 1993; Davies \\ea 1994; McGaugh \\ea 1995; FMG95; MG95.  The presence of LSB galaxies affects many areas of extragalactic astronomy. In particular, the luminosity functions hitherto derived from the local observable galaxies have a strong underrepresentation of LSB galaxies. As MG95 points out, even a small number of observed LSB galaxies implies a large underlying population because of the small volume sampling when detecting them.  There may also be a large population of dwarf galaxies escaping the magnitude limits of present surveys. The goal of this work is to quantify the effect of faint and low-surface-brightness galaxies on the extragalactic background light and to examine whether or not existing observational limits of the EBL constrain any proposed models of faint galaxy properties.  In addition, the basic ingredients which affect the surface brightness of the extragalactic component of the sky are reviewed. ",
        "conclusions": "The main results of the work can be summarized as follows:  (1) Models which predict number counts that are consistent with observations can have clearly different $I_{\\rm EBL}$ levels and spectral shapes.  (2) An EBL value of $I_{\\rm EBL} \\sim 1 \\cdot$ \\cggs appears to be of the correct order of magnitude. All the models considered without specific selection effects give a value at or just above this level in the optical and dropping off to $\\sim 0.4$ by $K$-band.  These could be divided roughly into two, the first group having a flat SED of EBL (in $f_{\\lambda}$) in the UV to blue and the second group having an enhanced UV EBL, up to 2 in the above units. Models with the $I_{\\rm EBL}$ dropping towards the UV (\\eg Yoshii \\& Takahara 1988) -- which would constitute a'third group' in the above distinction -- are found to be inconsistent with observed galaxy  counts.   (3) The distribution of EBL in different magnitude ranges exhibits a strong dependance of cosmology, galaxy population, and evolutionary model; universes producing same $I_{\\rm EBL}$ with different galaxy populations can in principle be seperated using both the SED of EBL and its intensity as a function of cut-off magnitude.  Most importantly, different models predict very different EBL levels beyond the current (and some future) magnitude limits.   (4) The situation changes when selection effects due to isophotal detection, low surface brightness effects, {\\em and} LSB galaxies are included. In principle the LSB galaxies could help to produce a very high-intensity EBL. Even considering more realistic ideas about the properties of LSB's, one can still produce an $I_{\\rm EBL}$ of about $2 - 3 \\cdot$ \\cggs. It is with the LSB models that the present upper limits (around 5--9 in the same units) of the observed EBL start providing constraints.  It is sobering to see how much the  large uncertainties in the surface brightness characteristics of galaxies can affect the EBL and the galaxy  counts.  If all galaxies could be seen, then the EBL and galaxy counts would not give independent results (apart from non-galactic contributions to the EBL; see below). However, especially if there is a large population of LSB galaxies, the EBL -- as a function of wavelength and cut-off magnitude -- provides a powerful tool for observational cosmology to complement galaxy counts and redshift distributions.  Furthermore, assuming that we have a measured value for the EBL and that the galaxy model predictions provide an accurate prediction, then the difference would account for any previously unknown sources of radiation, such as decaying particles or any radiation of intergalactic or/and pregalactic origin."
    },
    "9604/astro-ph9604001_arXiv.txt": {
        "abstract": "We investigate, in a model-independent way, the conditions required to obtain a satisfactory model of extended inflation in which inflation is brought to an end by a first-order phase transition. The constraints are that the correct present strength of the gravitational coupling is obtained, that the present theory of gravity is satisfactorily close to general relativity, that the perturbation spectra from inflation are compatible with large scale structure observations and that the bubble spectrum produced at the phase transition doesn't conflict with the observed level of microwave background anisotropies. We demonstrate that these constraints can be summarized in terms of the behaviour in the conformally related Einstein frame, and can be compactly illustrated graphically. We confirm the failure of existing models including the original extended inflation model, and construct models, albeit rather contrived ones, which satisfy all existing constraints. ",
        "introduction": "The extended inflation scenario \\cite{LS,K91} offers the prospect of resuscitating the original idea of Guth \\cite{G81} that inflation \\cite{LINDE,KT,LLrep} could be driven by a metastable vacuum energy and end by the tunnelling of the associated scalar field to the true minimum of its potential. The strategy is to implement inflation in an extended theory of gravity, such as a scalar--tensor theory, in which the expansion rate induced by a vacuum energy is slower than exponential. Under such circumstances, one is guaranteed that the phase transition will be able to successfully complete, which proves not to be the case in Einstein gravity if sufficient inflation is demanded to solve the usual cosmological problems.  Moving to an extended gravity theory is an interesting way of generalizing existing inflation models, because such scenarios are much more highly constrained \\cite{W93} than alternative generalizations where extra scalar fields are added by hand. In particular, one knows that the theory must mimic general relativity to a high degree at the present epoch, and there are further strong constraints on the variation of the strength of the gravitational interaction going all the way back to the time of nucleosynthesis. These additional constraints naturally have the effect of making it more difficult to obtain a viable model.  When one chooses to end inflation via a first-order phase transition, where bubbles of true vacuum nucleate, expand and coalesce, this introduces further constraints, because it is possible for the earliest true-vacuum bubbles which nucleate to be caught up in the subsequent inflationary expansion and stretched to astrophysically large sizes \\cite{W89,LSB}. These can contribute both density perturbations and microwave background anisotropies over and above those caused by quantum fluctuations \\cite{PERTEI} as in all inflationary models \\cite{PERT}. Near the general relativity limit, the distribution of bubbles is scale-invariant (in the sense of equal volume residing in bubbles within a given logarithmic size interval) which is far from acceptable \\cite{W89,LSB}.  The original extended inflation model \\cite{LS} was implemented in the Jordan--Brans--Dicke (JBD) theory of gravity \\cite{BD,W93}, where the gravitational `constant' is replaced by a field $\\Phi$ whose variation is controlled by a coupling parameter $\\omega$. General relativity is obtained in the limit of large $\\omega$. In that model, it was quickly shown that the competing needs of staying close to the general relativity limit to match present observations ($\\omega>500$ \\cite{Retal,W93}), and of obtaining a satisfactory bubble distribution ($\\omega \\lesssim 25$ \\cite{W89,LSB}), are mutually exclusive. This became known as the big-bubble problem, and various strategies have been brought into play in an attempt to evade it. The simplest is to introduce a mechanism which invalidates the present-day bound on $\\omega$; this can for example be achieved by introducing a potential for the Brans--Dicke field $\\Phi$ which is negligible during inflation and which prevents its variation at the present epoch. Alternatively, one can move to a general scalar--tensor theory, in which $\\omega$ is allowed to depend on $\\Phi$, which allows one to exercise control over how closely the general relativity limit is attained at different epochs.  Both these strategies have more recently suffered further constraints, under the assumption that the quantum fluctuations during inflation provide the density perturbations which are responsible for large-scale structure and microwave background anisotropies. The results of COBE in combination with large-scale structure studies quickly led to the conclusion that the spectrum of density perturbations must not be too far from scale-invariant, if the observed structures are to be reproduced. However, in extended inflation models one expects that if one makes the necessary moves to {\\em break} the scale-invariance of the bubble distribution, then one will also destroy the scale-invariance of the density perturbation spectrum. This implies two opposing constraints, but now both to be applied during inflation. This extra consideration proved sufficiently stringent to exclude all models in the existing literature in which inflation ends by nucleation \\cite{LL,LLEI}\\footnote{There are however models such as hyperextended inflation \\cite{SA}, in which inflation is brought to an end through dynamical evolution, with bubbles then nucleating in the post-inflationary phase.}.  In the literature, a substantial number of models falling into the extended inflation class have been devised \\cite{MODELS,SA,BM}, and examined on a more or less case by case basis \\cite{Yun,LW2}. In this paper, we shall place the constraints in a more general framework, allowing one to see easily the problems of existing models. As a by-product, this will enable us to construct working models satisfying all present constraints, though as we shall see the constraints combine in such a way as to make such models appear extremely contrived. The prognosis for the extended inflation scenario therefore continues to look poor. ",
        "conclusions": "We have carried out a model-independent analysis of the constraints on extended inflation scenarios containing a trapped scalar field within an arbitrary scalar--tensor theory. We have shown that all of the constraints can be interpreted graphically, by considering the behaviour of the (logarithm of the) Einstein frame potential as a function of the number of $e$-foldings from the end of inflation. The principal competition arises from the need to keep the density perturbation spectrum adequately scale-invariant while suppressing the production of bubbles which finish with astrophysically large sizes.  Especially bearing in mind that our implementation of the constraints is quite conservative, we have been able to show how difficult it is to obtain a successful extended inflation scenario. Indeed, all models of this type which presently exist in the literature, in which bubble nucleation ends inflation, cannot evade the combination of constraints. Things seem particularly tough if one desires a `blue' spectrum of perturbations, which is a situation in which one would have hoped extended inflation might have fared well since such models are hard to implement in the chaotic inflation framework. Our graphical approach allows one to see exactly what is needed to obtain working models; for example, any model in which $\\omega$ monotonically increases with time will not work. We have devised models which are allowed, though they seem rather contrived.  To conclude, the extended inflation paradigm is an attractive one, because the new physics, that of extended gravitational theories, can be tested in a number of ways. Unfortunately, when one also adds the extra constraints brought on by demanding that inflation ends by a first-order transition, the scenario becomes so highly constrained that it is extremely hard to find any working models. Still, the fact that one can exclude inflationary models on the basis of observational data should be viewed as an encouraging situation, and one we shall hear much of in future years."
    },
    "9604/astro-ph9604034_arXiv.txt": {
        "abstract": "A recent observation of Steidel et al. indicates that a substantial fraction of giant galaxies were formed at an epoch as early as redshift $z>3-3.5$. We show that this early formation of giant galaxies gives strong constraints on models of cosmic structure formation. Adopting the COBE normalization for the density perturbation spectrum, we argue that the following models do not have large enough power on galactic scales to yield the observed abundance: (i) standard cold dark matter (CDM) models (where mass density $\\Omega_0=1$ and power index $n=1$) with the Hubble constant $h\\lsim 0.35$; (ii) tilted CDM models with $h=0.5$ and $n\\lsim 0.75$; (iii) open CDM models with $h\\lsim 0.8$ and $\\Omega_0 \\lsim 0.3$, and (iv) mixed dark matter models with $h=0.5$ and $\\Omega _\\nu \\gsim 0.2$. Flat CDM models with a cosmological constant $\\lambda_0 \\sim 0.7$ are consistent with the observation, provided that $h\\gsim 0.6$. Combined with constraints from large-scale structure formation, these results imply that the flat CDM model with a low $\\Omega_0$ is the only one that is fully consistent with observations. We predict that these high-redshift galaxies are more strongly clustered than normal galaxies observed today. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604172_arXiv.txt": {
        "abstract": "A considerable experimental effort is underway to detect the `Doppler peaks' in the angular power spectrum of the cosmic microwave anisotropy. These peaks offer unique information about structure formation in the universe. One  key issue is whether structure could have formed by the action of causal physics within the standard hot big bang, or whether a prior period of inflation was required. Recently there has been some discussion of whether causal sources could reproduce the pattern of Doppler peaks produced by the standard adiabatic theory. This paper gives a rigorous definition of causality, and a causal decomposition of a general source. I present an example of a very simple causal source which  mimics the standard adiabatic theory, accurately reproducing the behaviour of the local intrinsic temperature perturbations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604108_arXiv.txt": {
        "abstract": "We introduce Void Hierarchy as an important property of the Large--Scale Structure in the Universe and demonstrate how it can be used to interpret observations. Moreover the void hierarchy constraints any realistic galaxy and structure formation scenario. ",
        "introduction": "Voids were defined as low density regions or, alternatively, as regions completely devoid of a certain type of object. Mean void diameters listed in Table~1 demonstrate the dependence of the void size on the type of object used in the (second) void definition. Both definitions imply that voids are not completely empty. Thus, the question is meaningful whether the distribution of galaxies in voids is homogeneous or reveals any structure. For example, it was concluded that Blue Compact Galaxies (BCG) from the Second Byurakan Survey (SBS) or other peculiar galaxies occur isolated within voids (Pustil'nik {\\it et al.} 1995). Such questions are very relevant concerning scenarios of large scale structure and galaxy formation, but they are not conclusively answered up to now.  \\vskip 0.3truecm \\noindent{\\bf Table~1} \\ {\\small Mean diameters of voids surrounded by different types of object} \\vskip -0.4truecm {{$$\\vbox {\\tabskip=0.05truecm \\halign to \\hsize { \\hfil# & \\hfil# &  \\hfill# & \\hfil#\\cr \\noalign {\\smallskip} \\noalign{\\hrule} \\noalign{\\smallskip} type of object \\qquad & \\qquad mean void diameter  \\cr \\noalign{\\smallskip} \\noalign{\\hrule} \\noalign {\\medskip} rich clusters (Abell/ACO--Catalogue)     & 100 \\Mmpc \\cr poor clusters (Zwicky--Catalogue)         &  37 \\Mmpc \\cr \\qquad bright ($M \\le -20.3$) elliptical galaxies &  30 \\Mmpc \\cr galaxies brighter than $M = -20.3$  &  23 \\Mmpc \\cr galaxies brighter than $M = -19.7$  &  16 \\Mmpc \\cr galaxies brighter than $M = -18.8$  &  13 \\Mmpc \\cr \\noalign{\\medskip} \\noalign{\\hrule} }}$$}} \\vskip -0.3truecm  Using the second void definition we have studied the properties of voids surrounded by galaxies from three different luminosity (absolute magnitude $M$) limited samples. Three void catalogues have been compiled. Comparisons of voids from different catalogues revealed that voids form a hierarchical system (cf. Lindner {\\it et al.} 1995, A\\&A 301, 329) as it is visualized in Fig.~1a). In this hierarchical concept apparently isolated galaxies in voids may have faint close neighbors which are not detected because of selection effects as it is shown in Fig.~1b).  \\begin{figure} \\epsfysize=8.5cm \\vskip -2.7truecm {\\epsffile{voids_fig1ab.ps}} \\vskip -0.5truecm \\caption{Wedge diagrams of a slice of the Universe 6000 \\kms\\ deep and bordered by about $9^h < \\alpha < 15^h$ and $49^\\circ < \\delta < 57^\\circ$. {\\bf a)} The three circles indicate an example of hierarchically interlaced voids defined by galaxies of different luminosity limit. {\\bf b)} Additionally BCGs from SBS are shown (crosses). The circle indicates the distance to the nearest bright ($M < -19.7$) neigboring galaxy.} \\end{figure} ",
        "conclusions": "By now the concept of void hierarchy is established only for galaxies brighter than $M = -18.8$ in the nearby Universe (up to distance $60 h^{-1}$Mpc). The study of the radial distribution of fainter galaxies in voids along with nearest neighbor tests (Lindner \\et\\ 1996) suggests that this hierarchy continues to fainter magnitudes and therefore contradicts a homogeneous distribution of dwarf ga\\-laxies in voids claimed by some theories of galaxy formation (e.g. Dekel \\& Silk 1986). With second generation instruments attached to the VLT (e.g. VIRMOS) it will be possible to confirm the hierarchy of voids towards fainter luminosity limits and for more distant regions of the Universe. The void hierarchy itself will be helpful to devise new concepts for the study of the large scale structure in the Universe."
