{
    "9603/astro-ph9603054_arXiv.txt": {
        "abstract": "A brief introduction is given to some aspects of the statistical description of the luminous matter distribution. I review the features of the redshift surveys that arise in the statistical analysis of the galaxy clustering. Special topics include intensity functions, correlation functions, correlation integrals, multifractals and multiscaling. ",
        "introduction": "I am sure that, in spite of the title of the School, at the end of these two weeks, we shall have a less dark picture of the three-dimensional distribution of matter in the Universe. There are still a huge amount of unsolved problems regarding the origin and evolution of the observed large scale structure in the Universe. Although important developments have occurred during the last two decades, the task has revealed so elusive, that most of the students of this School will have interesting research projects on these topics in the following years.  The statistical study of the clustering patterns formed by the three-dimensional distribution of galaxies is one of the most important observational clues to learn about the physical processes that led to the large scale structure of the Universe. A detailed statistical description of the observed distribution of matter in the Universe is needed to confront theoretical predictions of models of structure formation, such as $N$-body simulations involving dark matter, against  observations.  The aim of this lecture is to introduce some statistical aspects of the description of the clustering in the Universe. This introduction will be followed by more detailed lectures given by Drs. Borgani and Coles. ",
        "conclusions": "We have given a short introduction to different aspects of the characterization of the observed surveys of galaxy clustering by means of different statistical techniques. Redshift surveys, when considered as point processes, have peculiar features which can be expressed through statistical tools. Obscuration by dust in our own Galaxy, truncation in luminosity and the use of selection functions for flux-limited samples have been discussed in some detail. The analysis of clumpiness is often done in redshift space, which has important distortions when compared to real space. Morphological and luminosity segregation is an important clue for testing galaxy formation theories. It is interesting to consider the galaxy distribution as a marked point process. We have illustrated the relationship of the two-point correlation function to other statistical quantities such as the intensity functions or cumulant quantities such as the correlation integral. The multifractal nature of the matter distribution comes from the scaling of the moments of the cell-counts. We have introduced the concept of multiscaling to provide a neat scheme for the explanation of the clustering of galaxies of different kinds and clusters with different richness. In this context, we have shown how the correlation dimension $D_2$ attains specific values for each kind of cosmic object, being a clear measure of their clustering.  \\ack  The author wishes to thank his collaborators S. Borgani, P. Coles, S. Paredes and M.J. Pons for countless discussions on all aspects of the clustering phenomenon. This work is partially supported by the EC Human Capital and Mobility Programme network (Contract ERB CHRX-CT93-0129), by the project number GV-2207/94 of the Generalitat Valenciana and by the Instituci\\'o Valenciana d'Estudis i Investigaci\\'o.  \\smallskip"
    },
    "9603/astro-ph9603112_arXiv.txt": {
        "abstract": "We have analyzed approximately 1100 days of Cygnus X-1 hard X-ray data obtained with BATSE to study its rapid variability. We find for the first time correlations between the slope of the spectrum and the hard X-ray intensity, and between the spectral slope and the amplitude of the rapid variations of the hard X-ray flux. We compare our results with expectations from current theories of accretion onto black holes. ",
        "introduction": "Accreting black-hole candidates (BHC) in X-ray binaries show three \\lq source states' that are distinguished by characteristic spectral and (correlated) fast-variability properties. They are called the \\lq low state', \\lq high state' and \\lq very high state'; the dominant parameter that determines the source state is likely to be the mass accretion rate (see Van der Klis \\cite{vdk:b}, and Tanaka \\& Lewin \\cite{Tanaka:Lewin} for recent reviews).  In the low state the X-ray spectrum is dominated by a very hard power law component extending to several hundreds of keV. The source intensity shows strong fluctuations, with broad-band r.m.s.\\ amplitudes as high as 40 percent. At frequencies above $\\sim\\! 1$ Hz the power density spectrum (PDS) of these variations follows a power law; at frequencies below a low-frequency cut off the PDS is flat. In several BHC the cut-off frequency $\\nu_{\\rm c}$ has been observed to vary by up to an order of magnitude, while the high-frequency part of the PDS remained approximately constant (Belloni \\& Hasinger \\cite{Belloni}; Miyamoto et al. \\cite{Miya:canon}). As a result, the power integrated over a frequency range below $\\nu_{\\rm c}$ is strongly anti-correlated with that cut-off frequency.  In the high state the X-ray spectrum contains an ultra-soft thermal component, with (bremsstrahlung) temperatures of order 1 keV. In some sources (e.g., LMC~X-3, see Tanaka \\cite{Tanaka:ESLAB}) the power-law spectral component is not detected, but in others (e.g., GS~1124$-$68, see Ebisawa et al.\\ \\cite{Ebisawa:Spec}) it is observed in combination with the ultra-soft component. In the 1--10 keV range the amplitude of the intensity variations is correlated with photon energy; this can be understood if the variability is connected with the power law spectral component, which is \\lq diluted' by the much lower variability of the ultrasoft emission (Van der Klis \\cite{vdk}).  In the very high state the PDS of BHC show 3--10 Hz quasi-periodic oscillations and branch structure in an X-ray color-color diagram reminiscent of the normal-branch (NB) state of Z sources (Van der Klis \\cite{vdk:a}), i.e., accreting neutron stars with magnetic fields of order $10^{10}\\!$~G. In view of the strong arguments that the accretion rate of Z sources is then close to the Eddington limit it has been suggested that in the very high state the BHC are likewise accreting near the Eddington limit (Van der Klis \\cite{vdk:b}).  In their low state the BHC are remarkably similar to atoll sources, i.e., accreting neutron stars with magnetic field strengths generally believed to be below $10^9\\!$~G (substantially weaker than those of Z sources). At low accretion rates, i.e., in their so-called \\lq island state', atoll sources have power law X-ray spectra (see, e.g., Barret \\& Vedrenne \\cite{Barret}). On average, these spectra appear to be not quite as hard as those of low-state BHC (see, e.g., Gilfanov et al.\\ \\cite{Gilfanov}); however, the distribution of the spectral slopes of BHC and atoll sources shows clear overlap (Wilson et al.\\ \\cite{Wilson}; Ebisawa et al.\\ \\cite{Ebisawa:Spec}; see also Van Paradijs \\& Van der Klis \\cite{jvp:vdk}). Also, the PDS of island-state atoll sources are very similar to those of low-state BHC. They follow a power law at high frequencies, and are flat below a cut-off frequency; the high-frequency part of the PDS of the atoll source 1608$-$522 was observed to remain approximately constant as the cut-off frequency varied, like for Cyg X-1 (Yoshida et al.\\ \\cite{Yoshida}).  As part of our attempt to gain a better understanding of the similarities between accreting BHC and neutron stars we are studying the variability of bright BHC using the almost continuous record provided by the Burst and Transient Source Experiment (BATSE) on the Compton Gamma Ray Observatory. We here report on a variability record of \\cyg, covering approximately 1100 days, and we show that the amplitude of its variations are strongly correlated with the slope of the hard X-ray spectrum, and that correlations exist between the amplitude of the variations and the total flux. ",
        "conclusions": "We have found a strong correlation between the slope of the high-energy (20-100 keV) X-ray spectrum of Cyg X-1 and both its high-energy X-ray flux and the variability thereof. During most of our observations Cyg X-1 showed strong variability with an amplitude that varied in anti-correlation with a cut-off frequency (Crary et al.\\ \\cite{Crarya}) similar to the low-state behavior in the lower-energy range described by Belloni and Hasinger (\\cite{Belloni}). It is therefore likely that we encountered Cyg X-1 mainly in the low state. During a 150 day interval beginning in 1993 September both the high-energy X-ray flux and its variability were extremely low; it is possible that Cyg X-1 had then entered a high state; this remains uncertain, due to lack of low-energy coverage.  The global correlations we have found between flux, variability and spectral hardness suggest to us that all three are determined by a basic system parameter; the mass accretion rate is the obvious candidate. A variety of models have been calculated for the structure of the accretion disks around black holes, to explain the two-component character of their X-ray spectra, and the suspected relation of these spectral components with accretion rate (Liang \\& Nolan \\cite{Liang}, and references therein; see also Haardt et al.\\ \\cite{Haardt:etal}; Chakrabarti \\& Titarchuk \\cite{Chakra}).  Most of these models invoke a very hot medium, e.g., a corona around the inner disk regions, that produces the hard power law component through upscattering of low-energy photons, and a standard Shakura-Sunyaev disk that provides the latter. Most of these models do not address source variability.  The fast variability of Cyg X-1 and other black holes has often been described in terms of shot noise models; the break frequency in the PDS reflects the decay time of the shots (Sutherland, Weisskopf, \\& Kahn\\ \\cite{Sutherland}; Miyamoto \\& Kitamoto \\cite{Miya:cyg}; Belloni \\& Hasinger \\cite{Belloni}). However, the X-ray spectral properties are usually not considered in these models.  Mineshige et al.\\ (\\cite{Mineshige:crit}, \\cite{Mineshige:disks}), proposed that accretion disks around black holes are in a self-organized critical state. Their qualitative arguments, aimed at an understanding of the low and high states, indicate that relatively hard X-ray spectra go with strong variability, and relatively soft X-ray spectra with weak variability.  Chakrabarti \\& Titarchuk (\\cite{Chakra}) recently argued that accretion disks around black holes contain a shock, at several tens of Schwarzschild radii.  In the low state, the post-shock region is quite hot, and the emergent spectrum is very hard ($\\alpha\\! <\\! 2.5$).  In this model, $\\alpha$ increases with disk accretion rate in the low state. Also in this context, Molteni, Sponholz, \\& Chakrabarti\\ (\\cite{Molteni}) found that over a range in mass accretion rates the cooling time of the flow inside the shock is comparable to the infall time scale. Under these conditions the location of the shock, and the X-ray luminosity, undergoes quasi-periodic oscillations.  The centroid frequency of these oscillations increases with the mass accretion rate, with typical values around 5 Hz for a $5 M_\\odot$ black hole.  This result may be related to the anti-correlation between $\\alpha$ and $f$ for Cyg X-1, by using the recent results of Van der Hooft et al.\\ (\\cite{vdh}) on the PDS of the black-hole transient GRO J1719$-$24. They found that the detailed shape of the PDS remained invariant under a frequency shift marked by the variation of a strong QPO peak (between $\\sim\\! 40$ and $\\sim\\! 300$ mHz), while also the power in a given (stretched and squeezed) \\lq rest frequency' interval remained constant.  This would suggest that the QPO frequency is proportional to a break frequency in the PDS of this source, and that the latter may be used as a frequency scaler as well. If we are allowed to generalize this result for GRO J1719$-$24 to Cyg X-1, at least the direction of the correlation between $\\alpha$ and $f$ found by us for Cyg X-1 would be accounted for by the model described by Chakrabarti \\& Titarchuk (\\cite{Chakra}) and Molteni et al.\\ (\\cite{Molteni}).  At super-Eddington disk accretion rates the model of Chakrabarti \\& Titarchuk (\\cite{Chakra}) predicts that the flow inside the shock is cooled by Comptonization of low energy photons from the disk, and the post-shock flow becomes predominantly radial (and converging). Comptonization in this region can still occur due to bulk motions, producing a hard tail with $\\alpha\\! \\sim\\! 2.5$. The value of $\\alpha$ increases weakly with the accretion rate; note however that in this model the X-ray luminosity may not be a good measure of the accretion rate as part of the internal energy in the disk is advected into the black hole (see also Narayan, McClintock, \\& Yi\\ \\cite{Narayan}).  In this regime the flux variability is determined by the variability of the illumination geometry of the converging inflow, and not by the change in the size of the post-shock region, as it is in the low state. The amplitude variation of the hard flux is expected to be smaller in this case (corresponding to the high state, since the increase in disk accretion rate leads to a increase in soft emission) than in the low state.  It appears, then, that these theoretical results are in qualitative agreement with the data from the low-flux episode that we have observed in Cyg X-1.  The correlations we have found for Cyg~X-1 would easily have escaped attention were it not for the 1100 days of continuous coverage provided by the BATSE all-sky monitoring capabilities. We are looking forward to an improved understanding of our results in terms of the source state framework for black holes (Van der Klis \\cite{vdk}) by combining the BATSE hard X-ray monitoring with that provided in the near future at low energies by XTE."
    },
    "9603/astro-ph9603106_arXiv.txt": {
        "abstract": " ",
        "introduction": "In recent years various techniques have been developed to isolate field halo stars in situ, so that increasingly, samples of these stars are becoming available with which to address questions about the kinematic structure of the halo. Radial velocity and distance measurements of these objects provide important constraints on models of the kinematics and more indirectly the distribution of the dark matter in which the visible galaxy is apparently embedded.  Flynn, Sommer-Larsen and Christensen (1994) have developed techniques using broadband photoelectric photometry and medium dispersion spectroscopy for isolating Blue Horizontal Branch (BHB) stars in the outer Galactic halo and have assembled a catalog in four fields of about 100 stars with line-of-sight velocities and distances (Flynn et. al. 1995).  Sommer-Larsen, Flynn and Christensen (1994, hereafter SLFC) developed a model of the kinematics of the outer halo, which fits the observations surprisingly well, given the simple nature of the model.  They assumed that the outer stellar halo is a round, non-rotating system with a density falloff with Galactocentric radius like $r^{-3.4}$ and that the Galactic rotation curve is flat to large Galactic radii ($r >> R_\\odot $) where $R_\\odot$ is the sun's distance from the Galactic center.  The Jeans equation was then solved for the radial and one-dimensional tangential velocity dispersions ($\\sigma_r, \\sigma_t$) as functions of $r$ by fitting the observed line-of-sight velocity dispersions as a function of line-of-sight distance.  Their results indicate that the halo kinematics in the outer halo are {\\it tangentially} anisotropic, whereas they are {\\it radially} anisotropic near the sun.  Specifically, near the sun, $\\sigma_r \\approx 140$ \\kms~and $\\sigma_t \\approx$ 90-100 \\kms, whereas SLFC find a major kinematic change in the outer halo $(r \\ga 10-20)$ kpc, where $\\sigma_r \\approx 80-100$ \\kms~and $\\sigma_t \\approx $130-150 \\kms.  The solution to the Jeans equation only provides the second moments of the velocity distribution in the halo. To demonstrate that the solution for $(\\sigma_r, \\sigma_t)$ is actually physical, one must show that the velocity dispersions can be realized in terms of a stationary phase-space distribution function $f$ which is everywhere non-negative in phase-space.  Although there are methods in the literature for recovering the distribution function from simple potential-density pairs and the velocity dispersions (see Binney and Tremaine (1987) pp255 and references therein), the extension of these methods to the case of the Galactic potential is a formidable analytic task.  Since we would nevertheless like to test that the SLFC model is physical, we have taken in this paper a less ambitious and more direct approach: we place a large number of test particles with the kinematic characteristics of the SLFC type model (using a simple assumption about their velocity distribution) into a realistic 3-dimensional model of the Galaxy's potential, and integrate the orbits over a Hubble time, allowing the system to phase mix.  We show in this way that a stationary system can be realized with the kinematic and spatial properties of the SLFC model.  In section 2 we develop a realistic model of the Galactic potential. In section 3 we describe our simulations of the SLFC kinematics, and we discuss the results and draw conclusions in section 4. ",
        "conclusions": "We have tested the simple model for the kinematics of the Galactic halo (in particular the outer halo) proposed by SLFC, by directly integrating particles in a realistic 3-D model of the Galactic potential, under the assumption of an initially Gaussian velocity distribution. We found that the particles relax somewhat into the potential with time, the kinematics becoming everywhere more isotropic.  We experimented with models where we increased the initial anisotropy compared to the SLFC model, and found a configuration which after a short relaxation period ($\\approx$ 1 Gyr) becomes a quite good and stationary fit to the SLFC model. Hence, the SLFC model, which shows a notable change in the velocity anisotropy from markedly radial at the sun to markedly tangential beyond about $r=20$ kpc, seems a tenable description of outer halo kinematics.  The origin of the change in the anisotropy with Galactocentric radius remains unclear. We are currently gathering further data in the outer halo at the poles and in the anti-center directions to test the model more directly. We are also currently involved in a series of hydrodynamical and N-body simulations of disk galaxy formation including star-formation in which the kinematics of successive generations of stars can be examined, and we expect the observed variation in the velocity anisotropy with $r$ of the halo to be an interesting constraint on the simulations."
    },
    "9603/astro-ph9603040_arXiv.txt": {
        "abstract": " ",
        "introduction": "The APM Galaxy Survey (Maddox \\etal\\ 1990a, b) includes about two million galaxies down to magnitude $b_J = 20.5$ over 4300 square degrees of the southern sky. It was the first machine-generated galaxy survey to cover an area of sky significantly larger than one Schmidt plate, and has proved an important survey for measuring galaxy number-magnitude counts over a wide magnitude range (\\refer{Maddox \\etal\\ 1990c}) and particularly for the most reliable measurement to-date of the angular correlation function of galaxies on large scales (\\refer{Maddox \\etal\\ 1990d}). This latter measurement was one of the first results to rule out the standard cold dark matter model of galaxy formation (eg. \\refer{Davis \\etal\\ 1985}).  Unfortunately, the APM Galaxy Survey, while complete to a faint magnitude limit of $b_J = 20.5$, is not very reliable for galaxies brighter than $b_J \\approx 16.5$. There are several reasons for this. Firstly, the surface density of galaxies brighter than $b_J \\approx 16.5$ is only about $1/20$ of the surface density of stars at the same magnitude limit even at the galactic poles. Therefore the selection of an uncontaminated bright galaxy sample requires an exceptionally reliable method of rejecting stars and merged images. Secondly, photographic emulsions have a limited dynamic range, and in order to detect images as faint as $b_J = 20.5$, the brighter images are necessarily saturated. Thirdly, bright stars have diffraction spikes and `ghost' images (\\refer{UKSTU handbook}) and large galaxies contain sub-structure. All of these factors prevent the standard APM image parameters, which were designed to classify small, faint images, from selecting a sufficiently reliable bright galaxy catalogue.  This paper describes the construction of a bright galaxy catalogue, complete to $b_J = 16.44$, using the same APM scans used for the faint survey. We developed a semi-automated method of star-galaxy separation, whereby most stellar images were rejected (losing only about 3\\% of galaxies) and the remaining images inspected by eye on a film copy of the photographic plate. The distinction is emphasized between a survey {\\em constructed} by eye, for example the \\refer{Zwicky \\etal\\ (1961-68)} catalogue or the Lick (\\refer{Shane and Wirtanen 1967}) survey, where the observer has to locate each image and {\\em then} decide whether it should be included in the catalogue, and a semi-automated survey like the present one where the observer is given the position of each image satisfying a magnitude limit, and then classifies it as a galaxy or star. It is much easier for the eye to distinguish a galaxy from a star than it is to select a complete magnitude or diameter limited sample, and so the semi-automated survey should be much more reliable.  The APM Bright Galaxy Catalogue (APM-BGC) covers almost the same area as the faint APM galaxy survey of Maddox \\etal\\ (1990a, b), including 180 out of the 185 fields of the fainter survey, an area of approximately 4,180 square degrees. Figure~\\ref{fig:fields} shows the distribution of the 180 survey fields in an equal area projection on the sky.  The construction of the present catalogue was first described by Loveday (1989). A similar survey has been carried out by Raychaudhury (1989) in the region towards the `Great Attractor', and there is an ongoing effort (Raychaudhury \\etal\\ 1994) to map out the galaxy distribution near the equator.  The layout of the paper is as follows. The construction of the bright galaxy catalogue, including star-galaxy separation and plate matching, is described in \\S\\ref{sec:constr}. Internal tests for uniformity, completeness and consistency are described in \\S\\ref{sec:int_tests} and comparison with other catalogues is made in \\S\\ref{sec:other_cats}. \\S\\ref{sec:photometry} describes the CCD calibrations used to check the APM to $b_J$ magnitude conversion, as a further test of photometric uniformity in the survey, and to define a second-order correction for photographic saturation. In \\S\\ref{sec:catalogue} we describe the catalogue data and present plots of the galaxy distribution. We compare the angular and spatial correlation functions of early and late type galaxies in the APM-BGC in \\S\\ref{sec:clust}. Finally, the properties of the catalogue are summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have described the construction of a catalogue of 14,681 galaxies brighter than $b_J = 16.44$ over a large fraction of the southern sky. The images were detected and parameterized by scanning 180 United Kingdom Schmidt Telescope plates with the Automated Plate Measuring system. Preliminary star-galaxy separation was carried out automatically using image profiles. All galaxy candidates were inspected by eye and assigned a morphological classification. A completeness of 97\\% was aimed for. By comparing image classifications in overlaps between plates (\\S\\ref{sec:ovlps}), we infer an actual completeness of 96.3\\%. As an external check on completeness, we correlated the ESO galaxy catalogue with the APM-BGC. We found that about 1.5\\% of ESO galaxies are too large to be detected by the APM machine (these galaxies are listed in Table~\\ref{tab:esobig}) and a further 3\\% of ESO galaxies were detected by the APM machine but not classified as a galaxy in the APM-BGC. Overall, we estimate that the APM-BGC is at least 95\\% complete. Most of the incompleteness is due either to high surface brightness galaxies with star-like profiles (Fig.~\\ref{fig:eso_sb}(d)) or low surface brightness galaxies which fall below our detection threshold. We plan to study surface-brightness selection effects in APM galaxy data in a future paper.  The reliability of the morphological classification in the APM-BGC was checked both internally using plate overlaps and externally by comparison with other catalogues. We conclude that the APM-BGC does not reliably distinguish between elliptical and lenticular galaxies; these classes should be combined in any statistical analysis. Time-dependent classification effects may produce an error in the type-dependent angular correlation function of $\\Delta w \\approx 0.05$. We classify fewer galaxies as late type at fainter magnitudes.  The photometry of the survey has been checked using CCD photometry of 259 galaxies. We fit a polynomial to the CCD versus APM magnitudes to correct for saturation and to define the magnitude zero-point. We find a scatter of 0.31 mag about this relation for individual galaxies. Comparison of the angular correlation functions calculated using intra- and inter-plate pairs of galaxies (\\S\\ref{sec:field-effects}) shows no evidence for significant plate-to-plate errors in photometry. The CCD$-$APM magnitude residuals are uncorrelated with each other and the plate zero-points, confirming the absence of large scale gradients in calibration.  The APM Bright Galaxy Catalogue is a reliable, new catalogue of bright galaxies which complements the fainter APM Galaxy Survey and the diameter-selected ESO survey. The catalogue is about 96\\% complete and has essentially zero contamination since every galaxy has been inspected by eye. It has proved to be a valuable source catalogue for the Stromlo-APM Redshift Survey (Loveday \\etal\\ 1992) and we hope that it will be useful for other followup work.  \\subsection* \\vspace{0.1in} \\begin{tabular}{rrr} \\hline \\hline Galaxy Type & Number & \\% \\\\ \\hline Elliptical \t\t&  1791 &  12.2 \\\\ lenticular         \t\t&  2648 &  18.0 \\\\ Spiral     \t\t&  8217 &  56.0 \\\\ Irregular/Peculiar \t&   627 &   4.3 \\\\ Unsure     \t\t&   164 &   1.1 \\\\ Merged with star \t&   975 &   6.6 \\\\ Multiple   \t\t&   259 &   1.8 \\\\ \\hline Total \t\t& 14681 & 100.0 \\\\ \\hline Holes drilled & 1456 & --- \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\clearpage \\tablecaption{\\rm The APM Bright Galaxy Catalogue for field 076} \\tablefirsthead{ \\hline \\hline \\multicolumn{1}{c}{\\rm Name} & \\multicolumn{3}{c}{\\rm RA} & \\multicolumn{3}{c}{\\rm Dec} & $b_J$ & maj & min &{\\rm p.a.} & {\\rm Cl}\\\\ \\hline} \\tablehead{\\multicolumn{12}{l}{\\rm Table~\\ref{tab:the_cat}, continued}\\\\ \\hline \\multicolumn{1}{c}{\\rm Name} & \\multicolumn{3}{c}{\\rm RA} & \\multicolumn{3}{c}{\\rm Dec} & $b_J$ & maj & min &{\\rm p.a.} & {\\rm Cl}\\\\ \\hline} \\tabletail{\\hline} \\begin{tt} \\begin{supertabular}{rrrrrrrrrrrr} \\label{tab:the_cat} 076-086-127 & 22 &48 &53.59 & -67 &41 &10.0 & 15.65 &  41 &  28 & 150 &   3 \\\\ 076-069-060 & 22 &46 &18.04 & -68 &57 &21.3 & 14.67 &  69 &  30 & 156 &   3 \\\\ 076-060-049 & 22 &44 &39.38 & -69 &10 & 1.4 & 14.28 &  64 &  45 & 177 &   3 \\\\ 076-055-049 & 22 &43 &38.60 & -69 &10 &14.4 & 16.14 &  48 &  20 & 115 &   2 \\\\ 076-039-099 & 22 &39 &49.52 & -68 &15 & 8.3 & 15.88 &  35 &  18 &  55 &   3 \\\\ 076-033-110 & 22 &38 &41.63 & -68 & 2 &49.9 & 15.59 &  43 &  29 &  61 &   2 \\\\ 076-015-079 & 22 &35 & 0.24 & -68 &38 & 8.3 & 16.36 &  35 &  24 & 169 &   2 \\\\ 076-009-052 & 22 &34 & 1.52 & -69 & 7 &57.7 & 15.68 &  36 &  24 & 175 &  13 \\\\ 076+022-124 & 22 &27 &41.23 & -67 &47 &50.9 & 16.17 &  31 &  21 &  11 &   3 \\\\ 076+058-111 & 22 &20 &30.09 & -68 & 0 &35.6 & 16.41 &  33 &  21 &   8 &   2 \\\\ 076+065-133 & 22 &19 &12.69 & -67 &35 &25.1 & 16.44 &  24 &  17 &   5 &   3 \\\\ 076+069-114 & 22 &18 &17.69 & -67 &56 &23.4 & 16.10 &  34 &  23 &  10 &   3 \\\\ 076+100-103 & 22 &12 & 6.23 & -68 & 6 &33.0 & 16.18 &  51 &  30 & 124 &   4 \\\\ 076+105-125 & 22 &11 &27.23 & -67 &41 &33.8 & 15.52 &  43 &  31 & 108 &   3 \\\\ 076+106-066 & 22 &10 &13.79 & -68 &46 &43.0 & 15.97 &  51 &  25 &  99 &   3 \\\\ 076+112-059 & 22 & 8 &42.68 & -68 &54 &31.7 & 14.56 &  62 &  43 &  72 &   2 \\\\ 076+117-126 & 22 & 9 & 4.61 & -67 &39 &45.6 & 15.64 &  39 &  33 & 171 &   2 \\\\ 076+121-120 & 22 & 8 &14.48 & -67 &45 &21.6 & 16.26 &  46 &  14 & 156 &   3 \\\\ 076-116+007 & 22 &57 &29.69 & -70 & 7 &17.7 & 16.11 &  31 &  22 & 150 &   8 \\\\ 076-113-015 & 22 &56 &20.68 & -69 &43 & 6.7 & 16.42 &  41 &  17 &  52 &   3 \\\\ 076-108-037 & 22 &54 &51.23 & -69 &19 &14.2 & 14.15 &  91 &  39 &  31 &   3 \\\\ 076-099-007 & 22 &53 &32.00 & -69 &53 &20.9 & 15.89 &  43 &  24 &  80 &   3 \\\\ 076-090+043 & 22 &52 &28.98 & -70 &50 &27.6 & 14.82 &  48 &  46 & 132 &   2 \\\\ 076-087+046 & 22 &51 &56.43 & -70 &53 &53.9 & 15.26 &  75 &  26 & 174 &   3 \\\\ 076-078-031 & 22 &48 &34.08 & -69 &29 & 8.8 & 15.68 &  34 &  27 & 155 &   3 \\\\ 076-060-032 & 22 &44 &43.85 & -69 &28 &37.5 & 15.40 &  59 &  29 & 165 &   3 \\\\ 076-058+021 & 22 &44 &58.11 & -70 &27 &38.5 & 15.57 &  70 &  19 &  11 &   3 \\\\ 076-055-040 & 22 &43 &39.90 & -69 &20 &13.8 & 16.15 &  56 &  14 & 124 &   3 \\\\ 076-044-006 & 22 &41 &31.72 & -69 &58 &52.4 & 16.07 &  32 &  28 &  81 &   3 \\\\ 076-026+041 & 22 &37 &56.04 & -70 &52 &18.7 & 16.11 &  31 &  23 &  18 &   9 \\\\ 076-024-008 & 22 &37 &10.65 & -69 &57 &20.6 & 15.75 &  46 &  21 &  40 &   8 \\\\ 076+001+023 & 22 &31 &43.20 & -70 &32 &25.9 & 14.69 &  60 &  45 &  41 &   3 \\\\ 076+003-009 & 22 &31 &25.89 & -69 &56 & 0.7 & 16.26 &  28 &  28 &  90 &   2 \\\\ 076+006-014 & 22 &30 &42.23 & -69 &50 &58.5 & 16.17 &  41 &  17 & 107 &   8 \\\\ 076+020+043 & 22 &27 &32.88 & -70 &54 &33.2 & 16.22 &  84 &   9 &  31 &   3 \\\\ 076+037+029 & 22 &23 &41.67 & -70 &38 &37.1 & 15.03 &  48 &  39 &  65 &   3 \\\\ 076+038-044 & 22 &24 & 5.04 & -69 &16 &45.3 & 16.42 &  33 &  24 &  45 &   2 \\\\ 076+039+036 & 22 &23 &15.83 & -70 &46 &13.9 & 16.35 &  65 &   9 &  86 &   3 \\\\ 076+044+039 & 22 &22 & 3.71 & -70 &49 &19.9 & 16.39 &  33 &  20 & 161 &   3 \\\\ 076+047-006 & 22 &21 &54.39 & -69 &58 &21.2 & 16.28 &  33 &  27 & 107 &   3 \\\\ 076+056+021 & 22 &19 &29.76 & -70 &27 &48.5 & 16.43 &  29 &  14 &  14 &   3 \\\\ 076+058+029 & 22 &19 & 0.17 & -70 &37 & 7.8 & 15.76 &  36 &  34 & 115 &   2 \\\\ 076+070+034 & 22 &16 &14.59 & -70 &42 & 0.2 & 16.06 &  38 &  24 &  69 &   3 \\\\ 076+097+053 & 22 & 9 &46.39 & -71 & 0 &52.9 & 16.28 &  26 &  25 & 117 &   3 \\\\ 076+099+045 & 22 & 9 &34.21 & -70 &51 & 7.6 & 15.37 &  43 &  26 &  80 &   8 \\\\ 076+099-022 & 22 &10 &48.25 & -69 &36 &55.1 & 15.04 & 104 &  17 & 106 &   3 \\\\ 076+100+022 & 22 & 9 &39.87 & -70 &25 &37.6 & 16.41 &  36 &  19 & 169 &   4 \\\\ 076+111+021 & 22 & 7 &21.16 & -70 &23 &54.8 & 15.58 &  63 &  14 &  21 &   3 \\\\ 076+110+033 & 22 & 7 &16.35 & -70 &37 & 2.2 & 16.21 &  42 &  14 & 102 &   8 \\\\ 076-099+103 & 22 &55 &56.05 & -71 &56 &36.4 & 16.12 &  32 &  17 & 148 &   3 \\\\ 076-058+109 & 22 &46 & 4.85 & -72 & 5 &52.3 & 16.11 &  49 &  15 &  84 &   3 \\\\ 076-053+090 & 22 &44 &36.86 & -71 &45 &56.8 & 16.42 &  37 &  21 &  77 &   2 \\\\ 076-050+077 & 22 &43 &54.56 & -71 &31 &10.3 & 16.12 &  68 &  10 &   9 &   3 \\\\ 076-049+120 & 22 &44 & 0.94 & -72 &18 &59.0 & 16.06 &  33 &  32 &  57 &   3 \\\\ 076-046+073 & 22 &42 &46.51 & -71 &27 & 2.5 & 16.01 &  45 &  32 &  70 &   3 \\\\ 076-044+080 & 22 &42 &29.15 & -71 &35 &10.7 & 15.32 &  45 &  35 &  50 &   2 \\\\ 076-030+087 & 22 &39 &17.03 & -71 &43 &23.9 & 15.86 &  34 &  27 &  33 &   3 \\\\ 076-020+075 & 22 &36 &42.77 & -71 &30 & 8.5 & 15.28 &  50 &  26 &  62 &   3 \\\\ 076-017+057 & 22 &35 &54.92 & -71 &10 &20.8 & 16.29 &  46 &  17 &  86 &   3 \\\\ 076+021+079 & 22 &27 & 4.40 & -71 &34 &35.9 & 16.34 &  54 &  12 & 103 &   3 \\\\ 076+024+086 & 22 &26 &28.53 & -71 &42 & 4.2 & 15.89 &  35 &  32 & 172 &   1 \\\\ 076+027+083 & 22 &25 &38.44 & -71 &38 &35.9 & 16.41 &  30 &  26 &  86 &   2 \\\\ 076+030+066 & 22 &24 &59.38 & -71 &19 &47.6 & 16.27 &  32 &  30 &  88 &   3 \\\\ 076+033+060 & 22 &24 &31.23 & -71 &12 &33.6 & 16.34 &  31 &  19 & 170 &   4 \\\\ 076+038+095 & 22 &22 &54.57 & -71 &52 & 7.0 & 15.85 &  59 &  16 &  91 &   3 \\\\ 076+060+089 & 22 &17 &45.33 & -71 &43 &57.2 & 16.43 &  34 &  20 &  23 &   3 \\\\ 076+060+074 & 22 &17 &57.49 & -71 &27 &30.9 & 14.93 &  76 &  25 & 165 &   8 \\\\ 076+082+120 & 22 &12 & 4.38 & -72 &16 &27.0 & 16.06 &  63 &  11 & 105 &   3 \\\\ 076+083+092 & 22 &12 &13.65 & -71 &45 &15.4 & 14.81 &  54 &  39 & 166 &   3 \\\\ 076+092+096 & 22 &10 & 8.55 & -71 &49 &27.6 & 15.79 &  60 &  16 & 172 &   3 \\\\ 076+095+106 & 22 & 9 & 3.54 & -71 &59 &38.4 & 15.57 &  44 &  35 & 122 &   4 \\\\ 076+100+072 & 22 & 8 &43.43 & -71 &21 & 3.7 & 15.22 &  59 &  21 & 117 &   3 \\\\ 076+106+104 & 22 & 6 &35.49 & -71 &57 & 3.6 & 15.96 &  59 &  15 & 179 &   3 \\\\ 076+108+086 & 22 & 6 &23.34 & -71 &35 &58.7 & 15.64 &  58 &  29 &  55 &   3 \\\\ \\end{supertabular} \\end{tt}  \\begin{table}[htbp] \\begin{center} \\caption{Angular correlation function results for all, early and late type galaxies in the APM-BGC.} \\vspace{0.5cm} \\label{tab:w} \\begin{math} \\begin{array}{lccccc} \\hline \\hline {\\rm Type} & \\gamma & A & \\Delta w & B & r_0\\\\ \\hline All & 1.85 \\pm 0.14 & 0.18 \\pm 0.03 & 2.9\\ten{-3} & 16.4 \\pm 1.5 & 4.5 \\pm 0.6\\\\ Early & 1.93 \\pm 0.16 & 0.35 \\pm 0.05 & 4.6\\ten{-3} & 40.6 \\pm 3.5 & 6.8 \\pm 0.8\\\\ Late & 1.79 \\pm 0.16 & 0.15 \\pm 0.03 & 1.6\\ten{-3} & 11.3 \\pm 1.0 & 3.9 \\pm 0.5\\\\ \\hline \\hline \\end{array} \\end{math} \\vspace{0.2cm}  Note.---Power-law fits ($ w = A \\theta^{1-\\gamma}$) were made over the range 0.1--$5\\dg$. The integral constraint $\\Delta w $ is estimated from the observed $w(\\theta)$.  The amplitude, $B$, and corresponding scale length, $r_0$ are for the spatial correlation function inferred from inverting Limber's equation. \\end{center} \\end{table}  \\clearpage"
    },
    "9603/astro-ph9603127_arXiv.txt": {
        "abstract": "We selected samples of late-type dwarf galaxies in the Virgo cluster with HI information. The galaxies were observed at the Wise-Observatory using several broad-band and H$\\alpha$ bandpasses. UV measurements were carried out with the IUE Observatory from VILSPA, and with the FAUST shuttle-borne UV telescope.  We describe the observations in detail, paying particular attention to the determination of measurement errors, and present the observational results together with published data and far-infrared information from IRAS. The sample will be analyzed in subsequent papers, in order to study star formation mechanisms in galaxies. ",
        "introduction": "\\label{sec_int}   The process of star formation is one of the most important process in galactic evolution. Generally, star formation is initiated by the gravitational collapse of gas clouds, followed by fragmentation into future individual stars. When stars enter the main sequence, radiation pressure and stellar winds push on the ambient gas and prevent it from collapsing further.   Despite extensive progress in understanding the star formation process, there are several issues still not fully understood, due to the complexity of this process. One can summarize the currently open issues, concerning the star formation processes in galaxies, in two major questions: (1) What are the mechanisms that govern the star formation process, and how do they depend on the galactic type and environment? (2) How do the star formation rate (SFR) and the initial mass function (IMF) depend on various galactic properties, such as interstellar gas density, morphology of the interstellar gas, metallicity, and the amount of dust in the interstellar medium?    In order to find the SFR of a sample of galaxies one needs to know the IMF characterizing the star formation in these galaxies. The IMF can be derived by fitting a number of observed properties, such as broad-band colors, to a set of models of different stellar populations, with different IMFs (population synthesis). By comparing the various color indices to the synthetic colors calculated from models one can, in principle, determine the IMF of the galaxies in the sample. This requires a database spanning as large a range in wavelengths as possible, to eliminate the degeneracy of results in several conditions, being able to distinguish between different conditions that yield the same value for one or more colors. It is important that one of the measured bands is in the UV regime, in order to trace the massive stars which contribute most of the energy emitted in this part of the spectrum.   Using these data together with a direct massive-star tracer enables one to draw a consistent picture of the star formation history of the sample galaxies. One of the most common tracers of young massive stars, adopted also here, is based on the H$\\alpha$ line intensity. The H$\\alpha$ line has been used by many (Kennicutt 1983, Kennicutt \\& Kent 1983, Gallagher, Hunter \\& Tutukov 1984, Pogge \\& Eskridge 1987, Kennicutt {\\it et al.} 1994), mainly due to its high intensity. The main deficiency of this method is the dust extinction. However, the extinction is moderate in late-type irregular galaxies like the ones used in this study, usually some 0.2--0.4$\\;mag$ (van der Hulst {\\it et al.} 1988). ",
        "conclusions": "\\label{sec_conc}  We presented observations related to star formation properties of a sample of late-type dwarf galaxies. The intention in focusing on dwarf galaxies was to exclude some of the star formation inducing mechanisms, assumed to account for star formation in large galaxies. In addition, we concentrate on Virgo cluster members, in order to test the effects of the cluster environment on the star formation properties of the galaxies.  A data base consisting of a number of broad-band colors and H$\\alpha$ line observations is important for determining the ongoing star formation process, as well as the star formation history of the sample. In a subsequent paper we will show that some of the galaxies show signs of a strong burst of star formation, while others completely lack signs of recent star formation activity.  The observational data are affected primarily by internal dust extinction in the galaxies, which complicates the interpretation of the data. This effect will be discussed in detail in the next paper.  Another interesting finding, concerning the dwarf galaxies in Virgo, is their velocity field. Our results suggest an infall of the galaxies towards the cluster core, which may explain partly the skewed recession velocity distribution of the dwarf galaxies, in terms of a Malmquist bias."
    },
    "9603/astro-ph9603061_arXiv.txt": {
        "abstract": "Lada et al. (1994) have described a method for studying the distribution of dust in dark clouds using infrared imaging surveys. In particular they show that the method provides some information about the structure of the gas (dust) on scales smaller than their resolution.  In the present work we clarify the nature of the information provided by their method.  We show that:  \\begin{itemize}  \\item the 3D density field of the gas is well described by a Log-Normal distribution down to very small scales;  \\item the power spectrum and the standard deviation of the 3D density field can be constrained;  \\item the origin of such a structure of the density field is likely to be the supersonic turbulence in the gas.  \\end{itemize}  In fact we find a qualitative and quantitative agreement between the predictions based on recent numerical simulations of supersonic turbulence (Nordlund and Padoan 1996; Padoan, Nordlund and Jones 1996) and the constraints given by the infrared dust extinction measurements. ",
        "introduction": "In a recent paper Lada et al. (1994) have illustrated the method of mapping the distribution of dust, and therefore gas, in dark clouds by using stellar extinction measurements in the near-infrared. The method is based on the use of multi-channel array cameras that allow the simultaneous determinations of the colors of hundreds to thousands of stars through a molecular cloud. The infrared color excess is proportional to the dust column density, and the dust-to-gas ratio is known to be nearly constant in interstellar clouds; therefore cloud maps can be obtained. The maps obtained measuring the infrared excess are considerably more accurate than the maps based purely on stellar counts.  The gas column density is given by  \\begin{equation} N(H+H_{2})=2\\times 10^{21}A_{V} cm^{-2} \\label{1} \\end{equation} where the visual extinction in magnitudes is \\begin{equation} A_{V}=15.9E(H-K) \\label{2} \\end{equation} and the color excess is \\begin{equation} E(H-K)=(H-K)_{observed}-(H-K)_{intrinsic} \\label{3} \\end{equation} and finally \\begin{equation} <(H-K)_{intrinsic}>=13 \\pm 0.01 mag. \\label{4} \\end{equation} (see equations 1-4 in Lada et al. 1994).  The extinction data are used in two complementary ways, one exploiting an ordered sampling (information on large scale structure), the other a random sampling (information on scales below the resolution of the map).  In the first method of analysis the data are spatially binned like in stellar count and millimeter wave observations. At any position a few stars are found so that an average extinction $A_{V}$ can be measured. The result is an extinction map that compares well with the stellar count map  and the CS map.  The second method is that of plotting the mean extinction, $A_{V}$, and its standard deviation, $\\sigma$, measured at any position. Lada et al. (1994) found that the dispersion grows with the average extinction, and realized that this behavior contains information about the structure of the extinction (therefore of the gas mass distribution) in the cloud, on scales smaller than the resolution of the extinction map. They give examples of mass distributions that would generate or not generate such a plot, but their interpretation of the plot does not go very far.  In this work we focus on the second method of using the extinction data, that is on the meaning of the $\\sigma-A_{V}$ plot as a tracer of structure on scales below the resolution of the map.  Fig.1, which is the equivalent of fig.7 in Lada et al. (1994), shows the $\\sigma-A_{V}$ plot obtained from the original data, kindly provided to us by the authors. The measurements are taken for a dark cloud complex near the young cluster IC 5146 in Cygnus. The cloud has been mapped in $^{12}$CO and $^{13}$CO by Dobashi et al. (1992), who named it `Cloud C'.  In sections 3 and 4 we show that the $\\sigma-A_{V}$  plot is due to the `intermittent' distribution of the dust (that is of the gas density field in the cloud), and we show how to constrain such distribution using randomly generated fields with given statistics and power spectra. Before giving such details, though, we present  in the next section the results of recent numerical simulations, concerning the density field in isothermal  random supersonic flows. It will be clear, in section 4 and in the following discussion, that random supersonic flows are in fact excellent candidates to interpret the extinction data and to explain the origin of the distribution of dust in dark clouds. ",
        "conclusions": "In the present work we have re-interpreted the observational results obtained by Lada et al. (1994), i.e. the fact that the mean stellar extinction at any given position in space increases together with the dispersion of the extinction, where the averages are taken among the stars found at that position in space.  The authors were able to conclude that:  \\begin{itemize} \\item structure must be present down to scales smaller that the extinction map resolution; \\item generic models for the cloud structure (eg uniform or in clumps) do not easily reproduce the $\\sigma-A_{V}$ plot; \\item the $\\sigma-A_{V}$ plot is a basic test for any model for the dynamics and structure of the cold interstellar medium. \\end{itemize}  We have simulated the observations by generating random density distributions. In this way we have been able to better define the information contained in the $\\sigma-A_{V}$ plot. We can therefore add that:  \\begin{itemize} \\item the statistics of the 3-D density field in the dark cloud is certainly very intermittent; in particular it is consistent with a Log-Normal distribution; \\item the standard deviation of the statistics is $\\sigma_{x,3D}=5.0\\pm0.5$; \\item the index of the power spectrum (assumed to be a power law), of the 3-D density field, is $\\alpha=2.6\\pm0.5$; \\item the relation between the rms Mach number of the flow and the standard deviation of the 3-D density field is about $\\sigma_{x,3D}\\approx0.5{\\cal{M}}$; \\end{itemize}  We therefore conclude that the scenario for star formation and for the dynamics of dark clouds proposed by Padoan (1995), Nordlund  \\& Padoan 1996), and Padoan, Nordlund, \\& Jones (1996) is fully consistent with the dust extinction measurements in `Cloud C' by Lada et al. (1994).  In that scenario the dynamics of dark clouds is characterized by supersonic random motions, which are responsible for fragmenting the mass distribution. In the cited works we showed that a MIller-Scalo stellar mass function is a natural consequence of that scenario. Here we have shown that the shape, the standard deviation, and the spectral index of the density distribution, predicted with numerical simulations of supersonic turbulence and used in that scenario, are consistent with the observations of stellar extinction."
    },
    "9603/astro-ph9603149_arXiv.txt": {
        "abstract": "Photometric redshifts have been determined for the galaxies in the Hubble Deep Field. The resulting redshift distribution shows two peaks: one at $z\\sim0.6$ and one at $z\\sim2.2$. Luminosity functions derived from the redshifts show strong luminosity evolution as a function of redshift. This evolution is consistent with the Babul \\& Rees (1992) \\nocite{br92} scenario wherein massive galaxies form stars at high redshift while star formation in dwarf galaxies is delayed until after $z=1$. ",
        "introduction": "\\label{sec:intro}  The Hubble Deep Field (HDF) optical images are the deepest yet obtained. Objects as faint as $I_{ST}$=28.5 \\footnote{for simplicity, $U_{ST}$,$B_{ST}$,$R_{ST}$ and $I_{ST}$ will be used to denote magnitudes in the F300W, F450W, F606W and F814W bands respectively. The ST zero-point system is used unless otherwise specified.} can be detected at the $10\\sigma$ level. At this point in time, only a few spectroscopic redshifts have been measured for the brighter galaxies and none for the faintest galaxies in these images.  Photometric redshifts (see, for example, Gwyn 1995; Connolly et al. 1995; Koo 1985) \\nocite{gwyn95,evil,pmzm} can be measured much faster and to much fainter magnitudes than their spectroscopic counterparts. Because the wavelength bin-size in photometry is generally so much larger than in conventional spectroscopy ($\\sim$1000\\AA~ {\\em vs.} 1--2\\AA), far shorter exposure times are required to measure redshifts (but with a sacrifice in accuracy).  Photometric redshifts have been calculated for the galaxies brighter than $I_{ST}=28$ in the Hubble Deep Field. This paper presents the redshift distribution for this sample. Also presented are luminosity functions as a function of redshift out to $z=5$. ",
        "conclusions": "\\label {sec:conc} Using photometric redshifts, a redshift distribution has been determined for the Hubble Deep Field. It shows two peaks: one at $z\\simeq0.6$ and another at $z\\simeq2.2$. Luminosity functions have been calculated using these redshifts. The LF's show strong evolution: the brightest galaxies are 4 magnitudes brighter than their present day counterparts and the faint galaxies are fewer in number. The double-peaked redshift distribution and the evolution of the luminosity function can be understood if larger galaxies form stars early at $z\\sim3$ and if star formation is delayed in the dwarf galaxies until after $z\\simeq1$."
    },
    "9603/astro-ph9603133_arXiv.txt": {
        "abstract": "Transport properties of degenerate relativistic electrons along quantizing magnetic fields in neutron star crusts are considered. A kinetic equation is derived for the spin polarization density matrix of electrons. Its solution does not depend on the choice of basic electron wave functions unlike previous solutions of the traditional kinetic equation for the distribution function. The density matrix formalism shows that one can always reach high accuracy with the traditional method by a proper choice of the basic functions. Electron Coulomb scattering on ions is considered in liquid matter, and on high-temperature phonons or on charged impurities in solid matter. In the solid regime, the Debye -- Waller reduction of phonon scattering can strongly enhance the longitudinal thermal or electric conductivity. An efficient numerical method is proposed for calculating the transport properties of electron gas at any magnetic field of practical interest. ",
        "introduction": "% \\label{sect1} Accurate transport coefficients in neutron star crusts are important for analysing the thermal evolution of neutron stars and evolution of their magnetic fields. In outer crusts of cooling magnetized neutron stars, the heat is mainly transported along the magnetic fields. There exist several competing heat transport mechanisms across the field, but the longitudinal currents are carried mostly by electrons through their scattering on phonons or charged impurities in the solid phase and on ions in the liquid phase. As a rule, the electrons in the crust are strongly degenerate and may be relativistic; the magnetic field can be easily quantizing. Transport properties of the crusts have been studied in a number of papers (e.g., Yakovlev \\& Kaminker 1994, and references therein).  In the present work, the most important problem of longitudinal electron transport in quantizing magnetic fields is studied with the use of the quantum density matrix formalism, instead of the traditional kinetic equation for the electron distribution function employed in previous studies. The main advantage of the present approach is that it is independent of the choice of the basis of electron states (basis states are not unique due to the electron spin degeneracy).  We consider three main electron scattering mechanisms. The first one is the Coulomb scattering on ions in the liquid or gaseous phase. The second one is the scattering on high-temperature phonons in the solid phase. In the latter case, we take into account the Debye -- Waller factor whose importance has been emphasized and proved by Itoh et al. (1984b, 1993) for the non-magnetic case. We show that the effect of this factor is much stronger in quantizing magnetic fields. The third mechanism is the Coulomb scattering on charged impurities in the solid phase, important much below the melting temperature.  The paper is composed as follows. In Sect.~\\ref{sect2} we describe the physical conditions of interest, electron scattering potentials, transport coefficients and their expressions in the non-magnetic case. In Sect.~\\ref{sect3} we derive a linearized kinetic equation for the density matrix and compare its numerical solutions with the traditional solutions employed in all previous works. Mathematical properties of the new equation are discussed in Appendix~A. The effect of the Debye -- Waller factor in quantizing magnetic fields is studied in Sect.~\\ref{sect4}. The results are summarized in Sect.~\\ref{sect5}. In Appendix~B we present new expressions for some intermediate integrals. These expressions ensure efficient computation of the transport properties for the case when many Landau orbitals are occupied. Previous results (Yakovlev 1984, Hernquist 1984, Schaaf 1988, Van Riper 1988) were restricted to 30 Landau orbitals at most. ",
        "conclusions": "% \\label{sect5} We have presented the theory of transport properties of degenerate electrons along quantizing magnetic fields in neutron star crusts. Our results are advanced, compared to the previously known ones, in three respects.  First, a kinetic equation for the spin polarization density matrix of electrons is derived. The solution of this equation provides a justification of the standard approach based on the kinetic equation for the electron distribution function. The present results are compared with two versions of the standard approach used previously by different authors.  For non-relativistic magnetic fields, $B\\la 10^{13}$~G, our results confirm the arguments of Yakovlev (1984) that the standard approach which employs basic functions with fixed spin $z$-projection is the most appropriate in the non-relativistic limit. The fixed-helicity basic functions used by other authors lead to small inaccuracies which however seem to be insignificant in astrophysical implications.  For stronger fields, $B\\ga 10^{13}$~G, the inaccuracies of the traditional approach increase up to 20\\% when density is rather low and the electrons occupy low-lying Landau levels. The density-matrix results allow us to choose the most appropriate version of the standard approach. If density is higher and the electrons populate many Landau levels, the difference between various approaches becomes negligible.  Secondly, we have taken into account the cumulative effect of the Debye -- Waller factor with the magnetic quantization. In the non-magnetic case, this factor can increase the thermal and electric conductivities by a factor of 3 just below the melting temperature. We show that the magnetic quantization can enhance the effect by an order of magnitude.  Thirdly, we have derived semiclassical expressions for some intermediate integrals which enter the system of equations either for the density matrix or for a distribution function. These expressions provide fast and accurate calculation of the relaxation time for large number of occupied Landau levels.  In this paper we have calculated the kernel function $\\Phi$ which should undergo further thermal averaging, Eq.\\,(\\ref{3.15}), to determine the longitudinal electric and thermal conductivities and thermopower. We shall consider this averaging and astrophysical implications of the developed theory in the subsequent paper."
    },
    "9603/astro-ph9603023_arXiv.txt": {
        "abstract": "We investigate the evolution of metal deficient stellar structures, presenting H-burning isochrones covering cluster ages from 800 Myr to 7 Gyr. Evolutionary evidences for selection effects in the metallicity distribution of very metal poor H-burning red giants are reported. The evolution of stars during central and shell He burning is further investigated, discussing the occurrence of He burning pulsators as a function of cluster age. ",
        "introduction": "\\par In order to extend the theoretical scenario concerning stars with $Z=10^{-5}$ down to cluster ages of about 1 Gyr, evolutionary tracks already presented in Paper I for the quoted metallicity have been implemented with new tracks for suitable choices of the stellar masses. All the computations have been performed adopting a cosmological abundance of He as given by $Y=0.23$. \\par As well known, the modality of He ignition significantly depends on the total mass of the evolving giant. For each given stellar population (i.e., for each assumed value of $Y$ and $Z$), one may define a critical mass $M_{HeF}$ as the upper mass limits for stars experiencing strong electron degeneracy of the He core during the H shell burning phase and - thus - igniting He through one or more violent He flashes. One finds that evolutionary features of red giant (RG) stars with masses around $M_{HeF}$ change in a remarkable way in a range of only few tenths of solar mass, an occurrence already known as \\lq{\\sl Red Giant Branch Transition}\\rq\\ (RGT) (see Sweigart, Greggio \\& Renzini 1989, hereafter SGR, Sweigart, Greggio \\& Renzini 1990, Castellani et al. 1992). \\par \\begin{figure} \\epsscale{.60} \\plotone{fig1.eps} \\caption{The luminosity of the RGB tip and the mass of the helium core at the helium ignition versus the total star mass when Z=$10^{-5}$.} \\end{figure}   The behavior of our Z=$10^{-5}$ models through the transition is shown in figure 1, which shows the dependence of $M^{tip}_c$ (the mass of the He core at the He ignition) and $L_{tip}$ (the star luminosity at the tip of RGB) on the stellar mass. According to Sweigart \\& Gross (1978) and SGR, the onset of the helium flash has been taken at the model where the contribution of $3\\alpha$ reactions to the energetic reaches $100L_\\odot$; for structures which quietly ignite helium, the He ignition has been alternatively fixed at the first appearance of a convective core. The sudden variation of $L_{tip}$ around M=$1.5M_{\\odot}$ indicates that this stellar mass is near the transition between low mass stars developing full degenerate Helium cores and more massive structures where electron degeneracy is progressively removed. As in previous investigations, if one defines the critical mass $M_{HeF}$ as the mass of the star having at the He ignition a He core mass equal to the average value between the He core of fully degenerated structures and the absolute minimum in $M^{tip}_c$, when Z=$10^{-5}$ one finds $M_{HeF}$ of the order of $1.45M_\\odot$. Table 1 reports selected evolutionary parameters for all the computed models, allowing a quantitative inspection of the RGB transition. \\par  \\begin{figure} \\epsscale{.60} \\plotone{fig2.eps} \\caption{The amount of extra-helium brought to the surface by the first dredge up for all the computed models. For the sake of comparison, the results concerning two different assumptions  about the stellar metallicity are reported.} \\end{figure} Figure 2 compares the amount of extrahelium brought to the surface by the first dredge up with similar data but for the larger metallicities investigated in Paper I. As already discussed in Castellani \\& Degl'Innocenti (1995), for each given metallicity one finds a stellar mass separating the regime of low mass stars where $\\Delta{Y}$ increases when the stellar mass is increased from more massive stars with opposite behavior. Such an occurrence as well as the dependence of $\\Delta{Y}$ on the star metallicity can be easily understood in terms of the discussion given by Castellani \\& Degl'Innocenti (1995). \\par \\begin{figure} \\epsscale{.60} \\plotone{fig3.eps} \\caption{The critical mass $M_{HeF}$ (in solar mass) (a) and the age (in Gyrs) of a cluster with $M_{HeF}$ at the He ignition (b) versus the global amount of heavy elements.} \\end{figure}  \\begin{deluxetable}{ccccc} \\scriptsize \\tablecaption{Selected evolutionary parameters at the He ignition: 1) the star mass, 2) the stellar luminosity at the tip of the RGB, 3) the mass (in solar unit) of the He core, 4) the amount of extra-helium brought at the surface by the first dredge up and 5) the age (in Gyrs) of the star.} \\tablehead{ \\colhead{$M/M_{\\odot}$} & \\colhead{$log(L/L_{\\odot})_{tip}$} & \\colhead{$M_c^{tip}$} & \\colhead{$\\Delta Y$} & \\colhead{$t_{HeF}$}} \\startdata 1.1 &3.083 &0.501 &0.015 &5.04 \\nl 1.3 &2.835 &0.460 &0.021 &2.89 \\nl 1.4 &2.743 &0.436 &0.024 &2.27 \\nl 1.6 &2.603 &0.397 &0.027 &1.47 \\nl 1.8 &2.338 &0.358 &0.023 &1.01 \\nl 2.0 &2.166 &0.339 &0.004 &0.72 \\nl 2.2 &2.158 &0.344 &0.000 &0.53 \\nl \\enddata \\end{deluxetable} Figure 3a shows the dependence of the critical mass $M_{HeF}$ on star metallicity. In this figure, present results have been implemented with similar data given by Cassisi \\& Castellani (1993) or by SGR for lower or larger metallicities, respectively. The dependence of $M_{HeF}$ on Z has been already discussed (see Cassisi and Castellani 1993) and this discussion will not be repeated here. As a relevant point figure 3b discloses the dependence on the metallicity of the cluster age at the Helium ignition in stars with mass M=$M_{HeF}$. It appears that when Z=$10^{-5}$ the transition requires ages of the order of about 2.2 Gyr, i.e., a much larger age than for the Z=$10^{-4}$ case. As a consequence, in a dwarf galaxy with star metallicities ranging from Z=$10^{-5}$ to Z=$10^{-4}$ and ages around 1 Gyr, the red giant branch is expected to be populated by the more metal rich stars only. As a result, one finds that the distribution of metallicity of RG cannot taken in all cases as a bona fide indicator of the distribution of star metallicity in \\lq{not-too-old}\\rq\\ metal poor systems. \\par Evolutionary models, as computed for the case Z=$10^{-5}$ allow us to extend toward lower ages the set of isochrones presented in Paper I. This is shown in figure 4 where we report selected isochrones for H burning stellar structures covering cluster ages from 0.8 to 7Gyrs. \\begin{figure} \\epsscale{.60} \\plotone{fig4.eps} \\caption{Cluster isochrones for H burning phases and for the labeled range of ages. The interval is 100 Myr for ages lower than 1 Gyr and 1 Gyr for larger ages.} \\end{figure} ",
        "conclusions": "\\par This paper investigates the evolutionary properties of relatively massive, metal deficient stellar structures. The first part of the investigation has been devoted to H burning stars, presenting selected isochrones for ages in the range 800 Myr - 7 Gyr, increasing the range of ages covered by previous investigations.  Selected sets of HB models have been computed under different assumptions about the ages of the stellar system. We confirm the results already given for larger stellar metallicities about the possible, and sometime probable, occurrence of anomalous, overluminous variable stars, in relatively young, metal deficient system. \\par Both evolutionary tracks and isochrones are available by electronic mail upon request to cassisi@astrte.te.astro.it."
    },
    "9603/astro-ph9603059_arXiv.txt": {
        "abstract": " ",
        "introduction": "The $2.2 \\mu $m luminosity functions (LFs) of Baade's window (BW) and the inner $2'$ of the Galactic center (GC) are markedly different at the bright end. The GC LF has a substantially higher fraction of stars brighter than $K_{0} = 5.5$, but it is nearly identical to the BW LF at the faint end for  $K_{0} >7.0$. This suggests that the GC contains a population of younger stars in addition to an older population of red giants similar to the one in BW (Haller 1992; Blum et al. 1996). Here, we present $H$ and $K$ photometry of  $\\sim 42,000$ stars in an area $\\sim 16' \\times 16'$ centered on the GC. In our analysis, we exclude the inner $2'$ and refer to the remaining area as the Perigalactocentric region (PGC). We construct the $K$ band LF of the PGC by individually dereddening each star using its observed $(H-K)$ colors. We find that the PGC does not contain any super bright stars of the type found in the GC. For $K_{0} < 6.0$, the GC LF has a significant excess over the PGC LF. In comparing the PGC with BW, we find that for  BW the contamination by the foreground distribution of $M$ giants is significant for $K_{0} < 7.0$, while for the PGC it is negligible. Even if the contamination in BW is neglected, there is a still a clear excess of stars in the PGC over BW in the range $4.0 < K_{0} < 7.0$. We conclude that the luminosity function  of the PGC is intermediate between that of the GC and Baade's window. ",
        "conclusions": "\\par  The stellar population of the GC has an excess of bright stars compared to the  the older population of BW (Haller 1992). The central pc of the Galaxy contains helium rich, luminous, blue, emission-line stars and Wolf-Rayet stars with estimated zero-age main sequence masses of up to $\\sim 100M_{\\odot}$ (Allen, Hyland \\& Hillier 1990; Krabbe et al. 1991; Blum, Sellgren, \\& DePoy 1995a; Blum, DePoy, \\& Sellgren 1995b; Libonate et al. 1995; Krabbe et al. 1995). A plausible scenario is a burst of star formation in the GC $\\sim 10$Myr ago (Krabbe et al. 1995). Krabbe et al. (1995) also conclude that the intermediate mass asymptotic giant branch stars were formed in another burst of star formation $\\sim 100$Myr ago.  \\par In the PGC, we do not find stars that are as luminous as the brightest stars in the GC. Nevertheless, there is a significant excess of stars with $K_{0} \\leq 7$ over the older population of the BW. This could imply the existence of a population of stars that is significantly younger than the old bulge population. The best way to confirm this hypothesis would be to take spectra of the bright stars in the PGC. To this end, we present in Table 5 a list of all stars with $K_{0} < 5$ in the PGC region."
    },
    "9603/hep-ph9603336_arXiv.txt": {
        "abstract": "Motivated by the supersymmetric interpretation of the CDF $ee\\gamma\\gamma + \\slashchar{E}_T$ event and the reported $Z\\to b\\bar b$ excess at LEP, we analyze the Higgsino as a cold dark matter candidate.  We examine the constraints as implied by the collider experiments, and then calculate its relic density.  We find that this Higgsino-like lightest supersymmetric particle is a viable cold dark matter candidate ($0.05< \\Omega h^2 < 1$), and we discuss its favorable prospects for laboratory detection. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603150_arXiv.txt": {
        "abstract": " ",
        "introduction": "The outstanding problem of modern cosmology is to understand the formation of galaxies, clusters of galaxies and large-scale structure in the expanding universe. A scenario in which these structures form by the growth by gravitational instability of small--amplitude primordial density fluctuations in a universe dominated by weakly interacting non--baryonic dark matter has been the standard framework within this problem has been discussed for many years. For most purposes the crucial theoretical quantity in the  gravitational instability picture is the primordial fluctuation spectrum, $P(k)$. In most theories, density perturbations originate as a quantum phenomenon in the very early universe with a power--law spectrum of the form $P(k)\\propto k^{n}$; $n$ very close to unity  is favoured by many versions of the inflationary universe picture, as well as by more general considerations. With the discovery of temperature anisotropies in the Cosmic Microwave Background Radiation (CMBR)  by the COBE satellite (Smoot et al. 1992), it is possible to fix the  amplitude of this spectrum in a relatively unambiguous way on very large  scales, of order 1000 Mpc. On smaller scales, the shape of the primordial spectrum is expected to evolve from its initial power--law form because of the action of various damping and dissipative processes. During several intermediate stages different components can have different spectra and, in some models, residual differences can still be present at the onset of non--linear stages.  In many respects the problem of explaining structure formation in the gravitational instability picture can  be reduced to that of finding a power spectrum  whose primordial form  matches the COBE--inferred amplitude on large scales, and whose evolved form simultaneously matches the  statistical properties of galaxies and clusters on smaller scales. This has proved to be a non--trivial task for the Cold Dark Matter (CDM) model, which, at least in its standard formulation ($\\Omega_0=1$, $n=1$, $h=0.5$ and Gaussian adiabatic fluctuations), is now generally accepted to be ruled out by the data (Wright et al. 1992; Taylor  \\& Rowan--Robinson 1992; Liddle \\& Lyth 1993; White, Efstathiou \\& Frenk 1993). The essential problem of this model is that, once it is normalized to match the measure CMB  temperature anisotropy on large scales, it has too large a fluctuation amplitude on small scales $\\mincir 10\\hm$. Despite its failure, CDM is nevertheless considered as a reference model, several modifications of it having been suggested in order to remedy its shortcomings. Among the ``second generation'' of CDM--based models is the Cold+Hot Dark Matter (CHDM) model, which assumes that part of the dark matter content is in the form of massive neutrinos of mass $m_\\nu \\sim 10\\,$eV (Valdarnini \\& Bonometto 1985; Bonometto \\& Valdarnini 1985; Achilli,  Occhionero \\& Scaramella 1985; Holtzman 1989; Schaefer, Shafi \\& Stecker 1992; Schaefer \\& Shafi 1992; Davis, Summers \\& Schlegel 1992; Holtzman \\& Primack 1993;  Liddle \\& Lyth 1993; Klypin et al. 1993). In this scenario, the small--scale power is suppressed by neutrino free--streaming by an amount which depends on the hot percentage; the number of neutrino species participating in the hot component also plays a role (Primack et al. 1995; Babu, Shafer \\& Shafi 1995). This model, with $\\sim 20$--30\\% to the density from the hot component appears to provide  a good description of structure  on a rather large range of scales, although stringent constraints on the exact amount of the hot component are provided by the abundance of high--redshift objects (e.g. Ma \\& Bertschinger 1994; Klypin et al. 1995).  In this paper we discuss a scenario  similar to the CHDM picture,  where dark matter has a cold component (CDM) and a further  volatile component (VDM), which has a phase space distribution resulting from the decay of a heavier particle species. A previous paper (Pierpaoli \\& Bonometto 1995, hereafter PB95) calculated the effects of such a particle species upon the evolution of fluctuations through the period of recombination and up to the present epoch. Results were presented in PB95 in the form of transfer functions for several examples. In the following analysis we consider different models with  respect to those discussed in PB95, so as to provide a wider sampling of the cold+volatile dark matter (CVDM) model parameter space. Furthermore, we go beyond the calculation of the transfer functions and submit the models to a number of explicit tests, by comparing them with observational data. The tests performed take into account: (i) the large--scale galaxy clustering as deduced from the analysis of volume--limited galaxy samples obtained from redshift catalogues; (ii) the behaviour of bulk velocities calculated using POTENT; (iii) the observed abundance of galaxy clusters; (iv) the observed abundance of high--redshift structures traced by damped Ly--$\\alpha$ systems. Such tests therefore refer to scales ranging from a hundred of Mpc's down to a fraction of Mpc.  The plan of the paper is as follows. In Section 2 we introduce the CVDM models and the corresponding linear power spectra. By only resorting to linear--theory approaches, we compare in Section 3 such models to observational data. In Section 4 we discuss the main results and draw  general conclusions from our analysis. ",
        "conclusions": "The results we have presented demonstrate that the CVDM hypothesis yields potentially interesting models of structure formation. The aim of this paper is to show that rather slight changes in the parameters of volatile dark matter can make a significant difference to the transferred power spectrum. This contrasts with the case of a cold component, where the physical properties of the candidate particle do not really matter at all, in that the physical origin of the hot particles may leave a detectable imprint in the clustering pattern. In this context it is important to verify up to which point the shape of the distribution function  causes differences compared to the standard scenario based on relic thermal neutrinos.  The ability to change $z_{nr}$ almost independently of $\\Omega_X$ is especially significant in this respect: the power spectra we have obtained display considerable variations at a fixed value of $\\Omega_X$. Indeed, although the CVDM class of models involves one more parameter than is the case for CHDM, we have shown that observational data nevertheless allow us to put rather stringent constraints on the permitted values of $z_{nr}$ and $\\Omega_X$, even at the level of linear--theory. The most stringent of these constraints comes from the simultaneous requirement for a model to satisfy the observed abundance of high--redshift DLAS and of galaxy clusters. As for DLAS, the rather large value of the HI gas fraction involved in the absorbing systems, $\\Omega_g$, implies a substantial amount of power on galaxy scales, so as to favour models with $\\Omega_X\\mincir 0.2$. A larger volatile component would be allowed only resorting to a high value of $z_{nr}\\simeq 2\\times 10^5 \\Omega_X$ (cf. Figure 4). On the other hand, models with small $\\Omega_X$ and/or large $z_{nr}$ behave too much like the standard CDM model, drastically overproducing clusters (cf. Figure 3).  Therefore, the overall result would be that models with $z_{nr} \\magcir 5\\times 10^4 \\Omega_X$ have a hard time, quite independently of $\\Omega_X$. Among the models inspected, the only model which passes all the tests, or at least which can not be confidently ruled out, is the one with $\\Omega_X=0.2$ and $z_{nr}= 4\\times 10^3$. It is worth recalling, however, that such a model with low $z_{nr}$ requires that volatile particles occupy at least 5 helicity states [cf. eq.(1)]. We recall that this can accommodated only if {\\bf (a)} $g^*=7$ is allowed by standard nucleosynthesis and {\\bf (b)} two neutrino species are sufficiently massive that they have already decayed at the nucleosynthesis epoch.   An alternative possibility, holding if the physics of the decay is quite different from the axino model suggested in PB95, is that the decay itself takes place after the nucleosynthesis epoch. This would make low $z_{nr}$ models compatible with all $N_\\nu$. It must, however, be remembered that changing $N_\\nu$ itself causes an alteration of the transfer function, and a straightforward extrapolation of the above results to greater $N_\\nu$ values is not allowed. Changing the relation between $N_\\nu$ and $z_{nr}$ opens the way to inspecting different models and, in this context, we should also bear in mind that our analysis has been based on assuming a scale--free primordial spectrum, while  variations around this model are allowed by some classes of inflationary schemes. For instance, taking $P_i(k)\\propto k^n$ with $n<1$ (Adams et al. 1993; Liddle \\& Lyth 1993 and references therein) decreases the amount of power on the cluster mass scale, so as to alleviate the problem of cluster overproduction displayed by ``colder\" models. However, the amount of this tilt can not be too large, in order not to conflict with CMB (Bennet et al. 1994) and large--scale peculiar motions constraints (Tormen et al. 1993). On the other hand, the case of ``antitilting'', with $n\\simeq 1.2$, has been recently advocated to alleviate some of the problems of the CHDM scenario (Dvali, Shafi \\& Schaefer 1994; Lucchin et al. 1995). However, the subsequent increase of power on small scales goes in the undesired direction as far as the cluster abundance is concerned (Pogosyan \\& Starobinsky 1995; Borgani et al. 1995).  As a final remark, we should stress that the analysis presented in this paper is only preliminary and is entirely based on linear calculations. In order to be more definitive we would like to extend it  in two main directions. Firstly to calculate more detailed properties of the CMBR fluctuations they produce: we anticipate a rather different signature on angular scales around a degree than in the standard CHDM models. Furthermore, we would also like to study the non--linear evolution of some of these model by performing numerical calculations using N--body and other procedures."
    },
    "9603/astro-ph9603144_arXiv.txt": {
        "abstract": "A scenario in which cosmic rays (CRs) above $10^{20}{\\rm eV}$ are produced by cosmological gamma-ray bursts (GRBs) is consistent with observations provided that deflections by the inter-galactic magnetic field (IGMF) delay and spread the arrival time of the CRs over $\\geq50{\\rm yr}$. The energy lost by the CRs as they propagate and interact with the microwave background is transformed by cascading into secondary GeV-TeV photons. We show that a significant fraction of these photons can arrive with delays much smaller than the CR delay if much of inter-galactic space is occupied by large-scale magnetic ``voids'', regions of size $\\gtrsim5{\\rm Mpc}$ and field weaker than $10^{-15}{\\rm G}$. Such voids might be expected, for example, in models where a weak primordial field is amplified in shocked, turbulent regions of the intergalactic medium during the formation of large-scale structure. For a field strength $\\sim4\\times10^{-11}{\\rm G}$ in the high field regions, the value required to account for observed galactic fields if the IGMF were frozen in the protogalactic plasma, the delay of CRs produced by a burst at a distance of $100{\\rm Mpc}$ is $\\sim100{\\rm yr}$, and the fluence of secondary photons above $10{\\rm GeV}$ on hour--day time scales is $I(>E)\\sim10^{-6}E_{\\rm TeV}^{-1}{\\rm cm}^{-2}$. This fluence is close to the detection threshold of current high-energy $\\gamma$-ray experiments. Detection of the delayed flux would support the GRB-CR association and would also provide information on the IGMF structure. ",
        "introduction": "Recent gamma ray and cosmic ray observations give increasing evidence that the sources of gamma ray bursts (GRBs) and of cosmic rays (CRs) with energy $E>10^{19}{\\rm eV}$ are cosmological (see \\cite{cos} for GRB observations review; \\cite{Fly1}, \\cite{AGASA2}, \\cite{Wb} for CRs). The sources of both phenomena, however, remain unknown. In particular, most of the CR sources discussed so far have difficulties in accelerating CRs up to the highest observed energies (e.g., \\cite{huge}). Although the source of GRBs is unknown, their observational characteristics impose strong constraints on the physical conditions in the $\\gamma$-ray emitting region (\\cite{scen1}, \\cite{scen2}), which imply that protons may be accelerated in this region to energies $10^{20} - 10^{21} {\\rm eV}$ (\\cite{Wa}, \\cite{Vietri}). In addition, the average rate (over volume and time) at which energy is emitted as $\\gamma$-rays by GRBs and in CRs above $10^{19} {\\rm eV}$ in the cosmological scenario is, remarkably, comparable (\\cite{Wa},b). These two facts suggest that GRBs and high-energy CRs may have a common origin.  An essential ingredient of the GRB model for CRs is the time delay due to intergalactic magnetic fields. The energy of the most energetic CR detected by the Fly's Eye experiment is in excess of $2\\times10^{20}{\\rm eV}$ (\\cite{Fly1}), and that of the most energetic AGASA event is above $10^{20}{\\rm eV}$ (\\cite{AGASA2}). On a cosmological scale, the distance traveled by such energetic particles is small: $<100{\\rm Mpc}$ ($50{\\rm Mpc}$) for the AGASA (Fly's Eye) event (e.g., \\cite{huge}). Thus, the detection of these events over a $\\sim5 {\\rm yr}$ period can be reconciled with the rate of nearby GRBs, $\\sim1$ per $50\\, {\\rm yr}$ in the field of view of the CR experiments out to $100 {\\rm Mpc}$ in a standard cosmological scenario (e.g., \\cite{rate2}), only if there is a large dispersion, $\\geq50{\\rm yr}$, in the arrival time of protons produced in a single burst (Note, that this implies that if a direct correlation between high energy CR events and GRBs, as recently suggested by \\cite{MU}, is observed on a $\\sim10{\\rm yr}$ time scale, it would be strong evidence {\\it against} a cosmological GRB hypothesis). The required dispersion may result from deflections of CR protons by the inter-galactic magnetic field (\\cite{Wa}).  The inter-galactic magnetic field (IGMF) has not been detected so far. Faraday-rotation measures set an upper limit of $\\sim10^{-9}{\\rm G}$ for a field with $1 {\\rm Mpc}$ correlation length (see \\cite{Kron} for review). Other methods have recently been proposed to probe fields in the range $10^{-10}$--$10^{-20}{\\rm G}$ (e.g., \\cite{Plaga}, \\cite{Olinto}, \\cite{Avi}). Theoretical considerations regarding the existence and strength of the IGMF are related to the formation of the observed $\\mu{\\rm G}$ fields in galaxies. Recent studies suggest that a galactic dynamo cannot produce the observed large-scale fields in galactic disks (\\cite{Anderson}) and that one must turn to alternative mechanisms, which typically rely on a pre-existing field. Galactic fields might be created, for example, by compression of much weaker fields in collapsing protogalactic regions. This mechanism requires a protogalactic field of strength $10^{-11}$--$10^{-10}{\\rm G}$ and correlation length of order $1{\\rm Mpc}$. Such fields could be primordial, in which case they would likely permeate all intergalactic space. However, this need not be the case. For example, such fields could be generated from a much weaker primordial field, $\\sim 10^{-20}$~G, due to the turbulence induced in the formation of large scale structure in the universe (\\cite{Kulsrud}). In this picture, the IGMF would ``trace the mass'', with high $10^{-11}$--$10^{-10}{\\rm G}$ fields in the high density (proto-galactic) regions of the large scale structure and very low fields in the intervening voids.  Most of the energy lost by the CRs as they propagate and interact with the microwave background is transformed by cascading into secondary GeV-TeV photons (e.g., \\cite{cas3}, \\cite{cas1}). In this {\\it Letter} we show that even though the CR time delay must be $\\gtrsim50{\\rm yr}$, a significant fraction of the GeV-TeV cascade radiation can arrive with much shorter delays, on the order of hours to days, provided that a large fraction of the inter-galactic space is occupied by magnetic ``voids'', regions of very low magnetic field ($<10^{-15}$ G). In \\S 2 we present a qualitative discussion of the development of electro-magnetic cascades in the presence of an IGMF. The expected high energy photon flux is calculated using detailed Monte-Carlo simulations in \\S 3. Implications for current and future high energy gamma-ray experiments are discussed in \\S 4. ",
        "conclusions": "Although a scenario in which CRs above $10^{20}{\\rm eV}$ are produced by cosmological GRBs requires the arrival time of CRs to be delayed with respect to the $\\gamma$-rays by more than $\\sim50{\\rm yr}$, we have shown that the delay of secondary, $0.01-100{\\rm TeV}$ cascade photons may be much smaller. A short delay is possible provided that a large fraction of the inter-galactic medium is occupied by large-scale magnetic ``voids'', regions of size $\\gtrsim5{\\rm Mpc}$ and field weaker than $10^{-15}{\\rm G}$. For a field strength of $\\sim4\\times10^{-11}{\\rm G}$ in the high-field inter-galactic regions, which would account for the observed galactic fields if it were frozen in the protogalactic plasma, the delay of CRs produced by a burst at a distance of $100{\\rm Mpc}$ is $\\sim100{\\rm yr}$. At the same time, the fluence of secondary $>1{\\rm TeV}$ photons on a $1{\\rm day}$ time scale would be $\\sim10^{-6}{\\rm cm}^{-2}$, provided that $\\sim80\\%$ of the $20{\\rm Mpc}$ region around the source is occupied by magnetic ``voids'' (The fluence is inversely proportional to the burst distance squared, since the photon time delay is independent of the burst distance). The integral photon number flux in the energy range $10{\\rm GeV}<E<10{\\rm TeV}$ is approximately $I(>E) \\propto E^{-1.0}$. The flux at higher energies is very sensitive to the infrared background intensity. If the intensity is near the lower bound of current estimates, the flux extends beyond $10{\\rm TeV}$, approximately as $I(>E) \\propto E^{-1.4}$. If the intensity is close to its current upper bound, it would completely suppress the $>10{\\rm TeV}$ flux from distances $\\geq100{\\rm Mpc}$.  A $3\\sigma$ detection of $\\geq1{\\rm TeV}$ photons by current high-energy $\\gamma$-ray experiments requires a fluence $\\sim10^{-6} {\\rm cm}^{-2} \\sqrt{t_{day}} E_{min,TeV}^{-1}$, where $t_{day}$ is the observation time measured in days and $E_{min,TeV}$ is the detector threshold energy in TeV (see \\cite{CYGNUS} for CYGNUS, \\cite{HEGRA} for HEGRA, \\cite{CASA} for CASA-MIA, \\cite{Whipple} for Whipple). This fluence is close to that expected from a burst at a distance of $\\sim100{\\rm Mpc}$. However, in a cosmological model, the rate at which GRBs occur in a $100{\\rm Mpc}$ sphere around us is low, $\\sim0.1{\\rm yr}^{-1}$. A factor of $10$ increase in the sensitivity of TeV detectors, as expected for example in the near future in the Whipple and HEGRA observatories, would allow the detection of the delayed flux from bursts occuring at distances up to $\\sim300{\\rm Mpc}$. The rate of such bursts is $\\sim2{\\rm yr}^{-1}$.  Lower threshold energy, space-based detectors such as EGRET may also detect the delayed flux. At $10{\\rm GeV}$, EGRET has an effective area of $\\sim10^3\\, {\\rm cm}^2$. Thus, for the $>10{\\rm GeV}$ fluence expected from a burst at a $100{\\rm Mpc}$ distance, $\\sim10^{-4}{\\rm cm}^{-2}$ in one day, the probability that EGRET detects a $>10{\\rm GeV}$ photon is $\\sim0.1$. This probability, although not negligible, is small. Therefore, detection of the delayed flux would probably require next-generation GeV instruments, such as GLAST (\\cite{GLAST}), that are expected to have order of magnitude better sensitivity [It should be noted, however, that EGRET has detected a $18{\\rm GeV}$ photon from the direction of one of the strongest BATSE bursts (second in fluence), with $\\sim1.5{\\rm hours}$ delay (\\cite{GeV})]."
    },
    "9603/atom-ph9603005_arXiv.txt": {
        "abstract": "We report on a numerical study of the density matrix functional introduced by Lieb, Solovej and Yngvason for the investigation of heavy atoms in high magnetic fields. This functional describes {\\em exactly} the quantum mechanical ground state of atoms and ions in the limit when the nuclear charge $Z$ and the electron number $N$ tend to infinity with $N/Z$ fixed, and the magnetic field $B$ tends to infinity in such a way that $B/Z^{4/3}\\to\\infty$. We have calculated electronic density profiles and ground state energies for values of the parameters that prevail on neutron star surfaces and compared them with results obtained by other methods. For iron at $B=10^{12}$ G the ground state energy differs by less than 2 \\% from the Hartree-Fock value. We have also studied the maximal negative ionization of heavy atoms in this model at various field strengths. In contrast to Thomas-Fermi type theories atoms can bind excess negative charge in the density matrix model. For iron at $B=10^{12}$ G the maximal excess charge in this model corresponds to about one electron.\\\\ \\\\ \\noindent {PACS numbers: 31.15.-p, 03.65.-w, 32.10.-f, 97.60.Jd} ",
        "introduction": "The properties of matter in magnetic fields of the extreme strength of $10^{12}$ Gauss and higher have been the subject of numerous investigations since the early seventies, a major impetus being the discovery of pulsars in 1968 and the resulting interest in magnetized neutron stars. We refer to \\cite{C92}, \\cite{M92}, \\cite{R94}, \\cite{LSY94a},\\cite{LSY94b} for general reviews on this subject and lists of references.  The standard Hamiltonian of atomic physics, \\begin{eqnarray} H_{N,B,Z}&=&\\sum_{i=1}^N\\left\\{ [(\\p^{(i)}+\\A(\\bm r^{(i)}))\\cdot{\\sigma}^{(i)}]^2 -Z\\vert \\bm r^{(i)}\\vert^{-1}\\right\\}\\nonumber \\\\ && + \\sum_{1\\leq i<j\\leq N}\\vert \\bm r^{(i)}-\\bm r^{(j)}\\vert^{-1}\\label{dm1} \\end{eqnarray} is usually taken as a starting point for the study of atoms in the atmosphere and outermost crust of neutron stars. Here $N$ is the number of electrons that move in the Coulomb field of a nucleus, localized at the origin with charge $Ze$, and in a homogeneous magnetic field ${\\bf B}=(0,0,B)$ with vector potential $\\A({\\bm r})=(1/2)(-yB,xB,0)$. The Hamiltonian (\\ref{dm1}) operates on antisymmetric wave functions $\\Psi\\in\\bigwedge_1^N L^2({\\bf R}^3;{\\bf C}^2)$ of space and spin variables, and $\\sigma=(\\sigma_1,\\sigma_2,\\sigma_3)$ is the vector of Pauli matrices. Units are chosen such that $\\hbar=e=2m_e=1$, $c=1/\\alpha\\approx 137$; the energy unit is then four times the Rydberg energy, i.e., 54.4 eV. Besides the atomic Hamiltonian (\\ref{dm1}) it is, of course, important to study the Hamiltonian for molecules and matter in bulk, but the present paper is only concerned with (\\ref{dm1}), more specifically with its ground state energy \\be E^{\\mbox{\\tiny Q}}(N,B,Z)=\\inf_{(\\Psi,\\Psi)=1}(\\Psi,H_{N,B,Z}\\Psi),\\label{dm2}\\ee and the ground state electron density is \\begin{eqnarray} &&\\uprho^{\\mbox{\\tiny Q}}_{N,B,Z}(x)= N \\sum_{s^{(i)}=\\pm\\mfr1/2} \\\\ &&\\times \\int\\left|\\Psi_0(x,x^{(2)},\\ldots,x^{(N)} ;s^{(1)},\\ldots,s^{(N)})\\right|^2 dx^{(2)}\\cdots dx^{(N)}\\label{dm3} \\nonumber \\end{eqnarray} where $\\Psi_0$ is a ground state wave function.  Previous works on matter in strong magnetic fields can roughly be divided into two classes. On the one hand the focus has been on light atoms, in particular hydrogen with $Z=1$, on the other hand on heavy atoms with high $Z$. The present contribution falls into the second class. Here $Z=26$ plays a special role because iron is believed to be the most abundant element in the surface layer of a neutron star \\cite{C92}, \\cite{M92}. For such heavy atoms it is reasonable to expect that important aspects can be extracted from an asymptotic analysis in $Z$, and since $10^{12}$ G is large even compared with the natural atomic unit $B_0=m^2e^3c/\\hbar^3=2.35\\times 10^9$ G, an asymptotic analysis in $B$ is equally called for.\\footnote{With our choice of units $2m_e=1$ and the magnetic field is actually measured in units of $4B_0=9.40\\times 10^9$ G.}  The asymptotic behavior of the energy (\\ref{dm2}) and density (\\ref{dm3}) as $N,Z\\to\\infty$, where $N/Z$ is fixed and the magnetic field $B$ is allowed to vary with $Z$ as well, has recently been rigorously studied by Lieb, Solovej and Yngvason \\cite{LSY92},  \\cite{LSY94a}, \\cite{LSY94b}. In these papers it was proved that the ground state properties of (1) can in this limit be evaluated {\\em exactly} by five nonlinear functionals corresponding to different physics at different scales of the magnetic field $B$ as measured by powers of $Z$. These five parameter regions are characterized as follows: {\\em Region 1}, $B\\ll Z^{4/3}$; {\\em Region 2}, $B\\sim Z^{4/3}$; {\\em Region 3}, $Z^{4/3}\\ll B\\ll Z^3$; {\\em Region 4}, $B\\sim Z^3$; {\\em Region 5}, $B\\gg Z^3$. Here $B\\ll Z^p$, $B\\gg Z^p$ and $B\\sim Z^p$ means respectively that the ratio $B/Z^p$ tends to $0$, $\\infty$ or a constant $\\neq 0$ as as $Z\\to\\infty$.  The asymptotic theories corresponding to Regions 1-3 are semiclassical theories of Thomas-Fermi type that have been extensively applied to neutron stars in the past, see, e.g., \\cite{K70},\\cite{MRS71},\\cite{C74},\\cite{FGP89},\\cite{AS91},\\cite{RFGPY93} and \\cite{SO84}. Salient features of atoms in region 5 were captured by a different density functional theory already in the papers \\cite{KK71b} and \\cite{KK71a}. However, the conditions on the surface of a typical neutron star correspond rather to region 4, and this asymptotic region is also the most interesting one from the mathematical point of view. In fact, in \\cite{LSY94a} it is shown that it can be described by a functional of a novel type, where the variable is not a density, but a function with values in density {\\em matrices}. Moreover, this theory covers regions 3 and 5 as limiting cases. We refer to it as the density matrix (DM) theory.  In view of the fact that the DM theory is an exact limit of quantum mechanics it is important to know its properties in some detail. Being an asymptotic theory it is clear that it does not encompass the same information as the full Hamiltonian at finite $Z$ and $B$. In particular the DM theory does not capture exchange-correlation effects, and it is a theory of very strong fields in the sense that all electrons are confined to the lowest Landau band. These features should not be considered as a shortcoming of the DM theory, however. In fact, the hardest part of the derivation of the limit theorems in \\cite {LSY94a} is precisely to prove rigorously that contributions from exchange and higher Landau bands vanish in the limit considered. The DM theory should be judged in its own merits: It is enormously more simple numerically than the full quantum mechanical problem (\\ref{dm2}) and it is a well defined starting point for more refined approximations.  In the present contribution we report on a numerical study of the DM theory for atoms. We have computed ground state energies and electronic density profiles over a wide range of parameters and compared them with results obtained by different methods. In particular we compare the DM theory to the semiclassical theory that applies in Region 3, the simple density functional theory for Region 5, and also to other density functional \\cite{J85}, \\cite{KM88} and HF \\cite{KOONIN87} calculations. The difference between DM and HF calculations of ground state energies is less than 2\\% where data are available so that comparison can be made. This is remarkable in view of the fact that for standard Thomas Fermi theory with $B=0$ the Scott term, which corrects for the rough treatment of the electrons close to the nucleus in TF theory, must be incorporated in order to achieve such a good numerical agreement, cf.\\ \\cite{E88}. Thus, at least at this field strength,  DM theory is closer to HF theory than might have been expected. A more precise statement requires an analysis of the next to leading order terms in the asymptotic expansion of the ground state energy. Such an analysis has yet to be carried out.  Another point where the DM theory differs from semiclassical theories is in the possibility of negative ionization. It is a general feature of Thomas-Fermi type theories, based on potential theoretical arguments, cf.\\ \\cite{L82}, that the number of bound electrons never exceeds $Z$. For the quantum mechanical problem the meaning of this is simply that the binding energy of an excess electron must necessarily be of lower order in $Z$ than the ground state energy.  On the other hand it is known that a magnetic field enhances binding, for instance the Hamiltonian (1) with $N=Z+1$ has infinitely many bound states for $B\\neq 0$ \\cite{AHS81}. In the limit of extremely strong fields in Region 5 the negative charge can even be as large as $2Z$ \\cite{LSY94a}. The only rigorous results on the DM theory concern this extreme limit, but our numerical computations clearly show negative ionization that increases with $B$. It seems, however, that in order to approach the $2Z$ value extremely strong fields are needed; even at fields as strong\\footnote{The computations at these extreme field strength were carried out mainly to test the mathematical properties of DM theory. It is clear that doubts about the applicability of the nonrelativistic Hamiltonian (\\ref{dm1}) can be raised in such extreme fields, even for very heavy atoms.} as $10^{18}$ G the excess charge for iron is \\lq\\lq only\\rq\\rq\\ about 23 \\%.  Our interest in negative ionization is also motivated by its relation to another question, the binding of atoms into molecules and chains. Although a rigorous mathematical theorem linking these two aspects of binding does not seem to exist, it is a fact that in regions 1-3, i.e., for $B\\ll Z^3$, molecular binding energies are vanishingly small compared to ground state energies, whereas in region 5 binding becomes extremely strong: For a diatomic molecule the binding energy is 6 times the ground state energy of an individual atom! The question whether iron is weakly or strongly bound at field strengths of the order $10^{12}$ G has been controversial over the past 25 years. The best HF results \\cite{KOONIN87} indicate weak or no binding, but the computations are difficult for they amount to subtracting one large number from another.  It is decisive to treat the molecules and the individual atoms consistently by the same numerical methods so that unavoidable errors cancel as far as possible. Since the DM theory is numerically much simpler than HF theory it is easier to achieve this in the former and we shall return to the binding question in a separate paper. The atomic computations presented here are a necessary preparation for the study of molecules and chains. ",
        "conclusions": "We have carried out a numerical study of the density matrix model that describes exactly the quantum mechanical ground state of atoms in a homogeneous magnetic field in the asymptotic limit when the nuclear charge $Z$, the electron number $N$ and the magnetic field $B$ tend to $\\infty$ with $N/Z$ fixed and $B/Z^{4/3}\\to\\infty$. The calculations demonstrate the following features of heavy atoms in high magnetic fields as the field strength increases: A transition from an approximately spherical shape to a highly elongated shape, accompanied by a decrease in ground state energy and increasing ability to bind excess electrons. We have also compared the DM model with the semiclassical Thomas-Fermi theory and the one dimensional density functional theory that describe respectively its low  and high field limits. When $(B/Z^3)^{1/5}$ is of order unity these simpler theories are numerically and conceptually wrong and the full DM theory should be used. The quantitative agreement of DM theory with Hartree Fock calculations is quite good in strong fields; for iron at $B=10^{12}$ G the difference in binding energies is less than 2 \\%. The DM theory, however, is much simpler computationally than HF theory for large atoms and appears suitable as a starting point for more refined approximations and for the study of molecular binding in strong fields. \\appendix"
    },
    "9603/astro-ph9603078_arXiv.txt": {
        "abstract": "We measure the mean frequencies of acoustic waves propagating toward and away from a sunspot employing a spot-centered Fourier-Hankel decomposition of $p$-mode amplitudes as measured from a set of observations made at the South Pole in 1991. We demonstrate that there is a significant frequency shift between the inward and outward traveling waves which is consistent with the Doppler effect of a radial outflow from the sunspot. For $p$-modes of temporal frequencies of 3 mHz it is observed that the frequency shift decreases slightly with spatial frequency, for modes with degree $\\ell$ between 160 to 600. From the $\\ell$ dependence of the frequency shift, we infer that the mean radial outflow within the observed annular region (which extends between 30 and 137 Mm from the spot) increases nearly linearly with depth, reaching a magnitude of about 200 m/s at a depth of 20 Mm. This outflow exhibits properties similar to flows recently reported by Lindsey, et al. (1996) using spatially sensitive local helioseismic techniques. ",
        "introduction": "In the past several years, observational and theoretical tools have been developed in the field  of local helioseismology to study the interaction of solar acoustic waves with localized structures within the solar interior. These have included the use of ``ring diagrams'' as probes of local convective motions (Hill 1988, Patron, et al. 1995), time-distance helioseismic measurements of sound speed and Doppler perturbations (Duvall et al. 1996, D'Silva et al 1996, Kosovichev 1996), scattering and absorption of waves by sunspots (Braun, Duvall and LaBonte 1988, Bogdan, et al. 1993, Braun 1995, Fan, Braun and Chou 1995), and acoustic power maps (Lindsey and Braun 1990, Braun, et al 1992, Toner and LaBonte 1993, Lindsey et al. 1996).  Recently, new evidence for subsurface flows associated with active regions has been presented. Duvall, et al. (1996) have used time-distance helioseismic techniques to construct spatially resolved maps of the $p$-modes travel times over the solar surface, using full-disk solar images obtained at the geographic South Pole in 1991.  They found considerable time travel decreases in acoustic rays propagating away from active regions, which they interpreting as being due to 2000 m/s downflows within the upper 2 Mm of the convection zone. Lindsey, et al. (1996) constructed acoustic power maps of the same South Pole data, filtered in the Fourier domain to show Doppler signals caused by horizontal flows. These maps showed the presence of outflows from active regions which, for mature regions appeared to be predominately concentrated at depths greater than 8 Mm below the surface.  An important analysis procedure employed in the exploration of $p$-mode -- sunspot interactions has been the decomposition of the solar oscillation signal, observed in an annulus surrounding a spot, into appropriate inward and outward propagating wave modes. This method has been termed ``Fourier-Hankel spectral decomposition'', since the annulus is usually chosen to be small enough so that the spatial (radial) form of the wave modes is described very closely with Hankel functions. The initial utility of this procedure was demonstrated by the exploration of the $p$-mode absorption qualities of sunspots and other solar activity by looking at the difference in amplitudes between the inward and outward wave components (Braun, Duvall, and LaBonte 1988, Bogdan et al. 1993, Braun 1995, Chen, et al. 1996). Braun, et al. (1992) and Braun (1995) have subsequently determined the scattering phase shifts and mode-mixing amplitudes due to the refractive properties of spots.  We suggest that Fourier--Hankel spectral decomposition methods are highly useful in studying horizontal flows associated with active regions. For example, a net radial inflow or outflow centered on a particular point (e.g. a sunspot) will produce equal but opposite shifts in the horizontal wavenumber (at constant temporal frequency) of $p$-modes traveling toward and away from the center. The signature of a radial flow would then be a shift in the position of the $p$-mode ridge observed in power spectra constructed alternately from the inward and outward propagating waves. This procedure has a well known analogy in the determination of internal solar rotation by the measurement of frequency differences between prograde and retrograde traveling global $p$-modes. Before proceeding, we point out one important distinction between the use of the Hankel decomposition method in determining horizontal flows and its use in probing the absorbing and refracting properties of solar active regions. The amplitude and phase differences which were the primary diagnostics in the latter analyses were used to probe absorption and scattering effects which were largely confined to a area {\\it within the inner radius of the annulus}. In contrast, the radial Doppler diagnostic we are proposing here is, to first order, only sensitive to the mean radial flow of the medium {\\it between the inner and outer radii of the annulus}. A consequence of this is that the magnitude of the observed frequency shift will depend not only on the magnitude of the true velocity flow, but also on the choice of the annulus dimensions. ",
        "conclusions": ""
    },
    "9603/astro-ph9603121_arXiv.txt": {
        "abstract": "We estimate the level of confusion to Cosmic Microwave Background (CMB) anisotropy measurements caused by extragalactic infrared sources.  CMB anisotropy observations at high resolution and high frequencies are especially sensitive to this foreground. We use data from the COBE satellite to generate a Galactic emission spectrum covering mm and sub-mm wavelengths.  Using this spectrum as a template, we predict the microwave emission of the 5319 brightest infrared galaxies seen by IRAS. We simulate skymaps of extragalactic infrared sources over the relevant range of frequencies (30-900 GHz) and instrument resolutions ($10'-10^{\\circ}$ FWHM). Analysis of the temperature anisotropy of these skymaps shows that a reasonable observational window is available for CMB anisotropy measurements. ",
        "introduction": "The COBE detection of large-angular scale Cosmic Microwave Background anisotropy (Smoot et al. 1992) has generated interest in measuring CMB anisotropy on all angular scales with the goal of determining cosmological parameters. Current anisotropy observations look at  sub-degree angular scales which correspond to observable structures in the present universe.  Improved instrumentation and the MAP (Microwave Anisotropy Probe) and Max Planck Surveyor (formerly COBRAS/SAMBA) satellite missions focus attention on angular scales between one-half and one-sixth of a degree.  Due to its large beam size, COBE was basically unaffected by extragalactic foreground sources (Banday et al. 1996, Kogut et al. 1994). Because the antenna temperature contribution of a point source increases with the inverse of the solid angle of the beam, observations at higher angular resolution are more sensitive to extragalactic foregrounds, including radio sources, the Sunyaev-Zel'dovich effect from galaxy clusters, and the infrared-bright galaxies examined here.  Previous work in this area (Toffolatti et al. 1994, Franceschini et al. 1989, Wang 1991) used galactic evolution models with specific assumptions about dust temperatures to predict the level of extragalactic foreground. We choose instead a phenomenological approach using the infrared-bright galaxies detected by the Infrared Astronomical Satellite (IRAS) and the Galactic emission detected by the COBE (Cosmic Background Explorer) satellite. Section 3 compares our results with those from galaxy-evolution models.  The FIRAS (Far-Infrared Absolute Spectrophotometer) instrument of COBE gives evidence for the existence of Cold ($<15$K) Dust in the Galactic plane (Reach et al. 1995). If the Milky Way has Cold Dust, then it is likely present in other dusty spirals, which comprise the majority of bright extragalactic infrared sources.  Some observations (Chini et al. 1995, Block et al. 1994, Devereux \\& Young 1992) indicate the presence of Cold Dust in other galaxies.  Neither galactic evolution models nor pre-FIRAS observations (see Eales et al. 1989) were able to set tight constraints on emission from Cold Dust, but the FIRAS observations do. Emission from dust close in temperature to the 2.73 K background radiation is difficult to separate from real CMB anisotropies. If Cold Dust is typically accompanied by Warm Dust in spiral galaxies, we can use the FIRAS information about the total dust emission of the Galaxy to overcome this spectral similarity. ",
        "conclusions": "Our usage of the Galactic far-infrared emission spectrum as a template causes systematic errors on a galaxy-by-galaxy basis. It is easy to place constraints on our results; if all galaxies had only 33K dust as is typical for active galaxies, the resulting anisotropy would be a factor of 100 lower.  This is highly unlikely, because we know that most IRAS galaxies are inactive spirals, and galaxies with colder dust will dominate the anisotropy at mm-wavelengths because of the selection effect favoring sources with flatter spectra.  A robust upper limit on microwave anisotropy from infrared galaxies can be set by assuming that our IRAS 1.2 Jy galaxies cause the full cosmological far-infrared background (for which Puget et al. 1996 claims a detection and Mather et al. 1994 gives an upper limit).  In this case we have underestimated the anisotropy by a factor of 100, but no predictions of the IR background expect these nearby galaxies to produce more than a few percent of it. A more realistic check on our results comes from Andreani \\& Franceschini (1995), who measured a complete sample of IRAS galaxies at 1300 $\\mu$m (240 GHz).  Their average flux ratio of 1300 $\\mu$m over 100 $\\mu$m is half that of the Galaxy, but one of their beam correction methods brings their ratio into agreement with the Milky Way.  They find that the 60 $\\mu$m emission of spiral galaxies receives enough contribution from a starburst dust component mostly absent in the Galaxy that including 60 $\\mu$m fluxes in our fits may have caused a factor of 2 overestimate. Combined, these corrections give us a possible systematic overestimate of anisotropy by a factor of 4.  If typical IR-bright galaxies have dust colder than the Milky Way, our results could instead be an underestimate by a factor of a few, but this appears less likely. Our total systematic error is probably less than a factor of 3, which is consistent with our good agreement with previous results discussed in Section 3.  The recently obtained spectral knowledge of our Galaxy has enabled us to take into account the possible presence of Cold Dust. Our predicted level of temperature anisotropy makes the extragalactic infrared foreground dominant over the Galactic foregrounds of dust, free-free, and synchrotron for angular resolutions near $10'$ and frequencies above 100 GHz. Below 100 GHz, radio sources are expected to be the dominant extragalactic foreground. The extragalactic infrared foreground will not be significant in comparison to CMB anisotropies around 100 GHz but will be dominant above 500 GHz. Despite the likely presence of Cold Dust in infrared-bright galaxies, our results leave a window at intermediate frequencies for the measurement of CMB anisotropies without significant confusion from extragalactic infrared sources."
    },
    "9603/astro-ph9603073_arXiv.txt": {
        "abstract": "Ly$\\alpha$ forests are usually associated with intergalactic gas clouds intercepting the quasar sight--line.  Using as an example the Ly$\\alpha$ forest of QSO 0055--26 ($z_{em} = 3.66$), we show that the probability of observing a line at the velocity difference $\\Delta v$ has an intermittent structure in the redshift space.  On small scales (70 km~s$^{-1} \\leq \\Delta v \\leq 1580$ km~s$^{-1}$) the signature of intermittency appears as a self--similar structure of spikes systematically visible at the same redshift on all scales.  This behaviour can be interpreted as due to the presence of clustering of the Ly$\\alpha$ lines, which appear as singularities of the density of the probability measure.  From a direct measurement of the generalized dimensions, we show that the intermittency can be described by a multifractal structure which is due to a Fourier phase correlation of the signal.  The multifractal structure disappears at scales larger than $\\Delta v \\simeq 1580$ km~s$^{-1}$. ",
        "introduction": "The numerous absorption lines observed in the quasar spectra, and commonly referred as the Ly$\\alpha$ forest, are usually ascribed to intergalactic gas clouds which intercept the quasar sight--line (Sargent et al., 1980). In this sense Ly$\\alpha$ lines represent good tracers of the intergalactic matter distribution. Up to now data analysis of Ly$\\alpha$ forests is essentially based on the traditional statistical methods, like for example the two--point correlation function. However it would be quite interesting to perform an analysis where the traditional methods are replaced, or complemented, by the investigation of the scaling properties of the Ly$\\alpha$ column densities. In fact, the two--point correlation function has failed to highlight without ambiguity the presence of clustering in the Ly$\\alpha$ forest on low resolution data (Sargent et al. 1980; Bechtold et al. 1987; Webb \\& Barcons 1991). High resolution observations (Chernomordik 1995; Cristiani et al. 1996) has revealed clustering up to scales of about 300 km~s$^{-1}$.  We attempt to apply a method where spurious effects due to the finite size of the sample and to the poorly known column density distribution, can be controlled. The aim is to unambiguously recognize the presence of clusters or voids in the sample without losing information on their localization. In fact gravitational clustering should manifest itself as structures localized in redshift, detectable on all the dynamically interesting scales. The statistical analysis which we would like to perform is based on the idea that clustering is due to a Fourier phase correlation, which generates intermittency. The scaling behaviour of a probability measure derived from the intermittent signal is singular and can be described by the multifractal geometry. A quite similar approach has recently been used also by Pando \\& Fang (1995) who analyzed the scaling behaviour in the structure of the QSOs Ly$\\alpha$ lines by using the wavelets decomposition analysis. Starting from the work by Frisch \\& Parisi (1985), multifractals have been invoked to describe many physical phenomena where anomalous scaling laws are present. Examples include chaotic dynamical systems, the rate of energy dissipation in fluid flows etc (Halsey et al., 1986, and the reviews by Paladin \\& Vulpiani, 1987; Meneveau \\& Sreenivasan, 1991, and references therein). Multifractal geometry in astronomy has been used mainly to describe the mass distribution of galaxies (among others see for example Martinez et al., 1990; Jones et al., 1992; Coleman \\& Pietronero, 1992; Borgani et al., 1993, 1994; Martinez \\& Coles, 1994), the apparent luminosity field of galaxies (Garrido et al., 1996), the cosmic microwave background radiation (Pompilio et al., 1995) and the intermittency in the Solar Wind turbulence (Carbone, 1993).  However the literature on the subject has increased very rapidly in the last few years, and the multifractal analysis is now a standard way to recover anomalous scaling laws, if any exist, from random signals (Borgani, 1993; Vainshtein et al., 1994). ",
        "conclusions": "In order to better understand the structure of the Universe and to find, if any, the scale at which it  becomes homogeneous, a large sample of objects with a large spatial distribution is required. In this respect, the QSO absorption spectra available to the scientific community represent a valuable test, being distributed in a much more extended region (which can reach a space coverage of several hundred $h^{-1}$ Mpc in comoving distances) with respect to the observed galaxy distribution, and forming a more unbiased sample because the absorbing objects are not selected according to luminosity criteria.  In the present paper we performed a statistical analysis of the Ly$\\alpha$ absorption lines based on the multifractal formalism. The main motivation to do this analysis is the fact that the behaviour of the column density $N(z)$ with redshift shows a highly intermittent structure where singularities can be detected. Since in the multifractal formalism different regions at a given scale are weighted in different ways, as $q$ is increased the regions where the strongest column densities $N(z)$ lie are weighted differently from the regions where $N(z)$ are small or absent. Through a scaling analysis we are then able to detect the signature of clustering due to gravitational effects. The presence of a multifractal structure in the behaviour of $N(z)$ should indicate the presence of a hierarchy of intergalactic Ly$\\alpha$ clouds, probably generated through contraction effects (M\\\"ucket et al.,1995). We found this structure at smaller scales, say in the range 70 km~s$^{-1} \\leq \\Delta v \\leq 1580$ km~s$^{-1}$ (corresponding to comoving distances of about $0.5 \\leq r \\leq 7.6\\ h^{-1}$ Mpc). At scales larger than $7.6\\ h^{-1}$ Mpc the multifractal structure disappears and the structure is homogeneous. It is interesting to note that the maximum scale of clustering we obtain is larger than that obtained from preliminary results with the two--point correlation function (Chernomordik 1995; Cristiani et al., 1996), which in fact found clusters up to a scale $\\Delta v \\simeq 300$ km~s$^{-1}$.  High resolution spectroscopy of quasars has recently received a boost by observations with the HIRES spectrograph of the Keck telescope (Fan \\& Tytler 1994; Hu et al. 1995). It is important to see whether the scaling behaviour is also present in other available spectra. Preliminary results have shown that the scaling properties persists in different quasar sight--lines. The generalized dimension for small scales is consistent with what found for QSO 0055--26, but the characteristic scale $\\Delta v_*$  shows a redshift dependence. A detailed discussion will be presented in a forthcoming paper (Savaglio \\& Carbone, 1996).  A comparison of our results with those obtained from the same analysis when applied to galaxy catalogues (CfA1, CfA2 and QDOT redshift survey) and to the Abell and ACO catalogs of galaxy clusters, should allows us to obtain a more unified picture of the mass distribution. In fact from the CfA survey one obtains significantly smaller values for the generalized dimensions, say in a three-dimensional embedding space $D_0^{(3)} \\simeq 2.1$, and $D_2^{(3)} \\simeq 1.3$ for a wide range of scales (the relation between the one--dimensional cut and the $d$-dimensional cut is simply $D_q^{(d)}=(d-1)+D_q^{(1)}$). On the contrary the Abell and ACO clusters (Borgani et al., 1994) show a scale--invariant multifractal structure only in a limited range of scales, for Abell the range is $15-60\\ h^{-1}$ Mpc, while for ACO it extends to smaller scales. In both cases $D_2^{(3)} \\simeq 2.2$, and the picture of a pure scale invariant fractal structure extending to larger distances is disproved by these analyses. Finally the multifractal scaling properties of the QDOT redshift survey (Martinez \\& Coles, 1994) extend over two well defined scaling ranges, for $10-50\\ h^{-1}$ Mpc and for $1-10\\ h^{-1}$ Mpc. In both cases $D_0^{(3)} \\simeq 2.9$ while $D_2^{(3)} \\simeq 2.77$ at large scales and $D_2^{(3)} \\simeq 2.25$ at small scales. These last results are not really different from those obtained in our case. As regards the CfA surveys and the galaxy clusters, our results would highlight the fact that galaxies and the Ly$\\alpha$ cloud have different behaviour in space. This would be natural in a universe where the baryonic dark matter is distributed in a homogeneous way, as most of the cosmological models imply and the COBE results on the cosmic microwave background radiation show.  So far we have calculated the values of $D_q$ for positive values of $q$, because this part of the spectrum emphasizes the regions characterized by high densities of the measure (clusters). Even if there is a lack of statistical significance, due to the fact that the statistics of rare events requires very long data sets, we found two scaling regimes with multifractal scaling for $\\Delta v < \\Delta v_{\\ast}$. By using the clustering paradigm (Martinez et al., 1990; Martinez \\& Coles, 1994), this implies the presence of clusters localized at these scales. On the contrary the part of the spectrum $D_q$ with negative values of $q$ characterizes the low density regions. In this way we should be able to recognize the  presence, if any, of voids. Unfortunately (Borgani et al., 1993) discreteness effects and lack of statistics heavily affect the evaluation of $D_q$ for negative $q$, and this effect is stronger than the lack of statistics which affects the positive part of the spectrum."
    },
    "9603/astro-ph9603009_arXiv.txt": {
        "abstract": " ",
        "introduction": "Because of the central role big bang nucleosynthesis (BBN) plays in the standard cosmology, it is crucial to understand how robust the BBN results are.  Consequently there has recently been intense scrutiny of different possible sources of uncertainty in the BBN analysis.  There has been attention to refining the theoretical calculations, particularly the \\he4 yields \\cite{wssok,he4yields}, as well as quantifying their uncertainties due to the errors in their input parameters \\cite{kr,skm,kk1,kk2,hata1}, and there has been examination of the observational procedures and results for possible systematic errors \\cite{osa,csta}. Finally, there has been a focus on the model-dependence that arises in making the link between the observables and the theory predictions; in this paper we explore in detail an analysis\\cite{fo} that minimizes this model-dependence.  It is useful to situate our present concerns in the context of the subject of BBN analysis, which consists of three conceptually independent features. (1) The theory itself rests upon first-principles calculations and/or detailed laboratory results; the only really free parameter to the theory calculation is the baryon-to-photon ratio $\\eta \\equiv n_{\\rm B}/n_\\gamma$, and the output for each value of $\\eta$ consists of the primordial abundances of the light elements D, \\he3, \\he4, and \\li7. (We will consider throughout only homogeneous BBN, and furthermore we will first study the ``standard'' model with $\\nnu=3$; later we will relax the latter assumption and look at possible constraints on \\nnu.) (2) The observations pertinent to BBN are determinations of the light element abundances in various astrophysical settings; for the most part, these environments are at low or zero redshift, i.e., contemporary to our own, with the exception of QSO absorption line system (QSOALS) D/H measurements which we discuss in more detail below. As the observable abundances are contemporary, one must invoke (3) the final facet of BBN analysis, namely some account of galactic chemical evolution that models the evolution of elemental and isotopic abundances from their primordial state to the present. In this paper we emphasize that the chemical evolution results needed for each of the light isotopes have very different degrees of model-dependence; consequently we will retain only the results which are the most model-independent, and pursue their implications, both for BBN and for chemical evolution.  It is necessary to distinguish between the general implications chemical evolution gives and its detailed results. Generally, metal abundances (i.e., abundances of elements with $Z \\ge 6$) increase with time, while the specific histories require the detailed, quantitative modeling of abundance evolution. Turning to the light elements, one finds that to infer the primordial D and \\he3 from local and recent abundances (i.e., solar system and the ISM) requires detailed modeling. Of the two, D has the fewest complications and depends only on the chemical evolution model features and not at all on stellar evolutionary data, since D/H is entirely destroyed in the pre-main sequence phases of stars. \\he3 is subject to considerable uncertainties regarding not only chemical evolutionary parameters but also its stellar processing; the possible production in low mass stars and the possible destruction through mixing effects in the giant branch stars. On the other hand, to infer primordial abundances for \\he4 and \\li7 requires no detailed modeling, not because their evolution is simpler, but because we may determine their evolution empirically. Specifically, we observe \\he4 and \\li7 in systems with different, and in particular low metallicities.  Since metals increase with time, we may trace the evolution and infer the primordial abundances by extrapolating the observed trends to zero metallicity. This relatively simple procedure is, however, complicated by the very likely presence of systematic errors in the data, and so we will address in detail the possible instances of and distributions of systematic errors.  Given that the inferred primordial D and \\he3 abundances suffer from large chemical evolution uncertainties, there are several approaches one may take when trying to constrain BBN.  The traditional strategy has been to use the chemical evolution results for primordial D and \\he3, and either trying to quantify the concomitant uncertainty this introduces \\cite{st92,vop,bdf}, or to use a generalized form of chemical evolution which is designed to be ``generic'' \\cite{ytsso,wssok,st95}. This approach has the particular appeal that, as it uses all of the light elements, it strongly overconstrains the one-parameter BBN theory.  On the other hand, there remains a concern about the chemical evolution framework itself, which plays a central role.  As no first-principles calculations exist for this, it is hard to be certain that the idealizations in the available models can capture all relevant aspects of the (unavailable) full solution.  This procedure breaks down when the data begin to be refined enough to show inconsistencies between theory and observation and the conclusions for BBN become overly sensitive to highly uncertain modeling.  An alternative means of addressing the chemical evolution uncertainties is to minimize the reliance on chemical evolution by using only its most general features \\cite{fo}. That is, one uses only the inferred primordial \\he4 and \\li7 abundances, rather than all four light element isotopes. Of course, one always has to use chemical evolution at some level, but for \\he4 and \\li7 we need only assume that metal abundances (specifically C, N, O, and Fe) increase with time. With only this assumption, and a prescription for including possible systematic errors, one may constrain BBN. Note that in the case of \\li7, which is observed in old low metallicity stellar populations, we also include an uncertainty on the degree to which these stars did not deplete their \\li7 abundance.  In this paper, we will first review the observational status of all four light element isotopes.  In the cases of D and \\he3, we will also review briefly the arguments which limit their implementation in an analysis of BBN. Thus we asses the validity of BBN  using the most model-independent data. We quantify the success of BBN using likelihood analysis, including theoretical uncertainties as well as the observational statistical and systematic errors. This analysis is discussed in detail, and attention is given to the significance of the amplitudes of the likelihood functions as indicators of the goodness of fit. With this analysis we find, using only \\he4 and \\li7, that BBN theory is consistent with the observations, and we determine the best values for $\\eta$. We also attempt to quantify the reliability of these predictions.  Though we are privileging \\he4 and \\li7, this is not to say that these abundances are known with absolute certainty.  We take as a starting point the abundances directly from the observations. These observations are in remarkable agreement for a baryon-to-photon ratio which is rather low, $\\eta \\simeq 1.8 \\times 10^{-10}$. A similar range has been also suggested in  \\cite{Dar}. We will also show the effect of modifying the primordial abundances.  For example, it is sometimes argued that the lithium abundance observed in halo stars is not the primordial value but represents a depleted primordial abundance \\cite{ddp,ddk}.  We will show how our results change when the effects of \\li7 depletion are included.  As we have noted, the traditional use of solar and ISM data on D/H requires the application of chemical evolution before one may use it to constrain BBN; however, there is a real chance that the recent measurements of D/H in quasar absorption systems \\cite{quas1,rh} are in fact true measures of the primordial D/H abundance.  Therefore, we will also show results of the likelihood analysis which includes \\he4, \\li7 and D. In fact we find that the agreement is remarkably good, suggesting perhaps that the putative D in quasar absorption line systems is real and very close to primordial levels.  Having determined the success of BBN, we may then use it as has long been done, as a powerful tool in particle astrophysics. On the astrophysics side, we predict the abundances of primordial D and \\he3 on the basis of the \\he4 and \\li7 data.  These results are (in addition to being available for testing against quasar absorption line systems) then useful as inputs in chemical evolution models. We also determine the potential for BBN to set limits to \\li7 depletion; these constraints become considerably stronger, when D is included in the BBN analysis. On the particle side, we let \\nnu\\ vary, and quantify the allowed extra relativistic degrees of freedom. As we will show, if we were to use BBN, to ``predict\" the value of \\nnu\\, we would find, again remarkably, that the best value is \\nnu\\ = 3.0, independent of whether or not we include D in the analyses. To wit, the uncertainty in this ``prediction\" depends primarily on the uncertainty in the \\he4 observations, which at the the 2 $\\sigma$ level translates to a BBN upper limit of about 4 on \\nnu.  The paper is organized as follows: The observational data on each of the four light element isotopes is discussed in section 2, as is the role of chemical evolution in deriving primordial abundances of each isotope. We will carefully describe our choices of the input abundances and their statistical and systematic uncertainties. In section 3, we describe our likelihood analysis method and present results for BBN consistency in section 4. Included in section 4, is our discussion of the potential for constraining \\li7 depletion in stars, and our analyses including D with an assumed primordial value taken from the QSOALS measurements \\cite{rh}. In section 5, we discuss the implication of these results on particle astrophysics constraints, particularly, the value of \\nnu\\ . Finally, in section 6, we draw our final conclusions. ",
        "conclusions": "We have argued that because D and \\he3 are observed only in the solar system or the local ISM, determinations of their primordial abundances are particularly uncertain and dependent on models of galactic chemical evolution.  It is therefore important to determine the status of BBN without them.  Using a likelihood analysis we find that the compatibility of the \\he4 and \\li7 constraints depends on the assumptions one makes about the size and distribution of the possible systematic errors. Nevertheless, for the most ``standard'' assumptions, i.e.\\ taking the data on \\he4 and \\li7 at face value, we find that these two isotopes are quite compatible and give BBN concordance to a high degree of confidence. The baryon-to-photon ratio selected by this concordance is centered around $\\eta_{10} \\sim 1.8$ and ranges up to 3.8 (or 4.3 if all errors are assumed to be Gaussian).  While the putative determinations of D in quasar absorption line systems are still in their infancy, the technique can be a powerful probe. Upon introducing into our analysis the relatively high D abundance reported for several systems, we find that the agreement with \\he4 and \\li7 is striking.  It is crucial to further observe and better understand these systems, which can provide a decisive constraint on BBN.  Having found BBN theory to be consistent with observation, we turn to its implications.  In the realm of astrophysics, we find that the needed primordial D and \\he3 abundances challenge our understanding of the chemical evolution of these nuclides, particularly \\he3.  We find that the presence of \\li7 depletion does not improve the BBN agreement, and further, if one adopts the high D abundance suggested by the quasar measurements, one can determine the primordial \\li7 abundance and thus tightly constrain the degree of depletion.  In the realm of particle physics, we revisit the \\nnu\\ counting argument, with a careful treatment of all the data giving a best value for \\nnu\\ = 3.0 and a 95\\% CL upper limit of $\\nnu \\la 3.9$.  \\bigskip  We would like to thank Craig Copi, David Schramm, and Jim Truran for useful discussions. This material is based upon work supported by the North Atlantic Treaty Organization under a Grant awarded in 1994. This work was supported in part by DOE grant DE-FG02-94ER40823."