    },
    "9604/astro-ph9604093_arXiv.txt": {
        "abstract": "We explore scenarios where the highest energy cosmic rays (HECR) are produced by new particle physics near the grand unification scale. Using detailed numerical simulations of extragalactic cosmic and $\\gamma$-ray propagation, we show the existence of a significant parameter space for which such scenarios are consistent with all observational constraints. An average fraction of $\\simeq10\\%$ $\\gamma$-rays in the total cosmic ray flux around $10\\,$EeV ($10^{19}\\,$eV) would imply both a non-acceleration origin of HECR and a large scale extragalactic magnetic field $\\lesssim10^{-11}\\,$G. Proposed observatories for ultra-high energy cosmic rays should be able to test for this signature. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604020_arXiv.txt": {
        "abstract": "We present an expression for the nonlinear evolution of the cosmological power spectrum based on following Lagrangian trajectories. This is simplified using the Zel'dovich approximation to trace particle displacements, assuming Gaussian initial conditions. The model is found to exhibit the transfer of power from large to small scales expected in self-gravitating fields. Some exact solutions are found for power--law initial spectra.  We have extended this analysis into redshift--space and found a solution for the nonlinear, anisotropic redshift--space power spectrum in the limit of plane--parallel redshift distortions. The quadrupole--to--monopole ratio is calculated for the case of power--law initial spectra. We find that the shape of this ratio depends on the shape of the initial spectrum, but when scaled to linear theory depends only weakly on the redshift--space distortion parameter, $\\beta$. The point of zero--crossing of the quadrupole, $k_0$, is found to obey a simple scaling relation and we calculate this scale in the Zel'dovich approximation.  This model is found to be in good agreement with a series of $N$-body simulations on scales down to the zero-crossing of the quadrupole, although the wavenumber at zero-crossing is underestimated. These results are applied  to the quadrupole--monopole ratio found in the merged QDOT plus 1.2 Jy {\\it IRAS} redshift survey. Using a likelihood technique we have estimated that the distortion parameter is constrained to be $\\beta>0.5$ at the $95 \\%$ level. The local primordial spectral slope is not well constrained, but the likelihood analysis suggests $n \\approx -2$ in the translinear regime.  The zero--crossing scale of the quadrupole is $k_0=0.5 \\pm 0.1 h{\\rm Mpc}^{-1}$ and from this we infer the amplitude of clustering is $\\sigma_8=0.7 \\pm 0.05$.   We suggest that the success of this model is due to nonlinear redshift--space effects arising from infall onto caustics and is not dominated by virialised cluster cores. The latter should start to dominate on scales below the zero--crossing of the quadrupole, where our model breaks down. ",
        "introduction": "The Newtonian analysis of linear growth of perturbations in an expanding universe is a well understood problem (Peebles 1980; Efstathiou 1990). The extension of this into the nonlinear regime has proven more difficult due to the strong mode coupling that arises in gravitational collapse, and most progress has been made through the use of $N$-body simulations.  Actual observations of galaxies in redshift space are further complicated (and made more interesting) by redshift distortions, caused by peculiar velocities of galaxies along the line-of-sight. Again, the linear problem is relatively well understood. Kaiser (1987) has shown that linear redshift distortions take their simplest form when expressed in Fourier space, at least if structure is far from the observer so that the distortions are essentially plane-parallel. Here the redshift--space Fourier modes are related to the real--space ones by \\be \\label{Kaiser} \\delta^s(\\k) = (1+\\beta \\mu_\\k^2) \\delta(\\k) \\ee where a superscript $s$ denotes a redshift--space quantity, $\\mu_\\k$ is the cosine of the angle between the wavevector $\\k$ and the line of sight, and $\\beta$, the redshift distortion parameter, is the dimensionless growth rate of growing modes in linear theory, which is related to the cosmological density $\\Omega$ by (Peebles 1980) \\be \\label{beta} \\beta \\equiv {\\Omega^{0.6} \\over b} \\ee in the standard pressureless Friedmann cosmology with mass-to-light bias $b$. It is through measuring the distortion parameter $\\beta$ that one hopes to measure the cosmological density parameter, $\\Omega$.   Kaiser's formula (\\ref{Kaiser}) is valid only in the linear limit, and for plane-parallel distortions, neither of which approximations is well satisfied in reality. Several authors have now addressed the problem of generalising Kaiser's formula to the case of radial distortions, while retaining the assumption of linearity (Fisher, Scharf \\& Lahav 1994; Heavens \\& Taylor 1995; Fisher et al. 1995; Zaroubi \\& Hoffman 1994; Ballinger, Heavens \\& Taylor 1995; Hamilton \\& Culhane 1996).   The main purpose of the present paper is to extend the analysis of redshift distortions into the nonlinear regime, retaining the plane-parallel approximation for simplicity. Our approach to the problem is motivated by the consideration that the density in redshift space may appear highly nonlinear even when the density in real space is only mildly nonlinear. For example, a region which in real space is just turning around, a mildly nonlinear condition, appears in redshift space as a caustic, a surface of infinite density, which is thoroughly nonlinear. This leads us firstly to work in Lagrangian space (Section \\ref{Lagsec}), and secondly to adopt the Zel'dovich (1970) approximation (Sections \\ref{Zeldsec} \\& \\ref{Redsec}). The Zel'dovich approximation is in effect linear theory expressed in Lagrangian space, inasmuch as it approximates the trajectories of particles as straight lines with (comoving) displacements growing according to linear theory.   Our approach follows that of Taylor (1993), who studied the nonlinear evolution of the power spectrum. Comparable approaches have been used by Bond \\& Couchman (1987) to evolve the galaxy angular correlation function, by Mann, Heavens \\& Peacock (1993) to evolve the real--space correlation function of clusters, and by Schneider \\& Bartelmann (1995) to evolve the real-space (unredshifted) power spectrum. Our approach differs from that of Hivon \\et (1995), who applied a perturbation expansion in Lagrangian space to second order to calculate the redshift--space skewness.   For simplicity, we assume throughout this paper that the density field is unbiased, $b = 1$. Generally, the effect of  evolution, if continuity is assumed (which is a fundamental assumption of this paper), is to tend to drive the bias factor towards unity. Continuity implies that the ratio of galaxy to matter density $(1+\\delta)/(1+\\delta_{\\rm M})$ remains constant in Lagrangian elements, and if (somehow) a linear bias $\\delta = b \\delta_{\\rm M}$ is established at some early time when $\\delta$ and $\\delta_{\\rm M}$ are both small, then the ratio of galaxy to matter density must be close to unity, $(1+b\\delta_{\\rm M})/(1+\\delta_{\\rm M}) \\approx 1$. It follows that the bias will be close to unity at later, nonlinear epochs when $\\delta_{\\rm M}$ is no longer small. Conversely, if bias is in fact important at the present, nonlinear epoch (as may well be the case), then it must be that the assumption of continuity must break down in the not too distant past. Indeed, it may be that continuity is violated on an ongoing basis. In the present paper we choose to ignore this thorny problem, and simply assume an unbiased density field.   We begin in Section \\ref{Realsec} by deriving equations which relate the Lagrangian and Eulerian descriptions of density, and we calculate a general expression for the evolution of the power spectrum. We invoke the Zel'dovich approximation and Gaussian initial conditions and consider some of the general features of the resulting nonlinear power spectrum. Some exact solutions for initially power--law spectra are derived. In Section \\ref{Redsec} we calculate the power spectrum in redshift--space, again in the Zel'dovich approximation, and obtain a number of analytic and numerical results for the observationally interesting ratio of quadrupole-to-monopole power. In Section \\ref{Obssec}, we compare the predictions of the Zel'dovich approximation with $N$-body simulations, and we apply our findings to analyse the quadrupole distortion measured in the QDOT plus 1.2 Jy redshift survey. We summarise our conclusions in Section \\ref{Sumsec}. ",
        "conclusions": "\\label{Sumsec}  The main aim of this paper has been to study the redshift distortion of the power spectrum in the moderately nonlinear regime. Our approach is motivated by the consideration that structures may appear in redshift space more nonlinear than they really are --- for example, regions which are just turning around in real space appear as caustics in redshift space. This leads us firstly to work in Lagrangian space, and secondly to use the Zel'dovich approximation, which is in effect linear theory expressed in Lagrangian space.  We started by deriving an expression relating the power spectrum to the Lagrangian displacement field. We used this to determine the evolution of the power spectrum in the Zel'dovich approximation, first in real space, then in redshift space in the plane-parallel approximation. We presented some analytic solutions for initially power-law spectra. In particular, we showed that a spectrum with index $n=-1$ gives rise to an incoherent Gaussian displacement field, producing nonlinear redshift distortions with the same form as the Kaiser-Peacock model.  We derived various analytic and numerical results for the observationally interesting ratio $R$ of quadrupole-to-monopole redshift power, whose value in the linear regime, $R_\\lin = (\\frac{4}{3}\\beta+\\frac{4}{7}\\beta^2)/ (1+\\frac{2}{3}\\beta+\\frac{1}{5}\\beta^2)$, yields a measure of the distortion parameter $\\beta$. In the Zel'dovich approximation, the amplitude of the ratio $R$ is set by the linear value $R_\\lin$, but the shape of $R$ as a function of wavenumber $k$ depends mainly on the spectral index $n$, and is insensitive to $\\beta$ at fixed $n$. The zero-crossing of the quadrupole power occurs at the point where the Zel'dovich power spectrum is a maximum as a function of time.  We have tested the Zel'dovich results against $N$-body simulations with initially power law spectra, in both $\\Omega=1$ and low-$\\Omega$ models. The simulations show that the Zel'dovich approximation ceases to provide a good approximation to the power spectrum at moderately nonlinear epochs. Remarkably however, the Zel'dovich approximation predicts rather well the amplitude and shape of the quadrupole-to-monopole ratio $R$ on scales down to the zero-crossing of the quadrupole, when $R$ is scaled to the scale of the zero-crossing. The Zel'dovich approximation underestimates the wavenumber at zero-crossing by a factor of about two, although it predicts correctly the way the zero-crossing scales with $\\beta$ and $n$.  We have applied these findings to estimate the distortion parameter $\\beta$, the local spectral index $n$, and the variance $\\sigma_8^2$ of counts in $8\\, h^{-1} {\\rm Mpc}$ spheres, from the quadrupole-to-monopole ratio $R$ measured in the merged QDOT plus 1.2 Jy redshift survey. We find that the distortion parameter is constrained to $\\beta>0.5$ at the 95\\% level. The spectral index is not well constrained, but is consistent with $n\\approx -2$ at translinear scales. The clustering amplitude $\\sigma_8$, inferred from the scale of the zero-crossing of $R$, is $\\sigma_8=0.7\\pm0.05$, consistent with other estimates.  The success of the Zel'dovich model in describing the quadrupole-to-monopole ratio $R$ suggests that departures from the linear value $R_\\lin$ at translinear scales are caused mainly by infall on to clusters, not by virialised cluster cores.   \\bigskip \\noindent{\\bf ACKNOWLEDGMENTS} \\bib\\strut  \\noindent ANT is supported by a PPARC research assistantship, and thanks Alan Heavens and John Peacock for stimulating discussion. AJSH thanks George Efstathiou for the hospitality of the Nuclear and Astrophysics Laboratory at Oxford University, where AJSH enjoyed the support of a PPARC Visiting Fellowship during 1994/5, and he thanks John Peacock for hospitality during a visit to ROE, where this collaboration began. AJSH appreciates support from NSF grant AST93-19977, and NASA Astrophysical Theory Grant NAG 5-2797.  While in the process of completing this work we became aware of a similar paper by Fisher \\& Nusser (1995) who reach similar conclusions.   \\bigskip \\noindent{\\bf REFERENCES} \\bib \\strut   \\bib Abramowitz M. and Stegun I.A., 1968, Handbook of Mathematical Functions, Dover, New York  \\bib Bond J.R., Couchman H.M.P., 1987, preprint  \\bib Ballinger W. E., Heavens A. F., Taylor A. N., 1995, MNRAS, 276, L59  \\bib Cole S., Fisher K.B., Weinberg D.H., 1994, MNRAS, 267, 785  \\bib Couchman H.M.P., 1991, ApJ, 368, L23  \\bib Efstathiou G., 1990, in Peacock J.A., Heavens A.F., Davies A., eds, Proc. 36th  Scottish Universities Summer School in Physics, {\\em Physics of the Early Universe}, Adam Hilger, Bristol  \\bib Fisher K.B., Scharf C.A., Lahav O., 1994, MNRAS, 266, 219  \\bib Fisher K.B., Lahav O., Hoffman Y., Lynden--Bell D., Zaroubi S., 1995, MNRAS, 272, 885  \\bib Fisher K.B., Nusser A., 1995, preprint  \\bib Gorski K.M., 1988, ApJ Lett, 332, L7  \\bib Hamilton A.J.S., 1992, ApJ, 385, L5  \\bib Hamilton A.J.S., 1995, in {\\em Clustering in the Universe}, Proc. XVth Rencontres de Moriond, in press.  \\bib Hamilton A.J.S., Culhane M., 1996, MNRAS, 278, 73  \\bib Heavens A., Taylor A.N., 1995, MNRAS, 275, 483  \\bib Hivon E., Bouchet F.R., Colombi S., Juszkiewicz R., 1995, A.A., 298, 643  \\bib Hui L., Bertschinger E., 1995, preprint  \\bib Jackson J.C, 1972, MNRAS, 156, 1p  \\bib Kaiser N., 1987, MNRAS, 227, 1.  \\bib Mann R.G., Heavens A.F., Peacock J.A., 1993, MNRAS, 263, 798  \\bib Monin A.S., Yaglom A.M., 1975, {\\em Statistical Fluid Mechanics}, Vols. 1 \\& 2, MIT Press, Cambridge.  \\bib Peacock J.A., 1992 in Martinez V., Portilla M., S\\'{a}ez D., eds, {\\em New Insights into the Universe}, Proc. Valencia summer school. Springer,Berlin, p.1  \\bib Peacock J.A., Dodds S.J., 1994, MNRAS, 267, 1020  \\bib Peebles P.J.E., 1980, {\\sl Large Scale Structure in the Universe}, Princeton University Press  \\bib Schneider P., Bartelmann, M, 1995, MNRAS, 273, 475  \\bib Taylor A.N., 1993, in Proc. Cosmic Velocity Fields, 9$^{th}$ IAU Conf, eds F. Bouchet M. Lachieze-Rey, Edition Fronti\\`{e}re, Gif-sur-Yvett, p585  \\bib Zaroubi S., Hoffman Y., preprint 1994  \\bib Zel'dovich Y.B., 1970, A\\&A, 5, 84."
    },
    "9604/astro-ph9604179_arXiv.txt": {
        "abstract": "Deep {\\it Hubble Space Telescope} (HST) observations with WFPC2 of the nearby globular cluster NGC 6752 have allowed us to obtain accurate photometry for the cluster white dwarfs (WD). A sample of local WDs of known trigonometric parallax and mass close to that of the cluster WDs have also been observed with WFPC2. Matching the cluster and the local WD sequences provides a direct  measure of the distance to the cluster: $(m-M)_\\circ=13.05$, with an uncertainty less than $\\pm0.1$ mag which allows a substantial reduction in  the uncertainty in the  age of the cluster. Indeed, coupling this value of the cluster distance to the cluster metallicity, helium abundance and $\\alpha$-element enhancement [$\\alpha$/Fe]=0.5 yields an age of 15.5 Gyr and 14.5 Gyr using evolutionary models that do not include or do include helium diffusion, respectively. The uncertainty affecting these age determinations is  $\\sim 10\\%$. The majority of the cluster WDs appear to be of the DA variety, while the color-magnitude location of two WDs is consistent with the DB type. This suggests a cluster DB/DA ratio  similar to that of WDs in the solar neighborhood. ",
        "introduction": "The age of the universe $\\tz$ is the obvious partner of the Hubble constant $\\hz$. Together they set a constraint on $\\oz$ if we believe the cosmological constant $\\Lambda$ to be zero, or on a combination of $\\oz$ and $\\Lambda$ if one is willing to accept  $\\Lambda\\ne 0$ cosmologies. By general consensus, globular cluster ages provide  potentially the most accurate estimate of $\\tz=\\tgf+\\tgc$, $\\tgf$ being the age of the universe when the Galaxy formed and $\\tgc$  the present age of galactic globular clusters. Since presumably $\\tgf\\simeq 1-2\\,{\\rm Gyr}\\ll\\tz$, then $\\tz\\approx\\tgc$, and  $\\tgc$  provides a strict lower bound to $\\tz$. The age of Galactic globular clusters can be most accurately estimated by using the theoretical relation between age and the luminosity of the main sequence turnoff (TO), other methods being undermined by uncontrollable systematic errors (Renzini 1991, 1993). For example, one can use a relation that fits the isochrones of VandenBerg \\& Bell (1985): $$\\log t_9\\simeq -0.41 + 0.37\\,\\mvto - 0.43\\, Y - 0.13\\,\\feh,\\eqno(1)$$ where $t_9$ is the age in Gyr units, $Y$ the helium abundance, $\\feh$ the iron abundance in standard notations, and $\\mvto$ the TO absolute visual magnitude. In turn,  $\\mvto = V^{\\rm TO} -$ mod, where $V^{\\rm TO}$ (the TO apparent magnitude) is the directly {\\it observable} quantity, and mod is the cluster distance modulus. This relation allows one to estimate the relative importance of the uncertainty in each of the four input quantities ($V^{\\rm TO}$, mod, $Y$, and [Fe/H]) in establishing the final uncertainty in the age determination. The current distances are typically affected by a $\\sim$1/4 magnitude error in the modulus -- $\\sigma($mod$)\\simeq 0.^{\\rm m}25$ -- which immediately translates into a $\\sim 22\\%$ error in the derived cluster age ($\\sim 3$ Gyr for an age of 15 Gyr). All other input quantities convey substantially smaller errors. The high photometric accuracy of CCDs now allows one to determine a cluster's $V^{\\rm TO}$ with an accuracy better than $0^{\\rm m}.1$, which translates into a $\\sim 9\\%$ error in age. The helium abundance is very well known, from either the R method, primordial nucleosynthesis, or empirical determinations of the {\\it pregalactic} abundance, which all indicate $Y=0.23-0.24$ (e.g., Boesgaard \\& Steigman 1985), and even a $\\pm0.02$ uncertainty in $Y$ gives a negligible 2\\% error in age. The metal content of the best studied clusters is uncertain by $\\sim$0.3 dex (most of it being systematic), which translates into a $\\sim 9\\%$ uncertainty in age. There is a problem with the {\\it composition} of metallicity (e.g. enhanced [O/Fe], or [$\\alpha$/Fe]), a point to which we shall return in Section 4. Clearly the first concern is the error in the distance of the clusters, and it is therefore instructive to recognize that distance determinations dominate the error budget not just of the {\\it kinematical} age of the universe (via $\\hz$) but also of globular cluster ages. For the comparison of the two ages to be  unambiguous, the error in each of them must be  reduced as much as possible. The {\\it Hubble Space Telescope} (HST) {\\it Key Project} is aimed at achieving  $\\sim 10\\%$ accuracy on $\\hz$ (Kennicut, Freedman, \\& Mould 1995). We report  here  our own attempt at using HST observations to achieve similar accuracy on $\\tgc$.  Using ground based observations, the distance to globular clusters has been  estimated with either the RR Lyrae or the subdwarf methods. Their limitations are extensively discussed by e.g., Sandage \\& Cacciari (1990) and Renzini (1991, 1993). Suffice it to mention here that both methods are semi-empirical in nature, relying heavily on theoretical models (e.g., pulsational, atmosphere, and stellar models), and both require the metallicity of the calibrating stars and of the clusters to be measured. Hence, the resulting estimate of the distance  is  affected by both systematic errors that are difficult to quantify, and by errors in metallicity which can dominate the  age error budget. In this paper we present the first attempt at determining the distance to a globular cluster by using the white dwarf (WD) method (Renzini 1991 and ref.s therein), that is essentially free from these limitations. ",
        "conclusions": ""
    },
    "9604/astro-ph9604098_arXiv.txt": {
        "abstract": " ",
        "introduction": "Shortly after the cosmic microwave background (CMB) was discovered~\\cite{Penzias} it became clear that this universal radiation field has profound implications for the astrophysics of ultrahigh energy cosmic rays (UHE CR) of energies above $10^{18}{\\,{\\rm eV}}$. For nucleons the most profound effect is photoproduction of pions on the CMB. Known as the Greisen-Zatsepin-Kuz'min (GZK) ``cutoff''~\\cite{GZK,Stecker1}, this effect leads to a steep drop in their energy attenuation length by about a factor 100 at around $6\\times10^{19}{\\,{\\rm eV}}$ which corresponds to the threshold for this process. The nucleon attenuation length above this threshold is about $10{\\,{\\rm Mpc}}$. Heavy nuclei with energies above about $10^{19}{\\,{\\rm eV}}$ are photodisintegrated in the field of the CMB within a few ${\\,{\\rm Mpc}}$~\\cite{psb}. One of the major unresolved questions in cosmic ray physics is the existence or non-existence of a cutoff in the UHE CR spectrum at a few $10^{19}\\,$eV which, in the case of extragalactic sources, could be attributed to these effects.  Therefore, there has been renewed interest in UHE CR research since events with energies exceeding $10^{20}{\\,{\\rm eV}}$ have been detected. The Haverah Park experiment~\\cite{Watson} reported several events with energies near or slightly above $10^{20}{\\,{\\rm eV}}$. The Fly's Eye experiment~\\cite{FE1,FE2} detected the world's highest energy CR event to date, with an energy $\\simeq3\\times10^{20}{\\,{\\rm eV}}$. Near the arrival direction of this event the Yakutsk experiment~\\cite{Efiego} recorded another event of energy $\\simeq1.1\\times10^{20}{\\,{\\rm eV}}$. More recently, the AGASA experiment~\\cite{AG1,AG2} has also reported an event with energy $1.7-2.6\\times10^{20}{\\,{\\rm eV}}$. It is currently unclear whether these events indicate a spectrum continuing beyond $10^{20}{\\,{\\rm eV}}$ without any cutoff or the existence of a cutoff followed by a recovery in the form of a ``gap'' in the spectrum~\\cite{slsb}.  There has been much speculation about the nature and origin of these highest energy cosmic rays (HECRs)~\\cite{Hillas,Sorrell,Sommers,Elbert,SSB}. Concerning the production mechanism one can distinguish between two broad classes of models: Within acceleration models, charged primaries, namely protons and heavy nuclei are accelerated to very high energies~\\cite{Blandford,Gaisser} in a ``bottom-up'' manner. Preferred sites are large-scale astrophysical shocks which occur for instance in radio galaxies~\\cite{Biermann}. Even there it seems barely possible to accelerate CRs to the required energies~\\cite{Hillas,Cesarsky,Norman}. Recently it has also been suggested that acceleration of UHE CRs could be associated with cosmological gamma-ray bursts (GRBs)~\\cite{Wax1,Vietri,Wax2}. In the second class of so called ``top-down'' models, charged and neutral primaries are produced at UHEs in the first place, typically by quantum mechanical decay of supermassive elementary X particles related to grand unified theories (GUTs). Sources of such particles at present could be topological defects (TDs) left over from early universe phase transitions caused by the spontaneous breaking of symmetries underlying these GUTs~\\cite{Hill,HSW,Bh0,Bh1,BR,bhs,BS,Sigl1}. The injection spectra in top-down models tend to be considerably harder (flatter) than in acceleration models.  The particle identity of the UHE CRs is not known either. The Fly's Eye analysis~\\cite{FE1} suggested a transition from a spectrum dominated by heavy nuclei to a predominantly light composition, i.e. nucleons or even $\\gamma$-rays, above a few times $10^{19}{\\,{\\rm eV}}$. However, this has not been confirmed by the AGASA experiment~\\cite{AG3}. Although there have been claims that the shower profile of the highest energy Fly's Eye event may be inconsistent with a primary photon~\\cite{hvsv} or even with a proton primary~\\cite{Gaisser:pc}, the situation is not settled because of many uncertainties which can affect the shower development in the atmosphere.  Other options discussed for the nature of the HECRs include heavy nuclei and even neutrinos~\\cite{hvsv}. Heavy nuclei have their own merits because they can be deflected considerably by the galactic magnetic field which relaxes the source direction requirements \\cite{FE1,AG1}. In addition, for shock acceleration, heavy nuclei can be accelerated to higher terminal energies because of their higher charge. However, one should note that the range for heavy nuclei is limited to a few Mpc as mentioned above. Neutrinos, on the other hand, do not lose much energy over cosmological distances~\\cite{Weiler,Yoshida}, but by the same token the probability for interacting in the atmosphere is small. Attributing the HECRs to neutrinos would therefore require a neutrino flux at UHEs which is much higher than the observed CR flux at the same energies. This poses severe constraints on the possible sources for these neutrinos~\\cite{sl}. In addition, neutrinos would be expected to give rise to predominantly deeply penetrating showers in the atmosphere.  