    },
    "9603/astro-ph9603114_arXiv.txt": {
        "abstract": "We have mapped the velocity structure of the cometary compact H~II region \\gtn\\ using long-slit echelle spectra of the \\brg\\ line. This technique detects line emission over a much wider area at the necessary spatial resolution compared to radio recombination line observations.  Significant structure in both the velocity centroids and the line widths is seen over the entire nebula.  Large line widths are seen ahead of the bow and in the tail which may be due to turbulent motions in shocked and interface regions respectively.  We construct analytic models of the density and velocity structure in order to attempt to distinguish between the bow shock and champagne flow models which have been put forward to explain the cometary morphology of many compact H~II regions. The bow shock model is unable to explain the large velocity gradient that we see right across the tail of the cometary region which can only be explained by the streaming motions towards low density regions in the champagne model. However, our approximation to the champagne model is also not able to fit all of the features of the data. More realistic versions of this model which include the effects of stellar winds and density gradients may be able to provide a better match to these data. ",
        "introduction": "Young hot stars have a dramatic effect on their surroundings when they first begin to ionize the ambient molecular gas to form H~II regions. These effects are perhaps most important in their relation to further star formation; do they trigger further collapse through compression (Elmegreen \\& Lada 1977) or ultimately halt the process due to dispersal of the molecular cloud (e.g. Franco \\etal\\ 1994)?  The most compact and heavily embedded sources have attracted particular attention since they have been presumed to be the youngest and therefore most relevant to high mass star formation. These small, dense regions are easily visible across great distances at radio wavelengths via their free-free emission and at far infrared wavelengths since the hot stars also heat the associated dust.  The study of ultracompact (UC) H~II regions in particular has been stimulated in recent years by the sparsely sampled VLA survey of Wood \\& Churchwell (1989).   One of the greatest puzzles of the Wood \\& Churchwell survey is the total number of such regions found.   By taking both the total number of O stars in the galaxy (simply from extrapolating the number of optically visible OB associations), and their lifetimes on the main sequence, they predicted the total number of such sources that might still be expected to be in the UC H~II phase.  Comparing this with the numbers actually found  led Wood \\& Churchwell to conclude that there was an order of magnitude more sources in their VLA survey than there should have been. Since there is no other evidence for the massive star formation rate being this much larger than expected,  the likeliest solution to this problem is  that some mechanism is constraining the outward expansion of the H~II region.  Wood \\& Churchwell estimate that the UC H~II  phase must last about ten times longer than the $\\sim10^4$ years predicted on the basis of simple Str\\\"{o}mgren sphere expansion.  Another key result from the Wood \\& Churchwell study is that about 20\\% of these regions have a `cometary' appearance.  A subsequent survey by Kurtz \\etal\\ (1994) and a VLA survey of much larger, optically visible, H~II regions by Fich (1993) also found a similar proportion of this type. The near parabolic shape of many of these sources lead Wood \\& Churchwell to develop a bow-shock interpretation for UC H~II regions which provides an explanation for both of the lifetime of the UC H~II region phase and the cometary morphologies. The basis of this model is simple.  An OB star moves supersonically through a molecular cloud ($v_*>0.2$~\\kms) and the stellar wind from the star supports a bow shock along its direction of motion. This shock can trap the ionisation front (IF), preventing it from expanding.  Such a situation is potentially stable, and the lifetime of the UC H~II phase in this model is simply the star's crossing time through the cloud, typically of the order of 10$^5$ years.  Cometary-like morphologies for H~II regions have been known for some time and were originally labelled `blisters' by Israel (1978) due to their propensity for being found near the edges of molecular clouds. Icke, Gatley \\& Israel (1980) developed the idea that if there is a density gradient in the ambient gas the H~II region will expand fastest in the low density direction and so become very asymmetric. Tenorio-Tagle and co-workers in a series of papers (Tenorio-Tagle 1979; Bodenheimer, Tenorio-Tagle \\& Yorke 1979; Tenorio-Tagle, Yorke \\& Bodenheimer 1979; Yorke, Tenorio-Tagle \\& Bodenheimer 1983, Yorke, Tenorio-Tagle \\& Bodenheimer 1984) examined the gas dynamics of this situation and found that the pressure gradient set up when the IF reaches the edge of the cloud causes a `champagne' flow of ionized gas away from the cloud with velocities of order 30 \\kms.  They also give predictions for line profiles, and for the expected radio continuum. This model does not answer the lifetime problem however since there is no constraint on the H~II region expansion.  Crucially, it also does not include the effects of the stellar wind from the OB star.  Turner \\& Matthews (1984) considered the latter problem in a uniform static configuration: they find the IF could be trapped at early times in the shell formed by the stellar wind sweeping up the dense ambient gas ($n_e\\sim10^5$~cm$^{-3}$). A combination of this effect and the blister geometry plausibly satisfies both the lifetime and morphology constraints.  The combined model would have a shell structure near the core, and diffuse emission near the tail where lower densities preclude the trapping of the IF.  The bow-shock and champagne models for cometary H~II regions can be distinguished by the velocity structure of the ionized gas. Several studies have been carried out using radio recombination lines to map out the velocity structure. Garay \\etal\\ (1986) concluded that the cometary region in the G34.3+0.2 complex exhibits a champagne flow, but also shows a large velocity gradient perpendicular to the symmetry axis which they attributed to rotation of the parent cloud. A lower resolution study of the same object by Gaume, Fey \\& Claussen (1994) found a similar velocity structure, but they reject both bow-shock and champagne flow models in preference to a picture involving interactions with outflows from other sources ahead of the bow. Similar velocity structures have been seen in two cometary regions in the Sgr B2 complex by Gaume \\& Claussen (1990). Garay \\etal\\ (1994) presented radio recombination line maps of more extended H~II regions and find evidence for a bow-shock in one and champagne flows in two other more clumpy sources. Wood \\& Churchwell (1991) carried out high resolution radio recombination line studies of the cometary UC H~II region G29.96--0.02. Their results were analysed by Van Buren \\& Mac Low (1992), who claim good agreement with their bow-shock model.  There are substantial problems with the radio recombination line approach.  These lines are intrinsically very weak, and since they arise from levels well above the ground state (Wood \\& Churchwell 1991 used H76$\\alpha$), are prone to many broadening effects and maser activity. The high spatial resolution interferometric observations only detect the lines near the head of the bow and resolve out the the weaker and more diffuse tail emission. The synthesised beam sizes required to detect the radio recombination lines in the tail are usually then too large to resolve the head region simultaneously.  By contrast, infrared recombination lines are intrinsically brighter, and the detector technology is such that we can map emission across much larger regions where the surface brightness of the line emission is lower at high and uniform spatial resolution. Even allowing for the large extinction in these objects (A$_{V}\\sim20-30$ is typical, and much larger values possible), mapping the HI Br$\\gamma$ or Br$\\alpha$ emission provides many potential benefits over further radio surveys.  We have therefore embarked on a series of IR observations of UC H~II regions selected from Wood \\& Churchwell (1989) and from Kurtz, Wood \\& Churchwell (1994).  In this paper we demonstrate the value of our method by presenting observations and analysis of the prototypical cometary UC H~II region: G29.96-0.02.  In future papers we will present further observations. ",
        "conclusions": "We have demonstrated the value of our infrared techniques in studying the kinematical structures of young, heavily extinguished H~II regions.  The data we present here for G29.96-0.02 clearly shows features not present in previous radio data that allows us to test the validity of the two proposed models in much greater detail.  In particular we have shown that the regions in which the bow shock model has the poorest fit to the data are not at all well represented in the radio recombination line data of Wood \\& Churchwell (1991). Therefore, their conclusions, and those of Van Buren \\& Mac Low (1992) and Afflerbach \\etal\\ (1994), that G29.96-0.02 is completely consistent with a bow shock interpretation must be questioned.   From our own data, and the analytic approximation to the bow shock model we have employed, we find that the greatest discrepancies lie near the `head' of the region, and along the `sides'.  For the `head' of the region, we find the broadest lines are {\\em ahead} of the point at which the line intensity reaches a maximum, whereas the bow-shock model predicts they should coincide.   This may be due to turbulent mixing processes in the interface between the ionised and molecular gas.  We do not therefore consider it a strong argument against the bow-shock model.  Similar arguments could also be applied to the existence of this gas in the champagne flow model we have described.   The one aspect of the data that the bow shock model cannot explain away through an appeal to turbulence is the velocity gradient seen along the outer edge of the H~II region in the `tail'. Our velocity centroid data show large deviations from the model. Another way to demonstrate this is to compare the maps of line flux and velocity centroid.  The bow shock model predicts that the observed velocities along the outer edge of the comet should be constant since we are seeing gas which is essentially comoving with the molecular cloud material, e.g. the lowest contour in the velocity centroid map in Fig.\\ 6 of Van Buren \\& Mac Low (1992) follows the outer edge of the emission measure. It can be seen that the opening angle of the velocity centroid is wider than that of the flux as was also apparent in the radio data (Fig.\\ 5 of Van Buren \\& Mac Low 1992). In the champagne flow model (e.g. Bodenheimer et al.\\ 1979), the velocity gradient behind the star is easy to explain:  however, other means are required to produce  the velocity gradient ahead of the star at the edge of the H~II region.  For both models presented here, it is this aspect of the data that presents the greatest challenge.  In summary, the motions in the tail of \\gtn\\ are highly suggestive of a a champagne flow, although our very simple model is still far from a good fit to all the data. Our results are consistent with the radio continuum maps of Fey et al.\\ (1995) who also argue in favour of a champagne flow model on the basis of the combined radio and near-infrared morphology of the region. As advocated by Gaume \\etal\\ (1994) and others there is an urgent need to investigate champagne flows which include the effects of stellar winds since we know the latter exist in these regions as well as taking more realistic density distributions into account. This will lead to a better understanding of how young massive stars affect their natal environments and further star formation."
    },
    "9603/astro-ph9603052_arXiv.txt": {
        "abstract": "Hot big bang cosmology says nothing about the topology of the Universe. A topology-independent algorithm is presented which is complementary to that of Lehoucq \\etal\\  1996 and which searches for evidence of multi-connectedness using catalogues of astrophysically observed objects. The basis of this algorithm is simply to search for a quintuplet of quasars (over a region of a few hundred comoving Mpc) which can be seen in two different parts of our past time cone, allowing for a translation, an arbitrary rotation and possibly a reflection. This algorithm is demonstrated by application to the distribution of quasars between redshifts of $z=1$ and $z\\approx4,$ i.e., at a comoving distance from the observer $1700 h^{-1}\\mbox{\\rm Mpc} \\ltapprox d \\ltapprox 3300 h^{-1}\\mbox{\\rm Mpc}.$ Two pairs of isometric quintuplets separated by more than {$300$\\hMpc} are found. This is  consistent with the number expected from Monte Carlo simulations in a simply connected Universe if the detailed anisotropy of sky coverage by the individual quasar surveys is taken into account. The linear transformation in (flat) comoving space from one quintuplet to another requires translations of {$353$\\hMpc} and {4922\\hMpc} respectively, plus a reflection in the former case, and plus both a rotation and a reflection in the latter. Since reflections are required, if these two matches were due to multi-connectedness, then the Universe would be non-orientable. ",
        "introduction": "Given the immense observational effort to search for evidence of the possible non-trivial curvature of the observable Universe, (in particular for an ``open'', i.e., negatively curved, $\\Omega_0<1$, Universe) it is surely equally important to search for evidence for that the observable Universe could have a non-trivial topology. Apart from the obviously fundamental significance of knowing that the volume of the Universe is finite (if the topology turns out to be compact), studies of evolution of galaxies, quasars and other objects at cosmological distances would be immensely advanced by the ability to observe the same object at significantly different epochs. General relativity and the standard hot big bang model of the Universe are independent of topology. According to several authors (e.g., Hawking 1984; Zel'dovich \\& Grishchuk 1984), quantum gravity with a ``no-boundary'' boundary condition requires the Universe to be compact. A multi-connected topology of the observable Universe would provide such compactness.  Lachi\\`eze-Rey \\& Luminet (1995) present an extensive review of what multi-connectedness of our Universe would mean and of ef\\-forts made so far to de\\-tect multi\\--connected\\-ness, so only a short introduction is given here.  A simple example of a 3-dimensional manifold with a non-trivial topology is the hypertorus. This can be constructed by identifying opposite faces of a cube (or more generally, a rectangular prism). A particle which one would otherwise expect to leave the cube at one surface ``reenters'' at the opposite face, without ``noticing'' that anything strange has happened. Many N-body simulations used to calculate the non-linear effects of gravity use a hypertoroidal topology in order to provide boundary conditions; this is termed using ``periodic'' boundary conditions. If the universe were hypertoroidal, then the identification of opposite faces would be physical, not merely a numerical technique.  A useful way of thinking about this in terms of a simply connected Universe is to imagine the cube repeated endlessly in a 3-dimensional grid. Each repetition consists of the same physical region of space, and photon paths (geodesics) can be calculated just as for the simply connected Universe. Our past time cone can then be thought of as before---but with the difference that we see the same chunk of the (whole) Universe at earlier and earlier epochs further and further away. This apparent space containing multiple copies of the Universe is called a ``covering space''. In the case of the hypertorus, and if objects (e.g., quasars) had visible lifetimes as great as that of the Universe and were all bright enough to be seen in magnitude/surface-brightness limited surveys, then an obvious grid pattern would appear.  However, the hypertorus is only one of many possible topologies. There is an immense variety of topologies of the many possible 3-manifolds. For example, there are ten flat, homogeneous, isotropic, multi-connected 3-manifolds of finite volume (``compact''), and the classification of negatively curved, homogeneous, isotropic, multi-connected, compact 3-manifolds is still an open area of research (e.g., Thurston \\& Weeks 1984). Because of evolutionary effects and the vast range of possible topologies, the grid pattern possible for a hypertorus would not necessarily be obvious to the eye, so subjective impressions of present data are certainly not sufficient to rule out a non-trivial topology.  Other topologies can be thought of by starting with a ``fundamental polyhedron'' and identifying pairs of faces. Using the conventions of \\LaLu, the shortest distance from an object to any of its ``ghost'' counterparts is labelled $\\alpha$, and the largest distance from an object to an adjacent ghost (where there is an adjacent ghost for each face of the fundamental polyhedron) is labelled $\\beta$. The choice of a fundamental polyhedron for thinking about a multi-connected topology is not unique; such a Universe (as considered here) is on average homogeneous and isotropic, without any borders in hypersurfaces of constant cosmological time.  The claimed periodicity of the galaxy redshift distribution in pencil-beam surveys (Broadhurst et al. 1990) could have been an indication of a non-trivial topology, but closer investigation of the data has not confirmed this. As argued thoroughly by \\LaLu, observations on scales larger than this are not yet extensive enough to rule out repetitions of a finite-volume multiconnected Universe above a scale of about $\\beta=600 h^{-1}$~Mpc ($\\alpha=300 h^{-1}$~Mpc).  \\Figmink  Intuitively, one might think that a simple way to test for a compact topology would be to note that in a Universe of finite size no objects could exist which are larger than the size of the Universe, so that the initial perturbation spectrum (e.g., as observed by COBE) should be zero at large length scales, or small wavenumbers $k.$ The flaw in this argument is that in the covering space, structures which physically occupy the same, or nearly the same, part of space can coincide at the ``faces'' of the fundamental polyhedron in such a way that a pseudo-structure which is larger than the Universe can exist in the covering space. In analysis which works in the covering space (which is much easier than starting with a hypothesised topology), such a pseudo-structure would violate the minimum-$k$ criterion. As a demonstration, Fig.~1 shows a simple example of a sky map due to such pseudo-structure in a hypertoroidal Universe which % is only a sixth of the horizon size. The huge structures which extend across the sky are simply due to the ``microwave background'' sphere cutting through the same (physical) structure at slightly different angles in adjacent ghost images of the fundamental polyhedron (cube).  While such a dramatic pattern should be easy to rule out, it is not clear what assumptions would be needed  about the initial perturbation spectrum and the topology in order to make the minimum-$k$ test useful. In the example here, the maximum size of the fundamental polyhedron is clearly visible in the direction perpendicular to that of the pseudo-structures, so that averaging over the two dimensions of the sphere should help in the derivation of a generalised minimum-$k$ test. However, without a rigorous derivation of such a test, detailed models with assumptions about structure formation and specific topologies are required. Some authors have indeed calculated detailed models for particular cosmologies. In particular, toroidal topologies have been shown to be inconsistent with the COBE data to the scales of several thousand $h^{-1}$~Mpc (Stevens et al. 1993) to around the horizon size (Starobinsky 1993; de Oliveira \\& Smoot 1995) on the basis of considering a few (toroidal) of the topologies possible. Others (Jing \\& Fang 1994) note that the COBE data is well fit by a non-zero low wavelength cutoff in the primordial fluctuation spectrum, which would be consistent with a finite-volume, multiconnected Universe at about $1-3$ times the horizon scale.  An alternative to a minimum-$k$ test modified in some way by prior assumptions is the topology-independent method for searching for multi-connectedness in present and future CMB data which has been proposed by Cornish {\\etal} (1996). This method is based on the property that a fundamental polyhedron of about the horizon size or smaller should intersect the last scattering sphere several times, in circles. Different copies of an intersection which represent the same subset of space-time should be seen in different parts of a CMB all-sky map. As long as the temperature of a point in the CMB is independent of the direction of the observer (e.g., the ``Doppler'' fraction of $\\delta T/T$ due to the movement of the baryon-photon fluid towards or away from the observer is small), the values of $\\delta T/T$ around corresponding circles should be identical apart from a rotation due to the particular multi-connected manifold explaining the identification. Note that the circles themselves are not of constant temperature: the pseudo-structures shown in Fig.~1 are not the same as the circles discussed by Cornish {\\etal}.   To put these tests into perspective, the reader is reminded that in comoving units, i.e., within the hypersurface of constant $t=t_0,$ the distance to an object at redshift $z$ in the hypersurface of constant $t=t_0$ is $d= 2(c/H_0) (1-1/\\sqrt{1+z}),$ the horizon is at $6000 h^{-1}$~Mpc. (Except where otherwise stated, all distances here are quoted in comoving units.) The diameter across the observable Universe is, of course, twice the radius, i.e., $12000 h^{-1}$~Mpc, so it could only be claimed that the observable Universe is simply connected if evidence were available up to a scale of twice the horizon size. ",
        "conclusions": "A topology-independent method, complementary to those of Lehoucq {\\etal} 1996 and Cornish {\\etal} 1996, which observationally searches for the effects of a non-trivial topology of the observable Universe has been presented. The method is simply to search for isometries of quasar quintuplets in different parts of our past time cone, and an algorithm which reduces the number of computations required to within practical computing cabilities has been explained.  This method is illustrated by application to the observed set of quasars available from NED, and the significance is tested by comparison to Monte Carlo simulations. Even by taking advantage of Lehoucq \\etal's (1996) (and others') results against the Universe having a multi-connected topology on a scale $\\alpha \\le 300$\\hMpc, the number of detections remains statistically consistent with that expected for a simply connected Universe.    Without a detailed physical model of quasar evolution, the method is only capable of finding evidence for multi-connectedness, not for showing simple-connectedness. If a statistically significant number of corresponding quintuplets were found in a future more complete catalogue of quasars, it should be followed by finding a fundamental polyhedron and linear transformations which explain the quintuplet matches found. These should be applied to the full quasar catalogue in order to find the transformed (``ghost'') quasar positions in the copy of the fundamental polyhedron in which galaxy and galaxy cluster positions are best known, i.e., in the region of the covering space containing the Galaxy. If the transformed positions of the quasars significantly corresponded to the centres of, say, bright elliptical galaxies, or cluster centres, this would be a confirmation of multi-connectedness, and an indication of the evolutionary link between quasars and galaxies. On the contrary, a limit on the absence of luminous objects at the transformed quasar positions would either be a refutation of the claimed multi-connected manifold or evidence that the luminosity of ``dead'' quasars falls below the detection limit."
    },
    "9603/hep-ph9603201_arXiv.txt": {
        "abstract": "The primordial abundance of long-lived heavy Majorana neutrinos is calculated from the full Boltzmann equation. Inclusion of scattering reactions drastically change the predicted abundance of a heavy neutrino species. This loosens the well known mass constraint on MeV neutrinos from Big Bang nucleosynthesis, and allows for the existence of a Majorana $\\tau$ neutrino with mass $m_{\\nu_{\\tau}} \\geq 11$ MeV. Further experimental efforts are therefore needed to investigate the range $11 \\text{MeV} \\leq m_{\\nu_{\\tau}} \\leq 24 \\text{MeV}$. Some interesting cosmological consequences of an MeV $\\nu_\\tau$ are also pointed out. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603028_arXiv.txt": {
        "abstract": "The calculation of distances is of fundamental importance in extragalactic astronomy and cosmology. However, no practical implementation for the general case has previously been available.  We derive a second-order differential equation for the angular size distance valid not only in all {\\em homogeneous\\/} Friedmann-Lema\\^\\i tre cosmological models, parametrised by $\\lambda_{0}$ and $\\Omega_{0}$, but also in {\\em inhomogeneous\\/} `on-average' Friedmann-Lema\\^\\i tre models, where the inhomogeneity is given by the (in the general case redshift-dependent) parameter~$\\eta$.  Since most other cosmological distances can be obtained trivially from the angular size distance, and since the differential equation can be efficiently solved numerically, this offers for the first time a practical method for calculating distances in a large class of cosmological models.  We also briefly discuss our numerical implementation, which is publicly available. ",
        "introduction": "The determination of distances is one of the most important problems in extragalactic astronomy and cosmology.  Distances between two objects X and Y depend on their redshifts~$z_{x}$ and $z_{y}$, the Hubble constant~$H_{0}$, the cosmological constant $\\lambda_{0}$, the density parameter $\\Omega_{0}$ and the inhomogeneity parameter~$\\eta$.\\footnote{When discussing the distance between {\\em two\\/} objects, one can always make a coordinate transformation such that the contribution from the $\\theta$ and $\\phi$ terms in Eq.~(\\ref{rwm-g}) vanish.  Then one simply needs the redshifts and cosmological parameters in order to determine the distance between them. When discussing the distances between several objects, for example QSOs with $\\alpha$, $\\delta$ and $z$ as coordinates, this is no longer possible.  In many cases, however, suitable geometrical approximations can be made so that the most complicated part of the problem is essentially a determination of a distance between two objects. This point is further discussed in Sect.~\\ref{numerics}.} Usually, smaller distances are determined by the traditional `distance ladder' technique and larger distances are calculated from the redshift, assuming some cosmological model.  Since the redshift is for most purposes exactly measurable, knowledge of or assumptions about two of the factors~(a)~Hubble constant, (b)~other cosmological parameters and (c)~`astronomical distance' (i.e.~ultimately tied in to the local distance scale) determines the third.  In this paper we discuss distances given the Hubble constant $H_{0}$,  the redshifts $z_{x}$ and $z_{y}$ and the cosmological parameters $\\lambda_{0}$, $\\Omega_{0}$ and $\\eta$. Traditionally, a simple cosmological model is often assumed for ease of calculation, although the distances thus obtained, and results which depend on them, might be false if the assumed cosmological model does not appropriately describe our universe. A general method allows one to look at cosmological models whether or not they are easy-to-calculate special cases and offers the possibility of determining cosmological distances which are important for other astrophysical topics once the correct cosmological model is known.  We stress the fact that the inhomogeneity can be as important as the other cosmological parameters, both in the field of more traditional cosmology and in the case of gravitational lensing, where, e.g.~in the case of the time delay between the different images of a multiply imaged source, the inhomogeneity cannot be neglected in a thorough analysis (Kayser \\& Refsdal \\cite{RKayserSRefsdal83a}). For an example involving a more traditional cosmological test, Perlmutter et~al.\\ (\\cite{SPerlmutter95a}) (see also Goobar \\& Perlmutter (\\cite{AGoobarSPerlmutter95a})) discuss using supernovae with $z \\approx 0.25$--$0.5$ to determine $q_{0}$; for $z$ near the top of this range or larger, the uncertainty due to our ignorance of $\\eta$ is comparable with the other uncertainties of the method.  The plan of this paper is as follows.  In Sect.~\\ref{cosmology} the basics of Friedmann-Lema\\^\\i tre cosmology are briefly discussed; this also serves to define our terms, which is important since various conflicting notational schemes are in use.  (For a more thorough discussion using a similar notation see, e.g., Feige (\\cite{BFeige92a}).) Section~\\ref{distances} defines the various distances used in cosmology. In Sect.~\\ref{rainer} our new differential equation is derived.  Similar efforts in the literature are briefly discussed.  Section~\\ref{numerics} briefly describes our numerical implementation and gives the details on how to obtain the source code for use as a `black box' (which however can be opened) for use in cosmology and extragalactic astronomy.  The symmetry properties of the angular size distance, analytic solutions and methods of calculating the volume element are addressed in three appendices. ",
        "conclusions": "\\label{summary}   After discussing cosmological distances with an emphasis on practical distance measures for general use in cosmology and extragalactic astronomy, we have obtained a new differential equation, which gives the angular size distance for a class of `on average' Friedmann-Lema\\^\\i tre cosmological models, that is, models described not only by $\\lambda_{0}$ and $\\Omega_{0}$ but also by $\\eta(z)$, which describes the clumpiness of the distribution of matter.  We have also developed a practical numerical method of solving this equation, which we have made publicly available.  Since the equation is valid for {\\em all\\/} cases, this offers for the first time an efficient means of calculating distances in a large class of cosmological models.  The numerical implementation (in {\\tt FORTRAN77}), user's guide and a copy of the latest version of this paper can be obtained from either of the following URLs: \\begin{quote} \\scriptsize\\tt http://www.hs.uni-hamburg.de/english/persons/helbig/ \\\\ Research/Publications/Info/angsiz.html \\end{quote} \\begin{quote} \\scriptsize\\tt \\tt ftp://ftp.uni-hamburg.de/pub/unihh/astro/angsiz.tar.gz \\end{quote}"
    },
    "9603/astro-ph9603046_arXiv.txt": {
        "abstract": " ",
        "introduction": "The \\gr sky as we know it today is a composite of the emission from point sources and diffuse emission (see Fig.1). The most prominent feature is the galactic plane, in which interactions between cosmic rays and the thermal gas lead to \\gr emission by $\\pi^0$-decay and bremsstrahlung. On top of this diffuse emission we see point sources all over the map, part of them being pulsars, another part distant AGN, and also a significant fraction of yet unidentified sources. But we also see that the diffuse emission extends out of the galactic plane up to the poles. There appears to be an isotropic emission component which is presumably extragalactic in origin and may be understood as blend of unresolved AGN, but there is also considerable galactic emission. At higher latitudes interactions between cosmic rays and thermal gas still play a role, but inverse-Compton scattering of ambient photons by cosmic ray electrons becomes increasingly important. This emission tells us about the physical conditions in the halo as seen by the cosmic ray particles, and it may also reveal previously hidden gas, e.g. baryonic dark matter.  \\begin{figure}[htb] \\psfig{figure=all-bw.ps,width=13.0truecm,clip=} \\caption{The EGRET sky above 100 MeV \\gr energy. The image is deconvolved by a Maximum-Entropy algorithm (for the method see Strong 1995). The grey scale is logarithmic with darkness indicating high intensity.} \\end{figure} ",
        "conclusions": ""
    },
    "9603/astro-ph9603100_arXiv.txt": {
        "abstract": "We have analysed archival {\\it ROSAT} PSPC data for M32 in order to study the x-ray emission from this nearest elliptical galaxy.  We fit spectra from three long exposures with Raymond-Smith, thermal bremsstrahlung, and power-law models.  All models give excellent fits.  The thermal fits have kT$\\approx$4 keV, the Raymond-Smith iron abundance is $0.4^{+0.7}_{-0.3}$ Solar, the power-law fit has $\\alpha$=1.6$\\pm$0.1, and all fits have $N_H$ consistent with the Galactic column.  The source is centered on M32 to an accuracy of 9$''$, and unresolved at 27$''$ FWHM ($\\sim$90 pc).  M32 is x-ray variable by a factor of 3--5 on timescales of a decade down to minutes, with evidence for a possible period of $\\sim$1.3 days.  There are two plausible interpretations for these results:  1) Emission due to low-mass x-ray binaries; 2) Emission due to accretion onto a massive central black hole.  Both of these possibilities are supported by arguments based on previous studies of M32 and other old stellar systems; the {\\it ROSAT} PSPC data do not allow us to unambiguously choose between them.  Observations with the {\\it ROSAT} HRI and with {\\it ASCA} are required to determine which of these two very different physical models is correct. ",
        "introduction": "M32 (NGC 221, PCG 2555) is not only the nearest (D$\\approx$700 kpc, \\markcite{cjfn}e.g.~Ciardullo et al.~1989) example of the low-luminosity end of the sequence of normal elliptical galaxies \\markcite{korm1985}(Kormendy 1985), but the nearest such elliptical of any luminosity.  This allows us to study it to much lower absolute luminosity and size limits than is possible for other more luminous ellipticals.  The Virgo cluster is $\\sim$20 times more distant than M32 \\markcite{p94}(e.g.~Pierce et al.~1994), thus observations at a given angular scale or flux limit probe $\\sim$400 times deeper for M32 than for Virgo Es.  As a result, our understanding of the properties of M32 provides us with a cornerstone for the investigation of more luminous and more distant ellipticals.  X-ray emission from faint early-type galaxies ($L_B \\la 10^{40}~{\\rm erg~s^{-1}}$) typically has a hard ($\\sim$5 keV) spectrum, with $L_X \\propto L_B$ \\markcite{kft92a}(Kim, Fabbiano \\& Trinchieri 1992a; \\markcite{efka1995} Eskridge, Fabbiano \\& Kim 1995).  This emission is interpreted as arising from population II stellar binary sources (low mass x-ray binaries, or LMXRBs).  The {\\it Einstein} data for M32 are consistent with this interpretation \\markcite{fkt1992}(e.g.~Fabbiano, Kim \\& Trinchieri 1992), however the {\\it Einstein} observation of M32 was not particularly deep; other models for the x-ray emission are not excluded.  In particular, optical imaging and spectroscopy both indicate that M32 contains a central black hole of $M_{\\bullet} \\approx$ 1--3$\\times 10^6M_{\\odot}$ (van der Marel et al.~1994; Lauer et al.~1992).  The observed x-ray emission may be due to low-level accretion onto this black hole.  We have analysed archival {\\it ROSAT} PSPC observations of M32 in order to study the nature of the x-ray emission with the best currently available data. In \\S 2 we discuss the available PSPC data.  In \\S 3 we investigate the spectral properties, extendedness, and time variability of the x-ray emission. We discuss the two possible physical interpretations of these data in \\S 4, and conclude with a discussion of crucial future observations in \\S 5. ",
        "conclusions": "We have presented arguments that the x-ray emission from M32 can be due to {\\it either} a small population of LMXRBs, {\\it or} accretion onto the massive central black hole.  Two observations that could resolve this ambiguity are {\\it ROSAT} HRI imaging and {\\it ASCA} spectroscopy.  A deep on-axis HRI image would determine the extent and position of the M32 x-ray source much more precisely than do the PSPC data.  It would also resolve the weak source to the NE.  The HRI on-axis PSF is $\\sim$6$''$ FWHM ($\\sim$0.17$R_e$) giving a spatial resolution of $\\sim$20 pc.  Emission from an AGN would thus be unresolved with the HRI, and would come from the center of M32.  Emission from a collection of LMXRBs would not have to be coincident with the center of M32, and could be extended at HRI resolution.  A deep {\\it ASCA} spectrum will have the spectral resolution to clearly distinguish between the thermal and power-law spectral models in an integration time of $\\sim$20 ksec.  Given the importance and proximity of M32, the opportunity to resolve the current puzzle should not be missed."
    },
    "9603/astro-ph9603010_arXiv.txt": {
        "abstract": "If \\g--ray bursts are cosmological in origin, the sources are expected to trace the large--scale structure of luminous matter in the universe. I use a new likelihood method that compares the counts--in--cells distribution of \\g--ray bursts in the BATSE 3B catalog with that expected from the known large--scale structure of the universe, in order to place a constraint on the distance scale to cosmological bursts. I find, at the 95\\% confidence level, that the comoving distance to the ``edge'' of the burst distribution is greater than $630~h^{-1}$~Mpc ($z > 0.25$), and that the nearest burst is farther than $40~h^{-1}$~Mpc. The median distance to the nearest burst is $170~h^{-1}$~Mpc, implying that the total energy released in \\g--rays during a burst event is of order $3\\times 10^{51}~h^{-2}$ ergs. None of the bursts that have been observed by BATSE are in nearby galaxies, nor is a signature from the Coma cluster or the ``Great Wall'' likely to be seen in the data at present. ",
        "introduction": "The origin of \\g--ray bursts is still unknown and is currently the subject of a ``great debate'' in the astronomical community. Do the bursts have a Galactic origin (\\cite{L95}) or are they cosmological (\\cite{Pac95})? And what is their distance scale?  In this Letter, I do not attempt to answer the first question, but rather, I show that {\\it if} one assumes that \\g--ray bursts are cosmological in origin, one can begin to answer the second question and place a constraint on the distance scale to the bursts. This is because cosmological bursts are expected to trace the large--scale structure of luminous matter in the universe (\\cite{LQ93}, hereafter, LQ). The constraint comes from comparing the {\\it expected} clustering pattern of bursts on the sky --- which will depend on their distance scale because of projection effects --- with that {\\it actually observed}. The observed angular distribution is in fact quite isotropic (\\cite{Brig96}); hence, only a lower limit to the distance scale can be placed because a sufficiently large distance will always lead to a sufficiently isotropic distribution on the sky.  Hartmann and Blumenthal (1989) first used the absence of a significant angular correlation function in a small sample of \\g--ray bursts to put a lower limit of $71~h^{-1}$~Mpc on the distance scale to the bursts.\\footnote{I follow the usual convention and take $h$ to be the Hubble constant in units of 100 km~s$^{-1}$~Mpc$^{-1}$.} This lower limit was subsequently improved to $150~h^{-1}$~Mpc (\\cite{BHL94}) using the larger BATSE 1B catalog (\\cite{Fish94}).  Here I use a powerful new likelihood method, which I had previously developed to analyze repeating of \\g--ray bursts in the BATSE 1B and 2B catalogs (\\cite{Q95}), to compare the observed counts--in--cells distribution in the new BATSE 3B catalog (\\cite{Meeg96}) with that expected for bursts at cosmological distances. I describe this method and calculate the expected counts--in--cells distribution in \\S ~2.1, using the angular correlation function computed in \\S ~2.2. I present my results on the burst distance scale in \\S ~3, and discuss some implications of this work in \\S ~4. This work was presented in preliminary form elsewhere (\\cite{Q96}). ",
        "conclusions": "If \\g--ray bursts are cosmological and trace the large--scale structure of luminous matter in the universe, and their  positional errors are as quoted in the 3B catalog, then the lack of any angular clustering in the data implies that the observed distance to the ``edge'' of the  burst distribution must be farther than $630~h^{-1}$ Mpc. Since there are 1122 bursts in the catalog, an effective limit on the {\\it nearest} burst to us can be placed by convoluting the likelihood as a function of $D$ (Fig. 2) with the nearest neighbor distribution of 1122 bursts inside a sphere of radius $D$. I find that the nearest burst must be farther than $40~h^{-1}$ Mpc at the 95\\% confidence level, and farther than $10~h^{-1}$ Mpc at the 99.9\\% level. At this level of confidence, then, none of the bursts that have been observed by BATSE are in nearby galaxies. A signature from the Coma cluster or the ``Great Wall'' ($\\sim 70~h^{-1}$ Mpc) is not likely to be seen in the data at present, since only a few bursts could have originated from these distances. Indeed, a search for such a signature (\\cite{HBM96}) found no compelling evidence for anisotropy in supergalactic coordinates.  The median distance to the nearest burst is $170~h^{-1}$ Mpc. Since the brightest burst in the 3B catalog has a fluence of $7.8\\times 10^{-4}~{\\rm ergs~cm}^{-2}$ in \\g--rays, this implies that the total energy released in \\g--rays during a burst event is of order $3\\times 10^{51}~h^{-2}$ ergs.  As the number of observed \\g--ray bursts keeps increasing, the distance limit will improve. In fact, LQ showed that, with 3000 burst locations, the clustering of bursts might just be detectable and would provide compelling evidence for a cosmological origin. If it is not detected, the redshift to the ``edge'' of the bursts would be put at $z\\sim 1$ or beyond."
    },
    "9603/astro-ph9603156_arXiv.txt": {
        "abstract": "We investigate the properties of hybrid gravitational/hydrodynamical simulations, examining both the numerics and the general physical properties of gravitationally driven, hierarchical collapse in a mixed baryonic/dark matter fluid.  We demonstrate that, under certain restrictions, such simulations converge with increasing resolution to a consistent solution.  The dark matter achieves convergence provided that the relevant scales dominating nonlinear collapse are resolved.  If the gas has a minimum temperature (as expected, for example, when intergalactic gas is heated by photoionization due to the ultraviolet background) {\\em and} the corresponding Jeans mass is resolved, then the baryons also converge. However, if there is no minimum baryonic collapse mass or if this scale is not resolved, then the baryon results err in a systematic fashion.  In such a case, as resolution is increased the baryon distribution tends toward a higher density, more tightly bound state.  We attribute this to the fact that under hierarchical structure formation on all scales there is always an earlier generation of smaller scale collapses, causing shocks which irreversibly alter the state of the baryon gas.  In a simulation with finite resolution we therefore always miss such earlier generation collapses, unless a physical scale is introduced below which such structure formation is suppressed in the baryons.  We also find that the baryon/dark matter ratio follows a characteristic pattern, such that collapsed structures possess a baryon enriched core (enriched by factors $\\sim 2$ or more over the universal average) which is embedded within a dark matter halo, even without accounting for radiative cooling of the gas.  The dark matter is unaffected by changing the baryon distribution (at least in the dark matter dominated case investigated here), allowing hydrodynamics to alter the distribution of visible material in the universe from that of the underlying mass. ",
        "introduction": "Hydrodynamics is thought to play a key role in the formation of the visible structures in the universe, such as bright galaxies and hot intracluster gas.  For this reason there is a great deal of interest in incorporating hydrodynamical effects into cosmological structure formation simulations in order to make direct, quantitative comparisons of such simulations to observed data.  In addition to gravitation, a cosmological hydrodynamical simulation must minimally account for pressure support, shock physics, and radiative cooling, as these are the fundamental physical processes thought to play a dominant role in the formation of large, bright galaxies (White \\& Rees 1978).  There is already a bewildering array of such studies published, including Cen \\& Ostriker (1992a,b), Katz, Hernquist, \\& Weinberg (1992), Evrard, Summers, \\& Davis (1994), Navarro \\& White (1994), and Steinmetz \\& M\\\"{u}ller (1994), to name merely a few.  In order to appreciate the implications of such ambitious studies, it is important that we fully understand both the physical effects of hydrodynamics under a cosmological framework and the numerical aspects of the tools used for such investigations.  Basic questions such as how the baryon to dark matter ratio varies in differing structures (galaxies, clusters, and filaments) and exactly how this is affected by physical processes such as shock heating, pressure support, or radiative cooling remain unclear.  It is also difficult to separate real physical effects from numerical artifacts, particularly given the current limitations on the resolution which can be achieved.  For example, in a recent study of X-ray clusters Anninos \\& Norman (1996) find the observable characteristics of a simulated cluster to be quite resolution dependent, with the integrated X-ray luminosity varying as $L_x \\propto (\\Delta x)^{-1.17}$, core radius $r_c \\propto (\\Delta x)^{0.6}$, and emission weighted temperature $T_X \\propto (\\Delta x)^{0.35}$ (where $\\Delta x$ is the gridcell size of the simulation).  In a study of the effects of photoionization on galaxy formation, Weinberg, Hernquist, \\& Katz (1996) find that the complex interaction of numerical effects (such as resolution) with microphysical effects (such as radiative cooling and photoionization heating) strongly influences their resulting model galaxy population.  In this paper we focus on separating physical from numerical effects in a series of idealized cosmological hydrodynamical simulations.  This study is intended to be an exploratory survey of hydrodynamical cosmology, similar in spirit to the purely gravitational studies of Melott \\& Shandarin (1990), Beacom \\etal\\ (1991), and Little, Weinberg, \\& Park (1991).  We will examine the effects of pressure support and shock heating in a mixed baryonic/dark matter fluid undergoing gravitationally driven hierarchical collapse.  This problem is approached with two broad questions in mind: how stable and reliable is the numerical representation of the system, and what can we learn about the physics of such collapses?  These questions have been investigated for purely gravitational systems in studies such as those mentioned above.  In those studies numerically it is found that the distribution of collisionless matter converges to consistent states so long as the nonlinear collapse scale is resolved.  Such convergence has not been demonstrated for collisional systems, however.  It is not clear that hydrodynamical simulations will demonstrate such convergence in general, nor if they do that the nonlinear scale is the crucial scale which must be resolved.  Hydrodynamical processes are dominated by localized interactions on small scales, allowing the smallest scales to substantially affect the state of the baryonic gas.  As an example, consider a collisional fluid undergoing collapse.  Presumably such a system will undergo shocking near the point of maximal collapse, allowing a large fraction of the kinetic energy of the gas to be converted to thermal energy.  In a simple case such as a single plane-wave perturbation (the Zel'dovich pancake collapse), the obvious scale which must be resolved is the scale of the shock which forms around the caustic.  However, in a hierarchical structure formation scenario there is a hierarchy of collapse scales, and for any given resolution limit there is always a smaller scale which will undergo nonlinear collapse.  The subsequent evolution of the gas could well depend upon how well such small scale interactions are resolved, and changes in the density and temperature of gas on small scales could in turn influence how it behaves on larger scales (especially if cooling is important).  In this paper we examine a series of idealized experiments, evolving a mixed fluid of baryons and collisionless dark matter (dark matter dominated by mass), coupled gravitationally in a flat, Einstein-de Sitter cosmology. The mass is seeded with Gaussian distributed initial density perturbations with a power-law initial power spectrum.  We perform a number of simulations, varying the resolution, the initial cutoff in the density perturbation spectrum, and the minimum allowed temperature for the baryons. Enforcing a minimum temperature for the baryons implies there will be a minimal level of pressure support, and therefore a minimum collapse scale (the Jeans mass), below which the baryons are pressure supported against collapse.  From the numerical point of view, performing a number of simulations with identical initial physical conditions but varying resolution allows us to unambiguously identify resolution effects.  By enforcing a Jeans mass for the baryons we introduce an intrinsic mass scale to the problem, which may or may not be resolved in any individual experiment.  The hope is that even if the gas dynamical results do not converge with increasing resolution in the most general case, the system will converge if the fundamental Jeans mass is resolved.  The effects of the presence (or absence) of a baryonic Jeans mass also raises interesting physical questions.  Though we simply impose arbitrary minima for the baryon temperatures here, processes such as photoionization enforce minimum temperatures in the real universe by injecting thermal energy into intergalactic gas.  The Gunn-Peterson test indicates that the intergalactic medium is highly ionized (and therefore at temperatures $T \\gtrsim 10^4$K) out to at least $z \\lesssim 5$.  Shapiro, Giroux, \\& Babul (1994) discuss these issues for the intergalactic medium.  The dark matter, however, is not directly influenced by this minimal pressure support in the baryons, and therefore is capable of collapsing on arbitrarily small scales.  Pressure support provides a mechanism to separate the two species, and since the dark matter dominates the mass density it can create substantial gravitational perturbations on scales below the Jeans mass. While there are many studies of specific cosmological models with detailed microphysical assumptions, the general problem of the evolution of pressure supported baryons in the presence of nonlinear dark matter starting from Gaussian initial conditions has not been investigated in a systematic fashion.  This paper is organized as follows.  In \\S \\ref{Sim.sec} we discuss the particulars of how the simulations are constructed and performed.  In \\S \\ref{Numresults.sec} we characterize the numerical effects we find in these simulations, and in \\S \\ref{Physresults.sec} we discuss our findings about the physics of this problem.  Finally, \\S \\ref{disc.sec} summarizes the major results of this investigation. ",
        "conclusions": "\\label{disc.sec} The results of this investigation can be broken into two broad categories: what is revealed about the physics of hierarchical collapse in a mixed baryonic/dark matter fluid, and what is learned about the numerics of simulations of this process.  We find that the dark matter converges to a consistent state on resolved scales, so long as the nonlinear collapse scales are well resolved.  Increasing the resolution of the experiment does not fundamentally alter the dark matter distribution, but simply yields more detailed information about the small scale collapsed structures.  This is in agreement with previous, purely collisionless studies, though we demonstrate this here including a collisional component.  The numerical story is quite different for the collisional baryonic gas. We find that in the case where we do not impose a fundamental physical resolution scale in the baryons, the simulation results do not converge with increasing resolution.  Rather, as the numerical resolution of the experiment is increased, the collapsed fraction of the baryons is systematically altered toward a higher density, more tightly bound, and less strongly shocked state.  The physical reason for this behaviour is the presence of shocks, which allows the evolution on small scales to affect the overall state of the baryonic mass.  With improving resolution the simulation is able to resolve the collapse of smaller structures at earlier times.  The smaller scale (and therefore earlier) the resolved collapse, the weaker the resulting shock is found to be.  This effect is most obvious in Figure \\ref{RhoTDist.fig}, where there is a systematic trend of higher density/more weakly shocked material with increasing resolution.  The fact that the dark matter converges in general with resolution, while the baryons do not, highlights a fundamental difference in the physics of these two species.  While both dark matter and baryon fluids react to the global and local gravitational potential, the baryons are additionally subject to purely local hydrodynamical phenomena -- most prominently shocking in this case.  Once strong shocking sets in these hydrodynamical effects can rise to rival the gravitational force on the baryonic fluid, allowing the baryons to be strongly influenced by interactions on small scales in ways which the dark matter is not.  This implies that such small scale interactions can be just as important as the large scale forces in determining the final state of the baryons.  In other words, for the dark matter there is no back reaction from small to large scales, whereas the baryons are strongly influenced by interactions on small scales.  In the coupling of these physical processes, gravitation dominates the large scale structure, but hydrodynamics affects the local arrangements and characteristics of the baryonic gas.  If we want the quantitative results of such studies to be reliable, we must have reason to believe that the localized hydrodynamical processes are adequately resolved.  This gloomy picture is alleviated by an important physical effect: the Jeans mass.  Introducing a minimum temperature (and therefore pressure support) into the baryons creates a fundamental length/mass scale, below which the baryons are supported by pressure against any further collapse or structure formation.  We find that once we introduce such a minimal scale into the baryonic component, the simulation results converge as this scale is resolved.  This convergence holds even though the dark matter component continues to form structures below the baryon Jeans scale.  Although the Jeans scale is dependent upon the local density, we find that the global Jeans scale defined using the background density is adequate to define the critical resolution necessary for the hydrodynamics to converge.  This therefore describes an additional resolution scale necessary for hydrodynamical simulations to meet, much as the nonlinear mass scale represents the crucial resolution necessary for purely gravitational systems.  Furthermore, our experiments indicate that equation (\\ref{MR.eq}) is a reasonable estimate of an SPH simulation's true mass resolution, since we find that the threshold $M_R \\lesssim M_J$ marks the point at which convergence is achieved.  In these experiments we have tested the effects of the Jeans mass in an idealized framework by simply imposing an arbitrary minimum temperature into our system, but there is reason to believe that such minimum temperatures should exist in the real universe.  Based upon observations such as the Gunn-Peterson test (Gunn \\& Peterson 1965), it is known that the IGM is highly ionized out to redshifts $z \\lesssim 5$, which implies a minimum temperature for the IGM of at least $T \\gtrsim 10^4$.  Assuming an Einstein-de Sitter cosmology, a minimum temperature of $T \\sim 10^4$ requires a minimum spatial resolution (via eq. [\\ref{LJ.eq}]) \\beq \\lambda_J \\sim 0.777 \\; (1 + z)^{-3/2} \\; \\lp \\frac{\\mu}{0.6} \\rp^{-1/2} \\; \\lp \\frac{T}{10^4 \\mbox{K}} \\rp^{1/2} \\; h^{-1} \\mbox{\\ Mpc}, \\eeq which equates to a baryon mass resolution of (eq. [\\ref{MJ.eq}]) \\beq M_R \\lesssim \\Sub{\\Omega}{bary} \\; M_J \\sim 6.82 \\times 10^{10} \\; \\Sub{\\Omega}{bary} \\; (1 + z)^{-3/2} \\; \\lp \\frac{\\mu}{0.6} \\rp^{-3/2} \\; \\lp \\frac{T}{10^4 \\mbox{K}} \\rp^{3/2} \\; h^{-1} \\ M_{\\sun}. \\eeq This limit can also be expressed in terms of a minimum circular velocity, which has the advantage of being independent of redshift.  The minimum circular velocity can found as a function of the minimum temperature by relating the kinetic energy necessary for dynamical support to the internal energy for equivalent pressure support (Thoul \\& Weinberg 1996), yielding \\beq \\Sub{v}{circ} = \\lp \\frac{2 k T}{\\mu m_p} \\rp^{1/2} \\sim 16.6 \\lp \\frac{\\mu}{0.6} \\rp^{-1/2} \\lp \\frac{T}{10^4 \\mbox{K}} \\rp^{1/2} \\mbox{km/sec}. \\eeq In our $M_J>0$ simulations if we choose to call the scale at which RMS mass fluctuation is $\\Delta M/M \\sim 0.5$ to be 8 $h^{-1}$ Mpc at the final expansion, then our box scale is $L=64 h^{-1}$Mpc and the minimum temperature corresponds to $\\Sub{T}{min} \\sim 10^6$K.  While there are some suggestions that the intergalactic medium could be heated to temperatures as hot as $10^6$K (through mechanisms such as large scale shocks of the IGM), clearly these simulations do not meet our criteria if we wish to consider photoionization as setting the minimum temperature.  It is also not clear that the current generation of large-scale hydrodynamical cosmological simulations meet this criterion, but it should be achievable.  It is still unclear whether or not in the case with no minimum temperature imposed the baryon distribution will eventually converge.  It is well known that in a purely gravitational system, as structure builds and smaller dark matter groups merge into larger structures, the dark matter ``forgets'' about the earlier small scale collapses as such small structures are incorporated into larger halos and disrupted.  This is why the dark matter results converge once the nonlinear mass scale is resolved.  While it is evident from studies such as this that the baryons maintain a longer memory of their previous encounters, it seems likely that as the baryon gas is progressively processed through larger scale and stronger shocks, at some point the previous evolution should become unimportant.  At exactly what level this transition is reached remains uncertain, however, as we see no evidence for such convergence here.  Radiative cooling must be accounted for in order to model processes such as galaxy formation, and the inclusion of radiative cooling can only exacerbate the non-convergence problems we find here.  The amount of energy per unit mass dissipated by radiative cooling is proportional to the density, and we have already noted that the trend with finite resolution is to underestimate the local gas density and overestimate the temperature. Given these tendencies, it is not difficult to envision problems for finite resolution simulations which will tend to underestimate the effectiveness of radiative cooling in lowering the temperature (and therefore pressure support) of the shocked gas.  This could lead to perhaps drastic underestimates of the fraction of cold, collapsed baryons for a given system, and therefore strongly influence the inferred galaxy formation. Evrard \\etal\\ (1994) note this effect when comparing their high and low resolution 3-D simulations.  They find that altering their linear resolution by a factor of two (and therefore the mass resolution by a factor of eight) changes the measured total amount of cold collapsed baryons by a factor of $\\sim 3$.  They attribute this change to just the sort of problems we discuss here.  Weinberg, Hernquist, \\& Katz (1996) report similar findings and interpretation for simulations with a photoionizing background.  It therefore seems likely it is all the more important to resolve the minimum mass scale set by the minimum temperature in systems with radiative cooling.  We find that the majority of the baryonic mass undergoes strong shocking so long as the nonlinear mass scale exceeds the Jeans mass.  At infinite resolution in the $M_J=0$ case, it is possible that all of the baryonic material undergoes shocking.  As anticipated from previous investigations, the highest density collapsed fraction is characteristically less shocked as compared with later infalling material from larger regions.  The underlying cause for this behavior is the fact that potential wells deepen as structure grows.  The highest density material is that which collapses earliest due to the smallest scale perturbations.  This material falls into relatively shallow potential wells, and is only weakly shocked.  As the structures continue to grow, progressively larger scales go nonlinear and collapse.  The potential wells deepen and infalling material gains more energy, resulting in stronger shocking and higher temperatures.  Hydrodynamics can also play an important role in determining the distribution of the baryon mass, particularly in collapsed structures.  In the absence of external mechanisms to heat the baryons (such as energy input from photoionization), during the linear phase of structure growth the baryons evolve as a pressureless fluid and simply follow the dominant dark matter.  Once nonlinear collapse sets in, the baryons fall to the potential minimum, shock, convert their kinetic energy to thermal energy, and settle.  In contrast, the dark matter simply passes though the potential minimum and creates a more diffuse structure supported by the anisotropic pressure of random velocities.  This difference gives rise to a characteristic pattern in the baryon/dark matter ratio.  Wherever the evolution is still linear, the baryons and dark matter simply remain at the universal mix.  With the onset of nonlinear collapse, the baryons fall to the minimum of the potential well where they form a baryon enriched core, surrounded by a dark matter rich halo.  We find that even in the absence of radiative cooling the cores of collapsed structures can become baryon enriched by factors of $\\Sub{n}{bary}/\\Sub{n}{dm} \\sim 2$ or more, though this value is likely resolution and dimension dependent.  If the thermal energy of the baryons is raised to the point that it rivals the potential energy during the collapse, the baryons will become pressure supported and stop collapsing at that point.  In all cases we find that the dark matter is relatively unaffected by the baryon distribution.  This is due to the fact that the dark matter dominates the mass density, and therefore the gravitational potential.  In general it appears that under a dark matter dominated scenario hydrodynamics can substantially alter the characteristics of the baryonic material (and therefore the visible universe), such that it does not directly follow the true mass distribution which is dominated by the the dark matter."
    },
    "9603/astro-ph9603138_arXiv.txt": {
        "abstract": "Neutrino-nucleus elastic scattering is reduced in dense matter because of correlations between ions. The static structure factor for a plasma of electrons and ions is calculated from Monte Carlo simulations and parameterized with a least squares fit.   Our results imply a large increase in the neutrino mean free path.  This strongly limits the trapping of neutrinos in a supernova by coherent neutral current interactions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603142_arXiv.txt": {
        "abstract": "A collection of modern, field-theoretical equations of state is applied to the investigation of cooling properties of compact stars. These comprise neutron stars as well as hypothetical strange matter stars, made up of absolutely stable 3-flavor strange quark matter. Various uncertainties in the behavior of matter at supernuclear densities, e.g., hyperonic degrees of freedom, behavior of coupling strengths in matter, pion and meson condensation, superfluidity, transition to quark matter, absolute stability of strange quark matter, and last but not least the many-body technique itself are tested against the body of observed cooling data. ",
        "introduction": "A forefront area of research, both experimental and theoretical, concerns the exploration of the properties of matter under extreme conditions of temperature and/or density and the determination of the equation of state (pressure versus density) associated with it. Its knowledge is of key importance for our understanding of the physics of the early universe, its evolution to the present day, compact stars, various astrophysical phenomena, and laboratory physics (for an overview, see, for example, \\cite{Greiner94a}).  On the earth, relativistic heavy-ion colliders provide the only tool by means of which such matter can be created and its properties studied.  On the other hand, however, it is well known that nature has created a large number of massive stellar objects, i.e., white dwarfs and neutron stars, which contain matter in one of the densest forms found in the universe.  Neutron stars are associated with two classes of astrophysical objects -- pulsars and compact X-ray sources.  Matter in their cores possess densities ranging from a few times the density of normal nuclear matter to about an order of magnitude higher, depending on star mass.  To the present day, about 600 pulsars are known, and the discovery rate of new ones is rather high. This is accompanied by an impressive growth rate of the body of observed pulsar data, like pulsar temperatures determined by the X-ray observatories Einstein, EXOSAT, and ROSAT \\cite{Truemper92a,Truemper93a,Oegelman95a}.  In this paper, we shall apply a broad collection of modern, field-theoretical equations of state (EOS) to the study of the cooling behavior of both neutron stars and their strange counterparts -- strange matter stars -- which should exist if 3-flavor strange quark matter is more stable than confined hadronic matter. This collection of EOSs was derived under numerous model assumptions about the behavior of superdense stellar matter.  To mention several are: the many-body technique used to determine the equation of state; the model for the nucleon-nucleon interaction; description of electrically charge neutral neutron star matter in terms of either only neutrons, neutrons and protons in generalized chemical equilibrium ($\\beta$ equilibrium) with electrons and muons, or nucleons, hyperons and more massive baryon states in $\\beta$ equilibrium with leptons; behavior of the hyperon coupling strengths in matter, inclusion of meson ($\\pi$, $K$) condensation; treatment of the transition of confined hadronic matter into quark matter; and assumptions about the true ground state of strongly interacting matter (i.e., absolute stability of strange quark matter relative to baryon matter).  The paper is organized as follows.  In section \\ref{sec:structure} we introduce the set of equations that govern the cooling behavior of massive stars. The collection of equations of state for neutron stars is discussed in section \\ref{sec:models}.  The physics of strange stars and their associated EOS is explained in section \\ref{sec:strangestars}. The phenomenon of superfluidity and the various neutrino emission processes are outlined in sections \\ref{sec:superfl} and \\ref{sec:neutrino}. Our results and conclusions are presented in sections \\ref{sec:res} and \\ref{sec:con}, respectively. We summarize the paper in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:con}  We have already pointed out in section \\ref{sec:observations} that the observed cooling data of pulsars marked with squares in our figures appear to be the most significant ones. Primarily these data should be reproduced by theoretical cooling calculations. As we have seen above (section \\ref{sec:standard}) the luminosities of a few pulsars, like 1951+32 and 1055-52, appear to be compatible with the standard cooling scenarios.  Depending on the model employed for the equation of state, this may also be the case for Geminga and pulsars 0656+14, 2334+61 (cf. fig. \\ref{fig:standard}). Others, however, like the famous Vela pulsar (PSR 0833-45) and PSR's 1706-44, 1929+10, which have rather low observed luminosities, may require the introduction of so-called fast (or enhanced) cooling processes, which are characterized by higher neutrino emission rates by means of which stars cool more rapidly. A final decision on this issue -- standard versus enhanced cooling \\cite{Oegelman95a} -- is presently hampered by the uncertainties in a number of quantities, including the modified Urca neutrino-emission rate, superfluid gaps, equation of state, and mass and age of the star. Nevertheless, subject to the enormous progress that is being made presently in exploring the behavior of superdense matter \\cite{Greiner94a}, one may feel confident that this will change in the foreseeable future.  Figure \\ref{fig:enhanced.nosf} compares the influence of different \\begin{figure}[tbp] \\centering \\psfig{figure=p0995a_g300_enh.ps,width=7cm} \\capt{Influence of several different enhanced neutrino-emission processes on the cooling of neutron stars with masses $M=1.4\\,M_\\odot$, except for the kaon-condensed model where the mass is $M=1.78\\,M_\\odot$ (The $1.4\\,M_\\odot$ models do not possess large enough central densities to overcome the threshold density for kaon condensation).  The observed data are labeled in fig. \\protect{\\ref{fig:standard}}.} \\label{fig:enhanced.nosf} \\end{figure} enhanced neutrino emission processes, i.e., direct Urca, pion and kanon condensation, and quark Urca, on the cooling behavior of a neutron star with mass $M=1.4\\,M_\\odot$. The nucleons and quarks are treated as non-superflid particles. One sees that the inclusion of any of these enhanced cooling processes reduces the star temperature too quickly in order to get agreement with the observed data points.  The only exception is PSR 1929+10. The most significant data points (squares), however, lie considerably above the enhanced cooling curves.  The enhanced cooling scenarios can be slowed down if one assumes that the neutrons in the cores of neutron stars are superfluid, as was the case in figs. \\ref{fig:dr.uvu}--\\ref{fig:k240b180}. The only problem left then concerns the horizontal ``plateaus'' in these figures at star ages between $10^2$ and $10^5$ years which tend to lie somewhat too high for stars with canonical mass, $M=1.4\\,M_\\odot$, to be consistent with the data. Note, however, that one gets agreement with some of the data points if the star mass is assumed to be different from the canonical value.  Another possibility to get agreement consists in varying the gap energy, $\\Delta_\\mathrm{sf}$, which, as mentioned in section \\ref{sec:superfl}, is not very accurately know. This is particularly the case for the high-density \\sfp~ superfluid. A reduction of this gap by a factor of two, for instance, moves the cooling curves up right into the region where the oberved data is concentrated, as can be seen in fig. \\ref{fig:redsf}. \\begin{figure}[h] \\centering \\psfig{figure=p0995a_g300_enhsf.ps,width=7cm} \\capt{Influence of changes in the superfluid \\sfp~ gap on the cooling of neutron stars constructed for \\glen.  The enhanced cooling processes are direct Urca in pion (solid curves) and kaon condensed (dotted curves) matter.  The observed data are labeled in fig. \\protect{\\ref{fig:standard}}.} \\label{fig:redsf} \\end{figure}  Of course there are other possibilities, besides reducing the gap energy, by means of which agreement with the observed data may be achieved. For example, the superfluid phase of neutrons might not reach the center of very massive neutron stars, as is the case for stars constructed for \\UVU~ (see fig. \\ref{fig:dr.uvu}). This is due to the smaller proton fraction and the resulting higher Fermi momenta of neutrons.  Another possibility is an additional cooling process not suppressed by superfluidity, as, for example, the superfluid pair-breaking process \\cite{Voskresenskii86,Schaab95b}, or so-called internal heating of neutron stars \\cite{Shibazaki89:a,Umeda93:a,Sedrakian93:a,Reisenegger95:a,VanRiper95:a}."
    },
    "9603/astro-ph9603004_arXiv.txt": {
        "abstract": "We discuss the application of a class of machine learning algorithms known as decision trees to the process of galactic classification.  In particular, we explore the application of oblique decision trees induced with different impurity measures to the problem of classifying galactic morphology data provided by Storrie-Lombardi et al. \\shortcite{st}.  Our results are compared to those obtained by a neural network classifier created by Storrie-Lombardi et al, and we show that the two methodologies are comparable.  We conclude with a demonstration that the original data can be easily classified into less well-defined categories. ",
        "introduction": "Decision tree algorithms have proven themselves useful as automated classifiers in a number of astronomical domains.  An oblique decision tree algorithm created by Murthy, Kasif, \\& Salzberg (1994) has demonstrated that decision trees can be generated to distinguish between stars and galaxies or to identify cosmic rays in Space Telescope images.  Typically, the trees produced by this algorithm possess accuracies up to and exceeding 95\\%, and can now be used to classify additional data \\cite{sa2}.  In this way, decision trees free researchers from the tedious task of object classification.  The next logical step for a decision tree classification algorithm is, of course, the classification of galaxies morphologically.  As Storrie-Lombardi et al. point out \\shortcite{st}, this has been attempted by a number of researchers with limited success; today ``morphological classification into ellipticals, lenticulars, spirals and irregulars remains a process dependent on the eyes of a handful of dedicated individuals.''  In an attempt to rectify this situation, Storrie-Lombardi et al. (hereafter referred to as SLSS) have applied a computing technique known as artificial neural networks (also known as neural nets or ANNs), to the problem of galaxy classification.  In particular, SLSS used their neural net to classify galaxies taken from the ESO-LV catalog \\cite{lau}.  In this paper, we compare results from the SLSS ANN with those from decision trees. ",
        "conclusions": "We have shown that decision trees can be used to determine the morphological classification of galaxies with reasonable success. Furthermore, our comparisons of a neural net classifier to that of a decision tree algorithm have produced similar results.  Errors made by both classifiers can be attributed to one or more of the following problems: 1) there are errors in the visual classification of the 5217 galaxies which comprise both the training and test sets, or 2) none of the attributes used to describe the data provides sufficient information for accurate classification.  SLSS account for this by noting that the classifications are based on plate material rather than CCD frames, and that the parameters used to describe the galaxies were chosen somewhat arbitrarily.  As can be seen in both the ANN results and those obtained by our decision trees (tables 1, 4, and 5), while non-neighbor classes can, potentially, be easily separated, neighbor classes cannot.  In other words, while trees grown to discriminate E-type galaxies from Sa+Sb-types might typically be very accurate, trees that distinguish between neighboring types such as E and S0 would have very poor accuracies.  In fact, as SLSS noticed, scoring accuracy in terms of nearest neighbor classifications results in roughly a 90\\% accuracy.  Table 6 demonstrates that multiple decision trees can, in fact, be generated to easily distinguish between different regions along the continuum of classifications.  The decision trees used to produced these results were trained on the small 1700 object set used above.  Clearly, extremely accurate and simple trees can be induced from this simplified data.  With these trees, galaxies can now be confidently classified to larger, overlapping regions.  For example, by using a majority vote among all six trees, a galaxy might be classified either to the E-S0 region or to the S0-Sa+Sb-Sc+Sd region. In one experiment, the trees in table 6 were manually assembled to produce an E-Sa+Sb-Irr classifier with a 90.7\\% accuracy.  \\begin{table} \\caption{Accuracy and tree size (in number of leaves) of non-neighbor decision trees.} \\begin{tabular}{lccclcc} {\\bf Tree}\t&Acc\t&Lvs\t&&{\\bf Tree}\t&Acc &Lvs\t\\\\ 1. E / Sa+Sb\t&96.4\t&3\t&&4. S0 / Sc+Sd\t&91.4   &2      \\\\ 2. E / Sc+Sd\t&97.3\t&2\t&&5. S0 / Irr\t&95.7   &2      \\\\ 3. E / Irr\t&95.7\t&2\t&&6. Sa+Sb / Irr&92.8\t&3\t\\\\ \\end{tabular} \\end{table}  In one last attempt to overcome the five-classification ``fuzziness'' of the data, we tried growing two new trees: one to separate E and Irr-type galaxies from spirals, and another to identify the S0, Sa+Sb, and Sc+Sd-types in the spiral subset.  The result was a modest increase in accuracy to 66\\%.  While this result could most likely have been achieved by increasing the number of random searches performed by OC1, by initially filtering out the spirals, we were able to direct the search in a direction we wanted to explore.  By reducing the search space in this manner, we also reduced processing time.  Finally, even though neither the decision tree nor the ANN produced remarkable results, global classifications by the two differ for less than 3\\% of the test set.  (Attempts have not yet been made to determine accuracy by comparing classifications example by example.)  Furthermore, misclassifications made by the oblique decision tree in table 5 match roughly 83\\% of the misclassifications made by the SLSS ANN.  The similar results obtained by these two very different classifiers does, perhaps, point to the existence of error in the original data.  At the very least, our results confirm the discovery made by the SLSS ANN that the distinction between neighboring classes appears to be poorly defined.  The two classification algorithms may ultimately be discovering a more accurate way by which to classify the original Lauberts \\& Valentijn data."
    },
    "9603/astro-ph9603096_arXiv.txt": {
        "abstract": "A key argument in favor of orientation based unification schemes is the finding that among the most powerful 3CRR radio sources the (apparent) median linear size of quasars is smaller than that of radio galaxies, which supports the idea that quasars are a subset of radio galaxies, distinguished by being viewed at smaller angles to the line of sight. Recent measurements of radio sizes for a few other low frequency samples are, however, not in accord with this trend, leading to the claim that orientation may not be the main difference between radio galaxies and quasars.  We point out that this ``inconsistency'' can be removed by making allowance for the temporal evolution of sources in both size and luminosity, as inferred from independent observations. This approach can also readily explain the other claimed ``major discrepancy'' with the unified scheme, namely, the difference between the radio luminosity--size correlations for quasars and radio galaxies. ",
        "introduction": "According to a  widely discussed unified scheme for powerful F-R II extragalactic radio sources (with luminosities above $\\sim 10^{25.5}$ W Hz$^{-1}$ at 1 GHz, where we take $H_0 = 50$ km s$^{-1}$Mpc$^{-1}$ and $q_0 = 0.5$ throughout this {\\it Letter}), narrow-line radio galaxies (RGs) are identified as quasars (QSRs), or even blazars, whenever their principal axis happens to be oriented within a certain critical angle, ${\\psi}$, from the line-of-sight (for recent reviews, see Antonucci 1993; Urry \\& Padovani 1995; Gopal-Krishna 1995, 1996). In this model, the parsec-scale nuclear core of such sources, consisting of a compact central engine ejecting relativistic jets of non-thermal continuum emission, and a broad line region (BLR), is believed to be surrounded by a dusty torus. In the case of RGs, the torus is believed to obscure the core in the visible through soft X-ray bands, so  that the BLR is not directly visible. Strong evidence for such tori comes from the detection of the BLR in the (scattered) polarized light (e.g., Antonucci \\& Miller 1985; Cimatti et al.\\ 1993; Draper, Scarrott \\& Tadhunter 1993). Various evidences for the relativistic beaming aspect of this picture include: extremely rapid variability of blazars over all bands; apparent superluminal motion in the bright radio cores; larger ratios of core to total radio emission, $f_c$, in QSRs than in RGs; correlation of $f_c$ with the polarized optical nuclear continuum and its anti-correlations with apparent radio linear size, the symmetry of core-to-lobe separations, and the equivalent width of the [O II]$\\lambda$3727 line (Urry \\& Padovani 1995; Gopal-Krishna 1995, and references therein). The unified scheme is further supported by the analysis of the radio luminosity functions of the postulated parent and aligned populations (Padovani \\& Urry 1992).   One of the cornerstones of the orientation-based paradigm for powerful radio sources has been the observation that in the low-frequency 3CRR sample (Laing, Riley, \\& Longair 1983), where the axes of the sources should be randomly oriented, the median linear extent, $\\ell$, of the extended radio emission from QSRs {\\it appears} significantly smaller ($\\stackrel{<}{\\sim}$50\\% for redshifts $z >$ 0.5 ) than that of RGs (Barthel 1989). However, the lack of such a behavior in the same sample at $z <$ 0.5 (Kapahi 1990; Singal 1993a), bolstered by similar trends reported recently for a few other low-frequency samples a few times deeper in flux density than the 3CRR, has evoked serious doubts about the unified scheme (Singal 1995, 1996; Kapahi et al.\\ 1995). Here we show that this apparently irrefutable inconsistency with the data can be resolved quantitatively by taking into account the available strong independent evidence for temporal evolution in both the sizes and luminosities of extragalactic double radio sources. ",
        "conclusions": "We emphasize that the temporal evolution incorporated here into the orientational unified scheme is exceedingly simple and expressed in terms of just a few parameters, all of which are constrained by observations.  Since this evolutionary scenario is inspired by observations and elementary theoretical considerations, the neglect of this factor in the past attempts to verify the unified scheme was a major shortcoming. Consequently, the claimed ``mismatch'' with the radio size data (Singal 1993a, 1995, 1996; Kapahi et al.\\ 1995) should not mandate dismissal of the paramount role assigned to orientation effects in the unified scheme.  The present explanation for the decreasing RG-to-QSR size ratio, $R$, towards lower radio luminosities also provides an explanation for an {\\it equivalent} observational result according to which QSRs and RGs exhibit different $\\ell$--$P$ correlations. Several authors have pointed out that the empirically derived positive $\\ell$--$P$ dependence for RGs ($\\ell \\propto P^x$, $x \\simeq +0.3$), contrasts with practically no $\\ell$--$P$ dependence found for QSRs ($x \\simeq 0$), and have argued that this ``puzzling'' result effectively rules out the unified scheme (Chyzy \\& Zieba 1993; Singal 1993a,b, 1995, 1996; Kapahi et al.\\ 1995). However, this difference too, can be readily understood in our picture.  The values of $x$ corresponding to the three computed model curves shown in Fig.\\ 1 are  $+0.24 \\pm 0.07$ for RGs and $0.09 \\pm 0.08$ for QSRs, as deduced from the linear fits to the computed median sizes at different  radio luminosities for the three sets of input parameters; for each fit the slope $x$ is found to be greater for RGs than for QSRs. As a refinement to the present analysis, the use of full linear size distributions (rather than the median values) would be a useful step.  In summary, we conclude that the linear size measurements of radio galaxies and quasars reported up to this point do not  violate the basic tenet of the unified scheme, viz, that radio galaxies are oriented closer to the sky plane than are quasars."
    },
    "9603/astro-ph9603119_arXiv.txt": {
        "abstract": "Microlensing observations have now become a useful tool in searching for non--luminous astrophysical compact objects (brown dwarfs, faint stars, neutron stars, black holes and even planets). Originally conceived for establishing whether the halo of the Galaxy is composed of this type of objects, the ongoing searches are actually also sensitive to the dark constituents of other Galactic components (thin and thick disks, outer spheroid, bulge). We discuss here the present searches for microlensing of stars in the Magellanic Clouds and in the Galactic bulge (EROS, MACHO, OGLE and DUO collaborations). We analyse the information which can be obtained regarding the spatial distribution and motion of the lensing objects as well as about their mass function and their overall contribution to the mass of the Galaxy. We also discuss the additional signals, such as the parallax due to the motion of the Earth, the effects due to the finite source size and the lensing events involving binary objects, which can further constrain the lens properties. We describe the future prospects for these searches and the further proposed observations which could help to elucidate these issues, such as microlensing of stars in the Andromeda galaxy, satellite parallax measurements and infrared observations. ",
        "introduction": "The bending of light in a gravitational field predicted by general relativity provided one of the first verifications of Einstein theory \\cite{ed19}. The value of the deflection angle which results has now been tested at the 1\\% level by observing the change in the apparent position of stars whose light is deflected by the Sun. In 1920, Eddington \\cite{ed20} noted that the light deflection by a stellar object would lead to a secondary dimmer image of a source star on the opposite side of the deflector. Chwolson \\cite{ch24} later pointed out that these secondary images could make foreground stars appear as binaries, and that if the alignment were perfect, the image of the source would be a ring. In 1936, Einstein \\cite{ei36} published the correct formula for the magnification of the two images of a very distant star, and concluded that this lensing effect was of no practical relevance due to the unresolvably small angular separation of the images and the low probability for a high amplification event to take place.  The following year, Zwicky \\cite{zw37} showed that, if the deflecting object were a galaxy instead of a single star, the gravitational lensing  of the light of a background galaxy would lead to resolvable images. This `macrolensing' effect would provide information about the mass of the intervening galaxy and allow observation of objects at much larger distances due to the magnification of their light. Furthermore, the probability that this effect be observed is a `certainty' \\cite{zw37b}.  It was actually the multiple imaging of a high redshift quasar by a foreground galaxy which provided the first observation of gravitational lensing \\cite{wa79}, and this area of research has now become an active field in astronomy, with the potential of giving crucial information for cosmology \\cite{sc92,bl92,re94}. For instance, the time delay among the multiple images of a quasar allows one to relate the lens mass to the Hubble constant. Another interesting example is that the mass distribution of a rich cluster can be reconstructed from the shapes of the images of thousands of background faint blue galaxies around it, which become elongated due to the effect of weak lensing by the cluster \\cite{ty95}. Also, gravitational lensing can strongly affect the observable properties of AGN and quasars, and may have to be taken into account when inferring their intrinsic properties.  While these developments were going on, the theoretical study of gravitational lensing  of stars by stars restarted in 1964, with the work of Liebes \\cite{li64} and Refsdal \\cite{re64}, who extended the formalism and discussed the lensing of stars in the disk of the Galaxy, in globular clusters and in the Andromeda galaxy.  The lensing effect of individual stars belonging to a galaxy that is itself macrolensing a background source was discussed by Chang and Refsdal \\cite{ch79}. When both the lensed source and the intervening galaxy are at cosmological distances, the passage of one of these stars close to the line of sight to one of the images  further deflects the source light by an angle which is typically of ${\\cal O}(\\mu$arcsec). The name `microlensing' (ML) then became associated with this process, and is now generally applied to any gravitational lensing effect by a compact object producing unresolvable images of a source and potentially huge magnifications of its light.  Press and Gunn \\cite{pr73} showed in 1973 that a cosmological density of massive compact dark objects could  manifest through the ML of high redshift sources. In 1981, Gott \\cite{go81} pointed out the possibility of detecting the dark halos of remote galaxies by looking for ML of background quasars. It was in 1986 that Paczy\\'nski \\cite{pa86} gave a new face to the field when he noted that, by monitoring the light--curves of millions of stars in the Large Magellanic Cloud (LMC) for more than a year, it would become possible to test whether the halo of our galaxy consisted of compact objects with masses between $10^{-6}$ and $10^2\\ M_\\odot$, i.e. covering most of the range where baryonic dark matter in the form of planets, Jupiters, brown dwarfs or stellar remnants (dead stars, neutron stars or black holes) could lie. It was later realized \\cite{pa91,gr91b} that the Galactic bulge stars also provided an interesting target to look at, since at least the lensing by faint stars in the disk should grant the observation of ML events, and observations in the bulge could also allow one to test the dark constituency of the Galaxy close to the Galactic plane.  With these motivations, several groups undertook the observations towards the LMC (EROS and MACHO collaborations) and towards the bulge (MACHO and OGLE), obtaining the first harvest of ML events in 1993 \\cite{al93,au93,ud93}. As will be discussed in this review, the observations at present are already providing crucial insights into the dark matter problem, into the non--luminous contribution to the mass of the different Galactic components and they are helping to unravel the morphology of the inner part of the Galaxy. The continuation of these programs, as well as the new ones entering the scene (DUO collaboration looking towards the bulge; AGAPE and the Columbia--VATT collaborations looking at Andromeda and followup telescope networks such as PLANET, GMAN and MOA looking for fine details in the microlensing events) will certainly allow us to reach a much deeper knowledge of these fundamental issues in the near future. ",
        "conclusions": "The idea of using ML to study the dark constituents of our Galaxy has proved to be extremely useful. In the last few years, several experiments started searches for ML and many surprises have resulted.  The observations in the direction of the LMC returned rates which are significantly smaller than those expected from a standard halo composed of objects with masses in the range $10^{-7}$--$10^{-1}M_\\odot$, but anyway larger than those expected from the faint stars present in the known stellar populations, so that something new has certainly been found. If the compact objects belong to the halo, these results could imply that the halo has a large fraction of heavy objects ($m>0.1M_\\odot$), or a lot of gas, or is a mixture  with non--baryonic dark matter, or alternatively that it deviates significantly from the standard halo models due, for instance, to the disk being close to maximal and the rotation curve actually falling with distance. Another plausible explanation is that a large number of brown dwarfs and/or stellar remnants is present either in the Galactic spheroid or in a thick disk, with the halo being allowed to be completely non--baryonic. Increased statistics will clearly allow  these facts to be put on a more solid basis and to discriminate between the different proposed solutions. Continuing the observations will also allow one to gain sensitivity to longer event durations, i.e. to heavier lens masses.  Observations towards the bulge returned on the contrary more events than  were initially expected, implying that probably the bar in the inner Galaxy was rediscovered using ML. Continued searches will allow one to further constrain the geometry and mass of this Galactic component. The rates observed may also be indicative of large amounts of material in the disk, and this could also hint towards a smaller contribution of the halo to the rotation curve, so that bulge results may also be relevant for LMC searches. The large number of events observed has allowed identification of some particularly interesting cases of ML, such as events due to binary lenses and the observation of parallax  due to the Earth's motion. More events of these types, as well as proper motion measurements in the lensing of giant stars, will allow  better identification of the lensing agents. Also the mapping of the rates and event durations across the bulge are crucial for isolating the different contributions from disk and bulge lens and source populations.  Hence, microlensing has opened a new window for Galactic astronomy, allowing one to `see' dark objects which are otherwise unobservable. The  first results already had an important impact on our understanding of Galactic structure and of the composition of the dark halo, which still remains one of the fundamental open problems in physics. This field is rapidly moving, with new experiments joining the searches, so that one can expect that a deeper, and better established, knowledge of these fundamental issues, and possibly  some new surprises, will be obtained in the near future.  \\bigskip\\bigskip"
    },
    "9603/astro-ph9603025_arXiv.txt": {
        "abstract": "We present an analysis of single--site time--series photometry of the pulsating pre--white dwarf PG~0122+200.  We show the pulsations are consistent with a pattern of modes equally spaced in period; both the observed period range and spacing confirm that PG~0122 is a $g$--mode pulsator.  PG~0122 shows a pattern similar to that seen in multi--site observations of PG~2131+066 and PG~1159--035.  The measured period spacing, combined with the spectroscopic temperature, constrain the stellar mass much more precisely than the published measurement of its surface gravity.  Based on stellar models, the mass of PG~0122 falls in the range 0.66--0.72~$M_{\\odot}$.  Fine structure in the power spectrum indicates that PG~0122 rotates once every 1.6 days.  Future multi--site observation (e.g., using the Whole Earth Telescope) should increase the precision of these results and reveal detailed information on the internal structure of this variable pre--white dwarf star. ",
        "introduction": "The PG~1159 stars represent the ephemeral penultimate stage in the life of low-- and intermediate--mass stars.  Gaining knowledge of their properties provides insight into their predecessors---central stars of planetary nebulae and AGB stars---and their descendants---the white dwarfs.  The pulsations exhibited by some PG~1159 stars help us in this endeavor.  With the tools of asteroseismology in hand, the observed frequencies can be used to determine the total stellar mass and surface layer mass of the compositionally stratified pre--white dwarfs.  This then constrains their genealogy and structure.  The variable PG~1159 stars (GW Vir stars) are non--radial $g$--mode pulsators. Theory suggests---and observations show---that for such high surface gravity objects (log $g \\sim 7$), the power spectra should be rich but essentially well ordered: we expect to see patterns of modes equally spaced in period. High radial overtone $(n \\gg 1)$ modes of low spherical harmonic index $\\ell$ can show multiplet structure; rotation can split each one into $2\\ell+1$ peaks in the power spectrum.  Other effects, such as a stellar magnetic field, cause frequency splitting about the $m = 0$ mode in a given multiplet.  This splitting can be asymmetric depending on the field geometry.  Despite the potential complexity, only modes with low $\\ell$ have been identified in GW Vir stars.  Thus the power spectrum is not necessarily complex beyond comprehension.  Kawaler \\& Bradley (1994) showed that the average period spacing between modes of the same $\\ell$ and consecutive $n$ depends primarily on the total stellar mass, with lesser dependence on luminosity and only a very slight dependence on composition.  Periodic deviations from mean period spacing can reveal the existence of a composition interface (Kawaler 1988; c.f. Kawaler \\& Bradley, 1994).  Determination of the values of $n$, $\\ell$, and the azimuthal quantum number $m$ therefore reveals a wealth of information about the structure of the star (c.f. Winget et al. 1991 \\& Bradley 1994).  The task of decoding the power spectrum of a real star is complicated because not all possible modes are necessarily present.  Also, the periodic intervention of the Earth between the telescope and the star introduces aliases into the frequency spectrum which confuse mode identification.  The aliasing problem currently is addressed using the Whole Earth Telescope (WET, Nather et al. 1990). A major success of WET came in 1989, when observers at telescopes around the globe obtained two weeks of almost continuous data on the class prototype, PG~1159--035 (hereafter PG~1159).  This minimized the aliasing and allowed Winget et al. (1991) to conclusively identify pulsation modes.  They found an unbroken sequence of almost twenty multiplets---each corresponding to $\\ell=1$ and consecutive $n$---along with several $\\ell=2$ modes.  The total stellar mass, the rotation rate, and the envelope mass were determined with unprecedented precision, showing that asteroseismological models could be used with confidence to study the PG~1159 pulsation phenomenon.  Recently, Kawaler et al. (1995) analyzed a second PG~1159 star, PG~2131+066, again using WET data. PG~2131 displayed a much simpler frequency spectrum:  only a few consecutive $\\ell=1$ triplets, and no $\\ell=2$ modes, were identified.  The data from WET observations of other PG~1159 stars are currently in analysis.  The star PG~0122+200 (BB Psc, $m_{b} = 16.13$) was identified as a member of the PG~1159 spectral class by Wesemael, Green, \\& Liebert (1985).  Dreizler et al. (1995) report a surface gravity of log $g=7.5\\pm0.5$, and an effective temperature of $75,500\\pm5,000$~K, placing PG~0122 among the coolest PG~1159 stars.  Bond \\& Grauer (1987) discovered variability in PG~0122, reporting a power spectrum dominated by variations at 402.3~s and 443.7~s. Hill, Winget, \\& Nather (1987, hereafter HWN) observed the star in white light on four consecutive nights in late 1986 with the 2.1 m reflector at McDonald Observatory.  They tentatively identified eight pulsation modes between 300~s and 700~s, and suggested a mean period spacing of 16.4~s.  This implied a mass of $\\sim0.7 M_{\\odot}$, based on models developed by Kawaler (1987).  At the time HWN published their observations, rotational splitting had not been resolved in the power spectrum of any of the hot degenerate stars, so that identification of the value of $\\ell$ and $n$ for individual modes in these stars remained uncertain.  WET observations have since led to measurement of $\\ell=1$ rotational splitting for PG~1159 ($\\delta\\nu = 4.22$ $\\mu$Hz) and PG~2131 ($\\delta\\nu = 27.4$~$\\mu$Hz), implying rotation rates (see equation~[3] below) of 1.38 days and 5.07 hours respectively.  All of the high--amplitude variation in both stars is attributed to $\\ell=1$ triplets, with a series of low amplitude $\\ell=2$ modes discovered in PG~1159.  Armed with insight gained from these observations, we reanalyzed the HWN data. We hypothesized that the power spectrum of PG~0122 was comprised of $\\ell=1$ rotationally split triplets.  If PG~0122 had the complex mode structure of PG~1159, this analysis---based on single--site data---would have failed. Luckily, PG~0122 is a simple star that is remarkably similar to PG~2131 in the sparseness of its frequency spectrum.  This allowed us to test our hypothesis without multi--site observations.  This paper describes our reanalysis of the HWN data.  In the next section, we describe the observations and their reduction.  Section~3 outlines our analysis, including our efforts to separate peaks from aliases and the evidence for equal period spacing in the power spectrum.  In section~4, we use the periods to constrain the physical properties of the star, and we conclude with section~5. ",
        "conclusions": "Reanalysis of 1986 time--series photometry of the pulsating PG~1159 star PG~0122 revealed a wealth of new information about this pre--white dwarf star.  PG~0122 is a non--radial {\\it g-}mode pulsator with several $\\ell=1$ modes and one $\\ell=2$ mode present.  Fine structure in the power spectrum indicates that it rotates once every 1.6 days.  Comparison of the calculated period spacing of PG~0122 with that of stellar models implies a mass of 0.66--0.72~$M_{\\odot}$, given the spectroscopic constraint of its effective temperature.  PG~0122 is similar to PG~1159 in its pulsational structure and rotation rate and closely resembles PG~2131 in both mass and in the quantity of modes observed.  We found an insufficient number of consecutive--$n$ modes with which to analyze possible mode--trapping in PG~0122.  Future multi--site observation could provide a means to study such structure and to thereby measure the surface helium layer mass.  This measurement is crucial, since it is the surface layer thickness which primarily governs the subsequent evolution of these stars as they become white dwarfs.  However, with global parameters of three pulsating pre--white dwarfs now asteroseismologically constrained, we have begun to establish and to understand the general characteristics of this important evolutionary link between the AGB stars and the white dwarfs.  The authors wish to acknowledge an anonymous referee for many thoughtful suggestions and comments.  This work was supported in part by NSF Young Investigator Award AST-9257049 to Iowa State University (MSO'B, SDK, and BTD). Support for this work was also provided by NASA through grant number HF-01041.01-93A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.   \\clearpage  \\begin{table} \\begin{center} \\caption{Observing Log for the Archival McDonald 2.1~m Observations of PG~0122} \\vspace*{0.5in} \\begin{tabular}{cccc} \\tableline \\tableline & Start Date  & Start Time & Duration \\\\ Run Name        &   (UT)      &    (UT)    & (h:mm:ss) \\\\ \\tableline  RUN22 & 28 Nov 1986 & 01:47:00 & 5:58:25\\\\ RUN24 & 29 Nov 1986 & 01:33:50 & 2:20:40\\\\ RUN25 & 29 Nov 1986 & 03:54:43 & 3:44:50\\\\ RUN26 & 30 Nov 1986 & 01:19:00 & 2:47:40\\\\ RUN27 & 30 Nov 1986 & 04:08:10 & 3:14:50\\\\ RUN29 & 01 Dec 1986 & 02:49:30 & 5:22:30\\\\  \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\caption{Periodicities of PG 0122.  Numbers in parentheses show the mean consecutive $m$ splitting if the observed doublet is assumed to represent $m=\\pm1, \\ell=1$.  Commas separate possible $m$ identifications when constraint to a single value was not possible. Frequency and amplitude errors derive from a formal least squares analysis of the data. The amplitude units are modulation amplitude, ma = $\\Delta$I/I.} \\vspace*{0.5in} \\begin{tabular}{ccccccccc} \\tableline \\tableline Period&Frequency&$\\sigma_{f}$&Amplitude&$\\sigma_{A}$&&&$\\delta\\nu$ &$\\sigma_{\\delta\\nu}$\\\\ (s)&($\\mu$Hz)& ($\\mu$Hz)&(ma)&(ma)&$\\ell$&$m$&($\\mu$Hz)&($\\mu$Hz)\\\\ \\tableline &        &     &     &     &  &      & & \\\\ 612.4 & 1632.8 & 0.2 & 2.6 & 0.3 &? & ?    & &\\\\ &        &     &     &     &        &  &    & \\\\ 570.0 & 1754.4 & 0.2 & 2.5 & 0.3 &? & ?    & &\\\\ &        &     &     &     &        &  &    & \\\\ 466.4 & 2144.2 & 0.2 & 3.3 & 0.4 &2 & ?    && \\vspace{-.1in}\\\\ &        &     &     &     &        &      & 5.4 &0.3 \\vspace{-.1in}\\\\ 465.2 & 2149.6 & 0.2 & 2.3 & 0.4 &2 & ?    & &\\\\ &        &     &     &     &        &  &    & \\\\ 451.9 & 2213.1 & 0.2 & 4.8 & 0.4 &1 & $+1$ & &\\vspace{-.1in}\\\\ & & & & & & &7.9~($2\\times3.95$)&0.2 \\vspace{-.1in}\\\\ 450.2 & 2221.0 & 0.1 & 6.3 & 0.4 &1 & $-1$ & &\\\\ &        &     &     &     &        &  &    & \\\\ 401.6 & 2490.0 & 0.1 & 7.3 & 0.4 &1 & $+1$ & &\\vspace{-.1in}\\\\ &        &     &     &     &        &      & 3.6&0.2 \\vspace{-.1in}\\\\ 401.0 & 2493.6 & 0.2 & 3.0 & 0.4 &1 & $0$  && \\vspace{-.1in}\\\\ &        &     &     &     &        &      & 3.3&0.2 \\vspace{-.1in}\\\\ 400.5 & 2496.9 & 0.1 &12.3 & 0.4 &1 & $-1$ && \\\\ &        &     &     &     &        & &     & \\\\ 380.1 & 2631.0 & 0.2 & 1.9 & 0.4 &1 &$+1,0$&& \\vspace{-.1in}\\\\ &        &     &     &     &        &      & 3.3&0.3 \\vspace{-.1in}\\\\ 379.6 & 2634.3 & 0.2 & 2.1 & 0.4 &1 &$0,-1$& &\\\\ &        &     &     &     &        &  &    & \\\\ 337.1 & 2966.2 & 0.2 & 3.1 & 0.4 &1 & $+1$ & &\\vspace{-.1in}\\\\ & & & & & &      & 7.4~($2\\times3.70$)&0.2 \\vspace{-.1in}\\\\ 336.3 & 2973.6 & 0.1 & 5.1 & 0.4 &1 & $-1$ && \\\\  \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\begin{center} \\caption{Comparison of the period spectrum ($\\ell=1, m=0$) with a strict 21.2~s equal spacing model.  An asterisk indicates that the period represents the calculated center of a doublet splitting.  Numbers in parentheses show the effects of assuming a value of $m$ other than $m=0$ for the identified peak.} \\vspace*{0.5in} \\begin{tabular}{cccc} \\tableline \\tableline $\\Pi_{observed}$&$\\Pi_{predicted}$&$\\Delta$n&$\\Pi_{observed}-\\Pi_{predicted}$\\\\ ($m=0$)        &($\\Delta\\Pi$=21.2 s)&         &           \\\\ (s)     &  (s)            &        & (s)\\\\ \\tableline &                 &        &  \\\\ 612.4($_{-1.5}^{+1.4}$)  & 613.0   & $+10$    &  $-0.6$ \\\\ &                 &    &       \\\\ 570.0($_{-1.3}^{+1.2}$)  & 570.6   & $+8$    &    $-0.6$  \\\\ &                 &   &        \\\\ 451.0*   & 443.4   & $+2$    &  $+7.6$    \\\\ &         &                 &          \\\\ 401.0  & 401.0           &  $ 0$ &   \\\\ &         &                 &           \\\\ 379.6,380.1 & 379.8             &  $-1$   & $+0.2,-0.3$  \\\\ &         &                 &           \\\\ 336.7*  & 337.4             &  $-3$  & $-0.7$   \\\\  \\end{tabular} \\end{center} \\end{table}  \\clearpage"
    },
    "9603/hep-ph9603324_arXiv.txt": {
        "abstract": "We argue that an extension of the Minimal Supersymmetric Standard Model that gives rise to viable thermal inflation, and so does not suffer from a Polonyi/moduli problem, should contain right-handed neutrinos which acquire their masses due to the vacuum expectation value of the flaton that drives thermal inflation. This strongly disfavours ${\\rm SO}(10)$ Grand Unified Theories. The $\\mu$-term of the MSSM should also arise due to the vev of the flaton. With the extra assumption that $ m_L^2 - m_{H_u}^2 < 0 $, but of course $ m_L^2 - m_{H_u}^2 + |\\mu|^2 > 0 $, we show that a complicated Affleck-Dine type of baryogenesis employing an $LH_u$ $D$-flat direction can naturally generate the baryon asymmetry of the Universe. ",
        "introduction": "Thermal inflation \\cite{ti,Ed,old} provides the most compelling solution to the moduli (Polonyi) problem \\cite{problem,Dine,recent}. However, for a theory of the early Universe to be viable it must be capable of producing a baryon asymmetry \\cite{baryon} \\begin{equation} \\label{baryon} \\frac{n_B}{s} \\sim 3 \\times 10^{-11} \\end{equation} by the time of nucleosynthesis. Thermal inflation probably dilutes any pre-existing baryon asymmetry to negligible amounts, and the final reheat temperature after thermal inflation ($ T_{\\rm f} \\sim $ few GeV) is probably too low even for electroweak baryogenesis. Thus if thermal inflation really is the solution of the moduli problem, then it is likely also to be responsible for baryogenesis.  In Section~\\ref{ti} we explain why the flaton that gives rise to thermal inflation probably also generates the masses of right-handed neutrinos as well as the $\\mu$-term of the Minimal Supersymmetric Standard Model (MSSM). We also note the various ways in which a potential domain wall problem can be avoided. In Section~\\ref{AD} we describe how a somewhat complicated Affleck-Dine type mechanism can naturally generate the required baryon asymmetry after thermal inflation. In Section~\\ref{con} we give our conclusions. ",
        "conclusions": "\\label{con}  Right-handed neutrinos should acquire their masses due to the vacuum expectation value of the flaton that gives rise to thermal inflation, not some composite GUT operator. This will have important implications for GUT model building. In particular, $\\rm{SO}(10)$ GUT's are strongly disfavoured because the flaton would have to be in a {\\bf 126} representation which is difficult to derive from superstrings and one would have a flaton-125 splitting problem in addition to the usual doublet-triplet splitting problem.  The $\\mu$-term of the MSSM should also be generated by the vev of the flaton.  Our Affleck-Dine type mechanism generates a baryon asymmetry which is roughly estimated to be \\begin{equation} \\frac{n_B}{s} \\sim 10^{-10} \\theta \\left( \\frac{10\\,{\\rm eV}}{m_{\\nu_L}} \\right) \\left( \\frac{T_{\\rm f}}{{\\rm GeV}} \\right) \\left( \\frac{10^{11}\\,{\\rm GeV}}{M} \\right)^2 \\end{equation} where $\\theta$ is the lepton asymmetry per Affleck-Dine particle. $\\theta$ depends on the difference in phase between the soft supersymmetry breaking parameters of $W_{\\rm seesaw}$ and $W_{\\rm vev}$ (\\mbox{c.f.} Eqs.~(\\ref{wseesaw}) and~(\\ref{wvev})), as well as the detailed dynamics.  We also make the following prediction \\begin{equation} m_L^2 < m_{H_u}^2 \\end{equation} modulo renormalisation effects, where $ - m_{H_u}^2 $ is the soft supersymmetry breaking mass squared of $H_u$, and $ m_L^2 $ is the soft supersymmetry breaking mass squared of a lepton doublet.  \\subsection*{Acknowledgements} EDS thanks B. de Carlos and D. H. Lyth for helpful discussions. EDS was supported by a Royal Society Fellowship at Lancaster University during the early stages of this work and is now supported by a JSPS Fellowship at RESCEU. The work of EDS is supported by Monbusho Grant-in-Aid for JSPS Fellows No. 95209.  \\frenchspacing"
    },
    "9603/astro-ph9603041_arXiv.txt": {
        "abstract": "We present data from high-dispersion echelle spectra and simultaneous $uvby$ photometry for $\\gamma$~Doradus.  These data were obtained from several sites during 1994 November as part of the MUSICOS-94 campaign.  The star has two closely-spaced periods of about 0.75 d and is the brightest member of a new class of variable early F-type stars.  A previously suspected third period, very close to the other two, is confirmed.  Previous observations indicated that sudden changes could be expected in the spectrum, but none were found during the campaign.  The radial velocities rule out the possibility of a close companion.  The phasing between the radial velocity and light curve of the strongest periodic component rules out the starspot model.  The only viable mechanism for understanding the variability is nonradial pulsation.  We used the method of moments to identify the modes of pulsation of the three periodic components.  These appear to be sectorial retrograde modes with spherical harmonic degrees, ($\\ell, m$), as follows: $f_1$ = (3,3), $f_2$ = (1,1) and $f_4$ = (1,1).  The angle of inclination of the star is found to be $i \\approx 70^\\circ$. ",
        "introduction": "$\\gamma$ Doradus (F0V) is the brightest member of what appears to be a new class of pulsating variable stars.  In recent years, a number of late A- and early F-type dwarfs have been found to be variable with periods of the order of one day (see Krisciunas \\& Handler, 1995 for a list).  This is too long for membership of the $\\delta$~Scuti class, where the periods are just a few hours. The period is, in many cases, compatible with that expected for rotation, which suggests that the variability may be due to the rotation of a spotted star (Mantegazza, Poretti \\& Zerbi 1994).  However, most stars appear to be multiperiodic and require improbably large differential rotation.  The starspot model for $\\gamma$~Dor has been examined by Balona, Krisciunas \\& Cousins (1994b) who find that large overlapping spots would be required to give the observed light and colour amplitudes.  The most likely explanation for the variability in these stars is nonradial pulsation (NRP), since it is difficult to understand multiple periodicities in terms of the starspot model or as orbital motion.  Multiple periods are typical of NRP, but one needs to prove that the physical parameters determined from an NRP model are realistic and that the other models are unable to explain the observations. In NRP, the light variability is due to the periodic variation in temperature and, to a smaller extent, the change in the shape of the star during pulsation.  Pulsation with a period of about one day implies that the dominant restoring force is gravity ($g$-modes), rather than pressure ($p$-modes) as in the case of the $\\delta$~Sct stars.  Furthermore, because the stars show substantial light and radial velocity variations, the averaging effect over the visible hemisphere of the star must be quite small.  This implies that the spherical harmonic degree, $\\ell$, of the pulsation must be low, probably $\\ell < 4$.  Aerts \\& Krisciunas (1996) have analysed 9~Aur, a star belonging to this group, using photometry and the cross-correlation profile obtained with the CORAVEL instrument.  They deduce that 9~Aur is pulsating in two modes, both having spherical harmonic indices ($\\ell, m$) = (3, $|1|$).  The mechanism which excites the pulsation is not known at present.  The stars appear to be located at, or near, the red edge of the $\\delta$~Sct instability strip.  This indicates that the mechanism may be the same as for $\\delta$~Sct stars (the ionization of H and He), but perhaps modified by convection.  It would certainly be interesting to examine the line profile and light variations of these stars.  These data may enable us to distinguish between NRP and rotational modulation of a spotted star.  NRP produces line profile and light variations which, in principle, allow the determination of the spherical harmonic degree, $\\ell$ and the azimuthal number, $m$.  This can be done by examining the variation of the moments of a line profile and calculating a discriminant.  The value of the discriminant as a function of angle of inclination is used to determine the most probable mode ($\\ell, m$).  The method was developed by Balona (1986a, 1986b, 1987) and extended by Aerts (1993).  The two periods of $\\gamma$~Dor are stable and well determined from extensive photometric observations: $P_1 = 0.75701$ d, $P_2 = 0.73339$ d (Balona et al. 1994a).  More recently, evidence has been found for a third period at $P_4 = 0.67797$ d (Balona, Krisciunas \\& Cousins, 1994b).  (We call this period $P_4$ and not $P_3$ because $P_3$ was used in Balona et al. (1994a) to refer to a possible single period, later shown to be incorrect).  Since the star is bright ($V = 4.3$) and well placed, it was chosen as one of the targets for the MUSICOS 94 programme.  The aim was to obtain a large number of high signal-to-noise line profiles which, in conjunction with Str\\\"{o}mgren photometry, would decide which of the two models, starspots or NRP, is the correct one and to determine the spot geometry or pulsation modes.  Another factor in choosing to observe this star is the rapid change in the spectrum which appears from time to time (Balona et al. 1994a). $\\gamma$~Dor is an {\\it IRAS} source with a probable circumstellar dust cloud (Aumann 1985), which may be responsible for this effect.  In this paper we present details of the observations and reduction procedure and analyse the data in terms of the two models. ",
        "conclusions": "One of the most puzzling aspects of $\\gamma$~Dor is the sudden appearance, from time to time, of isolated absorption or emission features (Balona et al. 1994a).  Careful inspection of all the spectra obtained during the MUSICOS-94 campaign has failed to reveal such oddities.  The star appears to have a normal early-F spectrum, though broadened by rotation. There is no doubt, however, that the light variations of the star cannot be explained entirely by three periodic components.  There is a certain level of non-periodicity which is not yet understood (Balona et al. 1994a, Balona, Krisciunas \\& Cousins 1994b).  The spectroscopic observations during MUSICOS-94 show evidence for this effect on one night (JD~2449664), where the radial velocities are substantially smaller than expected.  Our new photometric observations confirm the existence of a third oscillation, $f_4 = 1.47447$ d$^{-1}$, previously suspected by Balona, Krisciunas \\& Cousins (1994b).  The existence of three frequencies implies substantial differential rotation on the basis of the starspot model.  This model also fails to account for the great differences which exist between the radial velocity to light amplitude ratios between the two main oscillations.  But the most serious objection to the starspot model is the fact that for $f_2$ radial velocity and light maximum coincide in phase.  One would expect a $90^\\circ$ phase difference for a starspot.  The small radial velocity amplitude also excludes the possibility of a binary companion.  We conclude that nonradial pulsation is the only viable explanation for the observations.  From the light curve and first four moments of the line profile of Fe~II $\\lambda$5018.450~{\\AA}, we show that $\\gamma$~Dor has a high angle of inclination, probably $i \\approx 70^\\circ$ and that the three modes may be identified as $f_1 = (3,~3)$, $f_2 = f_4 = (1,~1)$ with fair confidence. The pulsation parameters indicate quite a large value $v_f = 15.2$ km s$^{-1}$ for $f_1$.  Now the relative flux amplitude\\\\ $${\\Delta F \\over F} = {v_f \\over v \\sin i}$$\\\\ \\noindent which gives $\\Delta F/F = 0.32$ for $f_1$, implying that this mode produces a large temperature variation $\\Delta T/T \\approx 0.08$.  This may account for the significant phase dependence of the equivalent width for this mode, whereas the other modes show little or no phase dependence.  The fact that the amplitudes of the colour indices $c_1$ and $(u-b)$ are larger in this mode is also consistent with the large $v_f$.  The fact that $f_2$ dominates the radial velocity variation, but that $f_1$ does not, though they have similar light amplitudes, is due to the fact that $f_1$ has a relatively high spherical harmonic degree $\\ell = 3$, whereas $f_2$ has the lower value $\\ell = 1$.  The averaging effect in the latter case is much less, leading to a larger radial velocity amplitude.  On the other hand, the much larger value of $v_f$ for $f_1$ compensates in the light curve.  It is interesting that all three modes seem to be sectorial retrograde modes.  The frequency of pulsation in the co-rotating frame, $\\nu_0$, is related to the frequency of pulsation in the observer's frame, $\\nu$, by: $\\nu_0 = \\nu + m \\Omega$ where $\\Omega$ is the frequency of rotation. If $i = 70^\\circ$ and $R = 1.32 R_\\odot$, we find $\\Omega = 1.04$ d$^{-1}$ ($P_{\\rm rot} = 0.96$ d).  Thus, in the co-rotating frame, the three oscillations have the following frequencies: $f_{10} \\approx 4.44$, $f_{20} \\approx 2.40$ and $f_{40} \\approx 2.51$ d$^{-1}$.  $f_1$ and $f_2$ are substantially different and must have quite different radial orders in spite of the fact that they are nearly equal in the inertial frame.  On the other hand, $f_2$ and $f_4$, which are of the same degree, $\\ell$, probably have neighbouring radial orders.  This information could be used to extract seismological information, but the effect of rotation needs to be taken into account more accurately before this can done.  The success of mode identification in this star is encouraging.  The application of the method to other stars might well result in very interesting results, as in the case of 9~Aur (Aerts \\& Krisciunas 1996), but other types of pulsating star may be equally interesting. For example, an intensive campaign on a periodic Be star (and there are many of these which are bright) should result in a definitive test of the NRP/starspot models (Balona 1996).  All that would be required is an accurate determination of the radial velocity to light amplitude ratio."