The production spectrum of UHE CRs is modified during their propagation. There are many studies on nucleon propagation in the literature using analytical~\\cite{bg,akv,Rachen,Geddes} as well as numerical approaches~\\cite{Elbert,hs,yt,ac}, and the propagation of heavy nuclei has also been considered~\\cite{Elbert}. This was mainly motivated by the conventional acceleration models which usually predict UHE CR fluxes to be dominated by these particles. However, secondary $\\gamma$-rays and neutrinos can also be produced, for example as decay products of pions created by interactions with various radiation backgrounds at the source or during propagation~\\cite{yt}. Under certain circumstances their flux can become comparable with the primary flux~\\cite{Wolfendale1}. Furthermore, within TD models $\\gamma$-rays are expected to dominate to begin with~\\cite{abs}. A study on $\\gamma$-ray propagation in this context has been performed recently~\\cite{pj} using a quantitative treatment on the cascade initiated by UHE photons. In my opinion, however, it suffers from several unrealistic assumptions with respect to the injection scenarios considered. I improve on their treatment of the propagation of $\\gamma$-rays. Apart from that I find three reasons to explore UHE $\\gamma$-ray propagation in more detail in this paper:  First, due to the absence of threshold effects similar to photopion production which causes the GZK ``cutoff'' for nucleons, the $\\gamma$-ray spectrum is not expected to have a break around $10^{20}{\\,{\\rm eV}}$. Furthermore, $\\gamma$-rays can generate electromagnetic (EM) cascades while propagating rather than being absorbed right away. UHE electrons produced by pair production upscatter background photons and transfer most of the energy back to photons. This effect considerably increases the effective energy attenuation length of the ``cascade'' photons~\\cite{los,ls}. At a few times $10^{20}{\\,{\\rm eV}}$ this attenuation length may be even greater than that for protons which drops precipitously at the threshold for photopion production. Extragalactic $\\gamma$-rays could therefore have some potential to produce a recovery beyond the GZK ``cutoff''.  Second, in contrast to the case of nucleons, the propagation of $\\gamma$-rays is presently fraught by certain ambiguities which are mainly due to uncertainties in the intensity of the universal radio background and the strength and spectrum of the extragalactic magnetic field (EGMF). We hope that an application of the general framework presented here under different assumptions for such parameters could in turn provide some insights into their actual values once the UHE $\\gamma$-ray flux is known to some accuracy. This would be in some analogy to the method of using TeV $\\gamma$-ray observations to constrain or detect the universal infrared/optical background~\\cite{Stecker3}. In previous work~\\cite{los} it is shown that, depending on its strength, the large-scale EGMF could produce a feature in the $\\gamma$-ray spectrum which might be observable in the future.  Finally, the study of high energy cosmic and $\\gamma$-ray propagation can place stringent constraints on the nature and origin of UHE CRs. Such constraints can be obtained by computing the propagation modified spectra especially of lower energy $\\gamma$-rays expected within a certain scenario and comparing the predictions with the observed fluxes~\\cite{Wolfendale1,hpsv,SJSB}. At UHEs there are some experimental prospects to distinguish $\\gamma$-rays from other primaries in the future, possibly even on an event by event basis~\\cite{Cronin2}. This would allow comparing not only the total fluxes of UHE nucleons, heavy nuclei, and $\\gamma$-rays, but also their composition with model predictions.  This motivated the present comprehensive study of propagation of nucleons and $\\gamma$-rays and its application to models which attribute UHE CRs to top-down mechanisms within GUT-scale physics or associate them with cosmological GRBs. I explore the energy range of $10^8 < E < 10^{23} {\\,{\\rm eV}}$. The low end is chosen such that we can draw constraints by comparing the propagated spectra with existing measurements of the diffuse $\\gamma$-ray background around $100{\\,{\\rm MeV}}$~\\cite{Digel,Fichtel,Osborne}. The high end is chosen beyond the highest CR energies ever observed enabling us to study top-down models. I include not only the CMB but also the diffuse radio background which plays a big role at the highest energies and the infrared/optical (IR/O) background which influences the flux at somewhat lower energies. I also include the EGMF as a free parameter. The propagation of nucleons is also studied with special emphasis on the production of secondary $\\gamma$-rays, electrons, and neutrinos.  The rest of the paper is organized as follows: In Section 2, I present the general ingredients of calculating the propagation of extragalactic $\\gamma$-rays and nucleons. I discuss the role and nature of the low energy photon background and the EGMF, and explain in detail the implicit method used in solving the transport equations numerically. Section 3 is devoted to the treatment of the relevant interactions of $\\gamma$-rays and nucleons. I compare our analysis with other work in Section 4. Section 5 discusses the generic forms of the injection spectra and the source distribution for typical top-down models and the GRB scenario. Results and constraints from the spectra predicted at Earth are presented in detail. In Section 6, I summarize the findings and discuss future prospects. ",
        "conclusions": "I have performed detailed numerical simulations for the propagation of extragalactic nucleons, $\\gamma$-rays, and electrons in the energy range between $10^8{\\,{\\rm eV}}$ and $10^{23}{\\,{\\rm eV}}$. My goal was to explore constraints on various models of HECR origin from a comparison of predicted and observed $\\gamma$-ray fluxes at lower energies. The main focus thereby is on models which associate HECRs with GUT scale physics or with cosmological GRBs.  I find that at present the TD scenarios are primarily constrained by the observed $\\gamma$-ray background between $\\simeq 100{\\,{\\rm MeV}}$ and $\\simeq 10{\\,{\\rm GeV}}$ but not by the limit on the $\\gamma$ to charged CR flux ratio below $100\\,{\\rm TeV}$. The CEL approximation usually does not take the IR/O background into account, and thus may not be directly compared to the numerical calculation because the presence of the IR/O background may affect the $\\gamma$-ray flux level at $100\\,{\\rm MeV}$ by an order of magnitude. There is also a significant difference for the UHE spectrum between predictions by the CEL approximation and my numerical simulation. For an EGMF strength $\\la 10^{-11}\\,{\\rm G}$ the TD models yield the $\\gamma$-ray flux which is at about the same level as or below the current observed flux, depending on the adopted parameters. On the other hand, an EGMF stronger than $\\sim 10^{-11}\\,{\\rm G}$ stops the cascade at UHEs and the UHE end of the spectrum is suppressed significantly. Thus, the level of the $\\gamma$-ray flux at about $100\\,{\\rm MeV}$ is higher relatively, tightening the constraints. However, these results are rather insensitive to different models of the IR/O background \\cite{coaha}, although they are somewhat dependent on the poorly known universal radio background flux.  I conclude that TD scenarios with QCD motivated injection spectra up to energies $\\la10^{23}{\\,{\\rm eV}}$ are still viable if injection occurs uniformly or from a discrete source. This is in contrast to a recent claim in the literature~\\cite{pj}. In case of uniform injection this assumes an injection history motivated by energy release from a network of cosmic strings in the scaling regime or from monopole-antimonopole annihilation ($p=1$). Higher injection energy cutoffs are allowed for either a weaker source evolution or for injection spectra somewhat steeper than the QCD motivated spectra. For EGMF strengths larger than $\\simeq10^{-10}{\\,{\\rm G}}$, some of the predicted TD spectra have the potential to explain a possible gap in the HECR spectrum. The cosmological GRB scenarios recently suggested in the literature~\\cite{Wax1,Vietri,Wax2} are currently unconstrained by these limits.  With the arrival of the anticipated Pierre Auger Cosmic Ray Observatories~\\cite{Cronin2}, it is expected that the UHE end of the CR spectrum will be known with much better accuracy. Constraints derived from the influence of CR propagation on the observed spectrum will then be one of the most powerful tools in discriminating between models of HECR origin."
    },
    "9604/astro-ph9604051_arXiv.txt": {
        "abstract": "The velocity dispersion of galaxies on small scales ($r\\sim1h^{-1}$ Mpc), $\\sigma_{12}(r)$, can be estimated from the anisotropy of the galaxy-galaxy correlation function in redshift space (Davis \\& Peebles 1983). We apply this technique to ``mock-catalogs'' extracted from N-body simulations of several different variants of Cold Dark Matter dominated cosmological models, including models with Cold plus Hot Dark Matter, to obtain results which may be consistently compared to similar results from observations. We find a large variation in the value of $\\sigma_{12}(1 h^{-1} Mpc)$ in different regions of the same simulation.  We investigate the effects of removing clusters from the simulations using an automated cluster-removing routine, and find that this reduces the sky-variance but also reduces the discrimination between models. However, studying $\\sigma_{12}$ as clusters with different internal velocity dispersions are removed leads to interesting information about the amount of power on cluster and subcluster scales. We compute the pairwise velocity dispersion directly in order to check the Davis-Peebles method, and find agreement of better than 20\\% in all the models studied.  We also calculate the mean streaming velocity and the pairwise peculiar velocity distribution in the simulations and compare with the models used in the Davis-Peebles method. We find that the model for the mean streaming velocity may be a substantial source of error in the calculation of $\\sigma_{12}$. ",
        "introduction": "The velocity dispersion of galaxies on small scales ($r\\sim 1 h^{-1}$ Mpc), combined with cluster abundance data on intermediate scales and the COBE normalization and galaxy peculiar velocity information on large scales, provides a strong constraint on cosmological models by constraining the shape of the matter power spectrum. In this paper we investigate a method introduced by Peebles (1976; 1980) and Davis \\& Peebles (1983, hereafter DP83), which uses the anisotropy of the redshift-space correlation function to determine the pairwise velocity dispersion on small scales. We shall refer to this method as the Davis-Peebles method.  The galaxy-galaxy correlation function  $\\xi(r)$ is one of the canonical statistics used in studying large scale structure. A related statistic is the redshift-space correlation function, $\\xi(r_{p}, \\pi)$, which is a function of the components of the separation in redshift space perpendicular ($r_{p}$) and parallel ($\\pi$) to the line of sight. If the correlation function is isotropic in real space, it will be anisotropic in redshift space due to the peculiar velocities of the galaxies. Hence the degree of anisotropy of $\\xi(r_{p}, \\pi)$ is a measure of the moments of the peculiar velocity distribution.  The first moment of the pairwise velocity distribution, $\\overline{v_{12}}(r)$, is proportional to $\\Omega_{0}^{0.6}$ if galaxies trace mass and density fluctuations are in the linear regime, where $\\Omega_{0}$ is the density of matter in units of the critical density at the present epoch.  The second moment, $\\sigma_{12}$, is the velocity dispersion and measures the kinetic energy of the galaxy distribution. This quantity has been used in combination with the Cosmic Virial Theorem to estimate $\\Omega_{0}$ (DP83), although recently some authors have presented arguments that this calculation is not only plagued by extremely large uncertainties from cosmic variance (Fisher et al. 1994b), but that some fundamental assumptions in the usual formulation of the Cosmic Virial Theorem may also be incorrect (Bartlett \\& Blanchard 1995).  The first calculation of $\\sigma_{12}$ on a relatively large survey was done by Davis and Peebles (DP83). They calculated $\\sigma_{12}$ for the CfA1 redshift survey, a survey containing 1840 redshifts covering 1.83 steradians in the North galactic hemisphere (Huchra et al.  1983). In another paper (Somerville, Davis, \\& Primack 1996, hereafter SDP), we present a reanalysis of their work, in which we show that $\\sigma_{12}$ for this sample is extremely sensitive to the way in which corrections for infall into the Virgo cluster are applied, and that the value of $\\sigma_{12}$ for this survey is dominated by the clusters. The same calculation was done on the Southern Sky Redshift Survey(SSRS1) (da Costa et al. 1991), with results of $\\sigma_{12}(1 h^{-1} Mpc) \\equiv \\sigma_{12}(1) \\sim 300$ km/s (SDP; Davis 1988). Recently $\\sigma_{12}$ has also been calculated for the IRAS survey (Fisher et al. 1994b), the CfA2/SSRS2 survey (Marzke et al. 1995), and the Perseus-Pisces survey (Guzzo et al. 1995). These calculations have shown a range of values of $\\sigma_{12}(1)$ from 272 km/s for SSRS2 to 769 km/s for Perseus-Pisces (see Table 1 of SDP for a summary).  The values of $\\sigma_{12}$ usually quoted for simulations are calculated by measuring the dispersion of the pairwise peculiar velocity field directly using the full three-dimension position and velocity information for the halos (Davis et al. 1985; Gelb \\& Bertschinger 1994; Klypin et al. 1993). In real redshift surveys, not only are there errors introduced by edge effects and selection effects, but $\\sigma_{12}$ is extracted by fitting a model to the correlation function in redshift space, $\\xi(r_{p}, \\pi)$.  This quantity is quite noisy especially for samples with small numbers of galaxies. In addition the procedure involves a number of assumptions. It is a reasonable question, therefore, whether the values from simulations may be meaningfully compared to the observational values. Zurek et al. (1994) applied the Davis-Peebles method to mock redshift surveys extracted from simulations of a standard Cold Dark Matter (CDM) model, and found that this yields a rather large range of values for $\\sigma_{12}$. So far there has not been a comparison of $\\sigma_{12}$ calculated using the Davis-Peebles method on ``observed'' simulations of different cosmological models. In this paper, we investigate the robustness of $\\sigma_{12}$ using mock redshift surveys extracted from several different cold dark matter dominated models, and the ability of this statistic to discriminate between such models. We estimate the sky-variance and cosmic variance of the statistic and identify sources of error in the Davis-Peebles method. We also investigate the effects of removing clusters from the samples. ",
        "conclusions": "For the N-body simulations of three different cosmological models that we have analyzed, we find that the values of $\\sigma_{12}$ are considerably higher for CDM models than for CHDM models. The CHDM models, which give $\\sigma_{12}(1) \\sim 540 (440)$ km/s for CHDM2 to 740 (600) km/s for CHDM1 are perhaps more consistent with the body of observational values taken as a whole. (Numbers in paranthesis are for $F=0$). Although $\\sigma_{12}(1) = 647$ km/s for CfA2 North is marginally consistent even with unbiased CDM ($\\sigma_{12}(1) \\sim 1024 (880)$ km/s), $\\sigma_{12}(1)$ for CfA2 South (367 km/s) and SSRS2 (272 km/s) are not (Marzke et al. 1995).  However, because of the problems with galaxy identification discussed in Section \\ref{sec-sims}, we do not think that any model studied here should be ruled out on the basis of these results.  Any existing large volume N-body simulations would have similar problems. The point of this paper is precisely to suggest that it is premature to use $\\sigma_{12}$ to draw any strong conclusions about cosmological models.  However, studying the simulations has given us other interesting information.  We have estimated the expected sky-variance and cosmic variance of $\\sigma_{12}$ in our models. The following values are quoted for $\\sigma_{12}(1 h^{-1}Mpc)$. The sky-variance (calculated as the standard deviation over six mock catalogs) ranges from $\\sim 40$ km/s for CHDM2 to $\\sim 145$ km/s for CDM1. The cosmic variance (between two simulations with different initial conditions) for CHDM is $\\sim 200$ km/s. The errors usually quoted for $\\sigma_{12}$ are formal errors on the fit (for which we obtain typically $\\sim 40$ km/s) and in general are underestimates of the actual statistical errors.  We evaluate the accuracy of the Davis-Peebles method for extracting $\\sigma_{12}$ from redshift catalogs by comparing with the results of computing the velocity dispersion directly in our mock catalogs. We obtain agreement of better than 20\\% for all of our models. This leads us to interpret the large range of values of $\\sigma_{12}$ obtained from different redshift surveys as an intrinsic variation due to the sensitivity of the statistic to the clusters contained in the sample, rather than being due to errors in the method.  We have investigated the effects of removing clusters from the samples using an automated procedure. It was hoped that this might make $\\sigma_{12}$ a more robust statistic. However, we found that although this reduces the sample-to-sample variation in $\\sigma_{12}(r)$ by a small amount, it actually reduces the ability of the statistic to discriminate between the cosmological models we studied. This may be due to the fact that all of our simulations are of $\\Omega = 1$ models, and that once clusters are removed $\\sigma_{12}$ is really a measure of $\\Omega_{0}$. However, our study of the simulations suggests that the change in $\\sigma_{12}(1)$ as a function of the number of clusters removed may be an interesting quantity to study in future redshift surveys.  We find that an exponential form for the pairwise peculiar velocity distribution is an excellent approximation on small scales ($r<5 h^{-1}$ Mpc) in all the models studied. Measuring the mean streaming $\\overline{v_{12}}(r)$ directly from the simulations revealed that although the general form of the model used in the Davis-Peebles method does hold, on small scales the measured values may deviate from it considerably. We found that stable clustering is a reasonable approximation in the unbiased (b=1) CDM model but not in the CHDM models. The use of the BBGKY model for the mean streaming, especially with the assumption of stable clustering ($F=1$) could be a substantial source of error in the Davis-Peebles method.  We have shown that $\\sigma_{12}$ is very sensitive to both the number and the properties of the clusters in a sample. This makes $\\sigma_{12}$ a poor constraint on cosmological models given the current situation with regards to both simulations and observations. Even the largest existing redshift surveys do not represent a fair sample of rich clusters. Also, current N-body simulations do not simulate clusters realistically because cluster properties are probably sensitive to non-gravitational physics such as gas hydrodynamics, star formation, and supernova feedback --- effects which are impossible to include in large volume simulations with current computing capabilities. As larger redshift surveys become available and it becomes possible to simulate clusters in a cosmologically relevant volume, perhaps the robustness of $\\sigma_{12}(1)$ will improve. In fact, we have suggested a way in which the very sensitivity of $\\sigma_{12}$ to the properties of clusters could be used to define an interesting statistic for characterizing large scale structure in larger samples. In the meantime it is worthwhile to work on developing statistics which are discriminatory but less sensitive to the properties of clusters.  \\clearpage  \\begin{center} Acknowledgements \\end{center}"
    },
    "9604/astro-ph9604185_arXiv.txt": {
        "abstract": "We have proposed previously that Sgr A* is simply a scaled down AGN with a black hole, an accretion disk and a radio jet operating at a very low power. It appears as if M81* -- the nuclear source in the nearby galaxy M81 -- is an ideal laboratory to study a Sgr~A*-like source at a higher power level. The jet/disk model can explain M81* in great detail with no basic changes in the model parameters other than the accretion rate. Radio cores in other LINERs may be explained by the same model and they appear to be low-power counterparts to radio-loud quasar cores.  For Sgr A*, models without a supermassive black hole are facing difficulties -- some of which are discussed here, but a persistent puzzle in any scenario are the non-detections and low flux limits for Sgr A* from IR to x-rays. Especially the IR limits are a threat to accretion models. I discuss whether a thin molecular disk (as seen in NGC 4258) around Sgr A* could intercept infalling material before it reaches the black hole. ",
        "introduction": "The Galactic Center (GC) is a unique place in our galaxy, however, it is not necessarily a unique place in the universe. For this reason the GC has often been used as an analogy for other galaxies, and the GC can help us to understand what we do not understand in more distant places. But for some aspects of the GC itself, the GC is not necessarily the best place to look for answers.  The latter is especially true for the central compact radio nucleus Sgr~A*. While its basic radio properties are known for quite a while, the search for counterparts in other wavelength regimes has been largely unsuccessful. This is mainly due to the intrinsic weakness of the Galactic Center and the obscuration in the Galactic plane. Those difficulties have in part driven the developments of many new instruments and techniques -- the GC is often among the first objects to be observed with new cameras. And once in a while this has led to a detection of Sgr A* at frequencies inaccessible to radio astronomers (e.g. IR, NIR, X-rays, 511 keV line, etc.), but whenever the next generation of instruments provided higher sensitivity and resolution, it was shown that this emission was due to stellar objects and not due to the suspected supermassive black hole in the very center. This means that any successful model for Sgr A* has not only to be self-consistent but must also be stable against the annual variations in detections of Sgr A* which are a function of wavelength and spatial resolution. Fortunately, at least the evidence from dynamical estimates for the presence of a dark mass of $M_\\bullet=2\\cdot10^6M_\\odot$ in the center of the Galaxy has become more and more convincing (Genzel 1996, Rieke \\& Rieke 1996, Haller et al. 1996) in recent years and now seems well established.  The presence of such a large concentration of dark mass in the very nucleus of the Galaxy places the GC in line with many other galaxies where similar or even much higher dark mass concentrations have been found (Kormendy \\& Richstone 1995). Another similarity to other Galactic Nuclei is the presence of a compact flat spectrum radio core which is found in all radio loud active galaxies, as well as in many other active galaxies like radio quiet quasars, Seyferts, LINERs, in many elliptical galaxies, and also in some spiral galaxies. In this respect is the Galactic Center fairly typical and therefore should not be considered as an isolated case. In this paper I will therefore not only discuss possible explanations for Sgr A* and their difficulties, but also apply our Sgr A* jet/disk model to other weakly active galaxies, specifically to the nucleus of M81. ",
        "conclusions": "All the models proposed for Sgr A* have a certain appeal. The jet/disk model -- with and without monoenergetic electrons -- offers a scope that goes far beyond the GC and has survived a series of critical tests in a variety of very different source classes with compact flat spectrum cores, including Sgr A*. Advection-dominated and fossil disks may help to explain why the optical luminosity of Sgr A* is so low, and Bondi-Hoyle accretion is a process that seems to be unavoidable at a certain level. Applying all these concepts to the nuclei of nearby galaxies may help us to sort out which process dominates in which regime. Until then we perhaps could agree on a ``theorists-for-galactic-peace-model'' for Sgr A*: a jet of monoenergetic electrons, produced by an advection dominated disk coming from a fossil ring which is fed by Bondi-Hoyle accretion."