    },
    "9603/astro-ph9603107_arXiv.txt": {
        "abstract": "We study the radial pulsation frequencies of slowly rotating neutron stars in general relativistic formalism using realistic equations of state. It is found that the pulsation frequencies are always an increasing function of rotation rate. The increasing rate of frequency depends on the nature of equations of state. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603055_arXiv.txt": {
        "abstract": "We observed the Seyfert 1 galaxy NGC 3516 twice during the flight of Astro-2 using the Hopkins Ultraviolet Telescope aboard the space shuttle {\\it Endeavour} in March 1995. Simultaneous X-ray observations were performed with {\\it ASCA}. Our far-ultraviolet spectra cover the spectral range 820--1840 \\AA\\ with a resolution of 2--4 \\AA. No significant variations were found between the two observations. The total spectrum shows a red continuum, $f_\\nu \\sim \\nu^{-1.89}$, with an observed flux of $\\rm 2.2 \\times 10^{-14}~erg~cm^{-2}~s^{-1}~\\AA^{-1}$ at 1450 \\AA, slightly above the historical mean. Intrinsic absorption in Lyman $\\beta$ is visible as well as absorption from O~{\\sc vi} $\\lambda\\lambda 1032,1038$, N~{\\sc v} $\\lambda\\lambda 1239,1243$, Si~{\\sc iv} $\\lambda\\lambda 1394,1403$, and C~{\\sc iv} $\\lambda\\lambda 1548,1551$. The UV absorption lines are far weaker than is usual for NGC~3516, and also lie closer to the emission line redshift rather than showing the blueshift typical of these lines when they are strong. The neutral hydrogen absorption, however, is blueshifted by $400~\\rm km~s^{-1}$ relative to the systemic velocity, and it is opaque at the Lyman limit. The sharpness of the cutoff indicates a low effective Doppler parameter, $b < \\rm 20~km~s^{-1}$. For $b = \\rm 10~km~s^{-1}$ the derived intrinsic column is $\\rm 3.5 \\times 10^{17}~cm^{-2}$. As in NGC~4151, a single warm absorber cannot produce the strong absorption visible over the wide range of observed ionization states. Matching both the UV and X-ray absorption simultaneously requires absorbers spanning a range of $10^3$ in both ionization parameter and column density. ",
        "introduction": "Only 3--10\\% of Seyfert 1 galaxies show intrinsic UV absorption in the resonance lines of highly ionized elements (\\cite{Ulrich88}). Of these, NGC~3516 has shown the strongest and most variable absorption lines (\\cite{UB83}; \\cite{Voit87}; \\cite{Walter90}; \\cite{Kolman93}; \\cite{Koratkar96}). NGC~3516 is unusual for a Seyfert 1 in other respects as well. While as many as half of all Seyfert 1's show absorption by ionized material intrinsic to the source, characterized as a ``warm absorber\" (\\cite{NP94}), the only Seyfert 1 besides NGC~3516 with equivalent neutral hydrogen columns exceeding $5 \\times 10^{22}~\\rm cm^{-2}$ is NGC~4151 (\\cite{Kolman93}; \\cite{NP94}; \\cite{Yaqoob89}; \\cite{Yaqoob93}). In NGC~4151, these variable columns range from 1 to $12 \\times 10^{22}~\\rm cm^{-2}$ (\\cite{Yaqoob89}; \\cite{Yaqoob93}), and a similar range of variation has been observed in NGC~3516 (\\cite{Kolman93}; \\cite{NP94}; \\cite{Kriss96b}). Extended X-ray emission has been seen in the nuclei of Seyfert 2 galaxies (\\cite{Wilson92}; \\cite{Weaver95}), but among Seyfert 1's the only examples are NGC~4151 (\\cite{EBH83}; \\cite{Morse95}) and possibly NGC~3516 (\\cite{Morse95}). NGC~3516 is also one of the rare Seyfert 1's with an extended narrow-line region (NLR) having a biconical morphology (\\cite{UP80}; \\cite{Pogge89}; \\cite{Miyaji92}; \\cite{Golev95}). Again, NGC~4151 is the only similar counterpart (\\cite{Evans93}). This rare combination of strong UV and X-ray absorption and extended narrow-line emission suggests that they may be related phenomena.  Extended narrow-line emission in Seyfert galaxies is commonly associated with photoionization by a collimated source of radiation. Biconical morphologies are most often found in Seyfert 2 galaxies (\\cite{Pogge89}; \\cite{Evans94}; \\cite{SK96}), and they suggest that our line of sight to the central source of radiation is obscured. In the context of unified models of Seyfert galaxies (see the review by \\cite{Antonucci93}), Seyfert 1's present us with a direct line of sight to the broad emission-line region (BELR) and continuum source, whereas our line of sight in Seyfert 2's is blocked by a torus opaque from the mid-infrared to at least soft X-rays. If the torus collimates the ionizing radiation, then biconical morphologies should not be observed in Seyfert 1 galaxies. Given this line of reasoning, Evans et al. (1993)\\markcite{Evans93} proposed that the UV-absorbing material on our line of sight in NGC~4151, but not in the torus proper, might collimate the ionizing radiation.  Although the strength of the UV and X-ray absorption in both NGC~4151 and NGC~3516 suggests that the two absorbing mechanisms are related, it is not clear how. Kolman et al.'s (1993)\\markcite{Kolman93} simultaneous X-ray and UV observations of NGC~3516 were inconclusive due to a lack of variability. Common UV and X-ray absorption at much weaker levels in some active galactic nuclei has been successfully modeled with a single warm absorber (\\cite{Mathur94}; \\cite{Mathur95}), but the wide range of ionization states of the UV absorber in NGC~4151 is not compatible with the simplest warm absorber models (\\cite{Kriss95}).  The similarities of NGC~3516 and NGC~4151 in their UV and X-ray absorption and in their biconical NLR's prompted us to explore the far-UV spectrum of NGC~3516 shortward of 1200 \\AA\\ using the Hopkins Ultraviolet Telescope (HUT). Our goal was to search for further evidence that UV and X-ray absorbing gas in active galactic nuclei (AGN) is related to the collimation mechanism for the ionizing radiation. To improve our understanding of the relationship between the UV and X-ray absorbing gas, we also performed simultaneous X-ray observations using the Japanese X-ray satellite {\\it ASCA}. In this paper we present the far-UV spectrum obtained with HUT. A companion paper (Kriss et al. 1996) discusses the {\\it ASCA} observations. ",
        "conclusions": ""
    },
    "9603/astro-ph9603132_arXiv.txt": {
        "abstract": "We use $N$-body simulations to investigate the structure and dynamical evolution of dark matter halos in clusters of galaxies. Our sample consists of nine massive halos from an Einstein-De Sitter universe with scale free power spectrum and spectral index $n = -1$. Halos are resolved by 20000 particles each, on average, and have a dynamical resolution of 20-25 kpc, as shown by extensive tests. Large scale tidal fields are included up to a scale $L=150$ Mpc using background particles. We find that the halo formation process can be characterized by the alternation of two dynamical configurations: a {\\em merging} phase and a {\\em relaxation} phase, defined by their signature on the evolution of the total mass and root mean square ({\\em rms}) velocity. Halos spend on average one third of their evolution in the merging phase and two thirds in the relaxation phase. Using this definition, we study the density profiles and show how they change during the halo dynamical history. In particular, we find that the {\\em average} density profiles of our halos are fitted by the Navarro, Frenk \\& White (1995) analytical model with an {\\em rms} residual of 17\\% between the virial radius $R_v$ and $0.01 R_v$. The Hernquist (1990) analytical density profiles fits the same halos with an {\\em rms} residual of 26\\%. The trend with mass of the scale radius of these fits is marginally consistent with that found by Cole \\& Lacey (1996): compared to their results our halos are more centrally concentrated, and the relation between scale radius and halo mass is slightly steeper. We find a moderately large scatter in this relation, due both to dynamical evolution within halos and to fluctuations in the halo population. We analyze the dynamical equilibrium of our halos using the Jeans' equation, and find that on average they are approximately in equilibrium within their virial radius. Finally, we find that the projected mass profiles of our simulated halos are in very good agreement with the profiles of three rich galaxy clusters derived from strong and weak gravitational lensing observations. ",
        "introduction": "\\label{sec:intro}  Observational studies of galaxy clusters are providing ever more data that need theoretical interpretation in order to understand cluster formation and evolution. In current cosmological models the mass and dynamics of galaxy clusters are dominated by some kind of non-baryonic, {\\em dark} matter, which interacts with ordinary baryonic matter only through gravity. In studies focussing on the dynamics of galaxy clusters, they can thus be regarded as halos made of collisionless dark matter. From the theoretical point of view one can study the structure of dark matter halos both analytically and numerically. Much of the analytical work done so far is based on the {\\em secondary infall} paradigm, (Gunn \\& Gott 1972). The simplest version of this picture considers an initial point mass, which acts as a nonlinear seed, surrounded by an homogeneous uniformly expanding universe. Matter around the seed slows down due to its gravitational attraction, and eventually falls back in concentric spherical shells with purely radial motions. Calculations based on this model predict that the density profile of the virialized halo should scale as $\\rho(r) \\propto r^{-9/4}$. Self similar solutions were found by Fillmore \\& Goldreich (1984) and by Bertschinger (1985). Hoffman \\& Shaham (1985) applied an extension of this idea to the gravitational instability theory of hierarchical clustering. In their calculations they assumed a Gaussian random field of initial density perturbation with scale-free power spectra: $P(k) \\propto k^n$. They found that the virialized structures originating from density peaks should have density profiles whose shape depends on spectral index $n$ as $\\rho(r) \\propto r^{-(9+3n)/(4+n)}$.  In reality the collapse of an initial overdensity is not so simple. In particular, motions are not purely radial, and accretion does not happen in spherical shells (as assumed in the secondary infall model), but by aggregation of subclumps of matter which have already collapsed; a large fraction of observed galaxy clusters exhibit significant substructure (Kriessler et al. 1995). It is therefore very important to complement and compare these analytical studies with numerical simulations. These are not bound by such restrictions and so can tell if the gravitational collapse of a collisionless system eliminates all memory of the cosmological parameters which determined its initial conditions. They can also show whether scale-free universes, which have no characteristic scale, give rise to scale-free power-law density profiles. Work in this direction includes Quinn, Salmon \\& Zurek (1986), Efstathiou et al. (1988) and West, Dekel \\& Oemler (1987), but the somewhat conflicting results of these studies show that better numerical resolution is needed to settle the issue. Recent results from higher resolution simulations (Navarro et al. 1995 (hereafter NFW), Lemson 1995, Cole \\& Lacey 1996 (hereafter CL), Xu 1996), produced by different $N$-body codes with different setups for the initial conditions, seem finally to agree on the following results: \\begin{enumerate} \\item halo density profiles are curved, and are well approximated by a fitting formula governed by a single {\\em scale radius} $r_s$ and belonging to the family of curves: $\\rho(r) \\propto x^{-\\alpha}(1 + x^{\\beta})^{-\\gamma}$, $x = r/r_s$. NFW propose a model with $(\\alpha,\\beta,\\gamma) = (1,1,2)$; another candidate is the Hernquist (1990) (hereafter HER) profile: $(\\alpha,\\beta,\\gamma) = (1,1,3)$. \\item These fitting formulae provide a good model for halos formed in simulations of both scale-free and cold dark matter (CDM) universes. \\item The value of the scale radius $r_s$ depends both on the initial cosmology and on the mass of the halo in a way apparently related to the formation time of the halos. \\end{enumerate} Interestingly, a mass dependence for the scale radius was also found observationally by Sanders \\& Begeman (1994), who used the HER model to fit the dark matter component when modelling the rotation curves of a sample of spiral galaxies.  Despite the recent wealth of studies on this subject, the computational limits of present day machines are such that simulations of galaxy clusters are only now starting to reach a resolution sufficient to resolve reliably the dark matter structure of the central $\\sim$ 100 kpc. As a result several issues are still waiting for more detailed study. Among them: have the results presented so far converged? That is, are they independent of numerical limitations? Does the trend of the scale radius with halo mass depend on numerical resolution? Are the simulated halos in dynamical equilibrium? And, especially, do halo dynamics affects halos structure, for example, the shape of density profiles, and the scatter in the relation between the halo mass and the scale radius $r_s$?  More generally, the distribution of dark matter in the central regions of the halo is of particular interest; for example, the ability of a cluster to act as a gravitational lens, producing multiple magnified images of background galaxies, depends crucially on the mass content of the very central part (few tens of kpc) of the cluster, hence on the slope of the density profile at that scale. Further questions to ask are then: what is the structure of cluster-size halos in the central few tens of kpc? Are results of simulations compatible with recent lensing observations?  The purpose of the present paper is to address some of these points. Section 2 presents the simulations. Section 3 is dedicated to extensive numerical tests to establish the reliability of our results. In Section 4 the halo formation process is interpreted using a simple description in terms of its mass and {\\em rms} velocity. Section 5 presents our main results on halo density profiles, their dependence on the dynamical configuration of the system, analytical fits to them, and the dependence of these fits on halo mass. Section 6 discusses related topics, like the dynamical equilibrium of halos, and the comparison of simulations to dark matter observations from lensing studies. Finally Section 7 summarizes the results and presents some conclusions. ",
        "conclusions": "In this paper we have used high resolution $N$-body simulations to study the structure and dynamical evolution of the dark matter halos of galaxy clusters.  Firstly we performed extensive numerical tests of some parameters that define the simulation setup and determine the resolution of the results. These are the number of particles forming the final halo, the gravitational softening parameter $s$ which sets the small scale cutoff of gravity, and the tolerance parameter that determines the accuracy of the time integration of the system. From these tests we could assess the accuracy limits of our simulations, and so choose parameter values appropriate to the dynamical resolution we required. We then obtained a set of nine dark matter halos, resolved on average by $\\approx 20000$ particles each, with an effective force resolution of $\\simeq 25$ kpc.  We studied the formation process of these halos and discussed their dynamical equilibrium and departures from this equilibrium during their evolution. We analyzed the density and velocity field of the dark matter, especially in the central region of the halos, by using radial profiles of these quantities. We tried two analytical fits to these profiles, and estimated their performance and range of applicability. Under different approximations for the dynamics of the system, we tested the accuracy of the Jeans' equation in estimating the halo mass within different radii. We finally compared the dark matter profiles of the simulated halos to those inferred from recent observations of gravitational arcs, arclets and background distortions in rich clusters of galaxies.  Our main results are the following. \\begin{enumerate} \\item The halo formation process can be simply schematized as an alternation of merging phases and relaxation phases. Halos spend on average one third of their evolution in perturbed configurations, and the lasting two thirds in relaxed configurations. During merging the halo increases its total mass by $20\\%$ to $100\\%$; its velocity dispersion also increases by $\\approx 15\\%$ to $40\\%$ due to the velocity of the infalling lumps. After the first passage of a lump through the centre of the main halo, the latter starts to relax, the velocity dispersion decreases towards its equilibrium value, and substructure is erased. During these phases, halo densities can vary by up to a factor of two, circular velocities by up to 60\\% and radial velocity dispersions by up to 70\\%.  \\item The average configuration of a simulated halo is not isothermal. If we call $\\alpha$ the local logarithmic slope of the density profile: $\\rho(r) \\propto r^{-\\alpha}$, appropriate values for the simulations are $\\alpha \\simeq -1$ at the smaller radii, and $\\alpha \\simeq -3$ to $-4$ around the virial radius. The corresponding circular velocity profiles, $v_c(r) = (GM(r)/r)^{1/2}$, increase from the centre outwards, reach a maximum value and start decreasing before the virial radius, although only by a small amount, 10\\% to 25\\% of the peak value; a similar trend is shown by the {\\em rms} velocity. The velocity dispersion within halos is anisotropic; orbits are more elongated at larger radii, but almost isotropic at smaller radii. \\item The analytic models proposed by NFW and HER fit the dynamically averaged halo profiles with good accuracy: the {\\em rms} residuals for the density are 17\\% and 26\\% respectively. For the circular velocities they are 6\\% and 10\\%. The two models provide an even better fit to the simulations if one allows the halos to evolve until they erase most of their substructure. However, systematics in the residuals show that the fits are slightly too flat at small radii and too steep at larger radii.  \\item More massive halos have on average flatter density profiles than less massive ones, as indicated by the trend of the scale radius in the analytic models. The scatter in the relation is such that the same analytical profile can be the best fit for halos with mass differing by a factor of two at a $2\\sigma$ level. \\item The assumption of a static spherical system is always a very good approximation for estimating the halo mass using the Jeans' Equation. The true mass can be correctly recovered during most of the evolution, and although the scatter in the estimate reaches a factor of two in dynamically perturbed halos, the average estimated mass always agrees with the true value. In this sense, the halos are always in approximate dynamical equilibrium within their virial or Abell radius. A simple virial model provides a very good estimate of the Abell mass (but not of the mass at smaller radii) for halos not in a merging phase. \\item Finally, we found that our simulations produce dark matter halos that can match quite well the dark matter distribution of galaxy clusters A2218, 1455+22 and 0016+16, recovered from observations of gravitational lensing in these clusters, on scales from few kpc to $\\approx 1$ Mpc. \\end{enumerate}"
    },
    "9603/astro-ph9603060_arXiv.txt": {
        "abstract": "We extend our previous method to determine globular cluster ages using the luminosity function (Jimenez \\& Padoan 1996). We show that the luminosity function depends on both age and distance modulus and that it is possible to distinguish between the two. This method provides at the same time independent determinations of distance and age of a GC by simply counting the number of stars found inside specified luminosity bins. The main uncertainties in other traditional methods for determining GCs ages are absent (e.g. mixing length, color-$T_{\\rm eff}$ calibration, morphology of the color-magnitude diagram ).  The distance modulus is the biggest uncertainty in determining the age of GCs. Here we show that the age can be determined with small uncertainty for any value of distance modulus using the LF and that the LF allows a determination of the distance modulus itself. This is explained by the fact that the luminosity function is affected by a change in distance-modulus in a way that is different from its time evolution.  If GC stellar  counts with statistical errors not larger than $3\\%$ are available, the age can be determined with an uncertainty of about 0.4 Gyr (independent of distance modulus, mixing length and color calibration) and the distance modulus with an uncertainty of about 0.04 mag. ",
        "introduction": "GCs are still the best cosmological clock for measuring the age of the Universe. Since the dating of GCs is totally independent of the cosmological model adopted to describe the Universe, GCs serve as a constraint for different cosmology scenarios.  Despite the continuous effort carried out during more than 30 years to give a precise value for the age of GCs, the uncertainty in their age still remains of about 4 Gyr. The problem is particularly complicated because age and distance have the same effect on the morphology of the main sequence turn off point (MSTO).  In order to tackle this problem, some alternative methods to avoid the influence of the distance modulus have been proposed (Jimenez et al. 1996). Despite of this it is still interesting to find more precise methods that give both distance modulus and age with an accuracy better than 5\\% in order to be able to constrain cosmological models.  Until now it has been impossible to do so due to the degeneracy distance modulus - age. When the MSTO method is used (both isochrone fitting and delta V (Chaboyer, Demarque \\& Sarajedini 1996)) this degeneracy implies that a different distance modulus can be mimic with a different mass for the MSTO and therefore a different age. In general this leads to an uncertainty in the age of 3 Gyr. On top of this, some other uncertainties in the stellar physics input leads to another 1 Gyr (at least) in the uncertainty of the age.  The luminosity function seems the most natural observable to try to constrain both age and distance modulus in an independent and accurate way at the same time.  The luminosity function is a natural clock because the number of stars in a given luminosity bin decreases with time since more massive stars evolve more rapidly than less massive ones. The fact that small differences in stellar masses corresponds to large differences in evolutionary time explains the power of our clock, rather than being a source of uncertainty in getting GCs ages (as it is in the MSTO method).  The luminosity function is also a natural distance indicator because the number of stars in a given luminosity bin depends on the position of the bin. This dependence is different from the time dependence as it is illustrated by the fact that a time translation of an evolutionary track for a given stellar mass does not result in an evolutionary track of a star with different mass (see Fig. 1).  In this paper we illustrate how the luminosity function can be used to determine both the distance modulus and the age of a given GC. The paper is organised as follows: in section 2 we describe the luminosity function method, we continue in section 3 with the dependence of the luminosity function on distance and age. In section 4 we discuss the effect of the IMF. We conclude with a general discussion and the conclusions. ",
        "conclusions": "In this paper we have addressed the question of how to compute GC ages using the luminosity function. We have also shown how it is possible to break the age -- distance modulus degeneracy. The main conclusions of our work are the following:  \\begin{itemize} \\item A four bin luminosity function provides a new method to break the degeneracy age--distance modulus in GCs.  \\item The same four bin LF allows the determination of age with uncertainty of about 0.5 Gyr and distance modulus with uncertainty of about 0.05 mag, if observations are available such that the uncertainty in the stellar counts is not larger than 4\\% (which can be obtained with modern telescopes like HST and NTT). \\end{itemize}"
    },
    "9603/astro-ph9603036_arXiv.txt": {
        "abstract": "It is now clear that there is a substantial population of primordial binaries in galactic globular clusters and that binary interactions are a major influence on globular cluster evolution. Collisional interactions involving stars in binaries may provide a significant channel for the formation of various stellar exotica, such as blue stragglers, X--ray binaries and millisecond pulsars. We report on an extensive series of numerical experiments of binary--binary scattering, analysing the cross--section for close approach during interactions for a range of hard binary parameters of interest in globular cluster cores. We consider the implied rate for tidal interactions for different globular clusters and compare our results with previous, complementary estimates of stellar collision rates in globular clusters. We find that the collision rate for binary--binary encounters dominates in low density clusters if the binary fraction in the cluster is larger than $0.2$ for wide main--sequence binaries. In dense clusters binary--single interactions dominate the collision rate and the core binary fraction must be $\\ltorder 0.1$ per decade in semi--major axis or too many collisions take place compared to observations. The rates are consistent if binaries with semi--major axes $\\sim 100 AU$ are overabundant in low density clusters or if breakup and ejection substantially lowers the binary fraction in denser clusters. Given reasonable assumptions about fractions of binaries in the cores of low density clusters such as NGC~5053, we cannot account for all the observed blue stragglers by stellar collisions during binary encounters, suggesting a substantial fraction may be due to coalescence of tight primordial binaries. ",
        "introduction": "As the evidence for the presence of primordial binaries in globular clusters increases, it has become clear that the contribution of binary--single star and binary--binary scattering to stellar collisions and other stellar binary processes must be significant, at least in some clusters (see reviews by Hut \\etal 1992; Livio 1995; also, Sigurdsson \\& Phinney 1995, Davies 1995, Davies \\& Benz 1995, Leonard 1989, Goodman \\& Hut 1989). Of particular importance are tidal encounters, or stellar collisions, that occur during resonances that develop during hard binary--single and binary--binary scatterings. These may contribute significantly to the formation of blue stragglers (Leonard 1989, Leonard \\&\\ Fahlman 1991, Leonard \\& Linnell 1992), X--ray binaries, MSPs, CVs runaway stars and other exotica (Sigurdsson \\& Phinney 1995, Davies 1995).  Binary--binary scattering may be particularly important in low-density clusters where there may be a large number of primordial binaries and products of binary interactions such as blue stragglers (Hills 1975, Nemec \\& Harris 1987, Nemec \\&\\ Cohen 1989, Leonard 1989, Mateo \\etal 1990, Bolte 1991, Hills 1992, Leonard \\& Linnell 1992, Bolte \\etal 1993, Yan \\& Mateo 1994). Mass segregation effects in globular clusters will increase the binary fraction in the core compared to the rest of the cluster. Thus, even if the binary fraction in the whole cluster is low (say $\\sim 5$\\%) the fraction in the core may be much higher (see, for example, Leonard 1989, Hut \\etal 1992, McMillan \\& Hut 1994, but note also Sigurdsson \\& Phinney 1995).  Here we report the results of 100,000 numerical experiments of hard binary--binary scatterings, for a range of binary parameters appropriate to globular cluster interactions. Other studies of binary--binary encounters have been carried out (Mikkola 1983, 1984a,b, Hoffer 1983, Leonard 1989, Leonard \\&\\ Fahlman 1991, McMillan \\etal 1990, 1991, Hut \\etal 1992, Hut 1995, Rasio \\etal 1995). Mikkola considered a range of hard and soft binary scatterings looking at the final state and energy transfer, while Hoffer included mostly soft binary encounters. Leonard's work overlaps with ours, but does not present a systematic survey of cumulative cross--sections as reported here, and our work should be considered complementary to his. McMillan, Hut and Rasio have so far mostly reported studies of particular sets of encounters or encounters in particular models of clusters rather than surveys of cross--sections. We present a set of cumulative cross--sections for close approach during hard encounters for a range of mass ratios and semi--major axis ratios. We compute the relative event rate for the various possible outcomes of the encounters. We also present sample cross--sections for the change in semi--major axis during flybys and compare them to the one seen in encounters between binaries and single stars. We leave a detailed discussion of the subsequent evolution of the systems produced in encounters, such as triple-star systems, to a later paper.  The cross--sections for close approaches calculated here complement previous hydrodynamical calculations of the outcome of stellar collisions and strong tidal interactions in the context of hard binary encounters (Davies \\etal 1994, Davies \\& Benz 1995, Goodman \\& Hernquist 1991, Sigurdsson \\& Hernquist 1992). ",
        "conclusions": "Some care must be taken in considering the effectiveness of binary--binary collisions in globular clusters. The {\\it global} binary fraction at zero age in clusters is probably $0.5-1.0$, comparable with that seen in the field. However, this includes binaries from near contact, $a_i \\sim 0.01 AU$, to extremely wide binaries, $a_i \\gg 10^3 AU$. The former do not interact on short enough timescales to be of interest, except in core--collapsed clusters, and will in due course merge through their internal evolution; the latter are soft and have high encounter rates in all except the very lowest density clusters, and are broken up in a few dynamical timescales. The global binary fraction as a function of $a_i$ seems to have an initial distribution of approximately $0.1$ per decade in $a_i$, and that is the approximation we use above. However, the core population of binaries is modified by several processes, including mass--segregation, dynamical recoil, exchange and breakup. As a result the fraction of binaries in the core, per decade in $a_i$ and at different masses varies with time and cluster parameters. In calculating the expected number of blue stragglers above we made some effort to correct for the dominant processes in the different clusters considered.  We find collision timescales for plausible binary populations comparable to the lifetime of the clusters, and an expected number of blue stragglers sufficient to account for a large fraction of the low density blue straggler population, but overestimating the population in the denser clusters. This can be understood in terms of the dynamical evolution of the globular cluster binary population, as breakup and ejection decreases the core population of binaries.  We have refrained here from discussing in detail the properties of the final state of the binaries. In particular, parameters of interest include the final distribution of semi--major axis, not just for the flybys and exchanges, but also the breakups and triples, and the resultant cross--sections for energy transfer and recoil velocity distribution. Also of interest are the eccentricity distributions of the various final binary states. Of particular interest to us are the properties of the system after it undergoes an inelastic collision. Simulations of such collisions have been performed using SPH (Davies \\etal 1993, 1994, Goodman \\&\\ Hernquist 1991, Sigurdsson \\&\\ Hernquist 1993). Approximating the collision as a totally inelastic ``sticky particle'' merger, conserving momentum but not energy, allows a quick and reasonably accurate way of determining the properties of the merged systems, in particular whether they form a single merged star, or if the merged star is in a binary or even a triple, and if so what the orbital parameters and center of mass recoil velocity of the system containing the merged star is. An analysis of these properties is deferred to a second paper (in preparation).  It is clear that binary--binary interactions are significant for producing stellar exotica through collisions in globular cluster cores. Compared to binary--single interactions, the rates inferred suggest a modest global binary fraction in the cores of the denser clusters, in accord with previous estimates, with $f_b({\\rm all}\\  a_i) \\sim 0.2$ and $f_b(a_i) \\sim 0.05$ per decade in $a_i$, while in the low density clusters the blue straggler population is consistent with a somewhat higher binary population, with perhaps $ > 10 \\%$ of the turnoff mass main--sequence stars in the core being in binaries with $a_i \\sim 100 AU$. Binary--binary collisions most likely dominate binary--single collisions in many low density clusters as suggested by Leonard (1989) and may account for a significant fraction of the blue stragglers observed."
    },
    "9603/astro-ph9603022_arXiv.txt": {
        "abstract": " ",
        "introduction": "Fully self-consistent N-body simulations, where each galaxy is represented by a large number of particles, are a useful, albeit expensive, tool for studying the evolution of galaxy groups and clusters. However, for simulations of large clusters of galaxies, like the Coma cluster, the necessary computing time is prohibitive. As a substitute people have consi\\-dered explicit simulations, in which each galaxy is represented by a single point and the physics of the interactions is modelled by explicit prescriptions for merging conditions. In particular, a variety of recipes are explored for the conditions the two galaxies must fulfill in order to merge. In general, these merging conditions are based on self-consistent si\\-mulations of two-galaxy collisions, and do not include the tidal forces between the galaxies or collisions involving more than two galaxies. It is thus not a priori certain that they will perform well in simulations of group or cluster evolution. In some cases (Merritt, 1983; Richstone and Malumuth, 1983; Mamon 1987), the authors also introduce other effects like dynamical friction and tidal forces from the background. The main advantage of this type of approach is that it is inexpensive in computing time and therefore allows one to explore a wide parameter space. In any case, a considerable fraction of the results on the dynamics of galaxy groups are based on the explicit approach. We may cite works by Jones and Efstathiou (1979), Roos and Norman (1979), Aarseth and Fall (1980), Cooper and Miller (1981), Roos (1981), Roos and Aarseth (1982), Merritt (1983), Richstone and Malumuth (1983), Malumuth and Richstone (1984), Saarinen and Valtonen (1985), Mamon (1987), Navarro et al. (1987) and Schindler and B\\\"ohringer (1993).  Not many self-consistent simulations of groups with more than 10 galaxies can be found in the literature. We can cite the articles by Carnevalli et al. (1981), Ishizawa et al. (1983), Ishizawa (1986), Rhee and Roos (1990), Barnes (1992), Funato et al. (1993) and Bode et al. (1994). The first works of this kind used Aarseth's (1971) N-body code and a limited number of points, typically $10-20$, to represent each galaxy, and only recently it has become possible to use the order of 1000 particles per galaxy.  Our aim is to compare the two approaches to see whether, and under what conditions, one can use explicit simulations and have confidence in the results. For this purpose, we have evolved a set of initial conditions in two different ways. One way is to use an N-body code where physics is included explicitly, the other, to use self-consistent simulations and a treecode (Barnes and Hut 1986, Hernquist 1987 for a vectorised version), representing each galaxy either by $100$ or by $900$ points. In section 2 we describe our initial conditions and the different merging criteria used so far in the literature. In section 3 we compare the results of fully self-consistent numerical simulations to those of explicit simulations made with the various merging criteria, both without (section 3.1) and with dynamical friction (section 3.2). This comparison led us to propose a new merging criterion (section 3.3), whose performance we also compare with the fully self-consistent simulations. In this section we consider only groups with no common all-encompassing dark matter halo. Simulations including such a halo are presented in section 4, where again we compare the results of self-consistent and explicit simulations. We summarise and discuss our results in section 5. ",
        "conclusions": ""
    },
    "9603/astro-ph9603091_arXiv.txt": {
        "abstract": "Analysis results from {\\it ASCA} and {\\it ROSAT} observations of the narrow-line Seyfert 1 galaxy Mrk~766 are reported.  In the {\\it ASCA} observation we observed rapid variability with a doubling time scale of 1000 seconds.  A spectral variability event was observed in which the spectrum softened and hardened above and below $\\sim 1 \\rm keV$, respectively, as the flux increased.  The spectra could be modeled with 5 components: a power law, warm absorber, iron $K\\alpha$ line and soft excess component flux.  The spectral variability resulted from a highly significant change in the intrinsic photon law index from $\\Gamma \\sim 1.6$ to $\\sim 2.0$, an increase in the warm absorber ionization, and a marginally significant decrease in the soft component normalization.  A $\\sim 100\\rm eV$ equivalent width narrow iron $K\\alpha$ line was detected in the high state spectrum.  Spectral hardening during flux increases was observed in three {\\it ROSAT} observations.  The change in intrinsic photon index and disappearance of the soft excess component in the {\\it ASCA} spectra can be explained as a transition from a first order pair reprocessed spectrum to a pair cascade brought about by a sudden increase in the injected electron Lorentz factor.  The change in the ionization of the warm absorber, though model dependent, could correspond to the increase in flux at the oxygen edges resulting from the spectral index change. The {\\it ROSAT} spectral variability can be interpreted by variable intensity hard power law and a relatively nonvarying soft component, possibly primary disk emission.  These results are compared with those reported from other narrow-line Seyfert 1 galaxies. ",
        "introduction": "Mrk~766 is a bright ($F_{(2-10)}\\sim 2 \\times 10^{-11} \\rm erg\\, cm^{-2} s^{-1}$), soft ($\\Gamma_{0.1-2.4} \\sim 2.7$) X-ray source at redshift $z=0.012$.  The spectrum measured with the {\\it Einstein} IPC and MPC was complex and ultra-soft ($\\Gamma=1.77$; $kT=18.6 \\rm eV$; Urry {\\it et al.} \\markcite{34} 1990).  A shortest time scale of variability of 1000 seconds and a steep and variable power law index was found in a long observation using {\\it EXOSAT} (Molendi, Maccacaro \\& Schaeidt \\markcite{20} 1993).  During the {\\it ROSAT} All Sky Survey, Mrk~766 was bright ($F_{0.1-2.4} \\sim 1.5 \\times 10^{-10} \\rm erg\\, cm^{-2} s^{-1}$ (unabsorbed)) and variability by a factor of three with no accompanying spectral variability was observed in 10--12 hours (Molendi, Maccacaro \\& Schaeidt \\markcite{20} 1993).  Pointed {\\it ROSAT} observations revealed spectral variability that Netzer, Turner \\& George (1994) \\markcite{24} showed could not be explained by a change in ionization of a warm absorber, and Molendi \\& Maccacaro (1994) \\markcite{19} attributed to a change in the accretion rate.  Mrk~766 is a member of the X-ray narrow line Seyfert 1 (NLS1) galaxy class (Osterbrock \\& Pogge \\markcite{25} 1985; Goodrich \\markcite{13} 1989).  {\\it ROSAT} observations of NLS1s find soft 0.1--2.4 keV X-ray spectra and rapid, large amplitude soft X-ray variability.  The soft X-ray spectra of NLS1s are systematically steeper than the spectra of broad-line Seyfert 1 galaxies (Boller, Brandt \\& Fink \\markcite{3} 1996 and references therein). A harder high energy power law component generally was not observed in the relatively soft {\\it ROSAT} band. Only a few observations at higher energies have been reported.  The {\\it ASCA} spectrum of the NLS1 object IRAS~13224-3809 is dominated below $\\sim 2$ keV by a soft excess and from 2 to 10 keV by a hard ($\\Gamma \\sim 1.3$) power law (Otani \\markcite{26} 1995).  In contrast, a very steep spectrum with $\\Gamma_{(2-10keV)} \\sim 2.6$ was found from NLS1 object RE~1034+39 (Pounds, Done \\& Osborne \\markcite{39} 1995).  We report the results from December 1993 {\\it ROSAT} and {\\it ASCA} observations of Mrk~766. Timing analyses of two {\\it ROSAT} archival observations are also presented. In section 2 the data reduction is discussed briefly.  In section 3 timing analyses using normalized variability amplitudes and hardness ratios are presented.  In section 4 the spectral analysis of the {\\it ASCA} data is described.  The results are discussed in terms of standard models in Section 5 and compared with reported results from other NLS1s.  A summary and conclusions are given in Section 6. ",
        "conclusions": "We report analysis of {\\it ASCA} and {\\it ROSAT} observations of the narrow-line Seyfert 1 galaxy Mrk~766.  In the {\\it ASCA} observation rapid variability with doubling time scale of order $\\sim 1000$ seconds was observed, and dramatic spectral variability over as time period of less than $\\sim 10,000\\,\\rm s$ was discovered. Confined to a single event, during a 2--10 keV flux increase the spectrum above and below $\\sim 1$ keV softened and hardened respectively. The low and high flux spectra could be described with a model consisting of a power law, iron line, warm absorber and soft excess modeled as a black body.  The spectral variability was a result of a highly significant increase in the intrinsic power law index from $\\sim 1.6$ to $\\sim 2.0$ with the pivot point at $\\sim 9$ keV, a model dependent increase in the ionization of the warm absorber, and a marginal decrease in the soft excess component.  A $100 \\rm \\, eV$ equivalent width narrow iron line was detected in the high flux spectrum but not in the low flux spectrum, most likely because of poor statistics.  Variability on time scales as short as $\\sim 2400$ seconds was found in the {\\it ROSAT} data.  Because the variability in the softest {\\it ROSAT} band, below 0.4 keV, had relatively lower amplitude than the harder bands, spectral hardening during flux increases was detected on time scales as short as the orbital period of $ \\sim 6000 \\,\\rm s$.  The spectral index change, the disappearance of the soft component in the {\\it ASCA} band and the confinement of the spectral variability to a single event could be naturally explained in terms of non-thermal Comptonization models.  We postulate that the index change occurred through a transition from a first order pair reprocessed spectrum to a pair cascade spectrum brought about by a sudden increase in the Lorentz factor of the injected relativistic electrons.  The first order pair reprocessed spectrum observed in the low state as a soft excess disappeared in the high state cascade spectrum. Variations in the hard compactness resulted in pure flux variability.  The measured increase in the warm absorber ionization corresponds to the increase in flux near the oxygen edges resulting from the power law index change.  The spectral variability in the {\\it ROSAT} data was most naturally explained by a variable hard component and a nonvariable soft component which dominated the softest band and may be primary emission from an accretion disk perhaps implying that reprocessing is relatively less important in this object.  The flat and variable hard power law index observed in Mrk~766 is similar to that observed in NGC 4051 (Guainazzi et al. \\markcite{15} 1996), a Seyfert 1 with many properties common to NLS1s, but contrasts markedly with the very steep hard X-ray index $\\Gamma \\sim 2.6$ found in NLS1 object RE~1034+39 (Pounds, Done \\& Osborne \\markcite{39} 1995). Further hard X-ray observations of NLS1s using {\\it ASCA} are necessary to clearly understand the hard X-ray properties of these sources."
    },
    "9603/astro-ph9603058_arXiv.txt": {
        "abstract": "Rest frame Str\\\"omgren photometry (3500{\\AA}, 4100{\\AA}, 4750{\\AA} and 5500{\\AA}) is presented for 509 galaxies in 17 rich clusters between $z=0$ and $z=1$ as a test of color evolution. Our observations confirm a strong, rest frame, Butcher-Oemler effect where the fraction of blue galaxies increases from 20\\% at $z=0.4$ to 80\\% at $z=0.9$.  We also find that a majority of these blue cluster galaxies are composed of normal disk or post-starbursts systems based on color criteria. When comparing our colors to the morphological results from HST imaging, we propose that the blue cluster galaxies are a population of late-type, LSB objects who fade and are then destroyed by the cluster tidal field. After isolating the red objects from Butcher-Oemler objects, we have compared the mean color of these old, non-star forming objects with SED models in the literature as a test for passive galaxy evolution in ellipticals. We find good agreement with single burst models which predict a mean epoch of galaxy formation at $z=5$. Tracing the red envelope for ellipticals places the earliest epoch of galaxy formation at $z=10$. ",
        "introduction": "Observational astronomy has one major advantage over any other physical science, the phenomenon of lookback time. The combination of enormous distances plus the finite speed of light permits the study of the behavior of objects in our distant past and, with respect to extragalactic astronomy, this allows direct observation into the evolutionary history of galaxies, as opposed to just deducing the past based on their contents at the current epoch. This furnishes data for the investigation of the evolution of stellar populations, star formation history and the conditions of galaxy formation, all relatively new fields as our telescope collecting power has increased in recent years.  Unfortunately, the study of distant galaxies is inhibited by several difficulties primary being that their great distances imply small sizes and faint apparent luminosities. In addition, large distance also means the redshifting of regions of interest (e.g. the near-blue) into the technically and observationally troublesome near-IR.  The study of color evolution of galaxies has divided into three parts in recent years.  Foremost are optical and near-IR studies of extremely high redshift galaxies, mostly radio selected systems, to probe the conditions immediately after the time of galaxy formation (Eisenhardt and Chokshi 1990, McCarthy, Perrson and West 1992). The main purpose of these studies has been to search for protogalaxies, but they also place tight constraints on the the star formation history of galaxies. Second are color selected galaxy surveys in broadband colors or the 4000\\AA\\ break which have been used to test spectrophotometric model predictions out to redshifts of two (Hamilton 1985, Eisenhardt and Lebofsky 1987). Lastly are the numerous photometric studies on the fraction of blue galaxies in clusters (Butcher and Oemler 1984, Dressler, Gunn and Schneider 1985) relating to the rapid changes in cluster populations at only modest redshifts (0.2 to 0.5).  This study is a photometric analysis of galaxy cluster populations out to a redshift of one (10 Gyrs ago, $H_o = 50$ km sec$^{-1}$ Mpc$^{-1}$, $q_o=0$). We attempt a compromise between a full spectral analysis of distant clusters, versus simple, K-corrected broadband colors, by using narrow band blue filters ``redshifted\" to the cluster redshift ($\\lambda_{obs}=\\lambda_o(1+z)$). This use of narrow bandpasses, rather than direct spectroscopy, produces some loss in spectral resolution but gains in increased S/N per object and coverage of the entire cluster per exposure.  We present new photometry for 8 clusters from $z=0.6$ to $z=1$ (the limit of ground based, optical photometry). When combined with our previous results from our zero, low ($z=0.2$) and intermediate ($z=0.4$) redshift samples (Fiala, Rakos and Stockton 1986, Rakos, Fiala and Schombert 1988, Rakos, Schombert and Kreidl 1991, Schombert \\etal 1993), we can use the results to address four galaxy evolution questions:  1) the change in the fraction of red to blue galaxies with redshift (Butcher-Oemler effect), 2) the nature of blue cluster galaxies, 3) the mean colors of ellipticals as a function of redshift (color evolution) and 4) the redshift of galaxy formation. ",
        "conclusions": "The extragalactic meaning of the Str\\\"omgren colors are described in detail in our previous papers.  Briefly, the Str\\\"omgren colors are crude estimators of recent star formation, mean age and metallicity. The $uz$ filter is centered in the near-UV and measures the amount of recent high-mass star formation, $vz$ is centered of the CN-Fe blend at 4170\\AA\\ and is sensitive to metallicity effects and $bz$ plus $yz$ are centered on regions with no strong spectral features and serve as continuum measures. In previous papers, we have used $uz-vz$ as a measure of the 4000{\\AA} break, $vz-yz$ as a measure of global metallicity and $bz-yz$ as a measure of mean age of the underlying stellar population. However, precise understanding of these colors requires comparison to SED models with various assumptions (e.g. redshift of formation, mean metallicity and IMF) since varying star formation histories inhibit a unique interpretation of the colors. For example, $vz-yz$ colors are only sensitive to metallicity only for a homogenous age population. Recent star formation can also strongly influence $vz-yz$ (see Figure 1). To this end, we have convolved our filters with models from Guiderdoni and Rocca-Volmerange (1987) and restrict our interpretation to global properties of the sample and state the model dependence of our results specifically in our discussion and conclusions.  In addition, we have focused our analysis on four phenomenon, the changing fraction of blue to red galaxies in clusters (Butcher-Oemler effect), the nature of this blue cluster population, the color evolution of red galaxies (assumed to be the progenitor of present-day ellipticals) and estimating the epoch of galaxy formation from the red envelope (O'Connell 1987).  \\subsection{BUTCHER-OEMLER EFFECT}  The Butcher-Oemler effect is the strongest evidence of direct evolution of the stellar populations in galaxies that has been discovered to date. In its simplest form, the Butcher-Oemler effect is the observed increasing ratio of blue galaxies in a cluster as a function of redshift. There has long been an expectation that galaxy colors change with redshift since the gas depletion rates for many galaxy types are less than a Hubble time. Cessation of star formation naturally leads to a reddening of the integrated colors of the galaxy as massive blue stars in the underlying stellar population evolve to the red giant branch. Also, in practical terms, star formation rates are zero today for ellipticals and S0's, whereas they must have been non-zero at some point in the past to produce the current population. Thus, at some past epoch, their mean integrated colors must have been bluer. However, the Butcher-Oemler result is surprising because of the extremely rapid change in the fraction of blue galaxies from nearly zero at the current epoch to approximately 20\\% at only modest redshifts of 0.4, approximately 4 Gyrs ago (Butcher and Oemler 1984).  In contrast, even evolutionary models with late galaxy formation epochs predict that color changes were minor below a redshift of 0.5.  Our color data is unique from previous cluster studies in that we can determine the fraction of blue to red galaxies in rest frame colors without any K-corrections or model dependent parameters. Our color indices also remove any background or foreground contamination, a major inhibitor in broadband photometry surveys which require substantial background corrections or spectroscopic redshift confirmation. Our narrow band colors are also more finely tuned to test for star formation versus metallicity effects.  In addition, matching to the cluster redshift eliminates contamination from emission lines.  The original broadband definition of the fraction of blue to red galaxies, $f_B$, given by Butcher and Oemler (1984) is the fraction of galaxies 0.2 mags bluer than the mean color of the E/S0 sequence after K-corrections to the total number of galaxies in the cluster.  As discussed in Paper III, a sample of nearby spirals and irregulars is used to determine a value of $bz-yz$ and $vz-yz$ in our filter system that separates star-forming from quiescent galaxies. From this previous analysis, we have defined the fraction of blue galaxies, $f_B$, as the ratio of the number of galaxies bluer than $bz-yz$=0.20 or $vz- yz$=0.40 to the total number of galaxies. Since the mean $bz-yz$ color of a present-day elliptical is 0.37 (Schombert \\etal 1993) and $bz-yz$ maps into $B-V$ in a linear fashion with a slope of 1.33 (Matsushima 1969), then a cut at $bz-yz$=0.20 is effectively the same as Butcher and Oemler's 0.2 mag selection. To match limiting absolute magnitudes from cluster to cluster, we have only used galaxies with $yz$ mags greater than the magnitude of the 3rd ranked galaxy plus three (similar to Abell's definition of cluster richness). Since the lower redshift clusters were imaged on 1 to 2 meter class telescopes, and the clusters above $z=0.6$ on the KPNO 4m, we found the completeness of the data in terms of absolute luminosity was effectively balanced by S/N.  Both $bz-yz$ and $vz-yz$ are used to calculate $f_B$ and are listed in Table 1.  Our cluster data is shown in Figure 4 along with the original Butcher and Oemler (1984) data and a dotted line as our interpretation of the trend discussed below. Errors are $\\sqrt{N}$ the number of members in each cluster.  The value of $f_B$ from either $bz-yz$ or $vz-yz$ color produces basically the same distribution with redshift, as seen in Figure 4, even though $vz-yz$ is sensitive to metallicity effects and $bz-yz$ specifically tests continuum colors.  This would support our claim that foreground galaxies have not significantly contaminated our sample since foreground galaxies would produce bluer $vz-yz$ colors relative to $bz-yz$ and background galaxies produce very red $vz-yz$ colors and flat spectrum $uz-vz$ colors.  In either scenario, our color selection index, $mz$, would remove these objects from the sample as demonstrated in Papers I through III.  Our measurements of $f_B$ can be tested by comparison to previous studies.  For example, our sample has two clusters in common with the Butcher and Oemler (1984), A370 and CL0024.5+1653, with their values of $f_B$ of 0.21 and 0.16 respectively. Our values are 0.18$\\pm$0.10 and 0.17$\\pm$0.10, in good agreement. In addition, one cluster (CL0939.7+4713) in our sample was observed by HST and published in a morphological/color study of high redshift clusters (Dressler \\etal 1994).  Our color criteria of $bz-yz<0.20$ corresponds to $g-r<0.02$ in rest frame colors.  Applying a K-correction of 1.19 for a redshift of 0.04 (Schneider, Gunn and Hoessel 1983) gives a $g-r<1.21$ cutoff for blue galaxies. For galaxies with morphological classification brighter than $g$=22.0 (100 galaxies), the Dressler \\etal data produces a value of $f_B$ of 0.34, which is well within the errors of our value of 0.28$\\pm$0.10. The morphological information of the blue cluster galaxies is also insightful.  Of the 100 galaxies, 42 are described as early-type (E,L,S in their system) and 58 are classed as spiral or merger (A,B,C,D,M). Of the blue galaxies ($g-r<1.21$), only 3 are early-type and 33 are late-type. Of the red galaxies, 39 are early-type and 25 are late-type.  Thus, it appears that the blue cluster population contains few early-type galaxies (as expected), but the red population contains a mixture of ellipticals, S0's and spirals (see discussion below).  Beyond a redshift of 0.2, the data in Figure 4 shows that $f_B$ increases steadily to a value of 0.80 by a redshift of 0.9. This change in cluster populations is dramatic, not only in terms of the rapid pace of galaxy evolution as first discovered by Butcher and Oemler, but also in the extent of the blue galaxy population dominates the entire cluster population at high redshifts.  The largest value of $f_B$ from either Butcher and Oemler (1984) or Dressler, Gunn and Schneider (1985) was 0.36 at redshifts of 0.4.  The dashed line in Figure 4 represents Butcher and Oemler's interpretation of the trend with redshift from their sample of $z<0.4$ clusters.  A second line is draw displaying our interpretation from the high redshift sample. Our data suggests a slightly more rapid Butcher-Oemler effect than the lower redshift studies. In fact, many of our low redshift clusters have higher $f_B$ values than Butcher and Oemler's clusters at the same redshifts.  The present-day mixture of spirals in rich clusters is 20\\% Sa's, 16\\% Sb's and 4\\% Sc+Irr's (Whitmore, Gilmore and Jones 1993). The mean $B-V$ color of these types are 0.94, 0.87 and 0.69 respectively (Oemler 1991).  The 0.20 mag cutoff for the calculation of $f_B$ signifies that only galaxies with colors bluer than $B-V$=0.79 are selected. Thus, only Sc's and Irr's are counted in present-day clusters which agrees with Butcher and Oemler's value of 0.04 for low redshift clusters. However, our narrow band filters are less sensitive to reddening effects ($E(bz-yz)=0.22$ versus $E(B-V)=0.32$) and a 0.10 mag difference in internal reddening is sufficient to include Sa and Sb type galaxies into our measurements of $f_B$. This would explain why are values of $f_B$ for $z=0.2$ clusters are, on average, higher than the mean value from Butcher and Oemler clusters at the same redshift.  The relationship outlined in Figure 4 is not smooth nor clearly linear. Two clusters (A227 and A2317) display high fractions of blue galaxies at only $z=0.2$ ($f_B=$ 0.51 and 0.65). One cluster at $z=$0.66 (CL0128.8+0628) has a low value for its redshift of $f_B$=0.25.  The Butcher-Oemler effect in previous studies has shown similar behavior in that the observed pattern is more a deficiency of clusters dominated by red galaxies at high redshift, rather than an overabundance of blue dominated clusters or a correlated trend with redshift. In other words, at low redshift one can find clusters with both high and low values of $f_B$, but at higher redshifts we only find clusters with high $f_B$ values.  One possible explanation for the scatter at low redshifts is that various cluster types (irregular to compact) are being selected, whereas at higher redshifts only the richest, densest clusters are detected and cataloged. For example, Allington-Smith \\etal (1991) finds that the value of $f_B$ varies with cluster luminosity from a value of 0.30 for high luminosity clusters to 0.05 for low luminosity clusters. This would introduce a bias in that clusters detected at high redshift are found to be of much higher richness (i.e. cluster mass) than surveys of nearby clusters (the Scott effect). This might select clusters with stronger gas gradients at high $z$ relating to higher gas stripping rates or cluster galaxies with later infall times relating to a longer lasting blue population. This could be interpreted that the mechanism behind the Butcher-Oemler effect works best in dense clusters, or the Butcher-Oemler effect takes time to evolve and the richest clusters are the oldest (Mamon 1986).  \\subsection{THE NATURE OF BLUE CLUSTER GALAXIES}  The origin of the blue galaxies that result in the Butcher-Oemler effect has had two possible interpretations in the literature. The original papers by Butcher and Oemler propose that the blue cluster galaxies are unusual due to ongoing star formation. However, a different interpretation was proposed by Dressler and Gunn (1983) based on spectroscopic observations of Butcher-Oemler clusters. They found that some of the Butcher-Oemler galaxies did indeed have strong emission lines indicative of a protracted period of star formation, but there also exists an unusual number of objects with post-starburst signatures, such as strong Balmer absorption features (E+A galaxies), and with AGN features. For example, the 3C 295 cluster displayed a tenfold increase in the number of galaxies with high excitation emission lines compared to present-day clusters (Dressler and Gunn 1983).  Thus, the mystery lies both in the rapid timescale of the blue population's evolution, as related to the spectroscopic evidence of starbursts, and the identification of the ancestors of the blue population in present-day clusters.  The blue cluster population is selected based on $bz-yz$ continuum colors, although the same population is identified from $vz-yz$ colors. However, the 4000\\AA\\ break colors, $uz-vz$, are more relevant to testing the style and existence of star formation. Figure 5 displays the histograms of $uz-vz$ and $uz-yz$ color for the sample of clusters with redshifts less than 0.6 and greater than 0.6. There is no significant change in the $uz-vz$ colors between the high and low redshifts unlike the sharp change in mean $vz-yz$ or $bz-yz$ colors from Figure 3. The bottom panels of Figure 5 display the high redshift sample divided into red and blue galaxies based on the $bz-yz<0.2$ criteria. Although the blue galaxies have a slightly bluer median $uz-vz$ color than the red galaxies, the difference is not what would be predicted from changes in $bz-yz$ or $vz-yz$ (Schombert \\etal 1993).  The same galaxies are easily distinguished in $uz-yz$. In other words, the blue cluster galaxies at high redshift have redder $uz-vz$ colors than would be expected from their continuum colors and higher $uz$ fluxes compared to the red galaxies.  Redder $uz-vz$ color signals a strong 4000\\AA\\ break which, in turn, is also an indication of a hot star component due to recent star formation. This can be seen by considering a typical old population spectrum in the 4000 to 5000\\AA\\ region shown in Figure 1. Any hot, young star component will increase the near-UV flux as seen in the $uz-yz$ histograms, but will also sharply increase the flux on the red side of the 4000\\AA\\ break (i.e.  compare the 1.5 Gyr model to the 17 Gyr model in Figure 1), resulting in redder $uz-vz$ colors compared to a normal elliptical SED.  We rule out emission lines or AGN activity as a source of blue colors since our filters are specifically chosen to avoid all the significant emission features in the near-blue (O[II], H$\\beta$, O[III], etc.). If the blue galaxy population was due primarily to emission features, then they would not be distinguished in our narrow band colors as unusual.  Therefore, we interpret the trend in $uz-vz$ colors with redshift as evidence that the blue galaxies primarily have a young to intermediate age, post-starburst component dominating their colors rather than a strong, ongoing star formation episode.  The blue cluster population is a sample of galaxies with recent star formation, so there are three global possibilities for their evolution since the same population is absent in present-day clusters.  First, they have either evolved into some known red cluster galaxy type, such as cluster ellipticals and S0's. Second, they have faded from view and are simply not cataloged in our surveys of cluster populations (i.e. are of low surface brightness galaxies, LSB). Or, third, they are destroyed as identifiable units.  Considering the last option first, it is difficult to imagine a scenario where over 80\\% of a cluster population ($f_B$ values at $z=0.9$), or about 1/2 the total cluster luminosity, is destroyed with no remaining evidence.  There are strong tidal forces to play in rich clusters; however, interactions with other cluster members or a central cD galaxy occur at high velocities which is a condition that is not accommodating to tidal disruption (short interaction times) as much as tidal stripping (Merritt 1985, Malumuth and Richstone 1984). Even assuming some mechanism that only disrupts blue galaxies while ignoring red ones (perhaps the blue galaxies are low density structures that formed late), then the newly freed stellar population are still bound to the cluster potential. These stars would produce a luminous halo centered on the cluster core and, although giant halos have been detected around cD galaxies (Oemler 1976, Schombert 1988) and they have total luminosities which rival the brightest cluster members luminosity and not the total cluster luminosity necessary to explain the missing blue cluster population.  The second option is that blue cluster galaxies have faded from view. As the upper main sequence of a stellar population is depopulated after a burst of star formation, the red giant branch grows and the integrated colors redden. In addition, as high mass stars evolve into low luminosity white dwarfs, the luminosity per square parsec declines and the mean surface brightness of the galaxy decreases. For a disk population with a standard IMF, there is a change in one blue mag arcsec$^{-2}$ for every 0.23 mags change in $B-V$ color and, over a period of 5 Gyrs, a galaxy can drop 4 to 5 mags from its peak blue luminosity (Arimoto and Yoshii 1987). The peak luminosity of the blue cluster galaxies must be of order $M_B=-$19 to $-$20 for the blue population to be detected in the cluster surveys thus far. This amount of luminosity is only obtained from a starburst involving $4\\times10^9 M_{\\sun}$ of material.  However, this is a strong burst of star formation where a significant fraction of the mass of the galaxy is turned into stars and, after such an episode, the total number of stars has not decreased so its final surface brightness would remain high.  To see this consider a normal sized disk galaxy with an exponential scale length, $\\alpha$, of 2 kpc with total blue mags of $-$20. Its initial central surface brightness ($\\mu_o \\propto M_B + 5\\, {\\rm log}\\, \\alpha$) is 20.0 $B$ mag arcsec$^{-2}$ which would fade to 24.0 $B$ mag arcsec$^{-1}$ in 5 Gyrs, still quite visible in present-day cluster surveys.  In order for fading to be plausible, consider a larger galaxy with an $\\alpha=10$ kpc and the same absolute magnitude. It would have an initial central surface brightness of 22.5 $B$ mag arcsec$^{-2}$ which would fade to 26.5 $B$ mag arcsec$^{-2}$.  If we assume there were no high surface brightness bulge components (e.g., galaxy type Sc or later), then such a galaxy would be invisible on the Palomar Sky Survey prints and missing from any galaxy catalog.  This proposed LSB population would also explain the high fraction of AGN activity in high redshift clusters (Dressler and Gunn 1983). In a study of nearby, giant LSB galaxies, Knezek and Schombert (1994) found that 60\\% have low luminosity AGN activity. The weak emission is assumed to be due to the low surface density of gas in the cores of these systems, making a deficiency in fuel for the central engine. If the same event which triggers star formation in the blue cluster population also increases the core gas density, then the hidden AGN would increase in luminosity and the higher AGN fraction observed by Dressler and Gunn would be realized.  Unfortunately, there is presently no observational support for a hidden LSB population in present-day clusters. Photographically enhanced surveys of the Virgo cluster (Binggeli, Sandage and Tammann 1985, Impey, Bothun and Malin 1988) have achieved a depth of $\\mu_{lim} \\approx 27$ $B$ mag arcsec$^{-2}$ and have not detected a population of large, LSB disk galaxies. Field surveys for LSB galaxies (Schombert \\etal 1992) have found numerous examples of LSB counterparts to normal disk galaxies (scale length $\\alpha$ of 2 to 4 kpc) and large ($\\alpha > 10$ kpc) Malin galaxies, but none in a cluster environment (see Bothun \\etal 1993). Analysis of the structure of LSB galaxies demonstrates that their low luminosity densities reflect low surface mass densities and, thus, their is an expectation that such systems would not survive in a cluster environment (McGaugh 1992).  The remaining possibility for the fate of the blue cluster population is that they have evolved into some other kind of galaxy type. This galaxy type would also have to be red so as to reconcile the current distribution of galaxy colors in clusters to those in the past. Although there are highly reddened, dust-rich Sa's in clusters, the dominate red galaxy types are ellipticals and S0's. Ellipticals are poor candidates for the blue cluster population since they are composed of a single burst population of at least 12 Gyrs old with no evidence of recent star formation (Wyse 1985) plus have the correct morphological fraction from low to high redshift.  In the previous sections, we have shown that a large fraction of the cluster members have star forming colors by $z=0.9$.  Since 40\\% of present-day cluster galaxies are S0's and 40\\% are spirals, the change from 20\\% at $z=0.2$ to 80\\% at $z=0.9$ is suggestive that the increasing number of blue galaxies in clusters are star-forming S0's and early-type spirals. The decline in blue galaxies to the present epoch would then represent the gradual halt in star formation due to gas depletion.  The later passive evolution of the stellar population from a young, blue one to an old, red one would result in a quiescent S0.  The rapid reddening of blue cluster galaxies is also not unexpected in the context of spectrophotometric models. For example, Ellis (1988) showed that a 10\\% burst (i.e. a strong burst) on top of an old population can evolve in less than one Gyr to a galaxies whose colors are indistinguishable from a 16 Gyr population and blue galaxies at $z=0.5$ have sufficient time to evolve into the red galaxies. A plausible scenario is one where large bulge S0's have a burst or past episode of star formation in their disks. For redshifts from 0.4 to 0.9, the brightness of the disk dominates and the integrated color of the galaxy is blue. As the disk fades, the bulges dominates and the galaxies quickly becomes red. This hypothesis is supported by the near-IR colors of S0 disks which indicate that they have had their last episode of star formation 4 Gyrs ago or redshifts of 0.3 (Bothun and Gregg 1990).  It is tempting, then, to envision an evolutionary progression where spirals are converted to S0's and various mechanisms have been proposed over the past decades using ram pressure stripping or other gas depletion methods (Spitzer and Baade 1951, Gunn and Gott 1972, Larson, Tinsley and Caldwell 1980). However, the single greatest barrier to relating S0's to spirals has been that the distribution of bulge to disk ratios (B/D) for S0's is extremely difference from that of spirals (Dressler 1980). The Hubble sequence is also a sequence of increasing B/D from Sc to S0, where B/D is defined either by isophotal size or relative luminosity. If the Butcher-Oemler effect is a progression of spirals exhausting their gas supply with star formation to become red S0's, then the progenitors must be large B/D, early-type spirals.  Whitmore, Gilmore and Jones (1993) propose that the B/D ratio is also an indicator of formation time, with larger B/D galaxies forming first. Then, the succession from S0 to Sc is a chain of formation epochs where the older galaxies (i.e. future S0's) run out of gas first.  Thus, the cluster Sa's of today are the S0's of tomorrow and the separation of B/D between Hubble types is maintained, not by converting late-type spirals into S0's, but by merely adding early-type, large B/D galaxies to the S0 class in a steady fashion. This slowly shifts the distribution of B/D for S0's as a function of redshift, but maintains their high mean B/D value as compared to star-forming, gas-rich spirals in the cluster. If this conversion process is responsible for the Butcher-Oemler effect, then the blue galaxies in clusters at high redshift should be large B/D spirals where their blue component is a bright, star-forming disk.  Further insight to this scenario can be found if we return to the HST results on CL0939.7+4713 at $z=0.40$ (Dressler \\etal 1994) and examine the morphological types of the blue galaxies. Of the 100 galaxies brighter than $r=22$ with morphological classification, 20 are ellipticals, 22 are S0's (type S or L) and 58 are late-type galaxies (Sa's through Sd's plus seven interacting/merger systems). In contemporary clusters, the mean morophological mixture is 20\\% E's, 40\\% S0's and 40\\% spirals and irregulars (Oemler 1991) so already we can see that there is a deficiency of S0's and an overabundance of spirals in CL0939.7+4713.  The breakdown for spirals and irregulars in present-day clusters are 20\\%, 16\\% and 4\\% for Hubble types Sa, Sb and Sc+Irr (Whitmore, Gilmore and Jones 1993). However, in CL0939.7+4713, the fractions are 12\\%, 23\\% and 23\\%. There is a deficiency of S0's and Sa's and an overabundance of Sb's and Sc+Irr's. Using Dressler \\etal color values for the blue galaxies ($g-r<1.21$), only three were classed as E or S0 from the HST images, the remaining 33 are late-type systems Sa through Irr. Of the red galaxies, 39 were E/S0 and 25 were disk systems. There were no galaxies classed later than Sc in the red sample and the spiral sequence itself also shows a division from blue to red with only 3 blue Sa versus 9 red Sa's, 11 blue Sb's versus 12 red Sb's and 7 blue Sc's versus 4 red Sc's. This distribution has all the signatures of a gas depletion scenario since the ratio of HI mass to luminosity ($M_{HI}/L_B$) is also a decreasing function with galaxy type such that early-type spirals have lower current star formation rates and, therefore, redder optical colors.  However, the data from Dressler \\etal do not find large numbers of Sa's evolving into a population of S0's but, instead, an overabundance of late-type spirals whose B/D's are incompatible with conversion to S0's. Although there are proposals to build large B/D galaxies from late-type galaxies with recent star formation (see Pfenniger \\etal 1994) the stellar populations of present-day bulges do not support this hypothesis. In addition, the luminosity of Butcher-Oemler galaxies is not much brighter than normal cluster spirals and any fading to an S0 would produce a luminosity function for present-day cluster S0's that is fainter than field S0's, which is not found. All this would argue against a scenario where blue galaxies are transformed into present-day S0's; however, a B/D study of the HST images would further resolve this problem.  Our interpretation that the blue narrow band colors are due to disk systems with normal star formation rates or post-starburst galaxies, rather than an ongoing starburst event, is also confirmed by HST imaging. Dressler (1993) reports that the brightest blue cluster galaxies are interacting or merger systems; however, a majority are normal disk systems in appearance. In addition, the morphology of red objects at high redshift is homogenous (Couch \\etal 1993) such that they can not be highly reddened starbursts systems, put represent smooth, old population objects.  The colors of the late-type galaxies are also bluer than their present-day counterparts.  After taking the Dressler \\etal $g-r$ colors, applying a mean $K$-correction of 1.19 (variation by type was less than 0.05) and converting to $B-V$, we obtain means colors for Sa, Sb and Sc+Irr to be 0.70, 0.60 and 0.19.  Oemler (1991) reports that cluster values for Sa, Sb and Sc are 0.94, 0.87 and 0.69 respectively. The field values, also from Oemler (1991), are 0.77, 0.69 and 0.52. So even at a redshift of 0.4, the late-type galaxies have higher star formation rates then cluster spirals today, more in-line with field spirals colors. The Sc+Irr colors are too blue even for field spirals, but are similar to the colors of LSB disk galaxies (McGaugh 1992). Color evolution is not confined to just the Butcher-Oemler galaxies, but all galaxy types in the cluster.  Interestingly enough, when combining the above information from ground-based photometry and HST imaging, none of the three evolutionary scenarios for the blue cluster population are without significant drawbacks or contradictions.  The fading and destruction scenarios fail to match current observational limits for LSB galaxies in clusters or large luminous halos.  However, if the HST results for CL0939.7+4713 are indicative of all high redshift clusters, then neither is it plausible for a large fraction of the blue cluster galaxies to evolve into S0's since the progenitors are mostly small B/D late-type spirals. On the other hand, some fading must occur since the colors of the spirals in CL0939.7+4713 are much bluer than present-day cluster spirals and, when this episode of star formation ends, their mean surface brightnesses must decrease. If we eliminate the late-type spirals in this manner than the ratio of E/S0/Sa is roughly similar to present-day clusters implying that S0's had completed their star formation by redshift of CL0939.7+4713 ($z=0.4$) in agreement with the disk ages from Bothun and Gregg (1990). The fraction of ellipticals is slightly higher as compared to present-day values (37\\% versus 20\\%); however, bright ellipticals in rich clusters are merger products (Schombert 1988) and, therefore, their numbers will be decreased as dynamical friction produces numerous mergers among bright ellipticals in the cluster core.  An alternative to the above scenarios is to assume that shortly after the bright, star-forming period of its life, the blue cluster galaxies are destroyed while the stellar remnants evolve and dim, a hybrid of the destruction and fading scenarios for the evolution of the blue cluster population. If the blue galaxies have their origin as an infalling population of LSB disk galaxies, who are undergoing an enhanced phase of pressure-induced star formation (Evrard 1990), then shortly after their blue phase they will encounter the cluster core and be tidally disrupted, spreading the fading stellar population into the cluster potential. Since the population is undergoing a burst of star formation, it has an enhanced surface brightness to increase their detection at higher redshifts. In addition, the surface mass density is still as low as their present-day field counterparts making them more susceptible to tidal effects than normal disk systems. This hypothesis has the advantage of preferentially destroying these low surface density blue galaxies so as to explain why no giant LSB systems are seen in present-day clusters and also fading the remnant stellar population that would make up the large cluster halos.  There is some observational support for this hybrid scenario since HST images of high redshift blue galaxies indicates that many have a LSB appearance (Couch \\etal 1993) and LSB galaxies have also been offered up as candidates for the faint blue field population that plagues galaxy counts studies (McGaugh 1994).  The major drawback to this scenario that the combined effects of ram pressure induced star formation, enhanced AGN activity from increased fuel supply, fading then tidal destruction is a somewhat contrived scenario and too highly fine tuned to produce large numbers of red clusters galaxies at the present epoch.  \\subsection{COLOR EVOLUTION IN ELLIPTICALS}  Ellipticals are the best subjects for testing color evolution since they are relatively free of dust, nebular emission and non-thermal sources. They also represent the simplest test cases since all indications are that they are composed of a single burst system, i.e. a galaxy where all the stars formed at a single epoch of star formation near the time of formation and where subsequent supernovae removed the remaining gas by galactic winds. The color history of single burst objects becomes an exercise of composite stellar evolution. The variables in evolutionary models of this type are the form of the stellar mass function (IMF), metallicity distribution and history, redshift of galaxy formation and cosmological parameters such as $H_o$ and $\\Omega_o$ (see Buzzoni 1989).  An analysis of color evolution in ellipticals first requires a separation of star-forming, AGN and other active galaxies from the ``passive\" ellipticals. Our filter system uses multiple colors to distinguish blue from red galaxies and, therefore, a separate analysis can be done on the red objects as individuals. At low redshifts, cluster populations are dominated by red objects which are a mixture of ellipticals (20\\%) and S0's (40\\%) (Oemler 1991). At higher redshifts, the Butcher-Oemler effect begins to strongly influence the cluster population such that a over 80\\% of the cluster population is involved in some current (weak to strong) star formation.  However, analysis of the stellar populations in present-day ellipticals indicates that their epoch of star formation was at least 4 Gyrs before $z=0.9$ and that the duration of star formation was less than one Gyr to explain the metallicity dispersion (Rose 1985, Wyse 1985). Thus, this remaining 20\\% of red galaxies at $z=0.9$ must be the progenitors of the current epoch cluster ellipticals.  In our earlier papers we separated the ellipticals from other cluster members with a multiple color criteria (primarily $bz-yz>0.2$, $mz>-0.2$).  This criteria, based on the colors of present-day ellipticals and spirals, is conservative and our analysis was relativity insensitive to our selection process since only 20\\% or less of a cluster population was excluded up to redshifts of 0.4. However, beyond a redshift of 0.4, the number of blue galaxies increases rapidly introducing contamination problems from galaxies within the cluster itself. Also, by redshifts of 0.7, the mean color of an elliptical, as predicted by SED models, approaches our blue limit for galaxy formation redshifts between 5 and 10.  The end result is that our selection of the red population will be model dependent beyond $z=0.6$.  There are three avenues for analyzing the red population given the expected changes in mean color from the standard models. The first is to continue to apply our previous color selection based on present epoch galaxies.  These produce the values shown as open symbols in Figure 6 and, even applying this crude selection, the mean color of red objects, i.e. ellipticals, changes by 0.2 mags bluer from a redshift of 0.5 to 0.9 whereas the changes below $z=0.5$ were very small. This is the first evidence in our filter system that there is significant color evolution in ellipticals.  A second method is to use the predictions of the spectroevolutionary models to estimate a new color correction for each redshift. In this case we have use the UV-cold models of Guiderdoni and Rocca-Volmerangre (1987, hereafter GRV) for a redshift of formation of 5 (the best fit to the $z<0.4$ uncorrected clusters). Our original color criteria was 0.17 blueward of the mean elliptical color at the present epoch and we use the models to maintain this 0.17 mags difference at each redshift. The cutoff $bz-yz$ values are listed in Table 1, column 5 and the new mean colors using these cutoff values are found in columns 6 and 7 and plotted as solid symbols in Figure 6. The third method was to use a constant percentage of the cluster population by color. Guided by the morphological mixture in present-day cluster, we choose 20\\% of the reddest objects for each cluster beyond $z=0.6$.  This selection produced a distribution of mean colors identical to the model cuts above and, therefore, are not shown in Figure 6.  The mean colors of the red population from model cuts are shown in Figure 6 as the solid symbols (error bars are errors on the mean value for each cluster) along with the GRV models ($H_o= 50$, $\\Omega_o=0$ and a Miller-Scalo IMF for a single 1 Gyr star burst at $z_g=5$ and 10). Independent of the model tracks, there is a clear trend for increasing red colors to a redshift of 0.4 with a sharp blueward change from 0.6 to 0.9. There is little difference between the various methods of calculating the mean cluster color out to redshifts of 0.7.  There is also fairly good agreement with the GRV models in both $bz-yz$ and $vz-yz$ including the prediction of a red bump at $z=0.4$ (see Paper II). This small peak is due to a increased contribution from AGB stars 4 Gyrs ago and is not predicted by other models which do not include late stages of stellar evolution. This provides an example of how color observations can be used to test specific parameters in SED models, such as stellar tracks or global metallicity. On the other hand, a faster decrease in color is seen at high redshifts then predicted by the models.  Interpreting the data strictly within the context of these models indicates a redshift of elliptical formation in clusters cores as between 4 and 5.  However, we note that from redshifts of 0.6 to 0.9 there still exist truly red ($b-y > 0.35$) objects in all clusters and, assuming that these are not dust shrouded starbursts, this implies that cluster ellipticals did not all have the same epoch of formation (see \\S3.4).  The $bz-yz$ colors track the UV-cold models quite well for a formation redshift between 4 and 5. The $vz-yz$ colors are also in good agreement with these same models and formation redshifts. However, for individual cluster values of $bz-yz$, the $vz-yz$ colors are slightly bluer than model predictions by 0.05 to 0.10 mags between z=0.4 and z=0.9, although in agreement from z=0 to 0.4. There are two possible sources for this discrepancy between the continuum colors and the metallicity colors. One is that there is an increasing contribution from low metallicity stars are a function of redshift that is not accounted for in the SED models. Since the models are only for solar metallicity stars, a distribution of metallicities (similar to the bulge to halo distribution in our Galaxy) was not considered. The metallicity effects of line blanketing and the mean temperature of the giant branch as well as unusual stellar types from low metallicity populations such as blue horizontal branch stars are more than sufficient to cause this discrepancy and future SED codes that contain a full chemical evolution treatment can address these observations. The second possibility is that there is contamination from the blue galaxy population. The fact that the bluer $vz-yz$ colors begins at the same redshift where the blue cluster population begins to dominate is suggestive. In previous papers, we have pointed out that our study only samples the bright end of the luminosity function (down to about 1/2 $L^*$).  Normally for a cluster population, the bright end is dominated by ellipticals and S0's. However, if the blue cluster population is a population of star-forming S0's, then some post-starburst S0's may contribute to the mean cluster colors regardless of our selection criteria since the distinction between blue and red galaxies is not discrete. On the other hand, one expects a population of blue S0's to fade very quickly since present-day cluster S0's have large B/D ratios and when undergoing a star formation phase the disk light will dominate, but as the star formation ceases, the large bulge component will quickly replace the disk as the dominate source of integrated color. In either case, we can not ignore the possibly of contamination resulting in slight variations in the different colors with redshift.  \\subsection{EPOCH OF GALAXY FORMATION}  The redshift of galaxy formation is a critical constraint on cosmological models. For example, standard CDM with a biasing factor of $b=2.5$ is unable to reproduce large scale structure in a top-down hierarchy with a galaxy formation redshift greater than 3. Early formation redshifts ($z\\geq10$) have been suggested in numerous high redshift studies (Hamilton 1985, Steidel and Hamilton 1993, Hu and Ridgeway 1994). Their results can be summarized that up to redshifts of 3.5 one can still identify red, old population galaxies. Given our current understanding of the photometric evolution of a single burst population, a red galaxy at redshifts greater than three or, 15 Gyrs ago, requires an addition two Gyrs to evolve a dominate red giant branch. This additional evolution requires the epoch of first star formation to be greater than $z=5$.  To explore the earliest epoch of galaxy formation the concept of the ``red envelope'' was invented (O'Connell 1987). This idea selects out the reddest (i.e. oldest) objects at any particular epoch then traces the red edge in the color distribution of galaxies as a function of redshift. Since objects in this study are selected on a multicolor plane, it is unlikely that our red population is contaminated by highly reddened objects but must represent objects which are red due solely to their stellar population. To estimate the redshift of galaxy formation for ellipticals in clusters, we define the red envelope as the 3$\\sigma$ edge of the red galaxy distribution in color.  The data for the mean colors for red galaxies (i.e. ellipticals) in Figure 6 indicates a model dependent formation redshift approximately five. However, the red envelope data, shown in Figure 7 along with GRV SED models of various formation redshifts, indicates a a formation redshift of 10 for the oldest ellipticals in cluster cores.  At $z=1$, the typical elliptical has mean color indicative of a F star population (E+F type rather than E+A for post-starburst blue galaxies). Thus, we can rule out formation epochs less than $z=4$ for ellipticals.  This estimate assumes the reddest objects are the oldest objects, but there are several other factors which can effect the red envelope such as systematic errors in the photometry (i.e.  bias towards detecting red objects) or metallicity effects (i.e. galaxies with the highest mean $[Fe/H]$ have the reddest integrated colors).  Systematic errors in our photometry seem unlikely and work in the opposite direction since sky brightness increased in our reddest filters enhancing the detection of blue galaxies. Metallicity effects are also minor since the color-magnitude relation indicates very small changes in color (less than 0.002 mags in $bz-yz$, see Schombert \\etal 1993) over the range of luminosity sampled herein ($L > 1/2L^*$). However, we note that the $vz-yz$ colors are slightly bluer than the $bz-yz$ colors compared to model tracks and indicate a redshift of formation of 8. Since a metallicity distribution of stars within the galaxy was not specifically included in the models (see Arimoto and Yoshii 1987), then the interpretation of red envelope is more accurate using $bz-yz$ colors.  In addition, since the red envelope indicates a redshift of formation of 10, yet the mean colors produce a value of 5, we interpret this to imply that cluster ellipticals are not coeval.  In the same fashion as the red envelope, we can define a blue envelope to estimate the formation epoch of the blue cluster population. Rather then defining a lower 3$\\sigma$ envelope, we simply calculate the mean color of the blue population as shown in Figure 7. The track from these colors indicate a formation redshift between two and three. If the blue population is composed of LSB galaxies, then this would correct identify this first epoch of star formation, even though the galaxy itself may have formed as a quiescent gas cloud at higher redshifts. This scenario is consistent with the spatial distribution of LSB galaxies (Mo \\etal 1994) and quiescent gas clouds (Lacy \\etal 1993). If the blue population is composed of large B/D proto-S0's, then this track would not indicate the formation redshift, since the spheroidal components have similar ages to ellipticals, but rather the last epoch of disk star formation at $z=0.4$, a redshift consistent with S0 disk age estimates (Bothun and Gregg 1990)."
    },
    "9603/astro-ph9603003_arXiv.txt": {
        "abstract": "Cometary ultracompact \\HII\\ regions have been modelled as the interaction of the hypersonic wind from a moving star with the molecular cloud which surrounds the star.  We here show that a similar morphology can ensue even if the star is stationary with respect to the cloud material.  We assume that the \\HII\\ region is within a stellar wind bubble which is strongly mass loaded: the cometary shape results from a gradient in the distribution of mass loading sources. This model circumvents problems associated with the necessarily high spatial velocities of stars in the moving star models. ",
        "introduction": "The ultracompact \\HII\\ regions (\\UCHIIR) found deep within molecular clouds provide important information on the early phases of the interaction of massive stars with their natal environment.  The disruption of the cloud material by the hypersonic winds and UV radiation fields of these stars is a severe barrier to an understanding of the process of massive star formation -- only by studying the disruption process will anything be learnt about the innermost regions of the protostellar cloud.  The disruption process also involves many important problems of gas dynamics.  Considerable theoretical effort has gone into modelling the varied morphology of \\UCHIIR\\@.  Most theoretical attention has so far been given to the cometary regions, which comprise about 20 per cent of observed \\UCHIIR\\ \\cite{chur90}.  Perhaps the most detailed model, that of Van Buren \\& Mac Low (1992, and references therein), treats them as the steady-state partially ionized structures behind bow shocks driven by the winds of stars moving through molecular cloud material. Although it has been argued that some morphologies which are not apparently cometary (in particular, core--halo) can be explained as cometary structures viewed close to their axes \\cite{maclea91}, the shell and multiply peaked morphologies cannot, and so other models should also be investigated.  There are also a number of unresolved questions with regard to the cometary models.  First, the star is assumed to be moving through relatively homogeneous molecular cloud gas.  Yet it is well known that cloud material has a clumpy distribution down to very small scales. Secondly, current cometary models require rather high stellar velocities, characteristically 10--20$\\kms$ (Van Buren \\& Mac Low 1992; Garay, Lizano \\& Gomez 1994).  The origin of such a high velocity dispersion is problematic.  Moreover, as emphasized by Gaume, Fey \\& Claussen~\\shortcite{gafc94}, if these high stellar velocities are correct, the vast majority of \\UCHIIR\\ should be well isolated from other structures for most of their lifetime.  Yet they often appear closely associated with other radio continuum sources and molecular cores \\cite{galg94}.  Churchwell~\\shortcite{chur95} argues that the parameters used by Van Buren \\& Mac Low~\\shortcite{vbml92} have a fairly large range of possible values.  In particular, the ambient density used in these models is significantly smaller than has been observed ahead of some \\UCHIIR\\ \\cite{cesaea95}.  Since the shape of the bow shock is determined by the ram pressure of the impinging molecular material, such high densities will allow cometary regions to form for substantially smaller relative velocities.  However, such high densities are not common, and are likely to be confined to relatively small clumps.  Since these clumps will not be immediately destroyed on their passage through the bow shock, the model discussed in the present paper may in fact be a more appropriate description of the structure of these regions.  Finally, as more detailed observations become available, the similarity of shape to the structures predicted by Van Buren \\& Mac Low \\shortcite{vbml92} becomes less convincing.  Some of the cometary \\UCHIIR\\ observed by Gaume \\etal~\\shortcite{gaumea95} have tails that hook back around the centre of the nebula, while the bow shock model predicts structures with bright arcs and diffuse tails -- Gaume \\etal\\ indeed propose that such arc-like regions be considered a separate morphological class.  When observed at higher resolution, apparently cometary structures often seem to dissolve into collections of individual emission knots (\\eg{} Kurtz, Churchwell \\& Wood 1994), rather than remaining as smooth structures.  Dyson~\\shortcite{dyso94} and Lizano \\& Cant\\'o~\\shortcite{lizaea95} have suggested that the interaction of a massive star with clumpy molecular material provides a natural alternative model for \\UCHIIR, which circumvents the problem of the relatively long ($10^5\\yr$) lifetimes required by their high frequency compared to field OB stars \\cite{chur90}.  Clumps act as localized reservoirs of gas which can be injected into their surroundings by photoionization and/or hydrodynamic ablation.  The ionized region is bounded by a recombination front; that is to say, an \\HII\\ region is held at a constant radius by the high gas density which results from the mass loading.  The frictional heating of the flow in mass loading regions may explain the very high temperatures inferred for some fraction of the gas in the Sgr B2 F complex \\cite{mehrea95}.  There are many possible variants of this basic premise.  Dyson, Williams \\& Redman \\shortcite[1995 -- hereafter]{dywr95} have described models in which a strong stellar wind mass loads from clumps.  In the model presented in Paper I (as in the present paper), the ionized gas flow is everywhere supersonic.  Redman, Williams \\& Dyson \\shortcite[1996 -- hereafter]{rewd95} have discussed models where either the wind is so weak or the mass loading is so great that the interior flow shocks or is totally subsonic.  The models of Papers I and II are spherically symmetric, and reproduce well many observed morphologies (\\eg\\ shell-like and centre-brightened).  They cannot, of course, produce cometary morphologies.  We here describe how a simple geometrical distribution of mass loading centres can generate such morphologies.  This model has two important properties.  First, it again exploits the clumpy nature of clouds.  Secondly, the star can be at rest with respect to the cloud.  The structures predicted are dependent only in detail on the exact form of the mass loading distribution.  We consider here only the simplest case, namely where the flow in the \\UCHIIR\\ remains supersonic, and defer a discussion of the much more complex subsonic flows, and those in which an initially supersonic wind goes through a termination shock, to a later paper. ",
        "conclusions": "In this paper, we have shown how cometary morphologies may readily be produced in a mass loaded model for \\UCHIIR, with no need for relative motion between the central young OB star and the molecular cloud material, if mass loading sources are concentrated to one side. Velocity channel maps show a shell structure stronger than the integrated emission at small velocities relative to systemic, becoming centre-brightened at higher relative velocities.  We have assumed that the lines are broadened only to the thermal width: in reality, however, the velocity structure may be dominated by turbulence in the dense gas around individual mass loading sources.  We have assumed an exponential distribution in the rate of mass loading with distance, $x$, away from a plane containing the star. The general shape of the regions is not strongly dependent on the detailed form of the mass loading distribution, so long as it has a general decrease in one direction.  In one direction, the recombination front is trapped close to the star by a high rate of mass loading, while in the opposing direction the material remains ionized to a rather larger radius, or even out to infinity.  Such distributions might result not only near the edge of a molecular core, but also if a small relative motion carried clumps into the \\UCHIIR\\ which were then evaporated as they crossed it.  Related models can be considered for different distributions of mass loading centres: for instance, if the luminous star formed on a ridge of clumpy molecular material (or surrounded by a thick disc), the resultant \\UCHIIR\\ would have a bipolar morphology.  This can be pictured by imagining a mirror placed at the plane $x=0$ of Fig.~\\ref{f:lineprofs}, to produce a structure of two back-to-back cometary tails.  The model discussed here will give rise to observable velocity gradients in the head-to-tail direction when the symmetry axis is inclined to the plane of the sky.  Since the velocity of gas crossing the recombination front, $v|R$, is nearly constant around the comet (to within 50 per cent even when the tail is 50 scale lengths away from the star), the maximum observed difference in mean line-of-sight velocity will be $\\sim 2v|R\\sin i$, for inclination angle $i$.  Garay \\etal~\\shortcite{galg94} find differences of $8$--$12\\kms$ between the mean velocities of emission in the head and tail of three cometary \\UCHIIR, which can be fitted with a mild inclination of the cometary axis to the plane of the sky.  While Garay \\etal~\\shortcite{galg94} find no velocity gradients across the regions they observe, Gaume \\etal~\\shortcite{gafc94} list three well-defined cometary regions which have an asymmetric velocity gradient perpendicular to the continuum symmetry axis of the \\UCHIIR\\@.  This observation is problematic for the model discussed here, as it is also for the bow shock model.  These contradictory results may relate to morphology: Gaume \\etal's sample tend to have significantly longer and more complex tails, while Garay \\etal's sample are more arc-like.  We find that such arc-like morphologies are produced for milder asymmetries in mass loading.  They form a natural intermediate between shell and cometary morphologies.  Indeed, we would expect arc-like morphologies also to be found in bow shock models for \\UCHIIR, if the relative velocity of cloud and star were low: such a model might, however, have too low a gas pressure to correspond to observed \\UCHIIR\\@.  Kurtz \\etal\\ \\shortcite{kurcw94} find some observational evidence for a size--density relationship for spherical and unresolved regions, but not for cometary and core--halo \\UCHIIR\\@.  This suggests that a Str\\\"omgren relation might hold for the former class, but not for the latter.  However, given the sparse data on which these correlations are based (which do not allow, for instance, for correction for the luminosity function of the stars), we do not view this as a conclusive argument against the model discussed in the current paper.  In subsequent papers, we will address the effects of tangential pressure forces in mildly supersonic models and the structures expected when the anisotropic mass loading is sufficiently strong to produce a subsonic region of the flow."
    },
    "9603/hep-ph9603378_arXiv.txt": {
        "abstract": "We present the first fully non-linear calculation of inflaton decay. We map inflaton decay onto an equivalent classical problem and solve the latter numerically. In the $\\lambda\\phi^4$ model, we find that parametric resonance develops slower and ends at smaller values of fluctuating fields, as compared to estimates existing in literature. We also observe a number of qualitatively new phenomena, including a stage of semiclassical thermalization, during which the decay of inflaton is essentially as effective as during the resonance stage. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603017_arXiv.txt": {
        "abstract": "The goal of the second flight of the Medium Scale Anisotropy Measurement (MSAM1-94) was to confirm the measurement of cosmic microwave background radiation (CMBR) anisotropy made in the first flight (MSAM1-92). The CMBR anisotropy and interstellar dust emission signals from the two flights are compared by forming the sum and difference of those portions of the data with the same pointings on the sky.  The difference data are consistent with a null detection, while the summed data show significant signal.  We conclude that MSAM1-92 and MSAM1-94 measured the same celestial signal. ",
        "introduction": "Measurements of anisotropy in the Cosmic Microwave Background Radiation (CMBR) continue as a subject of considerable interest to the astrophysics community.  Future anisotropy measurements on scales of 0\\fdg1 to 1\\fdg0 will discriminate among early universe models and determine fundamental cosmological parameters (e.g. \\cite{hu96a}, \\cite{knox95b} and \\cite{jungman95}).  Measurements of anisotropy at angular scales near 0\\fdg5 have been reported recently by \\cite{ruhl95}, \\cite{netterfield96}, \\cite{gundersen94}, and \\cite{tanaka95}.  \\cite{wilkinson95} voiced a common concern when he pointed out that ``there are plausible systematic effects at levels comparable with the reported detections.'' To address this concern the 1994 flight of the Medium Scale Anisotropy Measurement (MSAM1) observed the same field as the 1992 flight to limit the possibility of systematic sources of the signal.  \\cite{cheng94} (hereafter Paper~I) reported observations of anisotropy in the CMBR from the first flight of MSAM1 in 1992 (MSAM1-92). \\cite{cheng95} (hereafter Paper~II) reported the results from the second flight in 1994 (MSAM1-94). A conclusion of the latter is that while a quantitative comparison was pending, there was good qualitative agreement between the two flights in the double difference data set, and that agreement was inconclusive for the single difference data set.  This Letter presents a quantitative comparison of the MSAM1-92 and MSAM1-94 data sets. ",
        "conclusions": "The same region of the sky was observed in the 1992 and 1994 flights to confirm the detection of a celestial signal.  It is clear from the statistical analysis that the same sky signal is measured in these two flights.  We conclude that at the level of our signal, our measurements are likely to be free from significant contamination from time-varying systematics such as sidelobe pickup or atmospheric contamination.  In addition to our own confirmation of the MSAM1-92 results, the Saskatoon experiment has recently observed this section of sky at lower frequencies, 36~GHz to 46~GHz (\\cite{netterfield96}).  They have compared their signal with the double difference CMBR signal from Paper I, and find good agreement.  This result, spanning nearly a decade in frequency, is strong evidence that we are observing CMBR anisotropies rather than some other astrophysical foreground source."
    },
    "9603/astro-ph9603151_arXiv.txt": {
        "abstract": "We discuss the possibility that a significant contribution of the hard X-ray Background is the integrated emission from a population of galaxies undergoing advection-dominated accretion in their nuclei. Owing to poor coupling between ions and electrons and to efficient radiative cooling of the electrons, the accreting plasma is two-temperature, with the ions being generally much hotter than the electrons and forming an ion-supported torus.  We show that the electron temperature then saturates at $\\approx 100 \\keV$ independent of model parameters. At this temperature the hard X-ray emission is dominated by bremsstrahlung radiation. We find that this physical model gives an excellent fit to the spectrum of the XRB in the 3-60 \\keV ~range, provided that there is some evolution associated with the spectral emissivity which must peak at a redshift $\\sim 2$. ~We estimate that such galaxies contribute only to a small fraction of the local X-ray  volume emissivity. The model implies a higher mean black hole mass than is obtained from the evolution of quasars alone. ",
        "introduction": "The puzzle of the origin of the X-ray-background (XRB)  remains unsolved after over 30 years of study. Thermal bremsstrahlung with a temperature of about 40 \\keV ~fits its spectrum between 3 and 60 \\keV ~very well (Marshall \\etal 1980). Despite this agreement, observations of the cosmic microwave background (CMB) exclude the possibility that hot intergalactic gas is a major contributor to the XRB. The excellent fit of the CMB to the blackbody function strictly limits the amount of hot diffuse gas between the current epoch and the last scattering surface at redshift $z\\approx 1000$; such gas  would distort the XRB spectrum through the Compton scattering. The upper limit to any such  distortion gives an upper limit to the contribution of the hot intergalactic gas to the XRB as $\\approxlt 3$ ~per cent (Mather \\etal 1990). \\par It is therefore likely that the XRB is due to a contribution of different types of discrete sources. Active galactic nuclei (AGN) are likely candidates (see Fabian and Barcons 1992 and references therein), especially since Seyferts and Quasars (QSO) provide a large fraction of the soft XRB below $\\sim$ 2 \\keV ~(Shanks \\etal 1991; Hasinger \\etal 1993). One of the major problems with an AGN origin of the XRB has been the apparent discrepancy of the XRB spectrum with the spectrum of resolved AGN. In the $2-10 \\keV$ ~band AGN spectra are too steep and subtraction of their contribution worsens the situation (this is the spectral paradox of Boldt 1987). Moreover detailed studies of the QSO X-ray luminosity function (Boyle \\etal 1994) and the source number count distribution have shown that QSOs are unlikely to form more than 50 per cent of the XRB, even at 1 \\keV. \\par Models based on the unified Seyfert scheme, in which a large fraction of the emission from AGN is absorbed by obscuring matter (Madau \\etal 1994; Celotti \\etal 1995; Comastri \\etal 1995)  seem to work. In these models though, it is difficult to obtain a smooth spectrum in the 1-10 \\keV ~band (Matt and Fabian 1993). Some spectral features are expected due to the matter in the vicinity producing an iron edge and emission line in the intrinsic spectrum. The requirement of a smooth spectrum for the XRB has been recently enphasized by ASCA observations which show no such spectral features (Gendreau \\etal 1995a, 1995b), implying that a new faint population with a very hard, smooth X-ray spectrum is required. \\par A promising new population of sources with harder mean X-ray spectra discovered  in deep ROSAT images (Hasinger \\etal 1993; Vikhlinin \\etal 1995; McHardy \\etal 1995; Boyle \\etal 1995; Almaini \\etal 1996) is emerging as a significant possibility for the missing hard component of the XRB. At present this population can account for only about 10 per cent of the XRB, but the source number counts are still climbing at the lowest detected fluxes. Recent deep ROSAT studies are beginning to resolve some of the population into  narrow emission-line galaxies (Boyle \\etal 1995; Griffiths \\etal 1996) with remarkably high X-ray luminosities, typically two orders of magnitude more than that of galaxies observed locally (Fabbiano 1989) despite comparable optical luminosities. Although it now seems clear that faint galaxies are emerging as a significant new X-ray population, questions on the origin and nature of their activity still remain unsolved. \\par The spectrum of the XRB, with its $\\sim 30 \\keV$ rollover, requires some spectral uniformity in its constituent sources. Since it resembles  thermal bremsstrahlung in the 3-60 \\keV ~range, a mechanism which produces such radiation in galaxy nuclei would be particularly appealing. A way of standardising the temperature is then  required. A possibility has emerged in relation to recent discussion of energy advection solutions (Shapiro, Lightman \\& Eardley 1976; Begelman 1978; Rees \\etal 1982; Abramowicz \\etal 1995; Narayan 1996; Chakrabarti 1996) for accretion disks by Narayan \\& Yi (1995a,b). Beginning with the work of Shapiro, Lightman \\& Eardley (1976) and Rees \\etal (1982), investigations of black hole accretion disks at low $\\dot{M} (\\le 0.01\\dot{M}_{\\rm Edd})$, have focussed on a class of optically thin solutions where the gas is significantly hotter than in the local Shakura-Sunyaev (1973) thin disk solution. Because of the poor radiative efficiency of the accreting gas, in the advection-dominated solution most of the accretion energy is stored within the gas and advected radially inward.  The accreting plasma in this solution is two-temperature; since the ions are much  hotter than the electrons they maintain a thick torus (which is supported by the ion pressure). This requires the gas density to be sufficiently low that ion-electron coupling via Coulomb collisions becomes weak and cooling via synchrotron and bremsstrahlung radiation is not very important. In addition, the gas is optically thin for Compton cooling to be modest. For these reasons, the ion-supported torus exists only at low mass accretion rates (Rees \\etal 1982).  A major attraction of this class of solution is that the plasma attains and maintains an electron temperature $kT_{\\rm e} \\approx 100 \\keV$ and, since the gas is optically thin, much of the X-ray cooling occurs through thermal bremsstrahlung. X-ray spectra from such sources, when integrated over redshift, should plausibly resemble the XRB. \\par In this paper we explore a model in which a population of galactic nuclei undergoes advection-dominated accretion, perhaps as their quasar activity diminishes, and so produces the XRB spectrum. Although we centre our model on  advection-dominated disks (ADD), any situation in which an optically thin, two-temperature magnetized gas occurs would suffice. \\par In section 2 of this paper we review the physical conditions in ADD, as discussed in Narayan \\& Yi (1995b). We determine the two plasma temperatures from balancing the heating and cooling, and deduce the process that dominates the X-ray emission. In section 3, we derive the spectrum of the XRB from our model. We then discuss our results and their implications in the last section. ",
        "conclusions": "The diffuse XRB in the few-100\\keV  ~band can be well explained as due to a population of apparently normal galaxies undergoing two-temperature advection-dominated accretion via an ion-supported torus. \\par Here, we have considered a physical model for the description of the high energy emission and the cosmological evolution of such sources. We find that a population of sources with X-ray emission  due to an optically-thin thermal plasma emitting bremsstrahlung at a temperature $\\theta_{\\rm e}\\approx0.2$  satisfies the background constraints. We show that the temperature requirement is plausibly achieved in ion-supported tori. \\par The spectral emissivity of such sources is most likely to be associated with (some) spectral evolution and the acceptable evolution parameters are roughly consistent with those of AGN. The predicted population of sources also satisfies the constraints of the local volume emissivity. In particular we have found that the predicted emissivity  contributes a small fraction to the local volume emissivity in the 2-10\\keV ~band. For the deeper ROSAT sample, the  spectral emissivity extrapolated in the 0.5-2\\keV ~more closely  resembles the value obtained with our model, suggesting that advection-dominated disks might be related to the activity of X-ray luminous, narrow-line galaxies.  We  note that the low radiative efficiency in the advection-dominated scenario may indicate that most galaxies might harbour a massive black hole as a remnant of earlier more active quasar epoch (see also Fabian \\& Rees 1995). In particular, if we suppose that QSOs contribute  about 20 per cent of the XRB at 1 \\keV ~and  have power-law spectrum  of energy index $\\sim 1$, we can  calculate a ratio of the total emission due to QSOs to that of ADD sources of $E_{\\rm QSO}/E_{\\rm ADD}\\approx 1/4$. If we now assume  that the  radiative efficiency $\\eta$ of QSOs and ADD tori to be about 10 per cent and $0.1$ per cent  respectively, we derive an upper limit for the ratio of the mass accretion rates, i.e., \\[ \\frac{\\dot{M}_{\\rm QSO}}{\\dot{M}_{\\rm ADD}}\\approxlt \\frac{1}{4}\\left(\\frac{\\eta_{\\rm ADD}}{0.1}\\right) \\approxlt \\frac{1}{400}.\\] This leads directly to the conclusion that  the mean  mass of `dead' black holes might be up to a factor of 100 greater than that estimated previously  from measurements of the luminosities. Moreover, given that most of the XRB could be due to a population of ion-supported tori, we can find limits for the total mass density, $\\rho$, of black-holes distributed in local galaxies (Soltan 1982 or Fabian \\& Canizares 1988). The energy density due to accretion is given by $\\varepsilon=\\rho\\eta c^{2}$  which is equivalent to the the total energy emitted by all of the model sources in a unit comoving volume; i.e. the model XRB emission. If we take the spatial density of host galaxies to be $\\approx 10^{-3}\\pmpccu$,  we find that most  normal galaxies should contain a central black hole with a mass $\\approx 10^{7} M_{\\odot}$. Lower space densities lead of course to higher mean masses."
    },
    "9603/astro-ph9603120_arXiv.txt": {
        "abstract": "IC 1396 is a relatively nearby (750 pc), large ($>$2\\arcdeg), HII region ionized by a single O6.5V star and containing bright-rimmed cometary globules.  We have made the first arcmin resolution images of atomic hydrogen toward IC 1396, and have found remarkable ``tail''-like structures associated with some of the globules and extending up to 6.5 pc radially away from the central ionizing star.  These HI ``tails'' may be material which has been ablated from the globule through ionization and/or photodissociation and then accelerated away from the globule by the stellar wind, but which has since drifted into the ``shadow'' of the globules.  This report presents the first results of the Galactic Plane Survey Project recently begun by the Dominion Radio Astrophysical Observatory. ",
        "introduction": "IC 1396 (S131, Sharpless 1959) is a nearby (750 pc; Garrison \\& Kormendy 1976), large ($>$2\\arcdeg), evolved HII region ionized by a single O6.5V star HD 206267 located near its center.  Within the HII region are many dark globules, some with bright rims facing toward the central ionizing star (Pottasch 1956, 1958a,b; Osterbrock 1957).  Since the pioneering work of Pottasch and Osterbrock, these globules have been studied in radio continuum, molecular emission, and optically (for a summary of these studies see Patel et al. 1995).  For example, Patel et al. (1995) have mapped the CO emission within the HII region, and have found that the bright-rimmed globules appear to trace a ring approximately 12.5 pc in radius which is expanding radially outward from the central star.  The expansion is apparently caused, not by the stellar wind or radiation pressure, but by a ``rocket effect'' (Harwit \\& Schmid-Burgk 1983) induced by the ionization of the inner-facing surface of the globules (the ``bright rims'').  Weikard et al. (1995) have also mapped the HII region in several transitions and isotopomers of CO and in HI, but at moderate resolution (several arcmin).  In this report we present the first high-resolution ($\\sim$1') images of atomic hydrogen toward IC1396, showing remarkable ``tails'' of HI associated with some of the globules. (A more detailed analysis, presenting all of the data, awaits a later paper.) These are the first results of the Galactic Plane Survey (GPS) recently begun by the Dominion Radio Astrophysical Observatory (DRAO)\\footnote{The Dominion Radio Astrophysical Observatory is operated as a national facility by the Herzberg Institute for Astrophysics of the National Research Council of Canada.}.    The GPS is being carried out by a consortium of Canadian and international astronomers, and will provide an image of the Galactic Plane in the longitude interval $75^{\\circ} - 145^{\\circ}$ and latitude range $-3^\\circ$ to $+5^\\circ$, yielding an atomic hydrogen (HI) spectral line data cube with 256 velocity channels and angular resolution of $1' \\times 1' {\\rm cosec}(\\delta)$.  At the same time, continuum images at 1420 MHz and 408 MHz are obtained, with full polarisation data at 1420 MHz.  Observations began in March 1995 and will continue for approximately 4 years. ",
        "conclusions": "The morphology of the HI associated with these globules is remarkable and is reminiscent of comets in the solar system.  One possibility for the origin of the HI ``tails'' is that, like solar system comets, they represent material which has been ablated from the globule through ionization and/or dissociation and then accelerated away from the globule by stellar wind from the central O star.  In this scenario, the HI should be moving outward from the globules radially away from the central star, at velocities greater than that of the globules. Patel et al. (1995) found the velocity of the CO emission from A (and associated globule 14) to range from -7.9 to -8.2 km s$^{-1}$, while the globules in the vicinity of B (including 10, 12 and 13) range from -4.8 to -5.6 km s$^{-1}$.  The HI associated with A is thus blue-shifted by $>$1 km s$^{-1}$, and that associated with B by $>$4 km s$^{-1}$.  The CO velocity of globule F is -3.2 km s$^{-1}$. In this case, the HI in the ``head'' is slightly red-shifted compared to the CO (by $\\sim$ 0.5 km s$^{-1}$), while is the ``tail'' is blue-shifted (by $\\sim$ 0.8 km s$^{-1}$). Patel et al. (1995) found that the CO globules appear to trace an expanding ellipse.  The radial components of the expansion velocities for globules A, B and F are toward the line-of-sight, so that if the HI is being blown away from the globules by the stellar wind, the velocity of the HI should be blue-shifted with respect to that of the CO, as we observe (except for the ``head'' of F). However, in this scenario we might also expect to see an increase in the velocity of HI with increasing distance from the central star if the stellar wind were to continue accelerating the gas, as we clearly do not see for the ``tails'' of A and B, and in addition one might expect the gas in the accelerated tails to be ionized.  A second possibility is that the HI ``comets'' are ambient material, perhaps predating the HII region, which lies within the ``shadow'' of the globules protecting it from ionization or acceleration.  In this case, since the globules have been accelerated by the ``rocket effect'' (Harwit \\& Schmid-Burgk 1983) while the HI should have been relatively undisturbed, the HI should be red-shifted with respect to the CO. Except for the ``head'' of F, this is not the case.  A third scenario is that the HI is material which, as in the first possibility above, has been ablated and accelerated from the globules, but the material has drifted into the shadow of the globules where it is sheltered from further ionization or acceleration.  Qualitatively, this scenario seems the most attractive of the three, since it can account for the blue-shifted HI relative to the CO and the lack of acceleration.  We can calculate the mass of atomic hydrogen in these HI ``comet-tails'' by assuming that the optical depth is small.  Then $N_{HI} = 1.823 \\times 10^{18} \\int T_B\\,dv$ cm$^{-2}$ (Kraus 1982). The mass of HI associated with globules A and B is then 22 M$_{\\odot}$ ($\\sim$ 4 M$_{\\odot}$ associated with A, $\\sim$18 M$_{\\odot}$ with B), which is $\\lesssim$ 5\\% of the total mass of molecular gas within these globules (Patel et al. 1995).  The mass associated with F is 1.5 M$_{\\odot}$ ($\\sim$0.25 M$_{\\odot}$ in the ``head'', $\\sim$1.25 M$_{\\odot}$ in the ``tail''), which is $<$2\\% of the molecular mass.  Is there, however, sufficient momentum flux in the stellar wind to have accelerated this material by several km s$^{-1}$?  Chlebowski \\& Garmany (1991) have determined the mass loss rate and wind terminal velocity of HD206267 to be \\.{M} = $7 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$ and $V_{\\infty} = 3.1 \\times 10^3$ km s$^{-1}$.  Thus the momentum flux over $4\\pi$ steradians is $\\approx 2.1 \\times 10^{-3}$ M$_{\\odot}$ km s$^{-1}$ yr$^{-1}$.  If we assume a 1pc diameter globule, roughly 12.5pc from the central star (the approximate current radius of the expanding ring (Patel et al. 1995)), then $\\Omega$/4$\\pi$ $\\approx$ 0.04 and the momentum flux on the globule is $\\approx 8.4 \\times 10^{-5}$ M$_{\\odot}$ km s$^{-1}$ yr$^{-1}$.  The dynamical ages of these tails, neglecting the inclination of the velocity vector to the line-of-sight, are 1.6 - 2.5 Myr, which is roughly the age of the HII region.  The material could have been accelerated for only a short fraction of that time, say $\\lesssim$5\\% or $\\sim 10^5$ yr. In the immediate vicinity of globule A is $\\sim$1 M$_{\\odot}$ of atomic gas. Over $\\sim 10^5$ yr, 1 M$_{\\odot}$ of material would be accelerated to $\\sim$8 km s$^{-1}$.  It is thus plausible that the tails are material which was initially ablated and accelerated from the globules, but is now in the shadow of the dense globules.  There remains the ``head'' of the HI comet associated with F, which is red-shifted rather than the expected blue-shifted.  According to the ``rocket effect'' model of Harwit \\& Schmid-Burgk (1983), it is the action of material being ionized/dissociated on the front surfaces of the globules which accelerates it away from the star.  The HI ``head'' of F might be the initially red-shifted ``rocket-exhaust'' before being accelerated itself by the stellar wind.  However, we see no such red-shifted emission on the front surface of globule A (Fig \\ref{sv_fig}).  Finally, the ring-like HI structure surrounding globule 18 (rim A) is intriguing.  No other globule seems to possess a similar structure.  The globule itself is unusual, with a central cavity or hole which can be seen both optically (Figure \\ref{Acomb_fig}a, Plate \\ref{Acomb_fig}) and in molecular emission (Wooten et al 1983; Nakano et al. 1989; Patel et al. 1995).  Inside the cavity are two stars, LkH$\\alpha$ 349 and LkH$\\alpha$ 349/c, which are young stellar objects, the former of which may be on its way to becoming a Herbig Be star (Hessman et al. 1995).  These stars are unlikely to have ionized or dissociated the gas in the cavity.  Instead, the cavity was likely evacuated during an earlier outflow stage of one or both of these stars (Nakano et al. 1989).  Near the outside edge of the ``backside'' of the globule, $\\sim$100'' south-west of LkH$\\alpha$ 349, is a B3V star VDB 142 (HD 239710, AG+57 1457). Optical images of the region (Osterbrock 1957; P. Boltwood, private communication) show that this star is surrounded by diffuse nebulosity, suggesting that this star is physically associated with the globule. A B3V star can create a small HII region and a larger HI photodissociation region (Roger \\& Dewdney 1992).  There is some indication of an enhancement of HI intensity toward this star.  Thus the ring morphology of this globule may the result of ionization and photodissociation on both the front and back surfaces of the globule, plus the evacuation of the central cavity."
    },
    "9603/astro-ph9603066_arXiv.txt": {
        "abstract": "We compare the radio and soft X-ray brightness as a function of position within the young supernova remnant Cassiopeia~A\\@.  A moderately strong correlation (r = 0.7) was found between the X-ray emission (corrected for interstellar absorption) and radio emission, showing that the thermal and relativistic plasmas occupy the same volumes and are regulated by common underlying parameters. The logarithmic slope of the relationship, $\\ln(S_{\\rm X\\!-\\!ray}) = 1.2 \\times \\ln(S_{\\rm radio}) + \\ln(k)$ implies that the variations in brightness are primarily due to path length differences.  The X-ray and radio emissivities are both high in the same general locations, but their more detailed relationship is poorly constrained and  probably shows significant scatter.  The strongest radio and X-ray absorption is found at the western boundary of Cas~A\\@.  Based on the properties of Cas~A and the absorbing molecular cloud, we argue that they are physically interacting.  We also compare ASCA derived column densities with $\\lambda21$~cm {\\sc H~i} and $\\lambda$18~cm OH optical depths in the direction of Cas~A, in order to provide an independent estimate of ISM properties.   We derive an average value for the {\\sc H~i} spin temperature of $\\approx 40 \\arcdeg K$ and measure the ratio OH/H$_2$ , which is nominally larger than previous estimates. ",
        "introduction": "The basic hydrodynamical structure of idealized young supernova remnants (SNRs) seems well-under\\-stood from a theoretical standpoint (Gull 1973a, Chevalier 1982)\\@.  We expect to find an outer shock, a contact discontinuity between the shocked circumstellar medium and the ejecta, and a reverse shock moving into and decelerating the ejected material. Each of these should give rise to radio and X-ray radiation, with different emissivities depending on the local physical processes.  Observationally, the situation is far from this ideal.  In the best studied young SNR, Cas~A,  none of these structures can be clearly identified (Anderson \\& Rudnick 1995, hereafter A\\&R)\\@.  There are also  questions about the nature of the outer shock, where the expected tangential magnetic fields are not seen (e.g., Kepler - Dickel et al.\\ 1989,   and  Tycho - Dickel et al.\\ 1991)\\@. Inhomogeneities on a variety of scales also complicate the observational as well as the theoretical  pictures (e.g.\\ Borkowski et al.\\ 1992, Cliffe \\& Jones 1994, Jun \\& Norman 1994)\\@.  We need to clarify the nature of actual SNR structures both to understand the hydrodynamics and also to begin addressing important physical issues such as magnetic field amplification and relativistic particle acceleration.  Although reasonable theoretical mechanisms exist for these processes (Gull 1973b, Reynolds \\& Ellison 1992), the observational signatures are far from clear ({\\it e.g.}\\ Anderson et al.\\  1994)\\@.  One fruitful approach to addressing such questions may be a careful examination of the relationship between the X-ray and radio emissivities within a remnant, because of the different physical processes involved.  The bulk of the X-ray emission at low energies results from thermal line emission (Becker et al.\\ 1979, Holt et al.\\ 1994, hereafter HGTN), depending primarily on the temperature ($\\approx 3$ {\\rm keV}) and density  ($\\approx 10~{{\\rm cm}^{-3}}$) of the plasma carrying most of the mass and momentum (Fabian et al.\\ 1980, hereafter F80)\\@. On the other hand, the radio emission is synchrotron radiation from relativistic electrons ($\\epsilon \\approx 0.05-8$ GeV) in magnetic fields of $\\approx 100 \\: \\mu$G (Cowsik \\& Sarkar 1980)\\@.  In one remnant, SN1006, X-ray synchrotron radiation is probably present at keV energies (Koyama et al., 1995), although this is an exceptional case.  Very little quantitative work has been done on the comparison of X-ray and radio emissivities in SNRs\\@.  The canonical wisdom is that the two are well-correlated on large scales, but show little correlation at smaller spatial scales (F80,  Matsui et al\\  1984, hereafter MLDG)\\@. MLDG studied these relations in Kepler, where they divided the remnant into twelve sectors and found a moderate correlation of the form $\\ln\\left( S_{\\rm X\\!-\\!ray} \\right) \\approx (1.1 \\; \\mbox{to} \\; 2.5) \\times \\ln(S_{\\rm radio})  + \\ln(\\mbox{scale factor}) $\\@.  We chose Cas~A for study  because of the availability of both high quality radio and X-ray data.  In the {\\rm cm} wavelength range,  Cas~A is the brightest object in the sky outside of the solar system.  At an age of 300 years, it is believed to be in a pre-Sedov phase, and is situated  $3.4^{+0.3}_{-0.1}$ kpc away (Reed et al.\\ 1995), at the far edge of the Perseus arm.  The column densities of hydrogen between here and Cas~A ($N_{\\rm H}\\approx10^{22}~{\\rm cm}^{-2}$) are such that the optical depths of 1-2 {\\rm keV} X-rays are of order unity.  Therefore, column densities inhomogeneously distributed across the remnant, such as due to structures local to Cas~A and the Perseus arm, can play a large role in determining the apparent soft X-ray morphology.  For these same reasons, Cas~A is an excellent choice as a background source for radio and X-ray interstellar medium (ISM) studies.  The ISM has been well-studied in $\\lambda 21$ cm absorption (Mebold and Hills 1975, hereafter MH75; Bieging et al.\\ 1991, hereafter BGW; and Schwarz et al.\\ 1996, hereafter SGK)\\@.  Although BGW's VLA study was the highest resolution and most detailed, it only covered the velocity range of the Perseus arm.  SGK's Westerbork study, though only at a resolution of 30\\arcsec, covered both the Orion and Perseus spiral arms.    In addition there have been numerous molecular absorption studies using Cas~A --- in H$_{2}$CO (Goss et al.\\ 1984), in CO (Troland et al.\\ 1985, hereafter TCH and Wilson et al.\\ 1993, hereafter WMMPO), NH$_{3}$ (Batrla et al.\\ 1984 and Gaume et al.\\ 1994) and OH (Bieging \\& Crutcher 1986, hereafter BC)\\@.  The absorption patterns in the various {\\it molecules} are similar to each other, but very different than that of the {\\sc H~i}\\@.  However, {\\sc C ii} seems to be correlated with the {\\sc H~i} instead of the molecules (Anantharamaiah et al.\\ 1994)\\@.  Rasmussen (1996, hereafter R96) studied the spatial dependence of X-ray model parameters using the ASCA satellite, resulting in a total column density ($N_{\\rm H}$) map.  In this paper, we compare the radio absorption data of BC and SGK to the $N_{\\rm H}$ map of R96\\@.  From this we measure the scaling relation between column density and equivalent line widths from the radio absorption measurements of SGK and BC\\@.  This allows measurements of the average {\\sc H~i} spin temperature and the $N_{\\rm OH}/N_{\\rm H_2}$ abundance of the ISM to be calculated.  \\begin{figure} \\epsscale{0.85} \\plotone{figures/images.eps} \\\\ \\caption{30\\arcsec\\ resolution images of Cas~A\\@. The histogram-equalization method of scaling was used to enhance the images.  The quantities represented are:  (A) the equivalent width of the $\\lambda21$~cm line; (B) the equivalent width of the $\\lambda18$~cm (OH) line; (C) the total column density as derived from images A and B; (D) the logarithm of the ROSAT HRI image; (E) the logarithmic HRI image corrected for absorption; (F) a logarithmic $\\lambda20$ cm VLA continuum map.  Histograms of the quantities shown in images A, B and C are shown in figure~\\ref{histograms}; image D ranges from -2.6 -- -0.8 $\\rm \\ln(cts~s^{-1}~beam^{-1})$; a plot of image E versus image F is shown in figure~\\ref{L.v.X}\\@. \\label{images} } \\end{figure}  \\begin{figure} \\epsscale{1.00} \\plotone{figures/histograms.eps} \\\\ \\caption{Histograms of $\\lambda21$ cm equivalent width, $\\lambda18$ cm equivalent width and the derived column density ($ N_{\\rm H} = D \\left( EW_{\\sc Hi}\\right) + E\\left(EW_{\\rm OH}\\right) + F$)\\@. \\label{histograms}} \\end{figure} ",
        "conclusions": "In this paper we have presented a technique to correct for spatially inhomogeneous absorption of soft X-rays in Cas~A using radio absorption data.  We find a good correlation between the soft X-ray and radio synchrotron emission from Cas~A, but with significant scatter.  The correlation is probably dominated by variations in path length, implying that the X-ray and radio emissions both occupy the same volumes.  However, we have no evidence for a more detailed relation between their emissivities.  A quantitative interpretation of these results requires more sophisticated modeling of both the X-ray radiative transfer and the relativistic plasma evolution in young SNRs.  Future X-ray\\discretionary{-}{}{/}radio comparisons of Cas~A should concentrate in at least the following two directions:  studies at higher spatial resolution with deeper HRI measurements and comparisons with ASCA's spatially resolved spectroscopy. With a deeper ROSAT HRI observation, it may be possible to separate emissivity from path length variations.  In addition, X-ray proper motions could be measured and compared with the radio proper motions. Since HGTN's paper, Cas~A has been used as a calibrator for ASCA, so more ASCA data have been obtained and ASCA's response functions have been refined.  This will enable image reconstruction techniques to be applied to narrower bandpass images and better quality spatially resolved spectral fitting.  Studying the correlations between the radio and X-ray morphologies as a function of X-ray energy will allow the different emission mechanisms and temperature and metalicity structures to be distinguished.   We have shown that Cas A is likely to be interacting with a dense cloud in the west.  This has affected both the properties of the remnant and the cloud.  Such interactions may play an important role both in SNR dynamics, and in the transfer of energy into the ISM.  We have also demonstrated a new technique for probing the ISM\\@.  By comparing X-ray and radio spectroscopic absorption measurements, the {\\sc H~i} spin temperature and molecular abundances ratios were measured.  Future studies of other radio and X-ray bright extended objects can significantly enhance our understanding of the ISM, by comparing spatially resolved column densities from either the ROSAT PSPC or ASCA with radio and far IR atomic and molecular line data."
    },
    "9603/astro-ph9603047_arXiv.txt": {
        "abstract": "This paper presents new and synthetic narrow band photometry of ellipticals, spirals, Seyferts and interacting galaxies in an attempt to identify the cause of the unusually high fraction of blue cluster galaxies in distant clusters (Butcher-Oemler effect).  The properties and distribution of the low redshift sample specifically points to starbursts as the origin of the blue narrow band colors in interacting Arp galaxies. Comparison between theoretical models and multicolor diagrams, particularly 4000{\\AA} break colors, indicates a photometric signature which differs from both normal disk galaxy star formation and nonthermal components. This photometric signature is absent for the Butcher-Oemler galaxies whose general color distribution, compared to present-day clusters, is consistent with a majority of the blue population involved in normal star formation rates (spiral- like) with the addition of a small fraction of bright, blue interacting/merger systems.  This photometric picture of the Butcher- Oemler galaxies is in agreement with the morophological evidence from HST imaging. ",
        "introduction": "Information gleaned from spectroenergy distributions (SEDs) of galaxies are key to our understanding of their age, metallicity, stellar population content and star formation histories.  Analysis of SED data has usually invoked either population synthesis, which uses a library of spectral templates (Pickles 1985), or evolutionary models, a method of convolving stellar isochrones with star formation rates and initial mass functions (Bruzual 1983, Arimoto and Yoshii 1986, Guiderdoni and Rocca- Volmerange 1987).  In this series of papers (Fiala, Rakos and Stockton 1986; Rakos, Fiala and Schombert 1988; Rakos, Schombert and Kreidl 1991; Rakos and Schombert 1995), a system of narrow band filters has been used to replace the SED sampling at the sacrifice of resolution for a gain in signal-to-noise.  And, just as stellar models are tested by predictions of integrated quantities such as total luminosity, mass, radius and effective temperature, population synthesis results can then be tested by comparison to predictions of integrated stellar content and metallicity of a galaxy as given by narrow band indices.  Thus, theoretical models, which produce spectral energy distributions as a function of age, can be convolved to produce synthetic colors for comparison to narrow band observations at various redshifts.  Our project has approached this problem through the use of a modified Str\\\"omgren filter system called $uz,vz,bz,yz$ to distinguish it from the original $uvby$ system.  Our primary goal was to investigate the color evolution of ellipticals, testing the hypothesis that ellipticals are the result of a single burst where all the stars formed at an epoch near the time of galaxy formation and later aged through passive stellar evolution.  However, in the previous paper in this series (Rakos and Schombert 1995), the data revealed an extended Butcher-Oemler effect (the increasing fraction of blue to red galaxies in distant clusters) out to redshifts of 0.9 and containing up to 80\\% of the cluster population as blue galaxies (versus the previous values of 40\\% at redshift of 0.4). Deep HST images of intermediate redshift clusters (Dressler et al. 1994a,b) indicate that, while the brightest blue galaxies have disturbed morphologies suggestive of an origin of their blue colors from tidally induced star formation, a majority of the Butcher-Oemler galaxies have normal, late-type morphology.  The goal of this paper is to combine the information obtained from narrow band photometry of nearby spirals, starburst, Seyferts and interacting galaxies in order to search for signatures of starburst activity that can then be linked to the $uz,vz,bz,yz$ photometry of Butcher-Oemler galaxies in distant clusters.  The procedure used herein is fourfold. First, the $uz,vz,bz,yz$ system will be linked to the original Str\\\"omgren system for stars to establish a baseline stellar sequence.  Second, synthetic colors for a range of galaxy types will be produced to map the behavior of the $uz,vz,bz,yz$ system on galaxies with composite stellar populations and varying star formation histories.  Third, photometry of interacting/merging galaxies will be presented to understand the signature of starburst phenomenon in galaxies.  And lastly, comparison of the above data will be made to high redshift clusters to examine the hypothesis that blue cluster galaxies are the result of similar starburst activity. ",
        "conclusions": "Rakos and Schombert (1995) demonstrated that Butcher-Oemler effect is much more dramatic than previously thought.  Early work on cluster galaxies found blue fractions from 30 to 40\\% in clusters with redshifts of 0.4.  Rakos and Schombert found fractions of 80\\% in clusters at redshifts of 0.9. This is a very dominate effect on cluster populations not only in terms of the rapid pace of galaxy evolution as first observed by Butcher and Oemler, but also in the that extent of the blue galaxy population prevails over the entire cluster population at high redshifts (80\\% would imply all present-day S0's are actively star-forming at these epochs).  In this paper we have derived and described the fundamental properties of the $uz,vz,bz,yz$ filter system. In particular, we have demonstrated the advantage of $mz$ index as a detector of starbursts in galaxies.  We then applied this index to establish the photometric signature of starbursts in interacting galaxies.  This method was then applied to high redshift clusters to demonstrate that the Butcher-Oemler galaxies are due to a mixture of star-forming spirals and irregulars plus an addition component from a small fraction of starburst galaxies.  The following is a summary of our primary results and interpretations:  \\noindent (1) The $uz,vz,bz,yz$ filter system is a robust tool for exploring the stellar populations in galaxies due to the placement of the filter centers in regions of clearer evolutionary interpretation.  The narrow width of the filters avoids contaminating emission lines so as to concentrate on the underlying continuum emission from stellar atmospheres or nonthermal emission.  \\noindent (2) The $uz,vz,bz,yz$ system is relinked to spectroscopy standards and the transformations from the original Str\\\"omgren system are derived.  The reddening effects are outlined and displayed in Figures 1 and 2.  \\noindent (3) Using spectrophotometry of ellipticals, spirals, Seyferts and starburst galaxies from the literature, synthetic $uz,vz,bz,yz$ indices are produced and plotted on multi-color diagrams.  Specific regions are occupied by various galaxy types and the $mz$ index is demonstrated as the best indicator of starburst activity.  \\noindent (4) Theoretical models from Leitherer and Heckman are used to define the behavior of the $uz,vz,bz,yz$ indices with time from the initial starburst.  The inadequacy of the 4000{\\AA} break color, $uz- vz$, is shown due to the contribution of intermediate mass stars to an aging burst.  Again, the $mz$ index is shown to be the most sensitive to starburst activity and for the longest amount of time post-starburst.  \\noindent (5) New photometry for a sample of Arp and Markarian galaxies is presented and confirms the results from the synthetic colors that starburst galaxies have a unique, detectable signature in the $uz,vz,bz,yz$ system.  Selection by disturbed morphology also confirms the dynamical arguments for starbursts in merging/interacting galaxies.  \\noindent (6) Direct comparison with the photometry of blue galaxies in distant clusters ($z>0.6$) yields the interpretation that the Butcher- Oemler effect is due primarily to normal star formation with a small component due to tidally induced starbursts.  The recent results from HST imaging (Dressler et al. 1994a,b) confirms this interpretation and suggests that the Butcher-Oemler effect is a indicator of dynamical evolution as well as color evolution."