    },
    "9604/astro-ph9604136_arXiv.txt": {
        "abstract": "We present an analysis of COMPTEL observations made between November 1991 and May 1994 of \\cg/, a bright \\gamms/-ray source located near the Galactic plane. At energies above 1 MeV, an excess consistent with the position of \\cg/ is detected in the sum of the observations, at flux levels which are a factor of $10-100$ below those published in the past. The detection significance of this excess, when the possible presence of underlying Galactic diffuse emission is neglected, is $6.6\\sigma$ for 3 degrees of freedom. The differential photon spectrum in the 1--30 MeV energy range can be described by a power law with a spectral index of $1.95^{+0.2}_{-0.3}$. Due to the uncertainties involved in modelling the Galactic-disk diffuse emission underneath the source, the absolute flux levels must be considered uncertain by a factor of two. They are consistent with the extrapolation of the time-averaged spectrum of \\cg/ measured with EGRET, thereby strengthening the identification. No significant temporal correlation between the \\gamms/-ray emission and the monitored radio emission of the possible counterpart radio source \\gts/ (showing a 26.5 day modulation) is found. ",
        "introduction": "With a flux above 100 MeV of $1.0\\times10^{-6}$ photons cm${}^{-2}$ s${}^{-1}$, the \\gamms/-ray source \\cg/ is one of the brightest unidentified high-energy Galactic-plane sources in the second COS-B catalogue (Hermsen \\etal/ 1977; Swanenburg \\etal/ 1981). \\cg/ attracted attention when its position was found to be consistent with that of the radio source \\gts/, which exhibits strong radio outbursts of a non-thermal character (Gregory \\& Taylor 1978). Long-term radio observations of \\gts/ subsequently revealed a 26.496-day periodicity in these outbursts (Taylor \\& Gregory 1982, 1984), of which the amplitude is possibly modulated by a 4-year period (Gregory \\etal/ 1989; Paredes \\etal/ 1990). The decay of the radio flux during outbursts is reminiscent of synchrotron emission from an expanding cloud of relativistic electrons (Taylor \\& Gregory 1984). In a high-resolution map obtained two days post-outburst from VLBI measurements at 6 cm, the object appears to consist of two components separated by $3.1\\times10^{13}(D/2.3\\hbox{ kpc})\\hbox{ cm}$ (Massi \\etal/ 1993). Optically, \\gts/ has been identified with \\lsi/, which has a spectrum typical of an early-type B star (Hutchings \\& Crampton 1981; Maraschi \\etal/ 1981; Paredes \\& Figueras 1986) and exhibits the near-infrared excess commonly observed for Be-type stars (D'Amico \\etal/ 1987; Hunt \\etal/ 1994).  Recent modelling of the JHK-band light curves of \\lsi/ has shown that the onsets of the radio outbursts roughly coincide with the inferred periastron passage (Mart\\'\\i\\ \\& Paredes 1995). The source has also been detected in X-rays (Bignami \\etal/ 1981; Goldoni \\& Mereghetti 1995). The X-ray spectrum is rather hard compared to those of normal B stars, but is similar to spectra observed for Be/X-ray binaries. \\medskip  In recent observations made with EGRET on board the Compton Gamma-Ray Observatory (CGRO), a source designated \\groname/ was detected on several occasions (Thompson \\etal/ 1995). This source, whose position is consistent with that of \\cg/ and which is roughly as bright, is generally assumed to be the COS-B source. The improved position of the \\gamms/-ray source established by the EGRET detection of \\groname/ (a 95\\% error radius of $33'$) is consistent with that of \\gts/ (located at $(l,b)=(\\decree{135}{.68},\\decree{1}{.09})$) and has strengthened the proposed identification with the radio source. In addition, the nearby QSO \\qso/ (located at $(l,b)=(\\decree{135}{.64},\\decree{2}{.43})$), which was also contained in the relatively large COS-B error region, is now firmly rejected (von Montigny \\etal/ 1993).  A detection of the 26.496 day periodicity in \\gamms/-rays would remove any remaining doubt about the identification with \\gts/, but neither COS-B nor EGRET has yet detected such flux variations. Note that the 26.496-day periodicity inferred from radio observations of \\gts/ has possibly been detected at infrared and optical wavelengths as well (Mendelson \\& Mazeh 1989; Paredes \\etal/ 1994), with smaller flux variations towards shorter wavelengths. At UV wavelengths and in X-rays, where the positional accuracy is good enough to allow a confident identification with \\gts/, orbital-phase related flux variations have not been detected up to recently (Bignami \\etal/ 1981; Howarth 1983; Goldoni \\& Mereghetti 1995). The factor of $\\sim10$ increase of the X-ray flux in the light curve measured with ROSAT during the monitoring of one orbital cycle (Taylor \\etal/ 1996) may be the first evidence of orbital-modulated X-ray emission. \\medskip  Despite the firm detections of \\gts/ up to X-rays and of \\cg/ above 100 MeV, the situation around 1 MeV is less clear. When we compare the spectral index measured for \\gts/ with ROSAT ($\\alpha\\approx1.1$: Goldoni \\& Mereghetti 1995) with that measured for \\cg//\\groname/ with EGRET ($\\alpha\\approx2.21$: Fierro 1995), it is evident that a spectral break must occur in between. This naturally also holds if the two sources are not related. A simple extrapolation of the measured power laws to the MeV energy range predicts fluxes in the range of $10^{-5}-10^{-4}\\photcmsmev$, which is close to the detection threshold of COMPTEL. A preliminary analysis of COMPTEL observations of the \\cg/ field, based on three observations made between November 1991 and August 1992, did not reveal evidence for such a strong source but showed the presence of a weak excess consistent with the position of \\cg/ (van Dijk \\etal/ 1994). ",
        "conclusions": "We have detected a COMPTEL source consistent with the position of \\cg//\\groname/ in the energy range 1--30 MeV. The flux at 1 MeV is roughly two orders of magnitude below the values reported in the past for a source in this celestial region. The COMPTEL spectral shape is consistent with the spectral break expected from extrapolations of the ROSAT and EGRET observations of \\gts/ and \\cg//\\groname/ respectively. \\gamms/-Ray observations with higher spatial resolution, or showing correlated variability, are necessary to  prove unambigously the association between the radio source \\gts/ and the \\gamms/-ray source \\cg/."
    },
    "9604/hep-ph9604403_arXiv.txt": {
        "abstract": "Neutrino oscillations in the presence of strong gravitational fields are studied. We look at very high energy neutrinos ($\\sim $1 TeV) emanating from Active Galactic Nuclei (AGN). It is observed that spin flavor resonant transitions of such neutrinos may occur in the vicinity of AGN due to {\\it gravitational} effects and due to the presence of a large magnetic field ($\\sim $1 Tesla). A point to note is that matter effects (normal MSW transitions) become negligible in comparison to gravitational effects in our scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9604/astro-ph9604122_arXiv.txt": {
        "abstract": "Globular cluster age estimates based on the absolute magnitude of the main sequence turn-off (\\mvto) are generally considered to be the most reliable from a theoretical viewpoint.  However, the difficulty in determining \\mvto\\ in observed colour-magnitude diagrams leads to a large error in the derived age.  In this paper, we advocate the use of the absolute magnitude of the point which is brighter than the turn-off and 0.05 mag redder (\\mvb) as a precision age indicator.  It is easy to measure this point on observed colour-magnitude diagrams, leading to small observational error bars.  Furthermore, an extensive Monte Carlo calculation indicates that the theoretical uncertainty in \\mvb\\ is similar to \\mvto.  As a result, ages derived using \\mvb\\ are at least a factor of 2 more precise than those derived using \\mvto. This technique is applied to the globular cluster M68 and an age of $12.8\\pm 0.3\\,$Gyr is derived (assuming $\\mvrr = 0.20\\,\\feh + 0.98$), indicating that M68 is a `young' globular cluster.  A homogeneous set of globular cluster age estimates with this precision would provide unprecedented insight into the formation of the Galactic halo. ",
        "introduction": "There are a number of different techniques which may be used to determine the age of a globular cluster (GC). All of these methods rely on comparing some aspect of theoretical stellar evolution models to the observations.  Thus, in order to evaluate the reliability of the various age indicators, one must be aware of the uncertainties in theoretical stellar evolution models.  The correct treatment of convection in stellar models is an area of active research (e.g.\\ Kim \\ea\\ 1995, 1996; Demarque, Guenther \\& Kim 1996a,b) and remains the largest possible source of error in stellar models.  For this reason, properties of the stellar models which depend on the treatment of convection are the most uncertain.  The main sequence and red giant branch stars in GCs have surface convection zones, and so the predicted radii (and hence, colours) are subject to large theoretical uncertainties.  The helium burning stars (horizontal branch, and asymptotic giant branches) are convective in the energy generation regions, and so even the predicted lifetimes and luminosities of stars in this phase of evolution are somewhat uncertain.  An additional consideration when considering the reliability of stellar models is that observed CNO abundances in stars on the red giant branch indicate that some form of deep mixing occurs in these stars, which is {\\em not} present in the models (e.g.\\ Langer et al.\\ 1983; Kraft 1994; Chaboyer 1995).  This indicates that the red giant branch models are in need of revision.  In contrast, low mass main sequence models are in excellent agreement with the observations. Indeed, inversions of solar models which use the observed $p$-modes indicate that the run of density and sound speed in solar models agree with the Sun to within 1\\% \\cite{basu}.  The relative reliability of the age-luminosity relationship for low mass stars is well known, and it is for this reason that the absolute magnitude of main sequence turn-off (\\mvto) results in GC ages with the smallest theoretical error (e.g.\\ Renzini 1991).  Operationally, \\mvto\\ is defined to be the magnitude of the bluest point on the main sequence.  (Since this definition involves the use of colour \\mvto\\ is not strictly independent of the uncertainties in stellar radii.) Unfortunately, the turn-off region has nearly the same colour over a large range in magnitude.  This leads to difficulties measuring \\mvto\\ observationally, due to the scatter in the observed points around the turn-off.  Observers typically quote errors of order 0.10 mag in \\mvto, which leads to an error in the derived age around $\\pm 1.5\\,$Gyr (e.g.\\ (Sarajedini \\& King 1989; Chaboyer, Demarque \\& Sarajedini 1996, hereafter CDS).  This large error in the derived age of any individual GC is a great obstacle in furthering our understanding of galaxy formation. This problem has lead Sarajedini \\& Demarque (1990) and VandenBerg, Bolte \\& Stetson (1990) to advocate the use of the difference in colour between the main sequence turn-off and the base of the giant branch (\\dbv) as an age indicator.  This method has the advantage that the colour of the turn-off and the base of the giant branch can be accurately determined in observed colour-magnitude diagrams (CMDs), and is independent of the distance modulus.  As a result, this method can lead {\\it in principle\\/} to very precise age estimates (of order $\\pm 0.5\\,$Gyr).  However, the theoretical colours are subject to large uncertainties and \\dbv\\ only yields reliable {\\em relative\\/} age differences between clusters of a similar metallicity (see, however, the case of M68 and M92 (\\S \\ref{sec3}) for which \\dbv\\ fails).  In this paper, we advocate the use of a point which is brighter than the turn-off, and 0.05 mag redder in B--V (hereafter referred to as \\mvb).  This point is easy to measure in observational data and has a small theoretical uncertainty.  As it still requires  knowledge of the distance modulus, \\mvb\\ complements the \\dbv\\ technique in providing precision age estimates for GCs.  In \\S \\ref{sec2}, a Monte Carlo set of isochrones is described and analyzed in order to estimate the theoretical error associated with \\mvb\\ and \\mvto.  The well studied GC M68 is used to illustrate the relative precision of ages derived using \\mvb\\ and \\mvto\\ in \\S \\ref{sec3}.  Finally, \\S \\ref{sec4} summarizes the results of this work and suggests that observers should quote \\mvb\\ in their papers in addition to \\mvto. Simple formulae are provided to determine GC ages, given \\mvb\\ and \\feh. ",
        "conclusions": "\\label{sec4} An extensive Monte Carlo analysis indicates that the theoretical uncertainty in \\mvb\\ is similar to \\mvto.  The sensitivity of \\mvb\\ to age changes is similar to \\mvto.  The objective fitting technique described in \\S \\ref{sec3} indicates that the error in measuring \\vb\\ in observational data is $\\sim \\pm 0.006\\,$mag, at least an order of magnitude smaller than the error typically quoted in V(TO).  Hence, \\mvb\\ is a superior age indicator to \\mvto.  We suggest that observers should measure \\vb\\ as outlined in \\S \\ref{sec3}, and provide this value as a routine part of the analysis of GC CMDs.  A calibration of age as a function of \\mvb\\ (for B--V data) is presented in eq.\\ 1 and Table \\ref{tab2}.  A similar calibration for V--I data is presented in appendix \\ref{app1}.  The use of \\mvb\\ as an age indicator requires a knowledge of the distance modulus.  We suggest the use of the absolute magnitude of the horizontal branch so that ages can be derived using the difference in magnitude between the horizontal branch, and V(BTO), \\dvb. This leads to an age for M68 ($\\feh = -2.1$) of $12.8\\pm 0.3\\,$Gyr, assuming $\\mvrr = 0.20\\,\\feh + 0.98$.  The error in the derived age is dominated by the error in metallicity.  This is an internal error, useful in comparing {\\em relative} ages, and does not include the error due to the uncertainty in the \\mvrr\\ calibration.  However, as shown by CDS, many statements concerning relative ages are true independently of the choice of \\mvrr.  For example, M68 is significantly younger ($\\sim 2.5\\,$Gyr) than the mean age of 17 other low metallicity GC.  This statement is true over the full range in \\mvrr\\ calibrations quoted in the literature.  If the age of M68 had been determined using the published value of V(TO), it would not be possible to state that M68 is a young GC, due to the large error in the derived age ($\\pm 1.3\\,$Gyr).  This demonstrates the unique advantage of using \\mvb\\ to probe relative GC ages,  which should lead to new insights into the formation of the Milky Way."