    },
    "9603/astro-ph9603053_arXiv.txt": {
        "abstract": "The optical and magnetic properties of dust grains are reviewed, as they relate to the problem of interstellar grain alignment. Grain geometry plays an important role in determining the optical properties, and scattering and absorption of starlight will produce radiative torques which may drive grains to suprathermal rates of rotation in interstellar clouds; these radiative torques appear likely to play an active role in the alignment process. The likely magnetic properties of grains are discussed, with particular attention to the imaginary part of the magnetic susceptibility. ",
        "introduction": "What physical processes are responsible for the observed alignment of interstellar dust grains? The answer to this question has proved remarkably elusive. The first major theoretical assault on this problem was the classic paper by Davis \\& Greenstein (1951), who examined many of the physical processes involved, and proposed that grain alignment was the result of magnetic dissipation within spinning grains. While other alignment mechanisms are possible -- in particular, alignment by gas-grain streaming (Gold 1952; Roberge \\& Hanany 1990; Lazarian 1994) -- the process of grain alignment by magnetic dissipation has continued to appear attractive, and has received continuing theoretical attention.  In this paper I review some optical and magnetic properties of grains, as they relate to the problem of grain alignment. There is little we can say for certain about the composition, sizes, and geometry of interstellar grains; our current state of understanding is summarized in \\S\\S\\ref{sec:graincomp},\\ref{sec:graingeom} At first sight it is not apparent how the optical properties can affect grain alignment; in \\S\\ref{sec:radtorq} a mechanism is discussed whereby radiation field anisotropy can result in substantial torques on interstellar grains. These radiative torques may produce both suprathermal rotation and grain alignment.  The efficacy of grain alignment by magnetic dissipation obviously depends on the magnetic properties of dust grains, reviewed in \\S\\ref{sec:magprop} Since we do not know what grains are composed of, their magnetic properties are necessarily uncertain; various possibilities are discussed. ``Ordinary'' paramagnetism appears to be marginally capable of producing grain alignment in suprathermally rotating grains. It seems quite possible, however, that at least some interstellar grains may contain ferromagnetic inclusions, which can enhance the rate of magnetic dissipation either by ``superparamagnetism'' or by subjecting adjacent paramagnetic material to a static magnetic field. ",
        "conclusions": "The following are the principal conclusions of this review: \\begin{itemize} \\item Irregular grains are subject to radiative torques when illuminated by anisotropic starlight. These torques are likely to result in suprathermal rotation, and to play an active role in the grain alignment process. \\item The magnetic properties of grains remain very uncertain. Normal paramagnetism by itself can bring about grain alignment if steady suprathermal rotation can be maintained for long periods. \\item The Fe content of interstellar grains is so high that at least part of the grain volume should be ferrimagnetic or ferromagnetic. \\item If the ferri- or ferromagnetic inclusions are not too large, they will be superparamagnetic, and can greatly enhance the rate of grain alignment by magnetic dissipation. Fe clusters with diameter $d\\ltsim90\\Angstrom$ will be effective. \\item Even if too large to be superparamagnetic, ferri- or ferromagnetic inclusions in grains may possibly enhance the rate of dissipation in surrounding paramagnetic material. \\end{itemize}"
    },
    "9603/astro-ph9603115_arXiv.txt": {
        "abstract": "We present complete, low resolution $IJHK$ spectroscopy of the ultracompact HII region, G45.12+0.13.  From the observed HI line strengths, we derive a near infrared extinction law that is  slightly steeper than the average. After correction with this extinction law, we find good agreement between the observed line ratios of HeI, Fe$^+$, Fe$^{2+}$, S$^+$ and S$^{++}$ and the available atomic data.   Our data show that the density within the core of G45.12+0.13 must be at least $10^4$cm$^{-3}$.  This is consistent with the known radio structure of the HII region and in considerable disagreement with previous work using mid and far infrared lines.    There must also be considerable opacity in the HeI 2$^3$P--2$^3$S transition, and we show how the observed strengths of the other HeI lines are consistent with this.   From modelling the photoionisation structure, we find good agreement with most of the observed data if the hottest star present has T$_{eff}\\le 42000$K.  Consideration of the helium ionisation state places a lower limit on this value so that we can also constrain T$_{eff}\\ge38000$K.  Discrepancies still exist between some of the observed and model line ratios, but the most obvious tend to be the mid-IR observations. ",
        "introduction": "Ultracompact (UC) HII regions are generally held to be unevolved nebula marking the site of a young OB star, or group of such stars.  The definition of the UC phase is usually taken to be $n_e > 10^4$cm$^{-3}$ and $r<0.1$pc. (Habing \\& Israel 1979).  Although masked in the visual by the extinction due to the molecular cloud in which they form, they are extremely luminous objects in both the infrared and radio wavebands.  Interest in this area has been renewed with the completion of a sparsely sampled VLA survey of the galactic plane by Wood \\& Churchwell (1989).  They found that the numbers of UC HII regions they detected were inconsistent with the simple picture in which the Str\\\"{o}mgren radius expands at the sound speed of the region, and hence expands to be larger than 0.1pc after $\\sim 10^4$ years (see, eg., Osterbrock 1989).  There are several more complex models that would explain the observed numbers. Van Buren et al.\\ (1990) have suggested that an O star may be formed with a peculiar velocity relative to its natal molecular cloud.  The stellar wind known to exist in main sequence OB stars  supports a bow-shock ahead of it as it passes through the molecular cloud. The ionisation front is trapped and the region is constrained from expanding by this motion.  In another scenario (e.g. Keto, Ho \\& Haschick 1987), material in the molecular cloud is still infalling into the HII region, again constraining the growth of the Str\\\"{o}mgren radius.  Hollenbach, Johnstone \\& Shu (1993) have proposed that a remnant circumstellar accretion disk may replenish the nebula with new material, hence intercepting much of the far UV light from the exciting star. Dyson, Williams \\& Redman (1996) advanced a somewhat similar model, in which the interaction of the stellar wind with a clumpy molecular cloud can lead to mass loading in the wind, creating the same effect as in the Hollenbach et al.\\ model. All of these models can satisfy the lifetime arguments, but all differ in the other physical properties of these regions (eg.\\ the dynamical structure).  The advent of near-infrared array spectrometers has made it possible to obtain spectra of bright nebular sources with a signal-to-noise that would have been impossible before.  Although compact HII regions have been observed widely in the near infrared in the past (see below), only very bright features with high equivalent width have been detected.  There are also many intrinsically fainter features that could be present.  By studying these we can accurately determine the extinction to the HII region, the conditions present in the regions of different ionisation, examine the impact of the large opacity in the HeI triplet 1.083$\\mu$m $2^{3}$P--$2^{3}$S transition on the weaker HeI lines present in the near-IR, and perhaps find transitions that may indicate the presence of bow-shocks (coronal lines are expected to be detectable in these models) or remnant accretion disks (strong neutral transitions or perhaps CO emission).  We have therefore carried out a series of observations of one of the most luminous sources in the VLA survey of Wood \\& Churchwell, G45.12+0.13. This ultracompact HII region has been well studied at many wavelengths, both prior to and since the Wood \\& Churchwell survey.  There have been observations of the hot excited molecular gas around the region (eg.\\ Churchwell, Walmsley \\& Wood 1990, Cesaroni et al.\\ 1991), of the continuum emission both at mm and sub-mm wavelengths (Hoare, Roche \\& Glencross 1991, Wood, Churchwell \\& Salter 1988, Wood et al.\\ 1988), and of the radio recombination lines (Garay, Reid \\& Moran 1985, Churchwell et al.\\ 1990).  Wood \\& Churchwell classify this source as `cometary' (with a bright bow-shaped ionisation front in the radio map), but, as Colgan et al.\\ (1991) note, it can also be modelled as a core-halo source, since there is considerable extended emission beyond the bright core observed by Wood \\& Churchwell (Matthews et al.\\ 1977, Wink, Altenhoff \\& Mezger 1982). Herter et al.\\ (1981) collected the results of all the near-infrared spectroscopy carried out on this source with the first generation of low resolution instruments.  Since then it has largely been overlooked in this waveband, though it has been observed in the mid and far-infrared using both IRAS and the KAO (Simpson \\& Rubin 1990, Colgan et al.\\ 1991).  The best models of G45.12+0.13 show a core with a density $n_e>10^4$cm$^{-3}$, and a lower density halo that may reflect the fact that the ultracompact HII region is embedded in a larger ionised zone (eg.\\ Colgan et al.\\ 1991). Wink et al.\\ (1982) determine the electron temperature to be T$_e=8000\\pm2000$K, in agreement with the data from Wood \\& Churchwell.   This paper presents our results for the 0.85--2.5$\\mu$m region.  In section 2, we detail the observations made, and the data reduction techniques involved. In section~3, we present the spectra of G45.12+0.13, and identify spectral features.  In section~4, we discuss the properties of the identified lines, and from these discuss the possible conditions within the nebula.  Finally, in section~5, we present the conclusions of this work. ",
        "conclusions": "We have presented high signal to noise  spectral data covering 0.9--2.5$\\mu$m for the  compact HII region G45.12+0.13.   All of the lines detected by us can be explained by normal recombination and collisional excitation processes. There is no evidence for hot shocked gas (eg.\\ Van Buren et al.\\ 1990) indicative of the bow-shock model, nor for the presence of CO emission features that may have indicated the presence of a remnant accretion disk (Hollenbach et al.\\ 1993).  We used the HI lines present in our data to estimate  the near-infrared extinction law in this source and found a  dependence on wavelength that is somewhat steeper than normally  observed.  After accounting for this extinction we were able to show that our data are also consistent with a dense core, with $n_e > 10^4$cm$^{-3}$.  This value is considerably larger than previous infrared studies (eg.\\ Colgan et al.\\ 1991) have found.  Most of the emission arises from within the central 3$''\\times3''$ region, though our Br$\\gamma$ image reveals considerable extent to the HII region that was not evident from the radio continuum map of Wood \\& Churchwell (1989).   There is tentative evidence that the higher ionisation species in our data trace higher electron densities, typical of a core-halo model for the HII region.  Using an approximation to the true radiative transfer problem, we showed how the observed HeI triplet lines indicated considerable opacity in the HeI $n^3$P--2$^3$S series, consistent with the high derived electron density. After consideration of all the possible effects we showed that helium must be fully ionised within the HII region, and that we could set  both upper and lower limits to the maximum stellar effective temperature of 38000K and 42000K respectively.  We show how these limits can be set independently of a full modelling process, and hope more effort will be put into determining the accuracy of this method theoretically.  Lastly, from a full photoionisation model,  we showed that we can model most of the observed data on G45.13+0.12 with an OB star cluster which has a restricted range of  stellar masses present.  In particular, a model where the upper mass cut-off is set to 50\\Msolar\\ and the lower mass cut-off to 20\\Msolar\\ gives a reasonable fit to the emission line spectrum with the general exception of the neon lines. We agree with Colgan et al.\\ that the asymmetry evident in the source will probably explain the other differences we found between the (spherically symmetric) model and the data."
    },
    "9603/astro-ph9603005_arXiv.txt": {
        "abstract": "Using the IRAM 30-m telescope, we have mapped the $\\lambda$ 1.2\\,mm continuum emission in the edge-on spiral galaxies NGC\\,891, NGC\\,5907 and NGC\\,4565.  Generally, the $\\lambda$ 1.2\\,mm continuum correlates remarkably well with the CO emission; the correlation with \\hi\\ is however different for the observed galaxies: in NGC\\,891, there is no obvious correlation; in NGC\\,5907 the continuum emission is extending a bit further out than the CO and seems to be correlated with \\hi\\ peaks.  In NGC\\,4565, however, the dust emission not only shows a central peak and an inner ring like the CO, but also, like \\hi{}, a weaker, extended plateau. Comparable to the \\hi , the 1.2 mm contours are warped near the NW edge of the galaxy.  The average dust temperature in this galaxy is 18\\,K near the center and 15\\,K in the \\hi\\ plateau. From the 1.2 mm continuum intensity and the \\hi\\ line integrated intensity, we derive a dust absorption cross section per H atom $\\sigma{\\rm _{1.2mm}^H}=5\\times 10^{-27}$ cm$^2$ in the plateau. This value is very close to that predicted for the local diffuse clouds. ",
        "introduction": "Cold dust represents most of the interstellar dust in normal galaxies and may be used as a tracer of both molecular and atomic gas (see Cox \\& Mezger 1989 and references therein).  Gu\\'elin et al.\\ (1993, 1995) mapped the \\mm continuum emission of two nearby spirals, NGC\\,891 and M\\,51. This emission there was found to correlate tightly with CO and poorly with \\hi{}. It was not even clearly detected beyond NGC\\,891's molecular `ring', in a region where \\hi\\ emission is still strong. The mean dust temperatures derived from the \\mm and FIR flux densities were found to be $\\leq 20$\\,K.  In order to further study the properties of the ISM, we observed two more edge-on galaxies of similar type, NGC\\,4565 (Neininger et al.\\ 1996) and NGC\\,5907 (Dumke et al.\\ 1996). In particular, NGC\\,4565 was chosen because of its weak CO emission: the dust emissivity per H atom is on the average 2 -- 4 times larger for the molecular clouds than for the \\hi\\ clouds. Thus, the dust associated with the atomic gas becomes predominant only when the H$_2$ column density becomes very small -- which is the case in the outer parts of NGC\\,4565. The edge-on geometry ($i \\ge 85^{\\circ}$) ensures long lines of sight in the disk which helps to detect weak components of the ISM.  \\begin{figure} \\vspace{-2.2cm} \\hbox{ \\hspace{-2mm}\\psfig{file=n4565bas.ps,height=95mm} \\hspace{-5mm}\\psfig{file=boloplot.ps,height=87mm,rheight=94mm} } \\caption{a) The \\mm continuum emission of NGC\\,4565 smoothed to a resolution of $20''$; contour levels are 3, 6, ... 21 mJy/beam. b) The brightness distributions of the integrated line intensities of atomic and molecular gas along the major axis of NGC\\,4565 together with the \\mm intensity and the 21-cm continuum.} \\end{figure} ",
        "conclusions": ""
    },
    "9603/hep-ph9603443_arXiv.txt": {
        "abstract": "A new mechanism of supernova shock revival is proposed, which involves resonant spin--flavor precession of neutrinos with a transition magnetic moment in the magnetic field of the supernova. The mechanism can be operative in supernovae for transition magnetic moments as small as $10^{-14}\\mu_B$ provided the neutrino mass squared difference is in the range $\\Delta m^2 \\sim (3 \\;{\\rm eV})^2-(600 \\;{\\rm eV})^2$. It is shown that this mechanism can increase the neutrino--induced shock reheating energy by about 60\\%. ",
        "introduction": "One of the most important problems in the theory of supernova explosions is to understand the physical mechanisms which eventually expel the outer mantle delivering the right amount of energy. Although the main ideas involved in the theory have received major confirmation after the detection of the neutrinos from SN1987A, the problem of accelerating the outward going shock wave which forms after core collapse is still unresolved. For many years, when all computer calculations were essentially done in one dimension, the majority of computations were unsuccessful since the shock would travel for about $300-500$ km and then stall, after losing energy by dissociating heavy nuclei in the envelope into nucleons. Many proposals have been made to solve this problem, starting from including general relativistic corrections, different equations of state, better neutrino transport description and new neutrino physics. Recent numerical calculations \\cite{Her94,Bur95,Janka} have shown that as one moves to more than one dimension one can easily get convective instabilities driven by neutrino heating which are very effective and fast in reheating the material behind the shock. The idea that convection would help the explosion is not new. It has been in the literature for some time, see, e.g.,  \\cite{Epst79} (see \\cite{Bur95} for an extensive list of references). The idea that multidimensional calculations might help is also not new. Several attempts have been made to model supernova explosion in more than one dimension; probably they failed because they did not contain the right combinations of other aspects of the physics involved. Although these most recent calculations make an important contribution to the subject we still believe that much work should be done to understand fully the mechanisms involved. The proposed convective instability relies heavily on the details of neutrino interactions with matter which control the energy transport process. The details of this process are still controversial, especially when the matter is at high density and is not spherically symmetrical. The question whether convective instabilities can revive the shock and lead to a successful supernova explosion is therefore still far from settled, and any new mechanism which could contribute towards the shock energy would be very welcome.  Recently there has been discussion of whether massive neutrinos, which seem to be necessary to reconcile the solar neutrino experiments with the standard model of the sun, might also help in accelerating the shock. Matter--enhanced neutrino oscillations (MSW effect, \\cite{Wolf78,Wolf79,MS85}) in supernovae might play an important role in reviving the shock after core collapse by increasing the amount of energy that neutrinos deposit behind the shock \\cite{MS86,Ful87,Ful92}. The idea is based on the fact that in the region between the neutrinosphere and the position of the stalled shock the matter density is such that flavor transformation of $\\nu_\\mu$ or $\\nu_\\tau$ into $\\nu_e$ is resonant for masses of the heavier neutrinos in the range $10-100$ eV; this transformation can be efficient even if the vacuum neutrino mixing angle is quite small, $\\theta\\aprge 10^{-4}$. Since the average energy of $\\nu_\\mu$'s and $\\nu_\\tau$'s at the neutrinosphere is about $20$ MeV whereas that of $\\nu_e$'s is about $10$ MeV, the electron neutrinos emerging as a result of the $\\nu_{\\mu} (\\nu_{\\tau})\\to \\nu_e$ transformation would have twice as high an energy as the originally emitted ones, and this extra energy would be available for heating the matter behind the shock. Electron neutrinos interact with matter with a larger cross section than muon or tauon neutrinos since they have charged--current interactions with matter in addition to the neutral--current ones. Therefore the $\\nu_e$'s produced in the $\\nu_{\\mu}(\\nu_{\\tau})\\to \\nu_e$ transformation will more efficiently deposit energy behind the shock. Fuller et al. \\cite{Ful92} have shown that the net effect is a $\\sim 60\\%$ increase in the supernova explosion energy.  In this paper, we show that a similar result can be obtained in the framework of the resonant spin--flavor precession mechanism. In this case one has to assume that the neutrino has a nonzero transition magnetic moment $\\mu$ by which it interacts with a magnetic field. Since we know that after a supernova explosion a pulsar is, in many cases, left over, the important role played by the magnetic field during and after core collapse is not in question. The strong magnetic field and the high density make the environment between the neutrinosphere and the position of the stalled shock suitable for spin--flavor conversion due to transition magnetic moments of neutrinos.  Very much as in the case of the MSW effect, spin--flavor precession due to a transition magnetic moment of neutrinos, in which neutrino helicity and flavor are rotated simultaneously \\cite{ShVa}, can be resonantly enhanced in matter \\cite{Akhm88,LM88}. This effect can explain the observed deficit of solar neutrinos with respect to the predictions of the standard solar model \\cite{Akhm88,LM88,ALP93,LNPul,KraOT,ALP95}. That requires the neutrino transition magnetic moment to be of the order of $\\mu\\approx 10^{-11}\\mu_B$ ($\\mu_B=e/2m_e$ is the electron Bohr magneton) provided the strength of the magnetic field near the bottom of the convective zone of the sun is of the order of a few tens of kG. The indicated value of the transition magnetic moment is to be compared with recent astrophysical upper bounds derived from the limits on the energy loss rates of white dwarfs ($10^{-11}\\mu_B$, ref. \\cite{Blin}) and helium stars ($3\\times 10^{-12}\\mu_B$, ref. \\cite{Raff90}, and $10^{-12}\\mu_B$, ref. \\cite{CastDInn}). The latter two values imply that an order of magnitude stronger magnetic field might be necessary to account for the solar neutrino problem in the framework of the neutrino magnetic moment scenario.  In the present paper we show that the spin--flavor precession of neutrinos may play an important role in supernova dynamics even if neutrino transition magnetic moments are far below the present astrophysical upper limits. In particular, it can be resonantly enhanced in the region between the neutrinosphere and the position of the shock for typical values of $\\mu\\approx 10^{-14}\\mu_B$, magnetic field strengths of $B\\approx 10^{12}-10^{15}$ G and neutrino masses which lie in the range $\\sim (3-600)$ eV. This strength of the magnetic field is natural in the context of supernovae if the explosion does give rise to a pulsar, the value of $\\mu$ is consistent with the prediction of the decaying neutrino hypothesis \\cite{Sciama1993book}, and the range of neutrino masses is the one which is relevant for cosmology and the decaying neutrino theory.  The idea that neutrino magnetic moments can play an important role in supernova dynamics was first put forward by Dar \\cite{Dar} in the context of the usual Dirac neutrino magnetic moments and transitions of active left-handed neutrinos into sterile right-handed ones. Our mechanism, based on the spin--flavor precession of neutrinos due to their Majorana-like transition magnetic moments, is different from Dar's. Resonant spin--flavor precession (RSFP) of neutrinos in type II supernovae has been studied earlier \\cite{AB,APS}; however the main goal of those papers was to explore possible consequences for the neutrino signal from supernovae, and no implications of this effect for supernova dynamics were discussed.  The plan of the paper is as follows. We start by reviewing the main features of RSFP in Sec. 2. In Sec. 3 we discuss the RSFP in supernovae, and in Sec. 4 consider the implications of this effect for supernova shock reheating. Sec. 5 contains our conclusions. ",
        "conclusions": "We have shown that resonant spin--flavor precession of neutrinos due to interaction of their transition magnetic moments with the strong magnetic fields inside supernovae may increase the energy deposited by neutrinos in the matter behind the shock by about 60\\% and thus help to re-accelerate it. For the process to be efficient in the supernova environment, the heavier neutrino mass should be in the range $\\sim (3 - 600)$ eV, and the transition magnetic moment $\\mu$ should be of the order of $10^{-14}\\mu_B$ provided that the magnetic field strength at the resonance position is of the order of $10^{12}-10^{15}$ G. All these values of the parameters are consistent with the available laboratory and astrophysical constraints on neutrino properties as well as our present ideas about supernova magnetic fields. In fact, the neutrinos with the masses and magnetic moments in the above range may be very interesting for cosmology and for the decaying neutrino theory \\cite{Sciama1993book}.  Our simple consideration gave only a rough estimate of the supernova explosion energy increase due to the RSFP conversion; whether or not this effect is sufficient to revive the shock leading to a successful supernova explosion can only be decided on the basis of a full-scale supernova dynamics calculation with the RSFP transition included, which goes beyond the scope of the present paper. However, our estimates show that the RSFP--induced neutrino conversion can result in quite a sizable increase of the supernova explosion energy, and we believe that this effect deserves further investigation.  E.A. is grateful to Alexei Smirnov for useful discussions and to George Fuller and Hiroshi Nunokawa for correspondence. S.T.P. would like to thank George Fuller and Yong Qian for discussions of the physics of supernovae. E.A. is grateful to ICTP and SISSA, where part of this work has been done, for hospitality and support. The work of A.L. and D.W.S. has been supported by MURST."
    },
    "9603/astro-ph9603030_arXiv.txt": {
        "abstract": "We present soft X-ray spectra of 74 BL Lacertae objects observed with the {~PSPC} detector on board of the {~ROSAT} satellite. The sample contains all BL Lac objects detected during the pointed observation phase as a target or serendipitously. We have investigated the soft X-ray and broad band spectral properties and discuss the consequences for the X-ray emission processes. For the first time a clear dependence of the  X-ray spectral steepness on the radio to X-ray spectral energy distribution is found: \\arx and \\ax are {\\em correlated} in the X-ray selected  (XBL) subsample and {\\em anticorrelated} in the radio selected (RBL) subsample. The objects with intermediate \\arx values thus do have the steepest soft X-ray spectra. Simulated {~PSPC} spectra based on a set of simple two component multifrequency spectra are in good agreement with the measurements and suggest a broad range of synchrotron cutoff energies.  We have calculated synchrotron self-Compton beaming factors for a subsample of radio bright objects and find a correlation of the beaming factors \\del with \\arx and \\ax. The most extreme RBL objects are very similar to flat spectrum radio quasars in all their broad band and X-ray properties. ",
        "introduction": "BL Lac objects are active galactic nuclei which by definition  do not have strong emission lines, are highly variable, and show strong polarization in the radio to optical emission. Commonly the properties of BL Lacs are explained by the concept that the emission in all spectral bands is dominated by  relativistic jets. Relativistic electrons emit synchrotron radiation and may scatter up either synchrotron photons (synchrotron self-Compton emission) or photons from other regions (e.g. from an accretion disc) to higher energies.  Originally most BL Lacs have been found as counterparts of flat spectrum radio sources (radio selected BL Lacs, RBLs), but an increasing number of less radio bright objects were discovered by the identification of X-ray sources (XBLs). Due to the lack of spectral features in the optical no complete optically selected samples exist. The total number  of BL Lac objects in the catalogue of \\vv (\\cite{veron}) is less than 200. A new catalogue consisting of 233 sources is being published by Padovani \\& Giommi (\\cite{PG95b}). The ongoing search for new BL Lacs from {~ROSAT} sources (Kock et al. \\cite{kock}, Nass et al. \\cite{nass}) will significantly increase the number of known objects.  Recently BL Lacs have gained great interest due to the detection of several objects in  high energy $\\gamma$-rays by EGRET on the Compton Gamma Ray Observatory (von Montigny et al. \\cite{montigny}).  There is some evidence that RBLs and XBLs form distinct subclasses as they show a bimodal distribution in the plane of the broad band spectral indices \\aro versus \\aox (e.g. Giommi et al. \\cite{giommi}). Regardless of the discovery waveband we will use these terms for a distinction of the subclasses based on the spectral energy distribution: RBL is used for {\\it radio} bright objects having \\arx$>$0.75 and XBL for {\\it X-ray} bright objects with \\arx$<$0.75.    There have been several attempts to explain the physical differences of the XBL and RBL objects and their relation to flat spectrum radio quasars ({~FSRQs}). The concept that parts of the continuum emission from radio loud AGN arises from jets of radiogalaxies more or less aligned with the line of sight is widely accepted.  Based on the study of number count relations and  luminosity functions several authors (e.g. Padovani \\& Urry \\cite{PU}) proposed that BL Lacs are the beamed subpopulation of FR I galaxies with RBLs having higher beaming factors than XBLs. Ghisellini \\& Maraschi (\\cite{GM}) discussed an ``accelerating jet'' model with lower bulk Lorentz factors $\\Gamma$ in the X-ray  emitting regions resulting in broader beaming cones for the X-ray emission and narrow radio cones. Celotti et al. (\\cite{celotti}) developed a ``wide jet'' model with geometrically wider  opening angles in the inner, X-ray emitting, parts of the jet. Assuming that RBLs have smaller viewing angles than XBLs, both models are able to explain both the relative numbers and different spectral energy distributions of XBLs and RBLs with an intrinsically uniform population of objects. Maraschi \\& Rovetti (\\cite{maraschi}) have extended these considerations on {~FSRQs} and propose that all radio loud AGN only essentially differ in viewing angle and intrinsic power of the central engine.  An alternative approach to explain the differences between XBLs and RBLs was made by Padovani \\& Giommi  (\\cite{PG95a}) with a ``different energy cutoff'' hypothesis. They argue that both types form a uniform class of objects spanning a wide range in the intrinsic energy distribution caused by  different cutoff frequencies of the synchrotron component.  The  X-ray spectra of RBLs in average were found to be significantly steeper than the  spectra of (higher redshifted) {~FSRQs} in {\\it Einstein} IPC (Worrall \\& Wilkes \\cite{ww}) and {~ROSAT} {~PSPC} (Brunner et al. \\cite{brunner}) investigations. Furthermore, the X-ray spectral indices of BL Lacs showed a broad distribution in both investigations. The mean X-ray spectra of XBLs and RBLs were not found to be significantly different in  {\\it Einstein} IPC (Worrall \\& Wilkes \\cite{ww}), EXOSAT ME + LE (Sambruna et al. \\cite{sambruna}), and {~ROSAT} {~PSPC} (Lamer et al. \\cite{lamer}) observations. Ciliegi et al. (\\cite{ciliegi}) found significant steepening of the mean spectrum of XBLs between the soft (0.2--4 keV) and medium (2--10 keV) energy X-ray band.  The {~ROSAT} data archive presently comprises the largest X-ray database for BL Lac objects collected by a single instrument. In this paper we present the analysis of X-ray and broad band spectra of 74 BL Lac objects with the X-ray data obtained from pointed {~ROSAT} observations. We find a strong interdependence of the X-ray spectral index \\ax with the radio to X-ray energy index \\arx which we interprete as the signature of two spectral components intersecting each other at different frequencies. ",
        "conclusions": "We find that the broad distribution of spectral slopes in the soft X-ray spectra of BL Lacs is due to a strong dependence of \\ax on the broad band spectral index \\arx. Objects with extreme values of \\arx exhibit flat X-ray spectra, whereas intermediate objects have steeper spectra. The symmetry of this dependence prevented the detection of significant differences between the mean X-ray spectra of XBLs and RBLs in previous investigations (Worrall \\& Wilkes \\cite{ww}, Sambruna et al. \\cite{sambruna}, Lamer et al. \\cite{lamer}).  By comparison with simulated {~PSPC} spectra we showed that a two component model with a hard power law ($\\alpha=0.7$) component and a steepening soft component is appropriate to explain the observed spectra. The frequency where the components intersect each other is below the soft X-ray band for extreme RBLs and crosses the energy band of the {~ROSAT} {~PSPC} with declining \\arx. In the framework of the SSC models this means that Compton emission causes the flat X-ray spectra of extreme RBLs, whereas the likewise flat X-ray spectra of extreme XBLs are due to synchrotron emission. The steep X-ray spectra of objects with intermediate spectral energy distribution ($0.6<$\\arx$<0.8$) represent the first direct evidence of the synchrotron high energy cutoff.  Padovani \\& Giommi (\\cite{PG95a}) explain the different spectral energy distributions of XBLs and RBLs by different energy cutoffs of the synchrotron spectra. They postulate the cutoff energy being intrinsic properties of the sources without discussing the physics of the emission processes. This scenario also is the most straightforward explanation for our findings, including the correlation of \\arx and \\ax for the XBL subsample. The wide range in synchrotron cutoff energies, and consequently the cutoff in  the energy spectrum of the relativistic electrons, has to be explained. The SSC cooling of the jet electrons may be more efficient in the more powerful jets of {~FSRQs} and RBLs than in the jets of XBLs. Ghisellini \\& Maraschi (\\cite{GM94}) proposed a more rapid cooling of jet electrons in {~FSRQs} by external UV photons compared to jets of BL Lac objects. It is conceivable that RBLs are intermediate objects between XBLs and {~FSRQs} regarding the ambient photon density.  The more physically motivated beaming models, such as the ``accelerating jet'' model (Ghisellini \\& Maraschi \\cite{GM}) and the ``wide jet'' model (Celotti et al. \\cite{celotti}) do not as naturally satisfy our data. The crossover frequencies of soft and hard components in the spectra calculated by Ghisellini and Maraschi (\\cite{GM}) do not span a sufficient range and do not move across the soft X-ray range when tilting the viewing angle, as required by the new data. Nevertheless, further tuning of the free parameters may provide spectra which are in accordance with the measurements. The {~ROSAT} spectra therefore are suitable to constrain the beaming models."
    },
    "9603/astro-ph9603083_arXiv.txt": {
        "abstract": "We propose that FU Orionis outbursts may play an important role in maintaining the slow rotation of classical disk-accreting T Tauri stars. Current estimates for the frequency and duration of FU Orionis outbursts and the mass accretion rates of T Tauri and FU Orionis stars suggest that more mass may be accreted during the outbursts than during the T Tauri phases. If this is the case, then the outbursts should also dominate the accretion of angular momentum. During the outbursts, the accretion rate is so high that the magnetic field of the star should not disrupt the disk, and the disk will extend all the way in to the stellar surface. Standard thin disk models then predict that the star should accrete large amounts of angular momentum, which will produce a secular spinup of the star.  We present boundary layer solutions for FU Orionis parameters which show that the rotation rate of the accreting material reaches a maximum value which is far less than the Keplerian rotation rate at the stellar surface. This is due to the importance of radial pressure support at these high mass accretion rates. As a result, the rate at which the star accretes angular momentum from the disk drops rapidly as the stellar rotation rate increases. In fact, the angular momentum accretion rate drops below zero for stellar rotation rates which are always substantially below breakup, but depend on the mass accretion rate and on the adopted definition of the stellar radius. When the angular momentum accretion rate drops close to zero, the star will stop spinning up. Faster stellar rotation rates will produce negative angular momentum accretion which will spin the star down. Therefore, FU Orionis outbursts can keep the stellar rotation rate close to some equilibrium value for which the angular momentum accretion rate is small. We show that this equilibrium rotation rate may be similar to the observed rotation rates of T Tauri stars; thus we propose that FU Orionis outbursts may be responsible for the observed slow rotation of T Tauri stars. This mechanism is independent of whether the disk is disrupted by the stellar magnetic field during the T Tauri phase. ",
        "introduction": "The FU Orionis stars are a class of accreting pre-main sequence stars which are observed to be T Tauri stars undergoing large outbursts (see Hartmann, Kenyon, \\& Hartigan 1993 for a review). During the outbursts, these systems brighten by $\\sim 5-6$ magnitudes over a period of $\\sim 1-10$ years. They reach luminosities of a few hundred $\\lsun$, with spectra indicating maximum temperatures $\\sim 6500-7000$ K. The outbursts last for $\\sim 100$ years, although this number is poorly known, since no outbursting systems have been observed to return to their pre-outburst state. Hartmann \\& Kenyon (1985, 1987) showed that FU Orionis stars contain accretion disks by demonstrating that spectra of these systems contain double-peaked absorption lines. Simple steady accretion disk models of two FU Orionis systems, FU Orionis and V1057 Cygni, were constructed by Kenyon, Hartmann, \\& Hewett (1988, hereafter KHH). The models which best fitted the observed broad-band spectral energy distributions of these systems had accretion rates $\\dot M \\sim 10^{-4} \\msyr$, and central stars with masses $M_* \\sim 0.3-1 \\msun$ and radii $R_* \\sim 4 \\rsun$. The presence of a disk in these systems suggests that their outbursts may result from a disk instability similar to those believed to cause similar outbursts in dwarf novae (Clarke, Lin, \\& Papaloizou 1989; Clarke, Lin, \\& Pringle 1990; Bell \\& Lin 1994; Bell \\et 1995).  Disk accretion also has important implications for the rotational evolution of the accreting star. In the standard picture of disk accretion (Shakura \\& Sunyaev 1973), the star accretes angular momentum from the disk. The angular momentum accretion rate is simply the Keplerian specific angular momentum at the stellar surface multiplied by the mass accretion rate. This quickly spins up the accreting star. Thus, in the standard picture, if T Tauri or FU Orionis stars have accreted even a small fraction of their mass through a disk, they should be spinning rapidly.  The rotation rates of FU Orionis stars are not well known, since the accretion luminosity is so high that it overwhelms the stellar luminosity. The rotation rates of T Tauri stars are more easily observable and have been studied in some detail (Bouvier \\et 1993, 1995). They tend to be about a tenth of breakup speed, on average, with rotation periods usually around 6--8 days. Since the angular momentum added by disk accretion should rapidly spin T Tauri stars up to breakup speed, some authors have hypothesized that the spinup of T Tauri stars is controlled by the interaction between the disk and a strong stellar magnetic field (K\\\"onigl 1991; Cameron \\& Campbell 1993; Hartmann 1994; Shu \\et 1994).  This mechanism for maintaining the spin rate of a star near some equilibrium value was first explored in detail in order to explain the spin behavior of some accreting neutron stars (Ghosh, Lamb, \\& Pethick 1977; Ghosh \\& Lamb 1979a,b). In this model, the stellar magnetic field disrupts the disk at some distance from the stellar surface. The location of this disruption radius is determined by the strength of the magnetic field and the mass accretion rate; a stronger magnetic field disrupts the disk farther away from the star, while a larger mass accretion rate moves the disruption radius closer to the stellar surface. The Keplerian rotation speed at the disruption radius determines approximately how fast the star will rotate. In order to produce the rotation speeds of most T Tauri stars, the disk would have to be disrupted around 5 stellar radii from the stellar surface, and for the accretion rates commonly ascribed to T Tauri stars, this would require a stellar magnetic field $\\sim$ 1 kG (K\\\"onigl 1991).  The magnetic hypothesis for explaining the slow rotation of T Tauri stars does not include the effects of angular momentum accretion during FU Orionis outbursts. Event statistics of these outbursts, in which the number of observed events is compared to the number of young stars in the observable volume, suggest that each star probably undergoes multiple outbursts (Herbig 1977; Hartmann \\& Kenyon 1985). Estimates for the time between outbursts vary widely from $\\sim 10^3-10^5$ years (e.g. Bell \\& Lin 1994; Kenyon 1995). If the accretion rate during this period is the typical T Tauri rate of $\\sim 10^{-7} \\msyr$, then $10^{-4}-10^{-2} \\msun$ will be accreted during the T Tauri phase. During the intervening FU Orionis outbursts, the accretion rate is $\\sim 10^{-4} \\msyr$ for $\\sim 100$ years, so $\\sim 10^{-2} \\msun$ should be accreted during the outburst. This suggests that as an accreting pre-main sequence star evolves through multiple cycles of FU Orionis outbursts and T Tauri quiescent phases, most of the mass accreted by the star may be added during the outbursts. If this is the case, then we should expect that the outbursts will also dominate the accretion of angular momentum.  During the FU Orionis outbursts, the stellar magnetic field should have little effect on the accretion disk. As noted above, in order to produce the observed rotation rates of T Tauri stars, the field should disrupt the disk at a radius of about 5 stellar radii. When such a star experiences an FU Orionis outburst, the mass accretion rate increases by a factor of order 1000. This should be sufficient to overwhelm the magnetic field, and allow the disk to reach all the way in to the stellar surface. In the standard picture of accretion disks, since the star accretes a large amount of mass during the FU Orionis outburst, it should also accrete a large amount of angular momentum, which will spin the star up substantially. This will in turn require that the star loses large amounts of angular momentum during the T Tauri phase in order to stay at the observed slow rotation rate. If the star accretes less mass in the T Tauri phase than it does in the FU Orionis outburst, it will be unable to lose sufficient angular momentum to spin back down, and over time, the star will continue to spin up toward breakup speed.  If the disk reaches the stellar surface, and the star is rotating slowly, a viscous boundary layer forms between the rapidly rotating disk material and the star. The boundary layer controls the flow of angular momentum between the disk and the star. Also, if the central star is not rotating, the boundary layer should produce as much luminosity as the disk. Clearly, if we wish to understand the angular momentum transfer and to compare disk models to observed spectra of FU Orionis systems, it is essential to include the boundary layer in our disk models.  Our initial investigation of the boundary layer structure examined the question of whether a star continues to accrete mass and angular momentum when it has been spun up to breakup speed (Popham \\& Narayan 1991, hereafter PN91). We found that as the star approaches the breakup rotation rate, the amount of angular momentum accreted per unit of mass accreted drops dramatically, and can even become negative, so that the star continues to accrete mass while losing angular momentum to the disk. Thus, there is some equilibrium rotation rate close to breakup for which the star accretes no angular momentum, so that it can continue to accrete mass without spinning up or down. This study used a very simple disk model with a polytropic relation between the disk pressure and density, so it was not clear that similar conclusions would hold when a more realistic treatment of the disk was used.  More recently, we have explored the structure of the boundary layer region for several types of accreting stars, including T Tauri and FU Orionis stars (Popham \\et 1993, hereafter PNHK) and cataclysmic variables (Narayan \\& Popham 1993; Popham \\& Narayan 1995), using a more realistic model which includes the energy balance and radiative transfer.  In this paper, we examine the angular momentum transfer between the star and the disk for boundary layer solutions calculated for parameters corresponding to FU Orionis systems. The boundary layer model we use is similar to the one used in our studies of cataclysmic variables (Narayan \\& Popham 1993; Popham \\& Narayan 1995) and in our previous study of pre-main sequence stars (PNHK). We find that many of the conclusions reached by PN91 continue to apply when our more detailed model is applied to FU Orionis systems. When the stellar rotation rate is plotted against the height of the disk at the stellar surface, the solutions fall on multiple solution branches. We find that the angular momentum accretion rate drops as the stellar rotation rate increases, and that solutions exist for a wide range of angular momentum accretion rates, including negative values. There is an equilibrium stellar rotation rate for which the star can accrete mass without spinning up or down. The major difference from the PN91 results is that for these FU Orionis-type solutions, the equilibrium stellar rotation rate is well below breakup.  On the basis of these solutions, we propose that FU Orionis outbursts play an important role in the spin evolution of accreting pre-main sequence stars. The large amount of mass accreted during an outburst can change the stellar rotation rate substantially. If more mass is accreted during the FU Orionis outburst phases than during the intervening T Tauri phases, then the stellar rotation rate will stay close to the equilibrium rate established during the FU Orionis phase. Although this equilibrium rate is rather uncertain, we will show that for reasonable choices of parameters, it is comparable to the observed rotation rates of T Tauri stars. Thus, angular momentum transfer during the FU Orionis outbursts can explain the slow rotation of T Tauri stars without invoking strong stellar magnetic fields. If T Tauri stars have weak fields which allow the accretion disk to extend down to the surface of the star and spin up the star, negative angular momentum accretion during the FU Orionis outburst phases can keep the star spinning slowly. Alternatively, if T Tauri stars have strong magnetic fields which disrupt the disk, the FU Orionis phases will still dominate the spin evolution as long as they dominate the mass accretion.  In \\S 2, we discuss the model used to calculate our boundary layer and disk solutions for FU Orionis parameters, and the input parameters for our solutions. We describe a typical solution in detail in \\S 3. We demonstrate the presence of multiple solution branches, and show how the character of the solutions varies along those branches. We also show a solution which accretes mass without accreting angular momentum. Finally, we show how variations in the mass accretion rate affect the solutions and solution branches. In \\S 4, we describe the choice of an additional condition which allows us to find a relation between the stellar rotation rate and the angular momentum accretion rate. We discuss the spin evolution of T Tauri/FU Orionis systems and the implications of our results in \\S 5. ",
        "conclusions": "\\subsection{Spin Evolution of T Tauri/FU Orionis Stars}  We have found that boundary layer solutions for parameters corresponding to FU Orionis outbursts have small and negative angular momentum accretion rates for fairly low stellar rotation rates. This suggests that angular momentum transfer during the FU Orionis outbursts can keep the rotation rate of T Tauri/FU Orionis stars near to an equilibrium rotation rate $\\omseqf \\simeq \\oms(j=0)$. If the star is rotating faster than $\\omseqf$ when it experiences an FU Orionis outburst, it will lose angular momentum in the outburst phase and spin back down to $\\omseqf$. If, on the other hand, the star is rotating slowly when an outburst occurs, it will have $j \\sim 1$ in the outburst phase and will spin up toward $\\omseqf$. We have found that for accretion rates $\\md = 5 \\times 10^{-5} - 10^{-4} \\msyr$, $j=0$ for spin rates $\\oms(j=0) \\simeq 0.66-3.04 \\xfs \\simeq \\omseqf$, corresponding to rotation periods of $2.4-11$ days. These are similar to the observed rotation periods of T Tauri stars (Bouvier \\et 1993, 1995); however, $\\oms(j=0)$ is sensitive to the value assumed for $\\vs$, as discussed below.  During the T Tauri phase, the stellar rotation rate will tend to move toward a different equilibrium value $\\omseqt$. If the stellar magnetic field is too weak to disrupt the disk, then the star will spin up with $j \\sim 1$ until it reaches $\\omseqt \\simeq \\omks$. If the stellar field does disrupt the disk, then $\\omseqt$ will be determined by the radius where the disruption occurs, and may be smaller or larger than $\\omseqf$. Therefore $j$ may be positive or negative, with a magnitude $|j| ~\\sles~ (R_d/\\rs)^{1/2}$ (Ghosh, Lamb, \\& Pethick 1977), where $R_d$ is the radius at which the field disrupts the disk. Models of magnetically disrupted disks in T Tauri stars suggest $R_d \\sim 4-5 \\rs$, which would give $j \\sim \\pm 2$.  The spin evolution of a T Tauri/FU Orionis star will depend on $\\omseqf$, $\\omseqt$, and on the amounts of mass accreted in typical FU Orionis and T Tauri phases, $\\dmf$ and $\\dmt$.  If $\\dmf > \\dmt$, as suggested by current estimates of the frequency and duration of FU Orionis outbursts, the stellar rotation rate $\\oms$ will move toward $\\omseqt$ during each T Tauri phase, but will return to $\\omseqf$ during each FU Orionis phase.  How large will the excursions in $\\oms$ be? A rigidly rotating star has angular momentum $J = I \\oms$, where $I$ is the star's moment of inertia, so $\\dot J = \\dot I \\oms + I \\dot \\oms$. If we assume that the moment of inertia changes slowly, so that $\\dot J \\simeq I \\dot \\oms$, then the star spins up and down at a rate $\\dot \\oms / \\omks = (j/k)(\\md / M)$, where we have defined $I \\equiv k M \\rs^2$. Since $k$ is substantially less than 1, the increase in $\\oms$ as a fraction of the breakup rotation rate $\\omks$ is several times larger than the fractional increase in the stellar mass $M$. If we assume $j=1$, $k=0.2$, then the accretion of $\\Delta M = 0.01 \\msun$ onto a $0.5 \\msun$ star would increase $\\oms$ by $0.1 \\omks$. If the typical change in $\\oms$ during the T Tauri phase is larger than the difference between the equilibrium spin rates $\\omseqf$ and $\\omseqt$, then $\\oms$ will reach $\\omseqt$ during each T Tauri phase, and remain there for the duration of the phase.  It is important to note that even though the FU Orionis phases may produce more mass accretion than the T Tauri phases, $\\dmf > \\dmt$, a T Tauri/FU Orionis star will still spend most of its lifetime with $\\oms$ moving away from $\\omseqf$, due to the short duration of the FU Orionis phases. This means that the mean value of $\\oms$ during the T Tauri phase will be offset somewhat from $\\omseqf$. The sign of this offset will depend on whether the star spins up or down during the T Tauri phase, i.e. on whether $\\omseqt$ is smaller or larger than $\\omseqf$. Also, the return to the T Tauri state after an FU Orionis outburst may be gradual, with the mass accretion rate declining slowly over a period of a few decades. At mass accretion rates intermediate between T Tauri and FU Orionis rates, the star should spin up, since the magnetic field may still be too weak to disrupt the disk, while the mechanism described in this paper will only operate at a large value of the equilibrium rotation rate. Despite these deviations, the high accretion rate of mass and angular momentum in the FU Orionis phase will return $\\oms$ to $\\omseqf$ quite rapidly, and as long as $\\dmf > \\dmt$, $\\oms$ will continue to be close to $\\omseqf$.  Note also that in general the equilibrium spin rate $\\omseq$ will differ slightly from the value of $\\oms$ which gives $j=0$.  Since $\\dot J = I \\dot \\oms + \\dot I \\oms$, we will have $\\dot \\oms = 0$ for $\\dot J = \\dot I \\oms$. This simply means that in equilibrium, $\\dot J$ must be sufficient to compensate for changes in the star's moment of inertia. If the moment of inertia increases in response to accretion, $j > 0$ at $\\omseq$. If the moment of inertia decreases, as it might in an accreting white dwarf, then $j(\\omseq) < 0$. For instance, if the star's radius of gyration stays constant, so that $\\dot I / I = \\md / M$, then $j(\\omseq) = k \\omseq / \\omks$, where again $I \\equiv k M \\rs^2$. If $\\omseq / \\omks$ is in the range 0.2-0.4 as suggested by our results, and $k \\simeq 0.2$, then $j(\\omseq) \\simeq 0.04-0.08$, and $\\omseq$ will be only slightly smaller than $\\oms(j=0)$.  \\subsection{Solutions with Negative Angular Momentum Accretion}  One of the most important results of this paper is that it establishes the existence of disk and boundary layer solutions with small and negative angular momentum accretion rates. In these solutions, the central star accretes mass but loses angular momentum. This is quite different from the standard thin disk formulation, in which the angular momentum accretion rate is always $\\dot J = \\md \\omks \\rs^2$, or $j=1$ in our units. The key to understanding this result is that the $\\O$ profiles of our solutions can be very different from the purely Keplerian $\\O$ assumed in the thin disk case. Most importantly, $\\O$ does not reach a maximum and then drop down to the stellar rotation rate $\\oms$. Instead, $\\O$ continues to increase all the way in to the stellar surface at $R=\\rs$. The fact that $\\O$ is monotonic means that there is no radius $R_{\\O_max}$ at which $d\\O/dR = 0$, where the viscous torque would vanish, and the angular momentum accretion rate would simply be the amount of angular momentum carried in the accreting material at that radius, $\\dot J = \\md \\O_{max} R_{\\O_max}^2$. Since for a steady flow $\\dot J$ must be constant with radius, this would constrain $\\dot J$ to be positive. In our low-$j$ and negative-$j$ solutions, there is no maximum in $\\O$, so that the viscous torque carries angular momentum outward at all radii.  If the torque exceeds the rate at which angular momentum is carried in by the accreting material, the net angular momentum accretion rate will be negative.  These solutions are very similar to the ones we found in an earlier paper (PN91). In that paper, we calculated the structure of the boundary layer using a simplified disk model with a polytropic pressure-density relation. We found that our solutions formed two branches in the $\\oms - \\hs$ plane for large values of $j \\sim 1$. We did not find the lower branch, but seems likely that it exists for these solutions as well. The PN91 solution branches formed sharp corners where they met, unlike the rounded transitions between branches seen in the current solutions. This difference probably arises from the very different disk thicknesses in the two studies; the PN91 solutions generally had $H/R < 0.01$. We were also able to find solutions with small and negative values of $j$; for these values only a single branch was present in the $\\oms - \\hs$ plane. Thus, despite the use of a different disk model, the qualitative behavior of these solutions is very close to that of the solutions described in this paper.  One important question is why the solutions presented here reach negative angular momentum accretion rates when the stellar rotation rate $\\oms$ is still well below the breakup rate $\\omks$. For the PN91 solutions, we defined the stellar radius as the point where $H/R = 0.1$ in the accretion flow, and found that with this definition, $j$ was very close to 1 for most $\\oms$, but dropped precipitously when $\\oms \\simeq 0.915 \\omks$. We have seen that the current solutions reach $j=0$ for $\\oms \\sim 0.2 - 0.5 \\omks$. As discussed above, solutions with small or negative values of $j$ must have $d\\O/dR < 0$ for all $R$. Thus, if $\\O \\simeq \\ok$ close to $\\rs$, the star must be rotating close to breakup, $\\oms \\simeq \\omks$, in order to allow a negative-$j$ solution. If $\\oms$ is much less than $\\omks$, $\\O$ will have to decrease close to $\\rs$, and there will be a maximum in $\\O$, so that $j \\simeq 1$.  In the FU Orionis solutions presented here, two factors combine to produce small and negative values of $j$ when $\\oms \\ll \\omks$. The first is the large radial width of the boundary layer. In the solutions with $j \\sim 1$, $\\O$ peaks fairly far from $\\rs$, so that the radial width of the boundary layer is comparable to $\\rs$. At the radius where $\\O$ peaks, $\\ok$ is substantially smaller than $\\omks$; for instance, at $R = 2 \\rs$, $\\ok \\simeq 0.35 \\omks$. As the star spins up, $\\oms$ only needs to reach this value to make $d\\O/dR < 0$ everywhere and allow $j$ to be small or negative. Second, pressure support plays an important role, and $\\O$ is substantially smaller than $\\ok$. This allows solutions with $d\\O/dR < 0$ everywhere for even smaller values of $\\oms$. This also explains why we reach $j=0$ at smaller $\\oms$ for larger values of $\\md$; as $\\md$ increases, pressure support plays a larger role, the disk becomes thicker, and $\\O$ is a smaller fraction of $\\ok$.  It is also worth noting that the total accretion luminosity increases substantially as $j$ decreases. In general, the accretion luminosity varies according to the expression \\[ L_{acc} \\simeq {G \\ms \\md \\over \\rs} \\left[ 1 - j {\\oms \\over \\omks} + {1 \\over 2} {\\oms^2 \\over \\omks^2} \\right] \\] (Popham \\& Narayan 1995). Thus, for large negative values of $j$, the accretion luminosity can be substantially larger than the standard value $G \\ms \\md / \\rs$. If the star spins up substantially during the T Tauri phase, then when an FU Orionis outbursts occurs, $j$ may reach fairly large negative values. For instance, Fig. 4 shows that $j = -1$ for $\\oms \\simeq 1.83 \\xfs$, which is only about a 6\\% increase in $\\oms/\\omks$ over the point where $j=0$. A 10\\% variation from $\\oms/\\omks \\simeq 0.3$ to 0.4, as discussed above, would give $j \\sim -2$. This would give an accretion luminosity of $\\sim 1.88 G \\ms \\md / \\rs$, which would make the outburst even brighter than one would expect from the increase in $\\md$. This additional luminosity would come from the rotational energy of the star, which is released as the star spins down.  Disk solutions with negative angular momentum accretion rates may be important in other types of accreting systems. In systems such as cataclysmic variables, the accretion rates are generally much smaller than those of FU Orionis systems. This produces a narrower boundary layer, and pressure support plays a smaller role, so we would expect that negative angular momentum accretion rates are only found for $\\oms \\simeq \\omks$. We are currently completing a study of negative-$j$ solutions for disks around cataclysmic variables with $\\md = 10^{-8} \\msyr$ which confirms these expectations (Popham 1995). In some other systems, such as embedded pre-main-sequence stars, or some X-ray binaries, the accretion rates may be large enough to produce solutions like the ones presented here, where negative angular momentum accretion rates are reached when the star is still spinning relatively slowly.  \\subsection{Assumptions}  The results described above depend on a number of assumptions. One very important assumption is that FU Orionis outbursts occur in classical T Tauri stars; i.e., that the T Tauri and FU Orionis phases occur during the same epoch of pre-main sequence stellar evolution. Models of FU Orionis outbursts based on disk instabilities generally assume that the outbursts arise in T Tauri star disks, so that the FU Orionis and T Tauri phases correspond to the high and low states observed in other types of accreting stars, such as dwarf novae. The observational connection between T Tauri and FU Orionis stars is rather tenuous, since only one of the current FU Orionis systems was observed spectroscopically before it went into outburst, and no known FU Orionis systems have yet returned to their pre-outburst levels. The one system for which a spectrum was taken prior to outburst, V1057 Cygni, resembled a typical T Tauri star (Herbig 1977). However, some of the observed characteristics of FU Orionis systems, such as their association with reflection nebulae and outflows, and their large far-infrared excesses, suggest that FU Orionis outbursts may arise in younger embedded sources rather than in classical T Tauri stars (Kenyon 1995). If this is the case, then the mechanism described here should serve to limit the stellar rotation rate during the embedded phase of pre-main-sequence stellar evolution. The rotation rate would then presumably be controlled by the stellar magnetic field during the T Tauri phase.  We must also make some assumption about the location of the stellar radius, which is difficult to define in our model (see \\S 4 for a discussion of this problem). We have chosen a particular value of $\\vs$ to define $\\rs$, in order to provide an example of the variation of $\\hs$ and $j$ as a function of $\\oms$. Unfortunately, the value of $\\oms(j=0)$ is fairly sensitive to the choice of $\\vs$. We have illustrated this in Fig. 4, where we have plotted lines corresponding to $\\vs = -500 \\cm \\pers$ and $\\vs = -2000 \\cm \\pers$. For these values of $\\vs$, $j$ reaches zero at $\\oms \\simeq 0.94$ and $2.19 \\xfs$, respectively. This means that for a factor of 4 variation in $\\vs$, $\\oms(j=0)$ varies by more than a factor of 2. Thus, even when we specify all the parameters of an FU Orionis system, we still have a substantial uncertainty in $\\oms(j=0)$ resulting from the uncertainty in the definition of the stellar radius. As discussed in \\S 4.5, $\\oms(j=0)$ also varies dramatically with $\\md$. Therefore we can only say that for what we believe to be reasonable choices of $\\vs$ and $\\md$, the resulting values of $\\oms(j=0)$ are comparable to T Tauri rotation rates.  Our results are not strongly dependent on the magnetic field of the star, as long as the field is not strong enough to disrupt the accretion flow in the high-$\\md$ FU Orionis phase. This would require a field substantially larger than the $\\sim 1$ kG needed to disrupt the disk during the T Tauri phase. The disruption radius varies as $R_d \\propto \\md^{-2/7} B^{4/7}$ (Pringle \\& Rees 1972), so to maintain the same $R_d$ requires that $B$ vary as $\\md^{1/2}$. Since FU Orionis accretion rates are a factor of $\\sim 1000$ higher than T Tauri rates, $B$ would have to be about 30 times larger to disrupt the disk, which would require $B \\sim 3 \\times 10^4$ G. A field this strong would disrupt T Tauri disks at $R ~\\sgreat~ 20 \\rs$, which can probably be ruled out by observations.  We have used stellar radii of 2,25, 3, and 4 $\\xec$, corresponding to 2.88, 4.31, and 5.75 $\\rsun$, in our calculations. These are larger than the radii generally found for T Tauri stars, which are in the range 1.5--2.5 $\\rsun$. These larger radii appear to be required by observations: smaller stellar radii would produce maximum temperatures in the inner disk which are substantially larger than those observed. This suggests that the accreting star may expand during the FU Orionis outburst.  Prialnik \\& Livio (1985) calculated the evolution of a fully convective 0.2 $\\msun$ star during the accretion of $2.5 \\times 10^{-3} \\msun$ of material at rates ranging from $10^{-1}$ to $10^{-10} \\msyr$. They also varied the fraction $f$ of the accretion energy carried into the star with the accreting material; $f$ ranged from 0.001 to 0.5. They found that for the accretion rates appropriate to FU Orionis outbursts, the star expanded stably if $f \\leq 0.03$. If $f \\geq 0.1$, the star expanded unstably. In our FU Orionis solutions, the central temperature of the disk at $\\rs$ is $\\sim 1-3 \\times 10^5$ K. The virial temperature for these parameters is $\\sim 10^6$ K, so a reasonable fraction $f \\sim 0.1-0.3$ of the accretion energy is advected into the star. This suggests that rapid stellar expansion may result from the high-$\\md$ accretion during an FU Orionis outburst, but it is difficult to know how the mass of the star, the presence of rotation, and the departure from spherical symmetry would affect the Prialnik \\& Livio results. This expansion, and the subsequent contraction during the T Tauri phase, could affect the spin evolution of the star, and other aspects of the star's interaction with the disk, and clearly deserves further study.  Our models also assume that the disk and boundary layer are in a steady state. This is clearly not true for FU Orionis systems, since they are experiencing outbursts. The viscous timescale is $t_{visc} \\simeq R^2/\\nu = \\alpha^{-1} (H/R)^{-2} \\ok^{-1}$. For our solutions with $\\alpha = 0.01$, $H/R \\simeq 0.3$, this gives $t_{visc} \\simeq 1000 \\ok^{-1}$. In the inner disk, we have $\\ok \\simeq {\\rm few} \\xfs$, so $t_{visc} \\simeq {\\rm few} \\times 10^7 ~{\\rm s} \\simeq 1 ~{\\rm yr}$. This is comparable to the rise times of the fastest-rising systems, FU Ori and V1057 Cyg, and much shorter than their decline times. Thus FU Orionis systems should be reasonably well-modeled by steady state solutions for this value of $\\alpha$.  \\subsection{Comparison with Observations}  Recent measurements of periodic photometric variations in T Tauri stars have clearly established that the classical disk-accreting T Tauri stars spin more slowly than weak-line T Tauri stars, which appear to lack accretion disks (Bouvier \\et 1993, 1995; Edwards \\et 1993). This has generally been interpreted to mean that the spinup of T Tauri stars is controlled by a stellar magnetic field strong enough to disrupt the disk. If FU Orionis outbursts regulate the spin evolution of T Tauri stars in the way described in this paper, it would also account for this observation. In both cases, disk accretion maintains the slow rotation speed of classical T Tauri stars. If FU Orionis outbursts cease late in the classical T Tauri stage, then the more rapid rotation of the weak-line T Tauri stars could be produced by accretion at the end of the classical T Tauri stage.  Observations of the spin evolution of T Tauri stars could be used to test the ideas discussed in this paper. The most useful data would be measurements of the spin period before and after an FU Orionis outburst. If the star spins down during an outburst, it would offer clear evidence that outbursts are regulating the rotation rates of T Tauri stars. If the star spins up, the result is less clear. If the amount of spinup is less than expected based on the amount of mass accreted during the outburst, it would suggest that the star is reaching the equilibrium rotation rate $\\omseqf$ at some point during the outburst. In this case the outbursts would still be limiting the rotation rate of the star. If, on the other hand, the star spins up as much as expected, it would suggest that the rotation rate must be limited by processes taking place during the T Tauri phase. One complication in interpreting the observations is the possibility that significant amounts of angular momentum might be carried off in the strong outflows which occur during outbursts.  Observations of variations in the rotation rates of T Tauri stars which have not experienced observed outbursts could also be useful. Over a short timescale, variations in $\\md$ could produce irregular period variations. Nonetheless, if FU Orionis outbursts are controlling the rotation rates of these stars, then their spin periods should vary secularly during the T Tauri phase as the rotation rate $\\oms$ moves from $\\omseqf$ toward $\\omseqt$. If instead the rotation rates of T Tauri stars remain constant over a long timescale, it would suggest that $\\oms$ has stabilized at $\\omseqt$, which would indicate that the spin evolution is being controlled by angular momentum transfer in the T Tauri phase.  The main sources of data which can be used to further constrain the solutions presented in this paper are the spectra of FU Orionis systems. These spectra yield valuable information on the disk temperatures and rotational velocities (Hartmann \\& Kenyon 1985, 1987; KHH). The solutions presented in this paper were selected to have luminosities and temperatures in approximate agreement with those derived from observations. Because our main aim in this paper has been to understand the spinup and angular momentum transfer in FU Orionis systems, we have not attempted to match our solutions with observations in any detail. We defer this to a subsequent paper (Popham \\et 1995), where we plan to compare the spectra and rotational velocities produced by our disk and boundary layer solutions to those observed in FU Orionis systems."