    },
    "9604/astro-ph9604064_arXiv.txt": {
        "abstract": "Current measurements of the Hubble constant $H_0$ on scale less than $\\sim100$ Mpc appear to be controversial, while the observations made at high redshift seem to provide a relatively low value. On the other hand, the  Hubble expansion is driven by the matter content of the universe. The dynamical analysis on scale of a few $\\sim10$ Mpc indicates that the  matter density $\\Omega_0$ is only $\\sim0.2$--$0.3$, which is significantly smaller than $\\Omega_0=1$ predicted in the standard inflation model. This might support the tendency of a decreasing Hubble constant towards distance. In this paper, we discuss the influence of a possible variant Hubble constant on two fundamental relations in astronomy: the magnitude-redshift ($m$--$z$) and the number-magnitude relations. Using a distant type Ia supernova at $z=0.458$, we show that the deceleration parameter $q_0$ or $\\Omega_0$ cannot be determined from the $m$--$z$ relation at moderate/high redshift unless the variation of the Hubble constant is {\\it a priori} measured. It is further demonstrated that the number density of distant sources would be underestimated when their local calibration is employed, which may partially account for the number excess of the faint blue galaxies observed at moderate/high redshift. ",
        "introduction": "The Hubble constant $H_0$ is an indicator of the global expansion of the universe.  Theoretically, this parameter is defined in terms of the Cosmological Principle, i.e.,  the overall matter distribution of the universe is assumed to be isotropic and homogeneous. The isotropy of the universe has been remarkably demonstrated by the measurement of the $3^{\\circ}$K microwave background radiation (e.g. Bennett et al. 1996). However, observations have indicated that there exists an inhomogeneous matter distribution in the local universe on scale up to $\\sim100$ Mpc, such as matter clumps, voids, filaments, etc., which may cause a deviation of the local Hubble expansion from the global one. The pioneering work to directly measure the variation of the Hubble constant with distance was carried out  by Sandage, Tammann and Hardy in 1972, using the first ranked E galaxies.  Two decades later, Lauer \\& Postman (1992) conducted another direct observation from the brightest cluster galaxies to $z=0.05$. These authors have essentially reached a similar conclusion that the ratio of the local Hubble constant $H_L$ to the global one $H_G$ is consistent with unity, indicative of a very minor effect of local matter clumps and voids on the expansion of the universe. Other evidences from the study of the Hubble diagram of various objects (e.g. Sandage \\& Hardy 1973; Tammann \\& Sandage 1995) supported that the Hubble constant is roughly invariant.  However, this situation has been recently challenged by several determinations of the Hubble constant utilizing Cepheid variables and the globular cluster luminosity function in some nearby galaxies, which yield a  large value of  $H_0$ ranging from 69 to 87 km s$^{-1}$ Mpc (Pierce et al. 1994; Freedman et al. 1994; Tanvir et al. 1995; Whitmore et al. 1995). These correspond to an age of the universe of $8-13$ Gyr in the frame of the standard cosmological model, while the oldest known stars in globular clusters of our Galaxy are estimated to be about $15.8\\pm2.1$ Gyr (Bolte \\& Hogan 1995). If these calibrations are correct rather than suffer from the systematic and logical errors as were argued by Sandage and collaborators (Sandage 1996; Sandage \\& Tammann 1996; reference therein), there are two possibilities to remove the apparent conflict over the age of the universe: (1) a nonzero cosmological constant $\\lambda_0\\equiv\\Lambda/(3H_0^2)$ [see Ostriker \\& Steinhardt (1995) for a summary] and (2) a local low-density region embedded in a globally flat universe (Turner, Cen \\& Ostriker 1992; Wu et al. 1995). The later requires a considerably variation of the Hubble constant with distance.  The claim for a deviation of the local Hubble flow from the global one is indirectly supported by the  following observations: From the time delay of 1.1 years revealed by over 10 years coverage of the monitoring of the gravitationally-doubled images QSO0957+561A,B at redshift of 1.41, one has found $H_G=48^{+16}_{-7}$ km s$^{-1}$ Mpc$^{-1}$ [see Wu (1996) for a recent review]. This global value is relatively low though large uncertainties may exist in the resulting $H_G$ due to the modeling of the deflectors. For instance, a new lensing model of QSO0957+561A,B based on the same data yields $H_G\\approx 82.5$ km s$^{-1}$ Mpc$^{-1}$ (Grogin \\& Narayan 1996). Another evidence comes from the measurement of Sunyaev-Zel'dovich (S-Z) effect, which is the spectral distortion of the cosmic background radiation due to the inverse Compton cooling of the hot X-ray gas in clusters of galaxies (Sunyaev \\& Zel'dovich 1972). It has been found that a relatively small value of the Hubble constant is derived from distant galaxy cluster, e.g. $H_0=65\\pm25$ km s$^{-1}$ Mpc$^{-1}$ for A2218 ($z=0.171$) (Jones et al. 1993; Birkinshaw \\& Hughes 1994), $H_0=47\\pm17$ km s$^{-1}$ Mpc$^{-1}$ for A665 ($z=0.182$)  and $H_0=41^{+15}_{-12}$ km s$^{-1}$ Mpc$^{-1}$ for Cl0016+16 ($z=0.545$) (Yamashita 1994). However, the detection of the S-Z effect for nearby Coma cluster provides $H_0=71_{-25}^{+30}$ km s$^{-1}$ Mpc$^{-1}$ (Herbig et al. 1995). It seems likely that the Hubble constant decreases with redshift. Yet, the errors in these measurements are still large and uncertainties from modeling of the hot X-ray gas in clusters of galaxies need to be further studied (Inagaki, Suginohara \\& Suto 1995; Rephaeli 1995).  On the other hand, it is crucial to realize that the Hubble constant depends on the matter distribution since the expansion is uniquely governed by the gravitational matter of the universe. There have been increasing observational evidences that the local matter density $\\Omega_L$ within a few 10 Mpc is only $\\Omega_L\\approx0.2$--$0.3$ (Bahcall, Lubin \\& Dorman 1995; references therein). However, the standard inflation cosmological model predicts an overall matter density of $\\Omega_G\\approx1$. This matter discrepancy implies a gradient of the Hubble constant along distance.   The present measurement of the Hubble constant as well as the density parameter is still a controversial issue. It turns out to be unclear whether or not the Hubble constant alters with distance. In this paper, we address the following question: If the Hubble constant is variant, what effect will it have on the fundamental relations in astronomy ?  In particular we will demonstrate this effect on the magnitude-redshift relation using the distant type Ia supernova (section 2) and on the source number counts using the faint galaxy population (section 3). In section 4 we will construct a simple analytic model for a locally low and globally high density universe instead of the usual simulation method (e.g. Turner et al. 1992; Suto et al. 1995; Nakamura \\& Suto 1995; Shi et al. 1995). A brief discussion and our conclusions will be given in section 5. ",
        "conclusions": "We have shown that if the matter density of our local universe within $\\sim100$ Mpc is $\\Omega_L\\approx0.2$--$0.3$, as indicated by the dynamical observations (Bahcall et al. 1995), and the global value is close to unity, as predicted by the standard inflation model, then the local Hubble expansion rate would be $\\sim1.3$ times larger than the global one. The current measurements of the Hubble constant do not exclude this variation with distance. As a consequence, we need to be cautious of using some fundamental relations.   Utilizinging a distant SN Ia at $z=0.458$ discovered recently (Perlmutter et al. 1995), we have demonstrated how the variation of $H_L/H_G$ affects the magnitude-redshift relation.  Given the global matter density $\\Omega_G$, we have placed the preliminary limit on $H_L/H_G$. It turns out that due to the large intrinsic dispersion for the SN Ia, the variation of $H_L/H_G$ has not been well constrained. The current limit is consistent with either a uniform expansion or a variation of $H_L/H_G$. Nevertheless, we have pointed out that there are two free parameters, $\\Omega_G$ and $H_L/H_G$, appearing in the $m$--$z$ relation at moderate/high redshift if the local calibration is used, which would invalidate the conventional method of determining $\\Omega_0$ from the difference in the Mattig relation unless observations are able to provide the relationship between $H_L$ and $H_G$. This is consistent with the recent work by Kim et al. (1996).   The deviation of the local Hubble expansion from the global one would have influence on the calibration of the number counts of distant sources. We have found that the number density of distant faint galaxies would be underestimated if $H_L>H_G$. This may partially remove the demand for a strong evolution scenario of galaxies population in order to account for the number excess of faint (blue) galaxies. As a result, the theoretically predicted number of giant arcs, which are the distorted images of distant galaxies by the gravitational potentials of foreground clusters of galaxies, would increases. This probably provides another explanation for the deficit of giant arcs in statistical lensing (Wu \\& Mao 1996).  We have only discussed the effect of $H_L/H_G$ on the magnitude-redshift relation and the number counts. Indeed, if the local expansion is larger than the global one, many aspects of astrophysics might need to be modified. The crucial point depends on the measurements of both the Hubble constant and the matter density at different redshift. The current status on the determinations of $H_0$ and $\\Omega_0$ appears to be unfortunate: The observations using the same method can result in very different results. Here we do not intend to be involved in the disputes. Instead, we point out that it would be of great significance if one would explore other possibilities which can be used to probe or constrain the variations of the Hubble constant and the density parameter."
    },
    "9604/astro-ph9604158_arXiv.txt": {
        "abstract": "We compute numerically the scalar- and tensor-mode induced Stokes parameters of the cosmic microwave background, by taking into account the basis rotation effects. It is found that the tensor contribution to the polarization power spectrum get enhanced and dominates over the scalar contribution for low multipoles in a universe with or without recombination. Furthermore, we show that all full-sky averaged two-point cross-correlation functions of the Stokes parameters vanish. We thus comment on the cross-correlation between the anisotropy and polarization of the CMB, and calculate the expected signal to noise ratio for the polarization experiment underway. ",
        "introduction": "The detection of the large-angle anisotropy of the cosmic microwave background by the {\\it COBE} DMR (\\cite{smo92}) is an important evidence for the large-scale space-time inhomogenities. Since then, several small-scale anisotropy measurements have hinted that the Doppler peak resulted from acoustic oscillations of the baryon-photon plasma on the last scattering surface seems to be present. It is believed that measuring CMB anisotropies on all angular scales can give invaluable information of the early universe, such as testing inflation, pinning down cosmological parameters, and differentiating scalar from tensor perturbations (see, e.g., \\cite{ste94}).  Complementary to the anisotropy, polarization is another important property of the CMB. The degree of CMB polarization has been extensively calculated (see \\cite{ng96} and references therein). In a universe with standard recombination, the polarization is about $10\\%$ of the anisotropy on small angular scales while the large-scale polarization is insignificant. If early matter reionization of the universe occurs, the large-scale polarization would be greatly enhanced to a $10\\%$ level with the large-scale anisotropy reasonably unaffected. Meanwhile, the small-scale anisotropy would be significantly suppressed. Thus, measuring CMB polarization might provide additional information about the early universe. So far, only experimental upper limits have been set (\\cite{lub83}; \\cite{par88}; \\cite{wol93}), and the signal level is expected to be achieved through the use of new technology and instrument design (P. T. Timbie 1995, private communication).  In previous calculations of the degree of CMB polarization, approximations have more or less been made. For instances, the separation angle in the two-point auto-correlation functions has been assumed small, or/and the basis rotation from the ${\\bf\\hat k}$-basis of the perturbation to the fixed basis of an experiment has been totally neglected. Although it is expected that these approximations would be valid for small-angle polarization calculations, they may break down on large-scale calculations. In this {\\it Letter}, we will compute the Stokes parameters by taking into account the basis rotation effects. Also, we will consider the two-point cross-correlation functions of the Stokes parameters. ",
        "conclusions": "We have calculated all whole-sky averaged two-point correlation functions of the Stokes parameters of the CMB. Incorporating the basis rotation effects, we have found that the tensor contribution to the polarization power spectrum get enhanced for low multipoles whereas the scalar contribution remains unchanged, and that all full-sky averaged cross-correlation functions vanish. A special case of $\\langle TQ\\rangle$ has been calculated, and used to show that the signal to noise of $\\langle TQ\\rangle$ is low for the polarization measurement underway.  We intend to calculate the exact correlation functions that would be important to future polarization experiments."
    },
    "9604/astro-ph9604081_arXiv.txt": {
        "abstract": "{\\baselineskip 0.4cm This review is a short introduction to numerical hydrodynamics in a cosmological context, intended for the non specialist. The main processes relevant to galaxy formation are first presented. The fluid equations are then introduced, and their implementation in numerical codes by Eulerian grid based methods and by Smooth Particle Hydrodynamics is sketched. As an application, I finally show some results from an SPH simulation of a galaxy cluster.} ",
        "introduction": "In current cosmological scenarios, the main matter components of the universe are some non baryonic Dark Matter (DM), which constitutes most of the universe, and a mixture of primordial gas (H, He). The DM component is decoupled from the rest of the universe, and interacts only through gravity. The gas component instead can be heated and cooled in several ways, and the physics involved is more complicated than in the pure gravitational case. Fortunately, while gravity is a long range force, hydrodynamic processes are only important on relatively small scales, so that on scales larger than a few Mpc the dynamics of structure formation can be studied with good accuracy even neglecting the gas component. Gas dynamics, and the related radiative processes, are instead fundamental on smaller scales, e.g. in the formation of galaxies, and in linking the matter distribution of the universe to the light distribution we actually observe$^{1)}$. In what follows I list and discuss very briefly the main gas processes relevant to galaxy formation.  \\subsection{Heating processes}  \\begin{description}  \\item [Adiabatic compression] is the easiest way to heat a gas. By the First Law of Thermodynamics, compression work is converted into internal energy: $dQ = dU + pdV = 0 \\Longrightarrow dU = - pdV$.  \\item[Viscous heating] is due to the small internal friction (viscosity) present in real gas. Velocity gradients in a gas cause an irreversible transfer of momentum from high velocity points to small velocity ones, with conversion of bulk velocities into random ones, i.e. into heat, and generation of entropy. In the context of galaxy formation viscous heating mostly occurs in shocks, which are discontinuities in the macroscopic fluid variables due to supersonic flows. These arise for example during gravitational collapse, or during supernova (SN) explosions.  \\item[Photoionization] takes place when atoms interact with the photons of some background of soft X-rays, or UV radiation emitted by QSO or stars: e.g. $\\gamma + H \\rightarrow e^- + H^+$. Observations show that at high redshift ($z \\gsim 2$) hydrogen in the IGM is indeed ionized (Gunn-Peterson effect$^{2)}$), and although the origin is not clear, this is thought to be caused by some early generation of QSO or massive stars. The photoionizing spectrum is usually approximated by a power law, with flux $J(\\nu) \\propto (\\nu / \\nu_L)^{-\\alpha}$, where $\\nu_L$ is the Lyman-$\\alpha$ frequency, corresponding to the hydrogen ionization energy, $13.6$ eV.  \\end{description}  \\subsection{Cooling processes}  Gas cooling is the key to galaxy formation. In fact, in our current understanding of structure formation, the dark matter component of the universe first undergoes gravitational collapse, forming dark matter halos. These provide the potential wells into which gas can fall and heat up by shocks, then immediately cool and form cold, dense, rotationally supported gas disks. In these disks conditions are favourable to trigger star formation, and eventually give rise to the galaxies we observe today$^{3)}$. The following cooling mechanisms are important in a cosmological context.  \\begin{description}  \\item [Adiabatic expansion] is the opposite process of adiabatic compression, with conversion of heat into expansion work.  \\item[Compton cooling] is electron cooling against the Cosmic Microwave Background (CMB) through inverse Compton effect: $\\gamma + e^-  \\rightarrow \\gamma + e^-$. The condition for this is that the temperature of the electrons is higher than the CMB temperature: $T_e > T_\\gamma$. The net energy transfer depends on the two densities, and on the temperature difference: $dE/dT \\propto n_e \\rho_\\gamma (T_e - T_\\gamma)$. Since $E\\propto n_e T_e$, the cooling time: $t_{cool} \\equiv E/\\dot{E}$ is in this case $\\propto \\rho_\\gamma^{-1}$. The CMB photon density decreases like $(1 + z)^4$ due to the expansion of the universe, therefore Compton cooling is only important at high redshifts (typically $z \\gsim 8$), when the cooling timescale is smaller than the Hubble time.  \\item[Radiative cooling] is the most relevant mechanism for the cooling of primordial gas. It is caused by inelastic collisions between free electrons and H, He atoms (or their ions). Assuming that the gas is optically thin and in ionization equilibrium, the cooling rate per unit volume may be written $dE/dt \\equiv \\Lambda(\\rho,T) = n_e n_i f(T)$, where $n_e$ and $n_i$ are the number densities of free electrons and of atoms (or ions), and $f(T)$ is called {\\em cooling function}. The main processes are:  \\begin{itemize}  \\item {\\em Collisional ionization:} the inelastic scattering of a free electron and an atom (or ion), which unbinds one electron from the latter, e.g. $e^- + H \\rightarrow H^+ + 2 e^-$. The net cooling for the system is equal to the extraction energy.  \\item {\\em Collisional excitation + line cooling:} the same situation as above, but the atom is only excited, and it then decays to the ground state, emitting a photon, e.g. $e^- + H \\rightarrow e^- + H^* \\rightarrow e^- + H + \\gamma$. This is the dominant cooling process at low ($10^4$~K$ \\lsim T \\lsim 10^6~$K) temperatures.  \\item {\\em Recombination:} e.g. $e^- + H^+ \\rightarrow H + \\gamma$. A free electron is captured by an ion and emits a (continuum) photon.  \\item {\\em Bremsstrahlung:} free-free scattering between a free electron and an ion, e.g. $e^- + H^+ \\rightarrow e^- + H^+ + \\gamma$. Its cooling rate grows with the temperature: $dE/dt \\propto T^{1/2}$; therefore, bremsstrahlung is the dominant cooling process at high ($T \\gsim 10^6$ K) temperatures.  \\end{itemize} \\end{description}  Figure 1 shows the cooling and heating functions in different cases.  \\begin{figure}[t] \\epsfxsize=\\hsize\\epsffile{cool.ps} \\caption{ \\baselineskip=0.4truecm {\\footnotesize Cooling and heating functions. Left panel: cooling function $\\Lambda(T)/n_H^2$ versus temperature, for a primordial gas. The different contributions to radiative cooling, and the total cooling curve are indicated. No photoionizing background is assumed. Right panel: a heating term $\\Gamma/n_H^2$ is added, due to a photoionizing UV background with spectral index $\\alpha = 1.5$. Its effect is both to change the ionization equilibrium (and so to change $\\Lambda(T)/n_H^2$), and to heat the gas. Thin lines correspond to a gas density equal to the mean background density at redshift $z=5$; thick lines correspond to a density 200 times larger. For low density gas the effect of a UV background is dramatic, and line cooling is suppressed. The temperature $T_{eq}$ where $|\\Lambda - \\Gamma|$ drops to zero is called equilibrium temperature. At $T > T_{eq}$ the gas is cooled, at $T < T_{eq}$ it is heated. This figure was kindly prepared by I.Forcada using the atomic rates provided by T.Abel.} } \\end{figure}  \\subsection{Other processes} Besides the heating and cooling mechanisms listed above, other processes may be relevant in a galaxy formation scenario. Among them:  \\begin{description}  \\item[Thermal conduction] is direct transfer of heat from regions at high temperature to regions at lower temperature, due to the energy transport of diffusing electrons. The induced change in internal energy per unit volume is $dE/dT = \\vec\\nabla\\left( \\kappa \\vec\\nabla T\\right)$, with (positive) thermal conductivity $\\kappa = \\kappa(T, p)$.  \\item[Radiation transfer] is important in an optically thick medium. Photons are absorbed by gas clouds, thermalized by multiple scattering and reemitted as Black Body radiation. This process depends on the optical depth $\\tau_{opt}$ of the gas.  \\item[Star formation] follows gas cooling as the next natural step in modeling galaxy formation. Our understanding of the detailed physics of star formation is still rather poor, so what is usually done is to use some empirical prescription to characterize the gas which is supposed to turn into stars. A good example of recipe$^{4)}$ is to require: {\\bf\\em i)} a convergent gas flow: $\\vec\\nabla\\cdot\\vec v <0$; {\\bf\\em ii)} a Jeans' instability criterion: the free-fall time of the gas cloud be less than its sound crossing time; {\\bf\\em iii)} a minimum number density of H atoms, e.g. $n_H > 0.1$ cm$^{-3}$; {\\bf\\em iv)} a minimum gas overdensity, e.g. $\\rho_g \\gsim \\rho_V = 178 \\bar{\\rho_g}$, where $\\rho_V$ is the virial density of the spherical collapse model, and $\\bar{\\rho_g}$ is the mean background gas density. If all these conditions are satisfied, the gas will form stars at a formation rate similar to the one observed in e.g. spiral galaxies.  Assuming some Initial Mass Function for the stars so formed, it is possible to compute the fraction of gas forming massive stars ($M \\geq 8 M_\\odot$). These stars will explode as Type~II~SN, each one releasing $10^{51}$~erg of energy back in the ISM, causing new shocks and gas heating, and triggering further star formation. It is also possible to include in this recipe gas release and metal enrichment from SN explosions. Unfortunately, at this point the physics of star formation is still poorly understood, and the resulting scenarios depend very much on the kind of assumptions and modeling made.  \\end{description} ",
        "conclusions": ""
    },
    "9604/astro-ph9604174_arXiv.txt": {
        "abstract": "We have mapped the \\hh\\ \\vone\\ ($\\lambda=2.1215\\, \\micron$) emission line along a 400~pc long strip and in a 50~pc region in the Galactic center. There is \\hh\\ emission  throughout the surveyed region. The typical de-reddened ($\\ak = 2.5$~mag) \\hh\\ \\vone\\ surface brightness, $\\sim 3 \\times 10^{-5}$ \\ergintensity, is similar to the surface brightness in large-scale photon-dominated regions in the galactic disk. We investigate two possible excitation mechanisms for the \\hh\\ emission: UV-excitation by photons from OB stars, and shock waves, and conclude that UV-excitation is more likely. The total \\hh\\ \\vone\\ luminosity in the inner 400 pc region of the Galaxy is 8000~\\lsun. The ratio of the \\hh\\ to far-IR luminosity in the inner 400 pc of the Galaxy agrees with that in starburst galaxies and ultraluminous infrared bright galaxies.  \\subjectheadings{Galaxies:ISM Infrared:ISM:Lines and Bands ISM:Molecules} ",
        "introduction": "\\label{sec:int}  Physical conditions in the interstellar medium of the Galactic center\\footnote{ We use here the term ``Galactic center'' to denote the inner several 100~pc region of our Galaxy. We adopt a distance of 8.5~kpc, with which 1\\arcdeg\\ corresponds to 148~pc.} are significantly different from those in the solar neighborhood. The thin disk of interstellar material in the Galactic center (size: $450 \\times 40$~pc) contains $\\sim10^{8}\\ \\msun$ of dense molecular gas, $\\sim$10~\\% of the Galaxy's molecular mass (G\\\"{u}sten 1989). The molecular clouds in the Galactic center have higher density, metallicity, and internal velocity dispersion than the clouds in the solar neighborhood (Blitz et al. 1993). Strong radio continuum radiation from giant \\hii\\ regions and extended, low-density (ELD) \\hii\\ (Sofue 1985), as well as far-IR radiation from dust (Odenwald \\& Fazio 1984), indicate that the UV radiation field is intense. The energetic conditions in the Galactic center can provide a unique view of the interaction between stellar UV radiation and molecular clouds, and a nearby example for the nuclei of other galaxies.  Ro-vibrational lines of \\hh\\ trace Photon-dominated Regions (PDRs) where far-UV photons excite the \\hh\\ and shocked regions where the \\hh\\ is thermally excited.  As a result, the central regions ($\\sim 1$~kpc) in starburst galaxies are powerful emitters of near-IR \\hh\\ emission (Puxley, Hawarden, \\& Mountain 1989; Joseph 1989; Lester et al. 1990; Moorwood \\& Oliva 1990). Vigorous star formation in these galaxies produces large numbers UV photons which can excite \\hh, while subsequent supernovae can shock excite the \\hh.  Gatley et al. (1984, 1986) and Gatley \\& Merrill (1993) have observed \\hh\\ emission from the inner 5 pc diameter (2\\arcmin) in the Galactic Center, a much smaller region than those observed in starburst galaxies. With the University of Texas Near-Infrared Fabry-Perot Spectrometer (Luhman et al. 1995), it is now possible to observe \\hh\\ emission over much larger angular scales. We describe here a program to map the Galactic center in \\hh\\ emission on a scale of several degrees (several hundred pc), and discuss the likely \\hh\\ excitation mechanism. We can then compare the central region of our Galaxy with those in other galaxies. ",
        "conclusions": "\\label{sec:dis}  \\subsection{ Extinction Correction } \\label{sec:extinction}  At 2.2 $\\mu$m, the emission from the Galactic center is attenuated by interstellar material in the foreground (``foreground extinction'', mostly 4-8 kpc from the Galactic Center) and by material in the Galactic center itself (``Galactic center extinction''). Catchpole et al. (1990) mapped the  extinction toward the Galactic center by observing the near-IR reddening of giant stars in the central few hundred parsecs.  Along our H$_2$ strip at b= --0\\fdg05 (for --0\\fdg6$<l<$+0\\fdg6), the extinction is fairly uniform with a value A$_K\\sim$2.5 mag.  Although most of this extinction is in the foreground, some of it could occur within the Galactic center since Catchpole et al. (1990) were able to identify patches in their maps with A$_K>$2.5 with known molecular clouds in the Galactic center (see plate 4 in Glass et al. 1987).  Based on this work, we adopt A$_K$=2.5 for the foreground extinction.  The Galactic center extinction greatly exceeds the foreground extinction.  Typical $^{12}$CO J=1$\\rightarrow$0 linestrengths along the strip we have surveyed in H$_2$ are 1500 Kkms$^{-1}$ (Oort 1977).  This line strength implies an A$_K$ of 10--40 mag, depending on the CO/H$_2$ and A$_K$/H$_2$ ratios in the Galactic center (Sodroski et al. 1994).  The extinction through individual clouds may also be substantial (A$_K\\sim$10--30 for a 10 pc long cloud with n$_{H_2}$=10$^4$ cm$^{-3}$). The relevance of the Galactic center extinction depends on the source of the H$_2$ emission.  Any H$_2$ emission originating within the clouds will be highly extincted.  If the H$_2$ emission arises on the cloud surfaces, however, we only miss the H$_2$ flux from the back side of each cloud.  Clouds lying in front of other clouds will further reduce the flux reaching us from the front surfaces. If, in the Galactic Center, the velocity integrated area filling factor of clouds, {\\it f}, is substantially greater than unity, extinction by overlapping clouds will reduce the H$_2$ flux observed from the front surfaces by a factor $\\sim1/f$ in addition to the foreground extinction and to the loss of emission on the opposite sides of the clouds. Typical clouds in the Galactic center disk have kinetic temperatures $\\sim$70 K and linewidths $\\sim$20 km s$^{-1}$ (G\\\"usten 1989). An ensemble of such clouds could produce the observed $^{12}$CO J=1$\\rightarrow$0 lines in the Galactic center with f$\\sim$1. We therefore conclude that the extinction of any H$_2$ emission from cloud surfaces facing the sun beyond the foreground extinction of A$_K$= 2.5 discussed above is not substantial.  Since the extinction of emission from within the clouds or from the sides facing away from us is difficult to estimate and since no correction is usally made for such effects in giant molecular clouds and galactic nuclei, we make no additional extinction corrections here.   \\subsection{ UV Excitation of  \\hh\\ } \\label{sec:UV}  If one ignores the region immediately around \\sgra$^*$, the de-reddened ($\\ak = 2.5$~mag, see Section~\\ref{sec:extinction}) \\hh\\ \\vone\\ surface brightness along the Galactic plane has a roughly constant value of $\\simeq 3 \\times 10^{-5}$ \\ergintensity. Any excitation mechanism for the \\hh\\ must be able to explain both the absolute intensity and the uniformity and extent of the emission. The excitation of the \\hh\\ $v=1$, $J=3$ state can result either from radiative decay from UV-excited electronic states or from energetic collisions. \\hh\\ can absorb 91--123 nm photons in the $B^1\\Sigma^+_u - X^1\\Sigma^+_g$ Lyman and $C^1\\Pi_u - X^1\\Sigma^+_g$ Werner bands.  About 90\\% of the time, the excited \\hh\\ decays to some ro-vibrational level of the ground electronic state.  The relative line intensities arising in UV-excited \\hh\\ are insensitive to density or to UV field strength if ${\\rm n_{H_2}} < 10^4$ \\cmv\\ (Black and van Dishoeck 1987). At densities $\\geq 10^5$ \\cmv, UV-excited gas can become hot enough that collisions populate states with $v=1$ (Sternberg and Dalgarno 1989).  Collisional excitation can also result from shocks which abruptly heat the gas to $> 10^3$~K (e.g. Hollenbach, Chernoff, \\& McKee 1989). Several observational results lead us to believe that UV excitation can explain the \\vone\\ emission in the Galactic Center.  The denser parts of clouds like Orion and NGC 2024 produce \\hh\\ emission with an intensity close to that observed in the Galactic Center. In Orion and NGC 2024, the degree-scale H$_2$ emission has a typical surface brightness $\\sim 6 \\times 10^{-6}$ \\ergintensity\\ (Luhman et al. 1994).  Along the molecular ridges in these clouds, the H$_2$ surface brightness is 3-5 times higher. Observations of \\hh\\ transitions arising from higher-lying states indicates that, in these sources, the \\vone\\ emission is a result of UV fluorescence.  If large-scale \\hh\\ emission arises in the surface layers of the clouds where UV photons can excite the molecules, the dust, which absorbs the bulk of the incident flux, ought to radiate in the far-IR continuum as well. Luhman \\& Jaffe (1996) have compared the \\hh\\ \\vone\\ observations of clouds in the galactic disk with IRAS far-IR continuum results and derived a relation between the \\hh\\ \\vone\\ line and far-IR continuum intensities. We can use this relationship and the measured far-IR intensities in the Galactic center to see if UV-excitation is plausible for our observed \\vone\\ emission. In most of the region along our Galactic center H$_2$ cut, the IRAS 100 $\\mu$m band detectors were saturated.  In order to compare the Galactic center H$_2$ data to far-IR continuum measurements with comparable angular resolution, we have combined the un-saturated IRAS measurements from the outer ends of our H$_2$ strip with the 40--250 $\\mu$m continuum measurements of Odenwald and Fazio (1984). To make the two data sets comparable, we have first converted the IRAS 60 $\\mu$m and 100 $\\mu$m fluxes into a total far-IR flux (the FIR parameter of Fullmer and Lonsdale 1989).  The IRAS total far-IR flux agrees well with the far-IR flux derived by Odenwald and Fazio in the regions where their data and the unsaturated IRAS data overlap. We then converted the combined datasets into integrated far-IR intensity for comparison with our H$_2$ strip. We used the Luhman \\& Jaffe galactic disk \\hh\\ dataset to re-derive their H$_2$/far-IR relation in intensity units.  We obtain, \\begin{displaymath} \\log(I_{\\rm H_2 v=1-0 S(1)}) = -4.65 + 0.39 \\log(I_{FIR}), \\end{displaymath} where both intensities are in \\ergintensity. The dispersion of the galactic disk cloud \\hh\\ intensities about this relation is log($\\sigma$) = 0.23.  If we de-redden the Galactic center H$_2$ observations by \\ak\\ = 2.5~mag ({\\it but otherwise do nothing to fit the data to the galactic disk relation}), the Galactic center \\hh\\ intensities have a dispersion log($\\sigma$) = 0.26 about this relation. The Galactic center results are therefore completely consistent with the empirical far-IR vs. \\hh\\ relationship derived for the UV-excited surfaces of clouds in the galactic disk.  We can also compare the H$_2$ line intensities predicted by models of photon-dominated regions to the observed intensities. The models use indirect observations of the far-UV field in the Galactic center (radio and far-IR continuum fluxes) as inputs. For the radio, we predict the far-UV field using emission from extended, low-density (ELD) \\hii\\ regions because the molecular cloud column densities, and therefore the extinction at the wavelength of \\hh, tend to be high (and uncertain) toward the discrete \\hii\\ regions. Away from discrete H~II regions, the typical 10.5 GHz flux density is 2.2 Jy in a 3\\farcm3 beam (Sofue 1985). Assuming $T_e = 10^4$~K, this flux density corresponds to $2.3 \\times 10^{49}$~sec$^{-1}$ Lyman continuum photons per second (Mezger, Smith, \\& Churchwell 1974), in the corresponding region (8.2~pc). For an ionizing source with an effective stellar temperature, $T_{eff} = 3.5 \\times 10^4$~K as the UV source, the 2.3$\\times$10$^{49}$ Lyman continuum photons imply $\\sim$2.3$\\times$10$^{49}$ photons in the range which can excite the \\hh\\ (91-123 nm), or a luminosity of $1.2 \\times 10^5$~\\lsun\\ (Black \\& van Dishoeck 1987). From our observations, the average \\hh\\ flux in a 3\\farcm3 beam is $2.4 \\times 10^{-12}$~\\ergflux. The corresponding total \\hh\\ luminosity in the 8.2~pc (3\\farcm3) region is $3.3 \\times 10^3$~\\lsun, if we correct for an extinction of $\\ak = 2.5$~mag and use the PDR model of Black \\& van Dishoeck (1987) to extrapolate to the \\hh\\ cooling in all lines ($I_{\\rm H_2 v=1-0 S(1)}/I_{\\rm H_2}$ $=$ $0.016$). The ratio of the near-IR \\hh\\ luminosity to the luminosity in the far-UV band that is effective in exciting \\hh\\ is 0.028, which is close to the value of 0.034 from an appropriate PDR model for the Galactic center (Model 19 in Black \\& van Dishoeck 1987, which has n$_H$ = 10$^4$ and a UV field I$_{UV}$ = 10$^3$).  The far-IR continuum intensities along our Galactic center strip are typically 0.8 \\ergintensity\\ (Odenwald and Fazio 1984).  If all of this emission arises from a single molecular cloud surface filling the beam, it corresponds to a far-UV flux $\\sim 2 \\times 10^3$ times the mean interstellar radiation field in the solar neighborhood (Draine 1978). Given the likely number of clouds along each line of sight and various geometric effects, the likely far-UV field is $500-1000$ times the solar neighborhood value.  For this range of UV field strengths and densities between 3$\\times$10$^3$ and 3$\\times$10$^4$ \\cmv,  Black and van Dishoeck (1987) predict \\hh\\ \\vone\\ line intensities in the range $1.2-4.2 \\times 10^{-5}$ \\ergintensity, bracketing our typical observed, de-reddened value.  The \\hh\\ emission from the circum-nuclear disk appears to be collisionally excited (I$_{\\rm v=2-1 S(1)}$/I$_{\\rm v=1-0 S(1)}$) $\\simeq 0.1$, Gatley et al. 1984). Gatley et al. suggest that shocks created by mass outflow from the Galactic nucleus might excite the \\hh. Such thermal line ratios can also occur, however, in UV-excited gas if the UV fields and densities are sufficiently high (Sternberg and Dalgarno 1989; Luhman et al. 1996). Since the typical hydrogen density in the circum-nuclear disk is large, i.e., $n_H \\simeq 10^5\\ \\cmv$, and the UV field is intense in the central 3~pc, (the number of total Lyman continuum photons absorbed by the gas is $\\sim2 \\times 10^{50}\\ \\lsun$,  Lacy et al. 1980), the strength and character of the \\hh\\ emission from the circumnuclear disk are also consistent with UV-excitation.   \\subsection{ Shock-Excitation }   \\label{sec:shock}  Shock excitation of the \\hh\\ \\vone\\ transition must take place, at some level, in the inner 400 pc of the Galaxy. A large variety of dynamical activity may give rise to shocks with appropriate characteristics.  Outflows around newly formed stars and shocks caused by supernova remnants impinging on molecular clouds in the galactic disk both produce \\hh\\ emission and should be observable in the Galactic Center. Bally et al (1987; 1988) surveyed the Galactic center region in the $^{12}$CO and \\thirteenco\\ $J=1 \\rightarrow 0$, and CS $J=2 \\rightarrow 1$ lines. The gas distribution is highly asymmetric about the center, and some negative velocity gas is seen at positive longitudes, which is ``forbidden'' to gas in circular orbits. This gas and other clouds with eccentric orbits may collide with material in more circular orbits. For example, in the \\sgrbtwo\\ complex, Hasegawa et al. (1994) suggested that a dense ($ n_{H_2} \\simeq 1.4 \\times 10^4\\ \\cmv $), $ 10^6\\ \\msun $  ``Clump'' has collided with the extended less dense ``Shell'' of the cloud complex producing a $\\sim 30\\ \\kms$ shock. Finally, the internal velocity dispersion of the molecular clouds is in the range of $\\Delta V = 20 - 50\\ \\kms$ (Bally et al. 1988). If the internal collisions efficiently dissipate the relative kinetic energy by radiative cooling following shocks, there should be \\hh\\ emission throughout the molecular clouds, much of it, however, heavily extincted.  Depending on the context, shock-excited \\hh\\ emission could result either from dissociative J-shocks (colliding clouds, supernova remnants), or from C-shocks (outflows, dissipation of turbulence). The J-shocks give rise to \\hh\\ \\vone\\ intensities in the range of 3$\\times$10$^{-5}$ -- 10$^{-4}$ \\ergintensity\\ with the intensity being fairly insensitive to density and shock velocity over the range 10$^4$ cm$^{-3} \\leq$ n $\\leq$ 10$^5$ cm$^{-3}$ and 30 km s$^{-1} \\leq$ v$_{shock} \\leq$ 150 km s$^{-1}$ (Hollenbach \\& McKee 1989). For A$_K$ = 2.5~mag, the predicted intensity matches what we observe in the Galactic center fairly well. In order to explain the distribution of observed \\hh\\ emission, however, the number of shock fronts times the area covered per beam must roughly equal the beam area along virtually every line of sight through the inner 400 pc of the Galaxy, an unlikely picture at best.  C-shocks can produce \\hh\\ \\vone\\ intensities in the range of those observe in the Galactic Center.  A single C-shock with n = 10$^4$ cm$^{-3}$ and V = 20 km s$^{-1}$ gives I$_{S(1)}$ $\\simeq$ 3$\\times$10$^{-5}$ \\ergintensity\\ (Draine, Roberge, \\& Dalgarno 1983).  The emergent intensity, however, is extremely sensitive to the shock velocity, varying (at n$_H$ = 10$^4$ cm$^{-3}$) by 3 orders of magnitude from V$_{shock}$ = 18 to V$_{shock}$ = 40~\\kms . A model making use of C-shocks to produce the observed uniform \\hh\\ \\vone\\ distribution would have to be somewhat contrived. While there may be some shock-excited \\hh\\ emission from the Galactic Center, it is difficult to argue away the expected PDR emission and then construct a reasonably simple shock model capable of explaining the observations. A reliable test of the excitation mechanism would be to observe \\hh\\ transitions arising higher above ground than the \\vone\\ line.    \\subsection{ Total \\hh\\ Luminosity }  To estimate the total \\hh\\ luminosity of the Galactic Center, we extrapolate from our 400 pc long strip by assuming that the scale height of the \\hh\\ emission equals that of the far-IR radiation ($h\\simeq 0\\fdg2$, Odenwald \\& Fazio 1984). For $\\ak = 2.5$~mag and $f \\leq 1$ (see Section~\\ref{sec:extinction}), the de-reddened \\hh\\ \\vone\\ luminosity in the inner 400~pc diameter of the Galaxy is $8.0 \\times 10^3$~\\lsun. Joseph (1989) gives ranges of \\hh\\ \\vone\\ luminosity in $> 1$~kpc regions for various classes of galaxies: (1) merging galaxies: $3 \\times 10^6 - 3 \\times 10^8\\ \\lsun$; (2) interacting galaxies: $10^5 - 10^7\\ \\lsun$; (3) barred spirals: $10^4 - 10^6\\ \\lsun$. Over its inner $\\sim 1$~kpc, our Galaxy most likely falls within the range for barred spirals.  In ultraluminous infrared bright galaxies ($L_{IR} \\gtrsim 10^{12}\\ \\lsun$), Goldader et al. (1995) show that log($L_{S(1)}/L_{FIR}$) $=$ $-4.95 \\pm 0.22$. For the nearby starburst M82. we can use \\hh\\ \\vone\\ measurements of the inner 60\\arcsec\\ (Pak \\& Jaffe, unpublished) together with far-IR continuum observations (D. A. Harper, as quoted in Lugten et al. 1986) to derive log(L$_{S(1)}$/L$_{FIR}$) = $-5.2$ for the inner 1 kpc. For the inner 400 pc of the Milky Way, the data presented here yield log(L$_{S(1)}$/L$_{FIR}$) = $-4.8$.  There is evidence in some high-luminosity galaxies that the \\hh\\ emission results from UV-excitation. In NGC~3256, a merging galaxy, the \\hh\\ \\vtwo /\\vone\\ line ratio in the 600~pc region ($3\\farcs5 \\times 3\\farcs5$) is $0.39 \\pm 0.06$, suggesting that UV fluorescence is responsible for at least 60~\\% of the \\hh\\ \\vone\\ emission (Doyon, Wright, \\& Joseph 1994) . If \\hh\\ in the Galactic center is UV-excited, as we suggest here, this mechanism could be shared by \\hh\\ emission from galaxies with an enormous range of nuclear conditions."
    },
    "9604/astro-ph9604026_arXiv.txt": {
        "abstract": "In this paper we apply the jet-disk symbiosis model developed for Sgr A* to M81* -- the nucleus of the nearby galaxy M81. The model accurately predicts radio flux and size of M81* for the observed bolometric luminosity of the nuclear source with no major free parameter except for the inclination angle. We point out that the usually applied free, conical jet emission model implies a longitudinal pressure gradient that must lead to a moderate acceleration of the jet along its flow direction. This, usually neglected, gradual acceleration naturally accounts for the inverted spectrum and the size/frequency relation of M81* and may be a general feature of radio cores.  M81* is so far the best case for a radio-loud jet nature of the compact radio core in the nucleus of a nearby spiral galaxy. The fact that one can account for Sgr A* and M81* with the same model by simply changing the accretion rate, strongly supports the jet-disk symbiosis model as an explanation for the compact radio cores of galaxies in general. ",
        "introduction": "Quite a few nearby galaxies seem to have compact radio cores in their nuclei, prominent cases are the Milky Way (Sgr A*), M31 and M81. Those radio cores resemble the cores of radio loud quasars, showing a very high brightness temperature and a flat to inverted radio spectrum that extends up to submm wavelengths. Several models have been developed to explain those cores in the context of black hole accretion: Melia (1992a\\&b) suggested a spherical accretion model for Sgr A* and, what he called, M31*. Falcke et al. (1993) and Falcke, Mannheim, \\& Biermann (1993, hereinafter FMB93) proposed an alternative model, where Sgr A* was explained as the core of a radio jet, fed by an underluminous, starving accretion disk (see also Falcke 1996a, for a review, and Falcke \\& Heinrich 1994 for M31*).  This model evolved into the jet-disk symbiosis approach (Falcke \\& Biermann 1995, hereinafter FB95; Falcke 1996b) that unified the explanation for radio cores of quasars (Falcke, Malkan, \\& Biermann 1995, hereinafter FMB95), galactic jet sources (Falcke \\& Biermann 1996) and Sgr A* into a single picture. The basic idea was to postulate that black holes, jets and disks form closely coupled systems with little variations from one system to another except for the accretion rate.  Recent VLBI (Bietenholz et al. 1996), multiwavelengths (Ho, Filipenko \\& Sargent 1996) and submm observations (Reuter \\& Duschl 1996) have now shown that the radio core in M81 -- which in analogy to Sgr A* and M31*, we will label M81* hereafter -- is very similar to Sgr A*. This wealth of data now makes M81* an excellent laboratory to test the jet-disk symbiosis in detail. ",
        "conclusions": "The jet-disk symbiosis emission model by FMB93 and FB95 which was initially developed to explain the radio core of Sgr A* and radio loud quasars can also explain the nuclear radio source in M81 in detail. Simply by scaling the accretion disk luminosity by several powers of ten to the observed value properly predicts flux, size, spectral index, and size index of the VLBI radio source.  If one inserts M81* into the universal $L_{\\rm disk}/$radio correlation presented in Falcke \\& Biermann (1996) it falls right onto the line that connects Sgr A* and radio loud quasars\\footnote{This was already shown in a preliminary way in Falcke (1994, Fig. 8.1).}. This suggests that all these sources are powered by a very similar (jet/disk) engine.  In this context it is quite interesting to note that M81* and Sgr A* are in spiral galaxies, yet appear radio loud in these diagrams. M81 also has recently shown double-peaked broad emission lines, a feature usually seen only in radio-loud galaxies (Bower et al. 1996).  Like in quasars and Sgr A*, one has to choose the most efficient, equipartition jet model to explain M81*, with high internal energy and a jet power that equals the disk luminosity. This leaves one with no free parameters for M81* other than the jet inclination angle which according to the current model is around $i\\sim 30-40^\\circ$. Some caution is, however, neccessary, as the PA of the VLBI component seems to change with frequency, something we have ignored here. It is not clear whether this indicates bending, helical motion, or an extrinsic effect and thus is difficult to interpret. Some of these effects could change the results slightly -- but not to an order of magnitude.  In this {\\it Letter} we have considered for the first time the effects of the longitudinal pressure gradient in the BK79 and FB95 jet model. Without any sophisticated mechanisms, this gradient alone will already lead to a moderate acceleration of the jet to bulk Lorentz factors of 2-3. Consequently, the Lorentz factor will increase towards lower frequencies and for a fixed viewing angle outside the boosting cone the flux will become Doppler-dimmed with respect to higher frequencies.  This can provide a natural explanation for the inverted spectrum seen in compact radio cores in general and for the size/frequency relation of M81*. Such a mildly accelerating jet may also be of interest for quasar and BL Lac radio cores or galactic jet sources. The biggest advantage, however, is that the jet velocity, which was previously a free parameter, now is fixed by a simple physical model.  Another interesting point of the presented model is the usage of an initial, quasi monoenergetic electron distribution for which we find $\\gamma_{\\rm e}\\sim200$. This not only reduces another free parameter (the electron distribution index) but also can explain the submm-IR cut-offs seen in Sgr A* and M81*. If the jet in M81* does not have strong shocks to re-accelerate those electrons into the ususal power-law, the jet would have size-dependent high-frequency cut-offs and thus explain the apparent absence of extended VLBI components in M81*."
    },
    "9604/astro-ph9604160_arXiv.txt": {
        "abstract": "We prove that the sonic surface of axi\\-symmetric meridional stationary flows is always attached to the accretor, however small, if the adiabatic index of the gas is $\\gamma=5/3$.\\\\ Using local expansions near a point-like accretor, we extend Bondi's classification of spherically symmetric flows to axisymmetric flows, introducing the possibility of angular sectors reached by no flow lines, and singular directions of infinite mass flux, in addition to the angular regions of subsonic and supersonic accretion. For $\\gamma<5/3$, we show the impossibility of subsonic accretion onto a point--like accretor when the entropy of the flow is not uniform. The special case $\\gamma=5/3$ is treated separately.\\\\ We analyse the influence of the adiabatic index and Mach number of the flow at infinity on the mass accretion rate of shocked spherical flows. We propose an interpolation formula for the mass accretion rate of axisymmetric flows as a function of the Mach number and the adiabatic index, in the range $9/7<\\gamma<5/3$. ",
        "introduction": "Numerical simulations of the Bondi--Hoyle--Lyttleton (hereafter BHL) accretion flow, in 2--D and 3--D, have not only enabled the determination of the mean rate of accretion of mass, linear and angular momentum, but have also significantly modified our understanding of its dynamics. They have revealed that the accretion column foreseen by Hoyle \\& Lyttleton (1939) widens to form a detached bow shock if the adiabatic index of the accreted gas is $\\gamma=5/3$ (Hunt~1971) or $\\gamma=4/3$ (Hunt~1979). The shock structure observed in simulations was used by Eadie \\etal (1975) to interpret the X--ray observations of Vel X-1.\\\\ Real situations like wind accretion in X-ray binaries  require taking into account the density and velocity gradients in the incoming flow (see Ishii \\etal~1993 and references therein), the effects of rotation (see Theuns \\etal~1996 and references therein) as well as the intricacies of radiative transfer (see an overview in Blondin~1994). For the sake of a better understanding, we restrict our present investigation to the axisymmetric meridional accretion of a gas with uniform adiabatic index onto a gravitating accretor with constant velocity. \\\\ Although this configuration is highly simplified, a completely unexpected non--axisymmetric instability was discovered in numerical simulations for supersonic flows (Fryxell \\& Taam~1988). This instability could explain the observed variability of the radio or X-ray luminosity in some astrophysical systems (\\eg Ruffert \\& Melia~1994).\\\\ However, its is not easy to establish the consistency of the results of the few published 3--D numerical simulations which depend on the details of the numerical method: no instability appeared in the 3--D SPH simulations by Boffin (1991) whereas it was observed in the 3--D Eulerian simulations by Matsuda \\etal (1992 and references therein) and Ruffert (1995 and references therein). Dome--shaped structures frequently appeared with Roe's method (approximate Riemann solver) of Matsuda \\etal (1991), but not in the nested grid Piecewise Parabolic Method (PPM) of Ruffert (1991). Since numerical viscosity appears to play an important role, a sufficient resolution of the simulation is crucial. In this respect, numerical simulations encounter the fo\\-llow\\-ing difficulty: the smallness of the accretor in simulations is limited  by the computational power, because the shortest timestep of the simulation is usually determined by the region of high velocities in the vicinity of the accretor. The ratio of scales between the accretion radius and the radius $\\rs$ of a compact star, moving with a supersonic velocity $v$ is typically of order: \\begin{equation} {r_{\\rm A}\\over \\rs}\\sim 10^{5}\\left({10^3{\\rm kms^{-1}}\\over v}\\right)^2 \\left({r_{\\rm Schw.}\\over \\rs}\\right)\\;,\\label{ratio5} \\end{equation} where $r_{\\rm Schw.}$ is the Schwarzschild radius. Note that in the case of accretion onto a neutron star, the relevant accretor boundary is the boundary of its magnetosphere, typically a factor $10-100$ larger than the Schwarzschild radius. Some of the most recent 3--D simulations (Ruffert~1992, 1994a) could allow a ratio up to $r_{\\rm A}/\\rs\\sim10^2$ (the full domain of the simulation extends up to $\\sim 16$ accretion radii), by using a PPM code with multiple nested grids. Such an accreting sphere is still more than 100 times larger than $r_{\\rm Schw.}$, and typically $10$ times larger than the magnetosphere of a neutron star. One would like to be able to predict, on an analytical basis, which changes are likely to occur if the size of the accretor is further decreased. The role of the boundary conditions at the surface of the accretor is closely related to the position of the sonic surface, and we shall use analytical arguments to obtain some insights on this question.  An analytical description of axisymmetric stationary flows is also a first step towards the analysis of their stability, which we postpone to a future paper. Bisnovatyi-Kogan \\etal (1979) found self-similar stationary solutions, and one would like to know to what extent any generic stationary solution would behave in the same way as their solution. More generality is gained by studying, in the same spirit as Theuns \\& David (1992) for the spherical case, the first order behaviour of all quantities near the accretor for an axisymmetric flow, and in particular their departure from sphericity.  Some analytic formulae have been proposed to fit the mass accretion rate observed in simulations (Hunt~1979; Ruffert~1994b), but none of them filled the gap between the well studied cases of axisymmetric accretion of dust (Hoyle \\& Lyttleton~1939; Bondi \\& Hoyle~1944) on the one hand, and spherical accretion of a gas (Bondi~1952) on the other, for an arbitrary value of the adiabatic index . Continuing in this direction, we add to these reference models the spherical accretion of a shocked gas with uniform adiabatic index, in order to see the effect of both the kinetic and the thermal energies, and extend this approach to axisymmetric meridional stationary flows.\\\\  We first recall in Sect.~\\ref{Sgeneral} the general equations gover\\-ning a sta\\-tion\\-ary flow and form\\-ulate the axi\\-symmetric problem with a single equation in cylindrical coordinates. The shape of the sonic surface is analysed in Sect.~\\ref{Ssonic}. We establish some properties of the flow by performing local expansions in the vicinity of the accretor in Sect.~\\ref{Sexpansions}.  We present in Sect.~\\ref{Sspherical} the analytical solution of the spherical accretion from a gas with constant kinetic and thermal energies at infinity. An interpolation formula for the mass accretion rate of axisymmetric meridional flows is proposed in Sect.~\\ref{Sinterp}. We summarize our conclusions in Sect.~\\ref{Sconclusion}. ",
        "conclusions": "}  By using a geometrical interpretation of the shape of the flux tube at the sonic radius, we have shown the particular significance of the radius $r_0$ (Eq.~\\ref{rspherique}), which depends only on the adiabatic index and the energy of the flow. It is not only the sonic radius of all spherical flows, but is also related through the property ${\\cal P}$ (Sect.~\\ref{SSprop}) to the sonic surface of axisymmetric flows.\\\\ By reducing the stationary flow equations to a single partial differential equation, we have performed a local analysis near the accretor and have obtained some general results which aim at clarifying the diversity of possible configurations. We have extended the classification introduced by Bondi (1952) to axisymmetric stationary flows. In addition to the subsonic and supersonic types of accretion, angular sectors without any accretion are also possible, as well as singular directions of infinitely high accretion flux. This is illustrated by the self-similar solution found by Bisnovatyi-Kogan \\etal (1979), which is a very particular singular case ($\\theta_{\\rm c}=0$).\\\\ The angular sectors without any accretion is subsonic for flows with a detached shock, so that $\\theta_{\\rm so}\\le\\theta_{\\rm c}$ in the notation introduced in Sect.~\\ref{notations}.\\\\ Among some of the properties that we established for flows with $\\gamma<5/3$ are the following: \\par(i) the sonic surface is likely to be attached to the accretor for nearly isothermal flows with high Mach numbers, \\par(ii) isentropic accretion proceeds mostly from the downstream hemisphere of the accretor, \\par(iii) subsonic accretion must be isentropic.\\\\ We have also found a series of properties for flows with $\\gamma=5/3$: \\par(i) accretion is always regular, \\par(ii) the sonic surface is always attached to the accretor, \\par(iii) the pressure distribution is spherically symmetric to first order, \\par(iv) the Mach number and the entropy at a point like accretor are algebraically related to the local mass flux, \\par(v) shocked matter is accreted mostly from the downstream hemisphere, where the density, velocity and Mach number are maximum, and the temperature and entropy are minimum, \\par(vi) isentropic accretion is spherically symmetric to first order for the velocity, temperature, density, mass flux, and Mach number.\\\\  3--D numerical simulations have shown that the strongest instabilities occur with the adiabatic index $\\gamma=5/3$. The fact that subsonic regions invariably reach the accretor for $\\gamma=5/3$ might be considered as a warning about the possible role, in the instability mechanism, of the boundary conditions at the surface of the accretor. For $\\gamma<5/3$ and sufficiently small accretors, it would be interesting to check for the presence of the instability in cases where the accretor is surrounded by a supersonic region (Type~F in the nomenclature of Sect.~\\ref{SStypes}).\\\\  We have also determined the departure from sphericity, and constrained the radial variation of the angle $\\beta$ for $r\\to0$. In particular, we have stressed the influence of the entropy gradient on the bending of the flow lines for supersonic flows with $\\gamma<5/3$ and $\\gamma=5/3$.\\\\  We have stressed the role of the geometrical shape of the shock in the determination of the entropy of the accreted matter, and contrasted it with the simpler case of shocked spherical flows. In this latter case, we have shown the relevance of the entropy-energy approach, and noted that the mass accretion rate scales like the Hoyle--Lyttleton formula for high Mach numbers. Moreover, we have outlined the particularity of nearly isothermal flows for which the mass accretion rate can be a factor $2/(\\gamma-1)$ higher than the Bondi mass accretion rate, when the kinetic energy at infinity is comparable to the thermal energy. This demonstrates the fundamental difference of the BHL flows that have $\\gamma$ close to 1 and those with $\\gamma$ close to 5/3. \\\\ We have shown that the total mass accretion rate equals the mass accretion rate through the sonic surface, for shocked regular flows with $\\gamma<5/3$. We have derived an interpolation formula for the mass accretion rate valid in the range $9/7<\\gamma<5/3$. It links spherical accretion to the case of highly supersonic accretion, with reasonable agreement with numerical simulations.\\\\  Analytical estimates of the following quantities, however, are still missing: \\par(i) the distance of the shock as a function of $\\gamma$ and $\\M_\\infty$, \\par(ii) the value of $\\theta_{\\rm so}$ and $\\theta_{\\rm c}$ as a function of $\\gamma$ and $\\M_\\infty$, \\par(iii) the mass accretion rate in the limit of infinite Mach numbers, as a function of $\\gamma$.  Nevertheless, our classification will help perform, in a forthcoming paper, a local stability analysis of all the configurations considered."
    },
    "9604/astro-ph9604095_arXiv.txt": {
        "abstract": "We present a schematic model for the formation of baryonic galactic halos and hot gas in the Local Group and the intergalactic medium. We follow the dynamics, chemical evolution, heat flow and gas flows of a hierarchy of scales, including: protogalactic clouds, galactic halos, and the Local Group itself. Within this hierarchy, the Galaxy is built via mergers of protogalactic fragments. Hot and cold gas components are distinguished, with star formation occurring in cold molecular cloud cores, while stellar winds, supernovae, and mergers convert cold gas into a hot intercloud medium. We find that early bursts of star formation lead to a large population of remnants (mostly white dwarfs), which would reside presently in the halo and contribute to the dark component observed in the microlensing experiments.  The hot, metal-rich gas from early starbursts and merging evaporates from the clouds and is eventually incorporated into the intergalactic medium.  The model thus suggests that most microlensing objects could be white dwarfs ($m \\sim 0.5 \\msol$), which comprise a significant fraction of the halo mass. Furthermore, the Local Group could have a component of metal-rich hot gas similar to, although less than, that observed in larger clusters.  We discuss the known constraints on such a scenario and show that all local observations can be satisfied with present data in this model.  The most stringent constraint comes from the metallicity distribution in the halo. The best-fit model has a halo that is 40\\% baryonic, with an upper limit of 77\\%. Our model predicts that the hot intragroup gas has a total luminosity $1.5 \\times 10^{40} \\ {\\rm erg} \\ {\\rm s}^{-1}$, and a temperature of 0.26 keV, just at the margin of detectability. Improved X-ray data could provide a key constraint on any remnant component in the halo. ",
        "introduction": "Recently there has been renewed interest in the nature of the dark matter in galaxy halos, motivated by the results of microlensing experiments.  Observations toward the Magellanic clouds (e.g., Alcock et al.\\ \\cite{macho95}; Aubourg et al.\\ \\cite{eros93}) and toward the Galactic bulge (e.g., Udalski et al.\\ \\cite{ogle94}) have detected gravitational microlensing and inferred the presence of dark, massive compact halo objects (MACHOs).  Furthermore, a recent binary detection in the direction of the LMC, along with average event durations of about 2.5 months, may imply masses of order $\\sim 0.5 \\msol$, suggestive of white dwarfs.  At the same time, X-ray observations of clusters (Mushotzky \\cite{cluster}) and groups (e.g., Mulchaey, Davis, Mushotzky, \\& Burstein \\cite{mdmb,mdmb93}; Pildis, Bregman, \\& Evrard \\cite{pbe}; Ponman, Bourner, Ebeling, \\& B\\\"ohringer \\cite{pbeb}) have discovered a large amount of hot, metal-rich gas. Where it is observed, this gas is a substantial fraction of the baryonic mass--it is the dominant baryonic component of clusters, and is comparable to the galactic component in groups. Indeed, it appears that most baryons in the universe are in the for of this hot X-ray emitting gas.  The relatively high metallicity ($Z \\sim 0.3 Z_\\odot$ for clusters, $Z \\sim 0.1 Z_\\odot$ for groups) of this gas is impressive and demands that the material has undergone a significant amount of stellar processing during an earlier epoch of star formation. Such an epoch would also produce remnants, mostly white dwarfs. If this epoch is a general consequence of the formation of the bulge and halo of spirals like the Milky Way, as well as the ellipticals of rich clusters, it could account for the observed microlensing objects.  Although white dwarfs are attractive MACHO candidates, there are important constraints on such objects and their formation (Ryu, Olive, \\& Silk \\cite{ros}). These include background light from the early evolution, the present luminosity of the halo, and the metal and helium content of the disk and halo stars. While these place important constraints on model parameters, they do not rule out a significant white dwarf halo population, as we shall show.  In modeling these galaxy aggregates and their hot gas components, one must account for the dependence of these systems on the morphology of the constituent galaxies. On the one hand, the hot gas in clusters and groups appears correlated with the luminosity of elliptical and S0 galaxies (Arnauld et al.\\ \\cite{arolv}). This suggests that the hot gas arises from the violent merging associated with these morphological types. On the other hand, it also seems well established (e.g., Rich \\cite{rich}) that the morphology of the bulge and halo of spiral galaxies is quite similar to that of ellipticals and S0's. This suggests that the bulge and halo of spirals may have experienced a similar epoch of star formation and outflow during their formation.  The difference in the morphologies may relate to the larger angular momentum or shallower gravitational potential of spirals (Zurek et al.\\ \\cite{zqs}), such that some gas survives halo formation, and settles afterwards into spiral arms.  With this background in mind, we find that a likely, and perhaps inevitable, consequence of the formation of the bulge and halo is the formation of a large remnant population in the Galactic halo, along with hot X-ray emitting gas in the Local Group and intergalactic (i.e., extragroup) medium. While our model should be widely applicable, in this paper we concentrate on the Local Group.  Given the detection of dark microlensing objects, as well as the need for significant amounts of dark baryonic matter {\\it somewhere} (Copi, Schramm, \\& Turner \\cite{cst}; Fields, Kainulainen, Olive, \\& Thomas \\cite{fkot}) halo white dwarfs are a very conservative candidate (e.g., Larson \\cite{lars}; Ryu, Olive, \\& Silk \\cite{ros}; Silk \\cite{silk}). This is particularly so since red and brown dwarfs are apparently excluded as halo candidates (Bahcall, Flynn, Gould, \\& Kirhakos \\cite{bfgk}; Graff \\& Freese \\cite{gf}). Here we make a specific though schematic model, and assess the plausibility of the white dwarf hypothesis.\\footnote{Others have suggested that the microlensing objects might be remnants from an early Population III; see Fujimoto, Sugiyama, Iben, \\& Hollowell (\\cite{fsih}).} As we will see, the model can be made to work but not without some assumptions (e.g., one must alter the halo initial mass function). In any case, the model is eminently testable, and perhaps has already been tested by X-ray observations in other groups. Indeed, should one find this model and the halo white dwarf hypothesis untenable, then it follows that the dark baryons and the MACHOs must take an even more exotic form. ",
        "conclusions": "\\label{sec:imp}  We have shown that, without violating constraints posed by luminosity and nucleosynthesis considerations, one may construct a plausible model in which the dark halo of the Galaxy contains a significant fraction of white dwarfs. These may have already been detected in halo microlensing events towards the LMC, and might also be detected via their luminosity function. The same bursts of star formation which produced the white dwarfs also led to hot, metal-rich intergalactic gas, some of which may still reside in the Local Group. This hot gas could be detectable via its X-rays, and by distortions in the cosmic microwave background radiation (Suto et al.\\ \\cite{smio}).  Thus, the predictions of the model are testable.  If the halo is comprised of white dwarfs then there must be a background of hot, X-ray emitting gas in the Local Group.  Conversely, if there is metal-rich hot gas in the Local Group, then a significant fraction of the halo mass must be in remnants. Clearly, further searches for both of these are warranted.  Furthermore, if our galaxy formation scheme is indeed universal, then hot gas production and ejection should be a ubiquitous aspect of halo formation. Consequently, X-ray observations of other systems could provide a key constraint on our model.  In particular, our model can be directly tested by observations which can unambiguously confirm or deny the presence of hot gas in other spiral-dominated groups. Also, if white dwarfs are ubiquitous in galactic halos, then they may lead to detectable infrared profiles in edge-on galaxies, which may already have been observed (Barnaby \\& Thronson \\cite{bt94}; Sackett, Morrison, Harding, \\& Boroson \\cite{smhb}; Lequeux, Fort, Dantel-Fort, Cuillandre, \\& Mellier \\cite{lfdcm}; Lehnert \\& Heckman \\cite{lh}). Finally,  even if our scenario turns out not to be applicable to spiral-dominated groups, it remains that ellipticals must eject gas in strong winds. Thus, our model may still be valid for clusters, which are elliptical-dominated.  This will be explored in a subsequent work.  We note as well that in our scenario, just as there is typically a large outflow from the halo, there is also a strong evaporative wind that ejects material from the Local Group. As a result, most baryons eventually reside in hot, intergalactic (as opposed to intragroup) gas.  If this gas stays hot, it could perhaps be the ionized intergalactic (as suggested by Gunn-Peterson limits on the neutral intergalactic medium). If it does cool, it presents serious problems, as it would lead to prodigious but unobserved absorption of extragalactic radiation.  Finally, we reiterate that stellar remnants and their associated hot ejecta are conservative candidates for both the halo microlensing objects and for the baryonic dark matter.  If these can be ruled out, then we are forced to conclude that the microlensing objects and the dark baryons are something stranger still."
    },
    "9604/astro-ph9604141_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  Since clusters of galaxies are the largest virialized structure in the universe, they should retain the initial conditions at their formation epochs fairly faithfully. This implies that detailed studies of the clusters at high redshifts, as well as at present, should provide important clues to the evolution of the universe itself.  Since X-ray identifications of clusters of galaxies are largely free from the projection effect which notoriously plagues optically selected cluster catalogues, X-ray observations are suitable for probing cosmological signatures from clusters of galaxies. Homogeneous samples of distant clusters of galaxies, which recent X-ray satellites such as {\\it ROSAT} and {\\it ASCA} are constructing, will uncover various statistical properties of clusters with higher reliability in the near future. Therefore, quantitative theoretical predictions are of great value in interpreting the observed data properly.  Most theoretical approaches in X-ray cosmology rely on either state-of-the-art numerical simulations, or simplified analytical formalisms. The former approach is limited by the dynamical range available on the present computer resources; a typical core radius of clusters is ($0.1\\sim 0.2 \\himpc$) while the mean separation of the Abell clusters (richness class 1) is $\\sim 55\\himpc$, where $h$ is the Hubble constant $H_0$ in units of $100 \\mbox{km}\\cdot\\mbox{sec}^{-1}\\cdot\\mbox{Mpc}^{-1}$. A small-scale resolution much below the core size is essential because a large fraction of the X-ray luminosity of clusters comes from the core. On the other hand, a large simulation box size is a prerequisite for statistical studies of clusters. Unfortunately, it is still hard to simultaneously satisfy these requirements even with the currently most advanced simulations (e.g. Kang et al. 1994; Bryan et al. 1994; Cen et al. 1995).  A major fraction of the X-ray luminosity from clusters is produced via a fairly simple process, thermal bremsstrahlung. Thus one may readily compute their temperature and luminosity functions at redshift $z$, $n_{\\rm T}(T,z)$ and $n_{\\rm L}(L,z)$, once the mass function $n_{\\rm M}(M,z)$ is given, where $T$, $L$ and $M$ are the temperature, luminosity and mass of the clusters, respectively. Although the Press--Schechter theory (Press \\& Schechter 1974, hereafter PS) is frequently applied for this purpose, it has a serious limitation in predicting the temperature and luminosity functions; PS theory predicts the number density of virialized objects of mass $M$ collapsed {\\it before} a given epoch $z$, but does not specify the formation epoch $z_{\\rm f}$ of the objects. In fact, the predictions of the spherical nonlinear collapse model (e.g. Peebles 1980) suggest that the temperature and luminosity of objects that virialize at $z_{\\rm f}$ should scale as $T(z_{\\rm f}) \\propto (1+z_{\\rm f})$ and $L(z_{\\rm f}) \\propto (1+z_{\\rm f})^{7/2}$ in the Einstein--de Sitter universe, for instance. The previous approaches based on the PS formula (e.g. Evrard \\& Henry 1991; Hanami 1993) have simply replaced $z_{\\rm f}$ by $z$ in computing $T$ and $L$ (see also eq.~[\\ref{tempfnps}] below). This procedure corresponds to implicitly assuming that $T$ and $L$ of individual clusters {\\it decline} with time as $T(z)/T(z_{\\rm f})= (1+z)/(1+z_{\\rm f})$ and $L(z)/L(z_{\\rm f})= (1+z)^{7/2}/(1+z_{\\rm f})^{7/2}$.  Numerical simulations, on the contrary, suggest modest evolution in the opposite direction (e.g. Evrard 1990; Suginohara 1994; Navarro, Frenk, \\& White 1995). While the above assumptions may not alter $n_{\\rm T}(T,z)$ and $n_{\\rm L}(L,z)$ for larger clusters most of which would have formed only recently ($z_{\\rm f} \\sim z \\ll 1$), it is likely to affect the results for less massive clusters.  This line of consideration motivates us to specify explicitly the formation epoch of virialized structures and their subsequent evolution in making statistical predictions for comparison with observations. A key quantity for this purpose is a distribution function of halo formation epochs proposed by Lacey \\& Cole (1993, hereafter LC) and Kitayama \\& Suto (1996, hereafter KS) in a similar but slightly different manner (see also Blain \\& Longair 1993; Sasaki 1994). In this paper, we combine these distribution functions with a simple model of cluster gas properties to make quantitative predictions on the temperature and luminosity functions of clusters of galaxies in cold dark matter (CDM) universes with/without a cosmological constant $\\lambda_0$.  The plan of this paper is as follows. Section 2 outlines the formalism we use in computing the temperature and luminosity functions. Section 3 describes a model of X-ray clusters, and our main results are presented in Section 4. Finally, Section 5 summarizes our conclusions. ",
        "conclusions": "\\setcounter{equation}{0}  We have examined statistical properties of X-ray clusters in CDM universes in a semi-analytic manner, which is complementary to numerical simulations. Our method is different from the previous approaches based upon the conventional PS theory (e.g. Evrard \\& Henry; Hanami 1993), in that we explicitly took account of the epochs of cluster formation adopting the halo formation epoch distributions proposed by LC and KS. As a result, our method can include the subsequent evolution of cluster temperature and luminosity.  Although the LC and KS proposals involve slightly different definitions of the halo formation epochs, we found that the resultant temperature and luminosity functions are remarkably similar. Deviations from the previous PS approach become larger for clusters that form at higher redshift and for stronger evolution.  The shape and amplitude of the temperature and luminosity functions vary sensitively with the cosmological density parameter $\\Omega_0$, if the fluctuation amplitude is fixed by the {\\it COBE} normalization. Given the qualitative nature of our simple model, however, we should not constrain the cosmological parameters so stringently. Rather we conclude that the low-density $\\Omega_0 \\sim 0.2 - 0.5$ CDM models with/without the cosmological constant are roughly consistent with the observed temperature and luminosity functions. Nevertheless we can argue against $\\Omega_0 =1$ CDM models at least from the present analysis only.   \\vskip1.2cm \\noindent We thank Takahiro T. Nakamura and Naoshi Sugiyama for useful discussions on the spherical collapse model and the {\\it COBE} normalization. We are grateful to Ewan D. Stewart for a careful reading of the manuscript. This research was supported in part by the Grants-in-Aid by the Ministry of Education, Science and Culture of Japan (07740183, 07CE2002).  \\vfill\\eject \\centerline{\\bf APPENDICES}  \\appendix \\renewcommand{\\thesection}{\\normalsize\\bf\\Alph{section}} \\renewcommand{\\thesubsection} {\\normalsize\\it\\Alph{section}.\\normalsize\\it\\arabic{subsection}. } \\renewcommand{\\theequation}{\\mbox{\\rm {\\Alph{section}.\\arabic{equation}}}}"
    },
    "9604/astro-ph9604155_arXiv.txt": {
        "abstract": "We report the characterization of bright, compact features in the cosmic microwave background radiation (CMBR) detected during the June 1992 and June 1994 balloon flights of the Medium Scale Anisotropy Measurement (MSAM1-92 and MSAM1-94, respectively). Spectral flux densities are determined for each feature at 5.7, 9.3, and 16.5~\\icm. No viable counterparts for these features were found in source catalogs at 5~GHz or at 100 \\micron. The measured spectrum of each feature is consistent with a temperature fluctuation in the CMBR. The existence of these features is consistent with adiabatic fluctuation models of anisotropy in the CMBR. ",
        "introduction": "Recent observations of the cosmic microwave background radiation (CMBR) have found statistically significant anisotropy at 0\\fdg5 angular scales (see \\cite{white94}, and references therein). In earlier Letters (\\cite{cheng94}, hereafter Paper~I, and \\cite{cheng95}, hereafter Paper~II), we reported on observations of CMBR anisotropy that included two particularly bright features. The presence of bright, unresolved sources suggested the possibility of foreground point-source contamination, or of non-Gaussian fluctuations in the CMBR. In this Letter, we present our analysis of these features and a discussion of other work studying their properties. ",
        "conclusions": "All three features are best described as temperature anisotropies in the CMBR, although the possibility remains that one of these is due to galactic foreground dust emission.  While we were initially surprised to find such bright, compact features in the data, we now understand that these are consistent with our assumption of Gaussian fluctuations.  Although our catalog searches were all null, we find the most compelling argument that the features are not true point-sources comes from the direct search by \\cite{church95}. Finally, the morphological agreement reported by the Saskatoon experiment, consistent with a CMBR anisotropy spectrum spanning three octaves in frequency from Q-band (1.2 \\icm) to the MSAM 9.3 \\icm\\ band, leaves little room for alternate explanations of the detection. Given the agreement with a CMBR temperature anisotropy spectrum, together with the implausability of other explanations, we conclude that at least the first two, and possibly all three of these detections, are cosmological in origin."
    }
}