    },
    "9603/astro-ph9603118_arXiv.txt": {
        "abstract": "The entire dataset of the GRANAT/SIGMA observations of Cyg X-1 and 1E1740.7-2942 in 1990--1994 was analyzed in order to search for correlations between primary observational characteristics of the hard X-ray (40--400 keV) emission -- hard X-ray luminosity $L_X$, hardness of the spectrum (quantified in terms of the best-fit thermal bremsstrahlung temperature $kT$) and the {\\em rms} of short-term flux variations.  Two distinct modes of the $kT$ vs. $L_X$ dependence were found for both sources (Fig.1). At low luminosity -- below the level corresponding approximately to the $\\gamma_1$ state of Cyg X-1 (Ling et al. 1987) -- the $kT$ increases as the $L_X$ increases.  Quantitatively it corresponds to increase of the temperature from 70 keV at $\\approx 0.5L_{\\gamma_1}$ to 150 keV at $\\approx 1.2L_{\\gamma_1}$.  Above the luminosity level of $\\approx 1.2L_{\\gamma_1}$ the spectrum hardness is nearly constant ($T\\approx 150$ keV) and does not depend on the luminosity.  In the case of Cyg X-1 (1E1740.7-2942 is not bright enough and is located in crowded Galactic Center region) the correlation of similar kind was found between the spectrum hardness and $rms$ of the short-term flux variations (Fig.2).  The increase of the $kT$, corresponding to the increasing branch on the $kT$ vs. $L_X$ diagram, is accompanied with increase of the $rms$ from $\\la$ few percent level to $\\approx 10-15$\\%. Further increase of the $rms$ is not accompanied with change of the $kT$ and does not correlate with changes in the luminosity. ",
        "introduction": "\\begin{figure} \\psfig{figure=f1.ps,height=9.8truecm} \\caption[]{The hardness of the spectrum versus hard X-ray luminosity (40-200 keV) for Cyg~X-1 (upper panel) and 1E1740-294 (lower panel).  The luminosity for each bin was calculated using corresponding best fit thermal bremsstrahlung spectrum. The broken constant best fit to Cyg~X-1 data is shown in both panels by the dashed lines.} \\end{figure}   In order to search for correlation between various characteristics of the hard X-ray emission from the source the original data were regrouped according to the source intensity in the following way. The entire range of the 40-150 keV flux variations was divided into number of bins of the same width. For each individual dataset the bin number was determined according to its intensity. The mean energy and power density spectra corresponding to each intensity bin  were calculated by averaging  over all individual datasets with intensity falling into the given bin intensity range.  For Cyg X-1 16 intensity bins were chosen covering the 1.9 to 6.9 cnt/sec/cm$^{2}\\times10^{-2}$\\ (0.5-1.8 Crab) intensity range. The regrouping procedure was applied to the data of individual SI exposures (4-8 hours long - the highest time resolution providing both spectral and timing information) each exposure being treated as a separate dataset.  In the case of 1E1740-294, having 5 to 20 times lower signal to noise ratio the data averaged over single observations (comprised of 1-6 SI exposures with total duration of 4-34 hours) were treated as individual datasets to be regrouped. The intensity range 0.3 to 5.6 cnt/sec/cm$^{2}\\times10^{-3}$\\ (80-1500 mCrab) was divided into 10 intensity bins. ",
        "conclusions": ""
    },
    "9603/astro-ph9603097_arXiv.txt": {
        "abstract": "We describe a new implementation of a parallel N-body tree code. The code is load-balanced using the method of orthogonal recursive bisection to subdivide the N-body system into independent rectangular volumes each of which is mapped to a processor on a parallel computer. On the Cray T3D, the load balance in the range  of 70-90\\% depending on the problem size and number of processors. The code can handle simulations with $>$ 10 million particles roughly a factor of 10 greater than allowed on vectorized tree codes. ",
        "introduction": "During the past decade, $N$-body tree codes have been applied successfully to various problems in galaxy dynamics, galaxy formation and cosmological structure formation (Barnes \\& Hut 1986; Hernquist 1987; Dubinski 1988). Computing time only scales as $N\\log N$ so they provide a relatively fast way to solve the general collisionless N-body problem. For the most part, the favored tree code in studies has been the Barnes-Hut (BH) code (1986) mainly because it was widely distributed by its authors but also because it has been easy to implement. Different tree-based algorithms have also been used successfully for various problems (Appel 1985; Jernigan \\& Porter 1989; 1993; Benz, Bowers, Cameron \\& Press 1990; Steinmetz \\& Muller 1993). Tree codes have also been linked up with smoothed particle hydrodynamics (SPH) (Hernquist \\& Katz 1989; Benz et al 1990; Steinmetz \\& Muller 1993). Tree code simulations running on the latest workstations and vector supercomputers are generally restricted to $N < 10^6$ because of memory and time limitations. Cosmological and hydrodynamical simulations require ever greater dynamic range, so there is a strong desire to increase these limits.  During the past few years, there has been a paradigm shift in supercomputing with the movement from vector machines to the new massively parallel machines.  Parallel supercomputers are collections of hundreds to thousands of independent smaller commodity processors interconnected with an internal high speed communication network. The NSF supercomputing centers support several machines: the Connection Machine 5, (CM-5), the Intel Paragon, the IBM SP/2, and the Cray T3D. Each processor operates independently but communication software allows messages containing data to be exchanged rapidly with other processors. In principle, a problem can be partitioned among the $N$ processors and a maximum $N$-fold increase in computational speed can be realized. In practice, the speed up is smaller because of the extra time needed for exchanging data in message-passing algorithms.  Parallel machines offer a new route to very fast computation if algorithms can be redesigned to conform to the message-passing paradigm and communication can be minimized. The new algorithms are often very different than their sequential counterparts and require a considerable effort to redesign. The key to a successful algorithm is good {\\em load balance}: both data and computational work must be distributed evenly among the processors. A common way of achieving load balance on parallel machines is through {\\em domain decomposition} (Fox, Williams \\& Messina 1994). The physical domain of the problem is partitioned into smaller subdomains and the physical quantities of these subdomains: either values of density, pressure and temperature at grid elements in an Eulerian fluid code or collections of particles and their attributes in an N-body code are assigned to each processor. The trick in designing a parallel algorithm is finding a way of partitioning the domain so that data and work are evenly distributed and communication is minimized.  There have been many efforts to parallelize the BH tree code. Hillis and Barnes (1987) and Makino and Hut (1989) ported the code to the Connection Machine-2 but at the time there was little gain over the existing vector machines. Salmon (1990) introduced a new parallel tree algorithm using a domain decomposition based on the orthogonal recursive bisection of the N-body volume into rectangular subvolumes for purposes of load balance. The code has been applied to dark halo formation with initial  density fluctuations based on different power spectra (Warren et al. 1992; Warren 1994; Zurek et al. 1994). Warren \\& Salmon (1993) have also recently designed a new algorithm which uses a load-balancing scheme based on a parallel hashed oct-tree (also Warren 1994). Dikaiakos and Stadel (1995) also have developed another variant of Salmon's algorithm in their parallel tree code PKDGRAV. Parallel N-body tree codes have also been discussed as a pure computer science problem (Bhatt et al. 1992; Pal Singh 1993; Pal Singh et al. 1995).  In this paper, we present a new version of Salmon's (1990) parallel tree algorithm. In \\S 2, we describe the BH tree code as implemented at the node level in the parallel code. In \\S 3, we describe a new portable implementation of Salmon's parallel algorithm using the Message Passing Interface (MPI) package to handle processor communication. In \\S 4, we examine the code's performance on the Cray T3D, checking the accuracy, speed, and load balance. The term ``node'' is often used to refer to a single processor on a parallel computer but it is also used to refer to a data structure in a tree. To avoid confusion, I will only use the term ``node'' in reference to tree structures and ``processor'' for a computational ``node''. ",
        "conclusions": "A parallel tree code based on Salmon's algorithm has been implemented  using the MPI message passing software on 3 parallel machines: the Cray T3D, Paragon and the  IBM SP/2. In principle, the code can also run on a small network of workstations although we have not had the opportunity to test it this way. The code incorporates an improved version of the BH tree algorithm including particle grouping for tree walks, non-recursive tree walking and a ``safe'' cell opening criterion. The code has been applied successfully to simulations of galaxy interactions (Dubinski, Mihos, \\& Hernquist 1996) and is currently being applied to a variety of other problems in galaxy dynamics.  The code works best with a small number of processors that are heavily loaded with particles. If $N < 10^6$, we find that  the number of processors should be less than 64 to get the best use of computing time. Bigger problems will of course require the added memory of more processors.  This code's main disadvantage is its excessive demand for memory. The memory needs become excessive when $\\theta$ becomes too small. The use of locally essential trees is the main culprit in this respect. Perhaps, this method of load balancing can be replaced by an asynchronous message passing scheme in which tree walks in various processors are done on demand by distributing particles among the processors. At present, the code is still marginally time limited and available memory will be greater in the next generation of machines.  We are now in the process of adapting the code for both cosmological simulations with periodic boundary conditions and smoothed particle hydrodynamics (Dave, Dubinski \\& Hernquist 1996). With this new code, we should be able to increase the typical number of particles used in current CRAY C90 TREESPH (Hernquist \\& Katz 1989) simulation by a factor of 10-100. The improved dynamic range should allow simulations of galaxy formation in a full cosmological context."
    },
    "9603/astro-ph9603125_arXiv.txt": {
        "abstract": "We present an analytical method to extract observational predictions about non linear evolution of perturbations in a Tolman Universe. We assume no a priori profile for them. We solve perturbatively a Hamilton - Jacobi equation for a timelike geodesic and obtain the null one as a limiting case in two situations: for an observer located in the center of symmetry and for a non-centered one. In the first case we find expressions to evaluate the density contrast and the number count and luminosity distance vs redshift relationships up to second order in the perturbations. In the second situation we calculate the CMBR anisotropies at large angular scales produced by the density contrast and by the asymmetry of the observer's location, up to first order in the perturbations. We develop our argument in such a way that the formulae are valid for any shape of the primordial spectrum. ",
        "introduction": "The study of large - scale structure formation represents one of the most exciting research fields in Cosmology. The currently accepted view is that the structures we observe today trace their origin to primeval density inhomogeneities, generated by zero point fluctuations in the scalar field responsible for Inflation. The evolution of these inhomogeneities, from certain initial values ($<<1$) to their present distribution in stars, galaxies, clusters of galaxies and so on, involves complex non linear hydrodynamic and gravitational processes. Due to this complexity, it has not been possible to study this evolution from a completely analytical perspective \\cite{Zeldovich},\\cite{Roman},\\cite{Shandarin}.  Given the impossibility at present (and possibly in principle \\cite{Hobill}) to find a general exact analytical solution to Einstein equations, our understanding of processes such as the clustering of galaxies is largely dependent on numerical techniques. Of the known exact solutions to those equations, most of them presuppose the presence of certain symmetries of space-time and/or energy density \\cite{MacCallum}, \\cite{Krasinski}. One of them is the Tolman solution, which describes an inhomogeneous Universe filled with pressureless matter and spherically symmetric around a point. This solution is completely characterized by two time independent functions $% f^2(r)$ and $F(r)$. This last one can be interpreted as the mass contained in a sphere and $f^2(r)$ as the mechanical energy of a shell, both of radius $r$.  In spite of its limitations, the Tolman solution has been fruitfully used to study a great variety of effects related to the formation of large scale structures: formation of large scale voids, anisotropies in the microwave background radiation, possible fractal distribution of galaxies, etc. \\cite {voids}, \\cite{Paczynski}, \\cite{Panek}, \\cite{Ribeiro}, \\cite {Moffat-Tatarski}, \\cite{Fullana}\\cite{langlois}. But all (or almost all) of these studies are based on the numerical integration of the corresponding equations (mainly the equation for null geodesics), and often simplified expressions for $f^2(r)$ and $F(r)$ are employed, sometimes with no other reason than to facilitate the numerical calculations.  The goal of this paper is to present an analytical method to study the non linear evolution of density perturbations in a matter dominated Universe, viewed as an instance of the Tolman Universe, laying the emphasis in a physically motivated choice for the Tolman functions. Concretely, we shall assume the Tolman functions to be such that at early times, in the linear regime, the model be equivalent to a spatially flat Friedmann - Robertson - Walker (FRW) matter dominated Universe with growing perturbations. Otherwise, we shall keep the Tolman functions completely general. Observe that in this way the Tolman functions are determined by the spectrum of perturbations to the original FRW model.  We shall focus on simulating in our spherically symmetric model the results of concrete cosmological observations such as the redshift - luminosity distance and the number count - redshift relationships. Since these observations rely on information carried by light, the analysis centers on the study of the null geodesics in the model. We device a perturbative method based on the Hamilton - Jacobi equation for a time like geodesic and obtain the null geodesic as a limiting case.  We apply our scheme to study two situations. In the first one we consider an observer located in the center of the Universe and evaluate the number count and luminosity distance vs redshift relationships. For the first cosmological test we assume the simplifying hypothesis that the luminosity of galaxies does not evolve, which is justified in view of the fact that the redshifts we consider are small ($z<0.08$). For larger redshifts however, the evolution of luminosity ought to be taken into account \\cite{Campos}.  In the second situation we consider an observer located away from the center and study the anisotropies in the cosmic background radiation due to the presence of perturbations in the photon path, from the last scattering surface to the observer's position.  Throughout the paper a unique matter component of the Universe is assumed.  We begin our discussion by briefly reviewing the theory of linearized scalar adiabatic perturbations to a FRW background, the main features of the Tolman solution to the Einstein equations and the matching procedure. In section {\\bf III} we review the cosmological observations $N$ vs $z$ and $d_\\ell $ vs $z$ and write their form for a Tolman Universe. In section {\\bf IV} we begin by sketching our method, we perform the explicit calculations for a centered observer and find the general expression for the mentioned observational tests. In section {\\bf V} we do the same as in section {\\bf IV} but for a non centered observer and evaluate the anisotropies in the CMBR temperature due to the different paths of the photon from the last scattering surface to the observer's position. To obtain an idea of the kind of results to be expected from concrete perturbation profiles, in section {\\bf VI} we apply our formulae to a particular Tolman Universe, corresponding to a scale invariant spectrum of perturbations. Finally we discuss our results in section {\\bf VII}.  The overall conclusion is that the method we propose affords a simple way to model the growth of nonlinear structures such as large scale voids, and may be used profitably in testing competing theories of primordial fluctuation generation. ",
        "conclusions": "In this paper we have developed an analytical method to study the non linear evolution of adiabatic perturbations in a matter dominated Universe. We build the initial profile of the fluctuation on a $t=const.$ surface (that corresponds to the beginning of the matter dominated era, when structure formation begins) as a superposition of all those modes whose wavenumber is smaller than a certain cut-off $k_0$. This cut-off may correspond to the mode that survived the Meszaros effect \\cite{meszaros}, Landau damping or free streaming \\cite{Peebles93}.  >From this point on we can follow the non linear evolution of the perturbations by means of the exact solution to Einstein equation for a pressureless, spherically symmetric Universe, namely the Tolman Universe. Since this solution is isotropic with respect to one point, the Tolman solution can be applied only to analyze the formation of a structure (for example a void or cluster) with the proper symmetry. This restriction in the applicability of the Tolman solution to the study of structure evolution is compensated by the advantages of having an exact solution at our disposal.  Formally, the basis of all calculations for observable quantities lays in the resolution of the equation for the null geodesics, which in general is not trivial and is carried out by numerical methods. In this work we presented a perturbative method to do this calculation analytically. We put our calculations in a physical framework by the way in which we build the perturbations. We carried out the calculations only to second order, which means a mild non linear evolution, but the extension to higher orders is straightforward.  In the last section we obtained a glimpse of the results to be obtained by assuming a scale free spectrum of initial perturbations. In view of the simplicity of the calculations involved, we find that our results give a remarkable approximation to the structure of some of the largest known voids, such as B\\\"otes'.Quantitatively, our results are accurate only up to an order of magnitude (in fact, the B\\\"otes void is surrounded by a wall whose density contrast is $\\Delta \\delta /\\bar \\delta \\simeq 4$, and its radius is of about $30Mpc$ \\cite{deLapparent}). The difference, of course, could be reduced by a better choice of the cutoff wavelength and, most importantly, by carrying the computation to higher orders.  With regard to future work, the most important feature of the method we propose, besides its analytical and relatively simple character, is that it may be applied for any form of the primordial spectrum. Thus it becomes only a matter of plugging in one's favorite theory of primordial fluctuation generation (as reflected in the particular form of $\\zeta _k$ (Eq. (\\ref {C2-superp})) to easily obtain (rough) testable predictions of that model. While we have used here a (unrealistic) scale invariant spectrum for demonstration purposes, the calculation is equally simple with red, blue, or more sophisticated alternatives.  We therefore believe the methods we describe in this paper will be an useful tool in the delicate task of sorting between the manyfold fluctuation generation scenarios now available."
    },
    "9603/astro-ph9603019_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\bigskip  We have used preliminary grainless model atmospheres to model the spectral distribution of cool brown dwarfs and found that these reproduce well the observed spectroscopic and photometric properties of Gl229B.   We confirm the presence of methane absorption bands and therefore the substellar nature of Gl229B. The absence of CO absorption in the 4-5~$\\mu$m range sets a secure upper limit of 1000~K on the effective temperature.  Gravity can be constrained at fixed temperature using the relative K to J band fluxes and this leads to values ranging from log~g=5.3 at 1000~K to log~g= 4.8 at $\\teff\\ = 900~K$.  A minimum temperature and gravity can be set by the presence of water vapor bands in Gl229B's spectrum to about 700-800~K where photospheric temperatures are expected to drop below the condensation temperature of H$_2$O.  While uncertainties in the age of the system and in the temperature of Gl229B yield some indetermination for the mass of Gl229B, we can derive upper limits from the effective temperature limit, namely $M\\le 0.065 \\mso$ for an age $t\\le$10 Gyr, and $M\\le 0.055 \\mso$ for a more likely age $t\\le$5 Gyr.  The most likely solution $M\\approx 0.04-0.055$ $\\msol$ is supported by the quality of the most reliable fit to the water bands for these parameters, and would be consistent with an age similar to that of our solar system.  Our analysis cannot however exclude a mass as low as 0.02 $\\mso$ if $\\te = 900$ K and $t=0.5$ Gyr.  One of the main sources of uncertainty in the present models is the absence of grain formation in the computed atmospheres. Recent work by Tsuji \\etal\\ \\markcite{tsuji96a}\\markcite{tsuji1996b} (1996a,b) shows that grain formation may lead to substantial heating of the photospheres of cool M dwarfs and brown dwarfs, possibly resulting in a much lower effective temperature for Gl229B than derived with the present grainless models.  The absence of predicted VO and FeH features in the red spectrum (Z band) of Gl229B indicates reduced abundances of grain forming elements, and therefore the presence of condensation and perhaps grain heating in this cool brown dwarf.  However homogeneous dusty models by Tsuji \\etal\\ \\markcite{tsuji96b} (1996b) failed to reproduce the spectral distribution of Gl229B while the present models reproduce the observed strengths of molecular features such as water vapor which would otherwise be substantially weakened by the greenhouse effect of grains.  This suggests that Gl229B's atmosphere is very reminiscent of that of Jupiter, in which condensed material and cloudtop layers were not found in the expected amounts by the Galileo atmospheric probe (D. Isbell and D. Morse, NASA press release 96-10).  Perhaps in both Jupiter and Gl229B, these condensates either formed in clouds or sank deeper into the atmosphere.  We thank T. Nakajima and colleagues for providing updated version of the Gl229B spectrum in electronic form as well as the referee for his valuable comments.  We are also indepted to J. M. Matthews and D. R. Alexander for proofreading the text.  This research has been partially supported by a NASA LTSA and ATP grants to ASU and an NSF grant AST-9217946 to WSU.  The atmospheric calculations have been performed on the Cray C90 of the San Diego Supercomputer Center and on the IBM SP2 of the Cornell Theory Center, supported by the NSF."
    },
    "9603/astro-ph9603110_arXiv.txt": {
        "abstract": "We present the results of a spectral and timing analysis of PSR~0656+14 based on the complete set of ROSAT observations carried out with the PSPC instrument in 1991 and 1992.  The present  analysis confirms the thermal origin of the bulk of the  emission in the soft X-ray band (Finley et al. 1992).  In addition, we find  strong evidence of a harder component, described equally well with a blackbody  at $T \\simeq 2\\times10^6$ K, or with a steep  power law with photon index $\\Gamma \\simeq 4.5.$ This bimodal emission is also supported by an analysis of the light curve shape as a function of the energy.  The  0.1--2.4 keV light curve of PSR~0656+14, with a pulsed fraction of about 9\\%, is interpreted with a simple model for the temperature di\\-stri\\-bution on the neutron star surface, coupled with the geometrical information derived from radio data. In this model, which includes the effects of relativistic light bending and gravitational redshift, the X--rays ori\\-gi\\-na\\-te from two thermal components resulting from neutron star cooling and blackbody emission released in the hotter polar  cap regions.  The observed modulation can be reproduced only if PSR~0656+14 has a relatively high dipole inclination ($\\sim30^o$) and $(1+z)~\\lsim$1.15. The overall pulsed fraction cannot be significantly increased by including the polar cap contribution, if its temperature and intensity are to be consistent with the observed spectra. ",
        "introduction": "ROSAT ob\\-ser\\-va\\-tions of isolated neutron stars with characteristic ages of  a few $10^{5}$ years (PSR~0656+14, PSR~1055-52, Geminga) have shown that the bulk of their X--ray e\\-mis\\-sion is thermal in origin and may result from neutron star cooling (see \\\"Ogelman 1995 for a review). The observed periodic modulation in the X--ray flux indicates the presence of large thermal gradients on the surface of these magnetized neutron stars. For a given surface temperature distribution, the degree of modulation depends on the intrinsic geometry of the source and on its orientation relative to the observer, which in the case of radio pulsars can be inferred from the shape, polarization and spectrum of the radio pulses.  PSR~0656+14, first studied in the  soft X--ray range with the Einstein Observatory (Cordova et al. 1989), is the brightest ''middle-aged''   neutron star      in this energy band. Its association with a  possible X--ray supernova remnant, the Gemini--Monoceros ring (Nousek et al. 1981; Thompson et al. 1991) has recently been criticized, on the basis of new results on the proper motion of the pulsar (Thompson \\& Cordova 1994). A candidate optical counterpart with V$\\sim25$ has been discovered by Caraveo, Bignami \\& Mereghetti (1994) within 1'' of the radio position of PSR~0656+14.  Three   pointings of PSR~0656+14 were carried out with the ROSAT PSPC detector (1991 March 26, 1992 March 24 and April 15). While the first observation ($\\sim$3200 s) has been presented in detail by Finley, \\\"Ogelman \\& Kizilo\\u{g}lu (1992), only the preliminary results of the two 1992 pointings have been reported (\\\"Ogelman 1995).   In this paper we analyze all the PSPC data of PSR~0656+14 and interpret them in the context of a simple model for its temperature distribution  coupled with the geometrical configuration derived from the radio data. On the basis of the new observational results, the paper carries out an analysis similar to the one presented by Page (1995a) on PSR~0656+14 and explores  the effect of polar cap emission on the light curves and spectra. ",
        "conclusions": "The ROSAT spectral and timing analysis of PSR~0656+14 indicate that its soft X--ray emission  is characterized by two separate components. The blackbody  temperature of the  softer component ($9\\times10^5~^o$K), independent of the model used for the hard tail, is similar to that derived by Finley et al. (1992).  The hard tail can be described equally well by a blackbody or by a steep power law. In the blackbody case, it is natural to attribute the hard component to localized emission from hotter regions of the star surface, e.g. the polar caps, as discussed above. In the case of a power law, the tail can be interpreted as a non-thermal emission of magnetospheric origin. The best fit power law index ($\\Gamma =4.5$) is greater than that observed in  PSR~1055--52 ($\\Gamma =1.5$ , \\\"Ogelman \\& Finley 1993) and Geminga ($\\Gamma =2.5$, Halpern \\& Ruderman 1993), the other   isolated neutron stars of comparable age and  with a similar two component soft X--ray emission. These two objects have also been clearly observed at $\\gamma$-ray energies (E$\\gsim$100 MeV), contrary to PSR~0656+14 for which only a marginal detection, at a much lower flux, has been claimed (Ramanamurthy et al. 1996). This fact might indicate a possible connection between the level of $\\gamma$--ray emission and the slope of the hard X--ray tail.  With a model for the anisotropic thermal cooling induced by the crustal, dipolar magnetic field (similar to our first component), Page (1995a)  derived for PSR~0656+14 pulsed fractions between 8 and 10\\%, smaller than that ($\\simeq$14\\%) reported by Finley et al. (1992). This result is confirmed by our analysis of the complete PSPC data set. We show that, even in the most favourable geo\\-me\\-tric configuration ($\\alpha \\sim 30^o$) among those compatible with the radio data, the observed modulation can be reproduced only for $(1+z)\\le$1.15 (corresponding to $R/R_{S}~\\gsim$4, where $R_{S}$ is the Schwarzschild radius). These low values of $z$ are consistent with a stellar radius between $12$ and $17$ km and a neutron star mass $M\\sim 1 M_{\\odot}$. For M = 1.4 $M_{\\odot}$, our limit on $z$ requires a radius  $\\sim 17$ km, a value     inconsistent with the most recent  neutron star models (Glendenning 1985; Wiringa et al. 1988). Masses below 1.4 $M_{\\odot}$ have radii sufficiently large to account for the observed modulation in the X-ray signal, only  if the equation of state is very stiff (Page 1995a).  The presence of higher order moments in the magnetic field, or the effects of a magnetized atmosphere causing anisotropic photon transport have been proposed as solutions for increasing the light curve modulation (Page 1995a, 1995b). Our analysis shows that the alternative possibility of obtaining a higher modulation by including the additional contribution of hot   polar caps does not solve the problem. The polar cap emission has little effect on the overall mo\\-du\\-la\\-tion, if its  temperature and intensity are to be consistent with the observed spectra.  \\vskip 20pt"
    },
    "9603/astro-ph9603042_arXiv.txt": {
        "abstract": "We investigate the dynamical evolution of the Galactic Globular Cluster System in considerably greater detail than has been done hitherto, finding that destruction rates are significantly larger than given by previous estimates. The general scheme (but not the detailed implementation) follows \\cite{AHO:88} (1988; AHO).  For the evolution of individual clusters we use a Fokker-Planck code including the most important physical processes governing the evolution: two-body relaxation, tidal truncation of clusters, compressive gravitational shocks while clusters pass through the Galactic disk, and tidal shocks due to passage close to the bulge. Gravitational shocks are treated comprehensively, using a recent result by \\cite{KO:95} (1995) that the $\\langle \\Delta E^2 \\rangle$ shock-induced relaxation term, driving an additional dispersion of energies, is generally more important than the usual energy shift term $\\langle \\Delta E \\rangle$. Various functional forms of the correction factor are adopted to allow for the adiabatic conservation of stellar actions in a presence of transient gravitational perturbation.  We use a recent compilation of the globular cluster positional and structural parameters, and a collection of radial velocity measurements. Two transverse to the line-of-sight velocity components were assigned randomly according to the two kinematic models for the cluster system (following the method of AHO): one with an isotropic peculiar velocity distribution, corresponding to the present day cluster population, and the other with the radially-preferred peculiar velocities, similar to those of the stellar halo. We use the Ostriker \\& Caldwell (1983) and the Bahcall, Schmidt, \\& Soneira (1983) models for our Galaxy.  For each cluster in our sample we calculated its orbits over a Hubble time, starting from the {\\it present} observed positions and assumed velocities. Medians of the resulting set of peri- and apogalactic distances and velocities are used then as an input for the Fokker-Planck code. Evolution of the cluster is followed up to its total dissolution due to a coherent action of all of the destruction mechanisms. The rate of destruction is then obtained as a median over all the cluster sample, in accord with AHO.  We find that the total destruction rate is much larger than that given by AHO with more than half of the present clusters ($52\\%-58\\%$ for the OC model, and $75\\%-86\\%$ for the BSS model) destroyed in the next Hubble time. Alternatively put, the typical time to destruction is comparable to the typical age, a results that would follow from (but is not required by) an initially power law distribution of destruction times. We discuss some implications for a past history of the Globular Cluster System, and the initial distribution of the destruction times raising the possibility that the current population is but a very small fraction of the initial population with the remnants of the destroyed clusters constituting presently a large fraction of the spheroid (bulge+halo) stellar population. ",
        "introduction": "Globular clusters are thought to be the oldest stellar systems in our Galaxy, and a history of attempts by theoreticians and observers to understand the keys of their evolution is as old as the discipline of stellar dynamics, with a comprehensive review provided by \\cite{S:87} (1987). Yet, these relatively simple stellar systems are not understood well enough to predict their future with desirable accuracy, in part due to lack of accurate observational values for the current dynamical state and in part due to a residual uncertainty concerning the complete catalog of relevant physical processes operating on these systems. More than that, we have almost no clues about their past, and in particular, whether what we see now is representative of the {\\it initial} Globular Cluster System, or just a small leftover after a great destruction battle that occurred earlier in the history of the Galaxy. In this paper we consider the evolution of the globular clusters and their ultimate disappearance, and propose a simple model for their initial distribution.  Pre- and post-core collapse evolution of an isolated cluster is relatively well understood (\\cite{S:87} 1987; \\cite{G:93} 1993). Significant progress in understanding of the evolution was achieved using Monte Carlo and Fokker-Planck calculations. But the galactic environment makes the clusters subject to external perturbations -- tidal truncation and the gravitational shocks due to passages close to the bulge and through the disk. The shock processes, although known to be important, have never been carefully included in the evolution of the system. We investigate the Fokker-Planck models including the shocks elsewhere (\\cite{GLO:96}; hereafter GLO) and show that the dispersion of energy of the stars, induced by the shocks (\\cite{KO:95} 1995; hereafter KO), is generally even more important for the evolution than the first-order energy shift. Another very important aspect of modelling globular clusters is the initial mass spectrum. Multi-mass clusters undergo core collapse much faster than in the single-mass case, and their destruction is much more efficient (see \\cite{LG:95} 1995, and the references in GLO). We restrict ourselves however to the single mass models in order to maintain clear physical understanding; but due to omission of this aspect of the problem, our results provide a lower bound to the rate of destruction of the globular clusters.  As long as we have (at least, approximate) understanding of the evolution of a single cluster we can turn to the study of the system of globular cluster in our Galaxy, and in external galaxies. \\cite{CKS:86} used a semi-analytical Monte-Carlo technique to estimate the importance of the different mechanisms acting upon the cluster. They considered two-body relaxation, tidal stripping of stars, and the first order tidal shocking effect - due to the crossing disk and interactions with giant molecular clouds (GMC). For each of those processes they calculated the cluster mass and energy changes associated with them to predict the evolution. They followed the cluster evolution only up to core collapse, and assumed a single-mass King model (\\cite{K:66}) for the internal cluster structure. Tidal heating due to the GMC was found to be negligible compared to the disk shocks. Note however, that the Galactic model assumed in that work (\\cite{BSS:83}; BSS) is strongly favorable to the disk shock for the clusters with small orbital radius since the surface density of the disk increases exponentially as the galactocentric radius decreases (see Section \\ref{sec:GalModel}). Chernoff et al. (1986) has concluded that many of the clusters located within inner 3 kpc from the Galactic center have undergone core collapse, and many of them may already have been destroyed. A number of authors have pointed out that the bulge and stellar spheroid themselves could be composed of remnants of the destroyed globular clusters, if prior destruction of globular clusters occurred at high enough rate.  Another mechanism for the mass loss is the stellar evolution. \\cite{CS:87} used a similar method to Chernoff et al. (1986) and included a power-law initial mass function for the cluster stars. Mass loss due to stellar evolution is important for the early evolution of the cluster, but then fades away because the mass loss is large only for massive stars whose lifetime is short. Comprehensive study including the effects of stellar evolution has been done by \\cite{CW:90}. They used a Fokker-Planck code with an extensive spectrum of stellar masses (20 species), which includes stellar evolution and relaxation processes. Multi-mass models evolve much faster and the evaporation rate is larger. Also, mass segregation (\\cite{S:87} 1987) speeds up the collapse. Mass loss during first $5 \\times 10^9$ yr is sufficiently strong to disrupt weakly concentrated clusters ($c<0.6$). Combined with the relaxation, it destroys many low mass and low concentration clusters within a Hubble time.  The present characteristics and the evolutionary state of the observed Galactic globular clusters were also investigated by \\cite{AHO:88} (1988; hereafter AHO). We will draw heavily on that paper and compare our results with AHO. AHO used a sample of the 83 Galactic globular clusters with the known structural parameters and line-of-sight velocities. They considered virtually all the important physical mechanisms (except the mass spectrum and mass loss due to stellar evolution) to calculate the present day destruction rates for the clusters in the sample. The rates were defined as the inverse time it takes for a given mechanism to dissolve a cluster, in units of a Hubble time, which is nominally defined as $10^{10}$ yr. They estimated the rates for the evaporation through tidal boundary, disk and bulge tidal shocks, and dynamical friction. The last process has not been widely investigated in its application for the clusters' evolution. Its effect reduces to the gradual spiraling of the cluster toward the Galactic center, as the cluster loses orbital energy due to continuous interactions with field stars and dark halo. AHO found that this mechanism is a relatively unimportant one, except for unusually massive clusters. They calculated a number of orbits associated with each cluster in the sample and took a median over the whole resulting distribution for a corresponding destruction rate. They investigated two Galactic models (OC and BSS; Section \\ref{sec:GalModel}), and two kinematic models for the globular cluster system: isotropic, which corresponds to the current distribution, and predominantly radial, which resembles that of the halo stars and might be closer to an {\\it initial} cluster distribution. The assumption concerning the velocity distribution is necessary as only one velocity component is observed. (Note however that the program to obtain the true space velocities of globular clusters is underway; see \\cite{CH:93} 1993). We will discuss the kinematic models in Section \\ref{sec:kinematics}.  The central bulge was found to be very efficient in destroying clusters on a highly elongated orbits. This leads to an {\\it isotropization} of the orbits, and even preferential survival of tangentially biased orbits in the Hubble time. Thus, an initially radial distribution better fits the present population, than the initially isotropic one. At the current time, evaporation is the most important destruction process. The difference between the two Galactic models was found to be relatively small, except for the disk shocking that obviously depends on a body of the surface mass profile. AHO introduced a weighting factor that accounts for the fact that some orbits are less probable because of the various destruction mechanisms acting upon the cluster. Such a weighting reduces the computed rates, but converges to a similar value of $0.05$ for all Galactic models and kinematic profiles considered. This implies that only 4 or 5 clusters are destroyed over a Hubble time for present conditions. But AHO treated clusters in a simplified fashion as King models and they (like other investigators) did not allow for the tidal shock relaxation phenomena. It is these defects that we remedy in the present paper.  Finally, a recent paper by \\cite{HD:92} used a statistical approach and approximate analytical estimates for the relaxation times for a sample of 140 clusters. They concluded that the current evaporation rate is $5\\pm 3~{\\rm Gyr}^{-1}$ (about ten times the rate found by AHO), which means that a significant fraction of the present day clusters would be destroyed in a next Hubble time.  In this paper we make a significant improvement over AHO by applying detailed Fokker-Planck calculations to the real globular cluster sample. We describe the sample in Section \\ref{sec:sample} and the Galactic models in Section \\ref{sec:GalModel}. Then we discuss the two kinematic models and compare the resulting properties of the orbits for our sample. In section \\ref{sec:processes} we describe the destruction mechanisms involved in the globular cluster evolution. We conclude section \\ref{sec:method} with the history of our code and the formulation of the numerical strategy. Section \\ref{sec:results} presents our results for all runs. Finally, we speculate on the possible past history and future fate of the Galactic globular clusters in Section \\ref{sec:discussion}. Section \\ref{sec:conclusions} sums up our conclusions. ",
        "conclusions": "\\label{sec:conclusions} We have used the Fokker-Planck code to investigate the destruction rate of globular clusters in our Galaxy. We applied two forms of the adiabatic correction for gravitational shocks and found that the median results do not depend much of the particular form of the correction. The current destruction rate for the sample is about $0.5-0.9$ per $10^{10}$ yr (depending on the Galactic and kinematic models), which implies that more than half of the {\\it present} clusters is to be destroyed within the next Hubble time. This estimate is approximately a factor of ten higher than that obtained by AHO. There are two principal reasons for the change. First, our Fokker-Planck detailed calculations for each cluster give systematically larger rates of two body relaxation, and core collapse than did the essentially time scale arguments of AHO. Secondly, the new tidal shock relaxation process described by KO further reduces lifetimes by a significant amount.  Trying to understand the {\\it original} population of Galactic globular clusters, we considered two possible models for the distribution of the cluster lifetimes. Both of them can be normalized to the median present destruction rate. We favor the power-law model on basis of a shape of the rate distribution as it naturally explains the fact that the current median time to destruction is comparable to the present mean cluster age and also because the predicted distribution of cluster destruction times provides a reasonable match to observations. The power-law distribution allows a much larger number of the clusters to have been formed initially than is currently observed, and allows the possibility that the debris of the early disrupted clusters might have formed the much of spheroid of our Galaxy.  \\cite{Su:95} has investigated the possibility of populating the stellar halo by remnants of the destroyed star cluster (open and globular). He comes to the conclusion that much larger number of low-mass ($10^3 - 10^4 M_{\\sun}$) clusters than observed now is required to match the mass of the spheroid. This conforms to our result, since low mass (and hence, weakly concentrated) clusters have short lifetime. All those non-observed clusters have to be destroyed before the present time.  Finally we note that the inclusion of the mass spectrum in the Fokker-Planck models would strongly enhance the relaxation and core collapse, and ultimately speed up the dissolution of the clusters. This could also increase the destruction rate by a significant amount.  The destruction rates for the individual clusters are available electronically upon request."
    },
    "9603/astro-ph9603038_arXiv.txt": {
        "abstract": "The light curves from a variety of celestial objects display aperiodic variations, often giving rise to red--noise components in their power spectra. Searching for a narrow power spectrum peak resulting from a periodic modulation over the frequency range in which these ``coloured'' noise components are dominant has proven a very complex task. Commonly used methods rely upon spectral smoothing or incoherent summation of sample spectra in order to decrease the variance of the power estimates. The consequent reduction in frequency resolution causes also a reduction of sensitivity to periodic signals.  We develop here a technique aimed at detecting periodicities in the presence of ``coloured\" power spectrum components, while mantaining the highest Fourier frequency resolution. First we {\\co introduce a simple approximation to the} statistical properties of the ``coloured'' power spectra from celestial objects, based on a few examples and the theory of linear processes. We then estimate the continuum components in the power spectrum through an {\\it ad hoc} smoothing technique. This involves averaging the spectral estimates adjacent to each frequency over a {\\co suitably chosen} interval, in order to follow steep red--noise features and produce estimates that are locally unaffected by the possible presence of a sharp peak. By dividing the sample spectrum by the smoothed one, a white--noise like spectrum is obtained, the approximate probability distribution of which is derived. A search for coherent pulsations is then carried out by looking for peaks in the divided spectrum, the chance probability of which is below a given detection threshold. If no significant peaks are found, an upper limit to the amplitude of a sinusoidal modulation is worked out for each searched frequency.  The technique is tested and its range of applicability determined through extensive numerical simulations. We present also an application to the X--ray light curves of V$0332$$+$$53$, a highly variable accreting X--ray pulsar, and GX$13$$+$$1$, a bright and variable accreting source in the galactic bulge. ",
        "introduction": "Since the prehistorical efforts aimed at developing the calendar, the detection and investigation of periodic phenomena has played a major role in astronomy. Crucial information is obtained through the observation and measurement of periodicities in many classes of celestial bodies encompassing all scales from comets and asteroids to the largest structures currently known at cosmological distances. In some cases periodic signals from the cosmos can be measured to an exceptionally high accuracy, that rivals that of atomic clocks (Kaspi, Taylor \\& Ryba 1995). Astronomical time series of increased statistical quality, time resolution and duration have become available over the last two decades for a variety of objects and over different bands of the electromagnetic spectrum. Power spectrum analysis is probably the single most important technique that is applied to these series in order to: (a) detect periodicities (or quasi--periodicities) by the presence of significant power spetrum peaks; (b) characterise the noise variability through the study of continuum power spectrum components. In particular, recent applications to high energy astronomical time series have been especially numerous and successfull, as a consequence of the pronounced  variability (often both periodic and aperiodic in nature) detected in many sources and the availability of long uninterrupted observations  (up to several days) of high signal to noise ratio. The continuum  power spectral components arising from noise variability usually increase towards  lower frequencies ({\\it red noise}), often  in a power law--like fashion.  Their study has proven to be a useful tool for morphological classifications and, sometimes, has provided  constraints on physical models (e.g. Stella 1988; Hasinger \\& van der Klis 1989; van der Klis 1995). The periodic modulations revealed in a number of high energy sources often arise from the rotation of compact magnetic stars, or the orbital motion of a binary system. The detection and accurate  measurement of these periods provides a tool of paramount importance. A variety of other periodic or quasi--periodic phenomena in X--ray sources have been discovered  over a range of timescales (from tens of milliseconds to years).  Astronomical observations rely more and more upon photon counting instruments; therefore, measurement errors are often dominated by the statistical uncertainties originating from the Poisson distribution of the counts. This translates into a white noise power spectrum component of known amplitude, which, after normalisation, follows a $\\chi^2$ distribution with 2 degrees of freedom ($\\chi^2_2$). Any intrinsic variability of the source, either resulting from periodic signal(s) or from noise(s), must possess significant power above the counting statistics white noise component in order to be detected.  Traditionally, the detection of periodic signals through peaks in the sample spectrum has been carried out either (i) by eye, in all those  cases in which the peak amplitude is so large that it is self--evident or (ii) by ruling out (at a given  confidence level) that a peak originates from an underlying white noise. The latter technique implicitely assumes that the power spectra do not possess any conspicuous ``coloured\" component above the white noise. As mentioned above, this hypothesis is not verified at least over a range of frequencies in many instances. Indeed the very presence of ``coloured\" continuum power spectrum components resulting from source variability noise makes the detection of significant power spectrum peaks a difficult statistical problem. In general, establishing whether or not a sample spectrum peak originates from a periodic modulation requires an evaluation of the peak significance with respect to the {\\it local} continuum of the sample spectrum which, in turn, can be dominated by the   aperiodic variability of the source. Techniques along these lines have been developed, which often rely upon smoothing or incoherent summing of sample spectra to decrease the variance of the power estimates and/or allow the use of relatively simple statistics. In this way, however, the frequency resolution and, correspondingly, the sensitivity of the searches is reduced (e.g. Jenkins \\& Watts 1968; van der Klis {\\co 1988}). Moreover, standard spectral smoothing does not allow to reproduce power law--like spectral shapes with acceptable accuracy.  This paper describes a new technique for detecting  power spectrum peaks arising from  a periodic modulation, in the presence of ``coloured\" power spectrum components, while preserving the Fourier frequency resolution. In this technique, the continuum components of the spectrum at the $j$--th frequency are evaluated based on an {\\it ad hoc} smoothing technique which involves averaging the spectral estimates adjacent to $j$--th frequency over a given logarithmic interval excluding the $j$--th frequency itself. The advantage of this type of smoothing is that, while it allows the continuum features of the power spectrum to be followed, it is locally unaffected (i.e.  for the same Fourier frequency) by the presence of sharp power spectrum peaks. By dividing the sample spectrum by the smoothed one a white--noise like spectrum is obtained, the approximate probability distribution function ($pdf$) of which is derived based on the characteristics of the sample spectrum. A search for coherent pulsations is then carried out by looking for peaks in the divided spectrum, for which the probability of chance occurrence is below a given detection level. If no significant peaks are found, an upper limit to the amplitude of a sinusoidal modulation is worked out for each searched frequency.  Our treatment assumes that the instrumental noise is due to Poisson statistics; as such it can be readily applied to observations with photon counting detectors in any band of the electromagnetic spectrum. The generalisation to the case of a Gaussian instrumental noise is straightforward. The time series are supposed to  be equispaced and continuous.  The paper is structured as follows: in section 2 we {\\co introduce a simple approximation to} the $pdf$ of the power estimates in the sample spectrum of cosmic sources characterised by ``coloured\" noise. Section 3 describes the smoothing algorithm that we devised in order to estimate the corresponding continuum power spectrum components, even in the presence of quite steep red--noises. In section 4 we derive the $pdf$ of the white--noise like spectrum that is obtained by dividing the sample spectrum by the smoothed one. The prescription for detecting significant periodic signals and deriving their amplitude, is given in section 5. This includes also how to obtain upper limits in the case in which no significant peak is found. Section 3--5 summarise also the results from the extensive numerical simulations that were carried out in order to assess the reliability of the technique. In section 6 an application to the ``coloured\" power spectra from the X--ray light curves of an accreting X--ray pulsar (V0332+53) and a bright galactic bulge source (GX13+1) is presented. Our conclusions are in section 7. ",
        "conclusions": "The power spectrum analysis technique that we developed for the detection of periodic signals in the presence of ``coloured\" noise components presents the following main advantages: $(i)$ it does not require any reduction of the Fourier frequency resolution; $(ii)$ it can be used also in the presence of the relatively steep red noise components (power law slopes as low as --2) which are commonly found in nature; $(iii)$ it takes into account the statistical uncertainties in the estimator of the continuum power spectrum components. Extensive numerical simulations were carried out in order to test the reliability of the technique and define the range of applicability of the adopted approximations. We found that very good results are obtained if the first and the last 5--6 Fourier frequencies of the sample spectrum are excluded from the analysis and the width of the smoothing is larger than 30--40 Fourier frequencies.  Though based on relatively simple statistics, the numerical evaluation of the peak detection threshold must be carried out separately for each Fourier frequency: this can be quite CPU--intensive. A way around this limitation consists in carrying out a much faster search for potentially significant peaks by using a preliminary and simpler detection threshold. The significance of candidate peaks is then reassessed on the basis of a complete application of the technique described in this paper. Computer programs based the technique described in this paper will be made available to the community through the timing analysis package {\\it Xronos} (Stella \\& Angelini 1992a,b)."
    },
    "9603/astro-ph9603152_arXiv.txt": {
        "abstract": "A {\\it ROSAT\\/} observation of the narrow-line \\ion{Fe}{2} QSO PHL~1092 shows rapid variability that requires an efficiency of at least 0.13, exceeding the theoretical maximum for an accretion disk around a non-rotating black hole.  Plausible explanations for this high efficiency incorporate anisotropic emission and/or accretion onto a rapidly rotating black hole, the latter recently suggested by Kwan et al. as a mechanism for generating PHL~1092's strong \\ion{Fe}{2} lines by mechanical heating in an accretion disk.  The soft X-ray luminosity of PHL~1092 had also increased by a factor of 21 over the weak {\\it Einstein} detection, to more than $5 \\times 10^{46}$~ergs~s$^{-1}$.  Its photon spectral index of 4.2 is among the steepest of any AGN.  These X-ray properties are characteristic of narrow-line Seyfert~1 galaxies, of which PHL~1092 is evidently a very luminous member.  Narrow-line QSOs also extend a significant correlation between X-ray luminosity and X-ray spectral index which we have found among a large sample of optically-selected, narrow-line Seyfert~1 galaxies observed by {\\it ROSAT}. ",
        "introduction": "Narrow-line Seyfert~1 galaxies (NLS1s, Osterbrock \\& Pogge 1985; Goodrich 1989) are defined by their optical emission-line ratios and widths: [\\ion{O}{3}]/H$\\beta < 3$ and FWHM H$\\beta~<~2000$~km~s$^{-1}$. NLS1s also tend to have strong permitted \\ion{Fe}{2}, \\ion{Ca}{2} and \\ion{O}{1}~$\\lambda$8446 emission lines (Persson 1988), as well as high-ionization lines that are typical of Seyfert~1 galaxies.  Their high X-ray luminosities were first noted by Remillard et al. (1986) and Halpern \\& Oke (1987).  With {\\it ROSAT}, NLS1s were discovered to have unprecedented X-ray properties (but first see Remillard et al. 1991).  Their X-ray spectra are much softer than those of ordinary Seyfert~1 galaxies (Brandt et al. 1994; Boller, Brandt, \\& Fink 1996; Pounds, Done, \\& Osborne 1995), and they display rapid, large-amplitude variability as well as extreme long-term changes (Boller et al. 1993; Brandt, Pounds, \\& Fink 1995; Grupe et al. 1995a,b).  New members of this class found in the {\\it ROSAT\\/} All-Sky Survey by Moran, Halpern, \\& Helfand (1996) all have \\ion{Fe}{2} emission to some degree. The latter paper further argues the need for a NLS1 class.  It has long been known that there are high-luminosity analogs of this class, the prototype of which is I~Zw~1 (e.g., Phillips 1976), which also have weak forbidden lines, narrow permitted lines, and strong \\ion{Fe}{2}.  One of the most extreme narrow-line \\ion{Fe}{2} QSOs is PHL~1092 (Bergeron \\& Kunth 1980, 1984; Kwan et al. 1995).  Its \\ion{Fe}{2}~$\\lambda$4570/H$\\beta$ ratio is 5.3, and its line widths are only $1800$~km~s$^{-1}$ (Bergeron \\& Kunth 1984).  Only a weak X-ray detection of PHL~1092 was made by {\\it Einstein} (Wilkes et al. 1994). Supported by the results of a {\\it ROSAT\\/} observation of PHL~1092, we argue that the NLS1s and their QSO analogs can be understood as a single phenomenon, the ``I~Zw~1'' objects.  In addition to sharing their extreme X-ray behavior, PHL~1092 might illuminate fundamental puzzles about this distinctive class. ",
        "conclusions": "We analyzed the {\\it ROSAT\\/} PSPC observation of the luminous \\ion{Fe}{2} QSO PHL~1092, and conclude that it deserves to be classified as a luminous narrow-line Seyfert~1 galaxy for the following reasons:  \\noindent 1) Its optical emission-line spectrum conforms with the NLS1 classification (Osterbrook \\& Pogge 1985; Goodrich 1989).  \\noindent 2) Its soft X-ray spectrum is extremely steep, with photon index $\\Gamma = 4.17 ^{+0.63}_{-0.50}$, very similar to the spectrum of the NLS1 IRAS 13224--3809.  \\noindent 3) The large amplitude variability of its X-ray emission over long time scales is similar to that of WPVS 007 (Grupe et al. 1995b) and PG 1404+226, both of which are NLS1s with ultra-soft X-ray spectra and strong \\ion{Fe}{2} emission.  \\noindent 4) Its rapid X-ray variability requires high efficiency, similar to the behavior of IRAS~13224--3809, and especially PKS~0558--504, which are also I~Zw~1 objects.  \\noindent We also found a correlation between X-ray spectral index and X-ray luminosity which may eventually help to constrain models of accretion in I~Zw~1 objects.  The possibility that rotating black holes may be implicated in the explanation of both the unusual X-ray and optical properties of this class is suggested."
    },
    "9603/astro-ph9603014_arXiv.txt": {
        "abstract": "The calibration of the Type Ia supernova distances using HST observations of Cepheid variables is discussed. A new maximum likelihood method of calibration is applied to derive the PL relation for a composite sample of Cepheids in the LMC and in the SNIa host galaxies NGC5253 and IC4182. Our results show that the calibration of the Cepheid PL relation is robust both to sampling error and to luminosity and period selection effects. Hence, the outstanding uncertainty in deriving estimates of $H_0$ from SNIa remains the dispersion of the SNIa luminosity function, and not unresolved systematic errors in its Cepheid calibration. ",
        "introduction": "Type Ia supernovae (henceforth SNIa) have long been regarded as useful cosmological distance indicators because they are observable to large velocity distances and their luminosity at maximum light displays a small intrinsic dispersion. In e.g. Sandage \\& Tammann (1993) the Hubble diagram of 34 SNIa in or beyond the Virgo cluster was found to have an observed V band dispersion of $\\sigma(M_v) = 0.36$ mag. Moreover, the linearity of the Hubble diagram indicated that these SNIa were not significantly affected by peculiar motions (after correction for Virgo infall) or luminosity selection effects. The mean absolute magnitude of these SNIa was found to be ${\\overline{M_v({\\rm{max}}}}) = -19.47 + 5 \\log (H_0/50)$. In order to estimate $H_0$ one must therefore determine independently the distance to one or more SNIa host galaxy.  HST has measured distances to IC4182 (host of SN1937C) and NGC5253 (host of SN1895B and SN1972E) from observations of Cepheid variables. These data yielded $H_0 = 52 \\pm 9$, from SN1937C alone (Saha et al. 1994) and $H_0 = 54 \\pm 8$, from the average of the three SNIa (Saha et al. 1995). In each case the SNIa were assumed to lie at the mean of the luminosity function. More recently Riess, Press \\& Kirshner (1995) have used the shape of the light curve to better constrain the luminosity at maximum light of SN1972E (the only one of the three SNIa with sufficient quality photometry to apply their method) and find evidence that SN1972E was significantly overluminous -- yielding $H_0 = 67 \\pm 8$.  In each of these analyses the Cepheid distances were determined assuming a distance modulus of $\\mu = 18.5$ for the LMC and fixing the slope of the PL relation in V and I to that obtained from a fiducial sample of LMC Cepheids (Madore \\& Freedman 1991). This was partly an attempt to avoid {\\em Malmquist bias\\/} -- i.e. a systematic error in the distance determinations due to V-band luminosity selection effects in each HST-observed galaxy (HOG), for which there appeared to be some evidence (Saha et al. 1994). Nevertheless, adopting the LMC slope left the results susceptible to two further (possibly systematic) uncertainties: sampling error due to the finite size and different period range of the LMC and HOG Cepheids; and a possible intrinsic difference in PL slope due to e.g. metallicity effects (c.f. Chiosi, Wood \\& Capitanio 1993). Although allowance has been made for these sources of uncertainty in the error budget for the published $H_0$ estimates, our aim in this work is to {\\em explicitly\\/} address their impact on $H_0$ by fitting PL relations to a composite sample of Cepheids in {\\em both\\/} the LMC and each HOG. ",
        "conclusions": "Using the Cepheid apparent distance moduli deduced from the fourth example calibration -- i.e. $\\mu(V) = 28.43 \\pm 0.08$ to IC4182 and $\\mu(V) = 28.16 \\pm 0.10$ to NGC5253 -- and assuming both SNIa to lie at the peak of the luminosity function we estimate $H_0 = 50 \\pm 9$ from SN1937C and $H_0 = 56 \\pm 11$ from SN1972E. The error estimate on $H_0$ is calculated by adding in quadrature the uncertainties on the Cepheid distance modulus and the SNIa apparent magnitude, and a further (conservative!) uncertainty of 0.15 mag. to allow for differential extinction between the Cepheids and SNIa in each HOG (see Saha et al. 1994). Note that in Saha et al. (1995) a value of $H_0 = 58 \\pm 9$ was obtained from SN1972E alone. Thus we find that our estimates of $H_0$ are very slightly reduced -- which ostensibly appears to be consistent with the general trend that luminosity selection effects tend to positively bias estimates of $H_0$. It is interesting to note, however, that if we adopt instead the distance moduli estimated by the second (smaller) calibration, using the same range of periods in the LMC and HOG, then our estimates for $H_0$ are 54 (IC4182) and 60 (NGC5253), which are both larger than the Saha et al. values. This comparison demonstrates that sampling error can have just as large an effect as selection bias on the value of $H_0$, causing it to be erroneously decreased {\\em or\\/} increased, although the important point of our results is that {\\em both\\/} effects are shown to be very small here.  If we apply the LCS correction of Riess et al. (1995) to SN1972E we instead find $H_0 = 65 \\pm 8$, which is also still in excellent agreement with the value deduced from the Cepheid distance to M96 (Tanvir et al. 1995) and the SN type II method of Schmidt et al. (1994) which completely by-passes the Cepheid distance scale.  Our analysis confirms that the Cepheid distances derived to these two SNIa host galaxies appear secure (at least provided that the LMC distance is secure) and accounting for possible sampling error and V-band luminosity selection does not significantly change the derived distance moduli. The outstanding uncertainty in estimating $H_0$ with SNIa is therefore the dispersion in the SNIa luminosity function at maximum light; this fact underlines the difficulty in making reliable statistical conclusions from only 2 or 3 data points. The LCS method offers one solution to this problem by reducing the dispersion, although the validity of luminosity--LCS correlations has recently been questioned in Tammann \\& Sandage (1995). Clearly measuring more distances to SNIa host galaxies would be a better solution. Indeed, since this conference took place distances to three more SNIa have appeared in preprints, and the issue of whether SNIa support a long or short distance scale should soon be resolved. Whatever the outcome, we conclude that the reliability of Cepheid distances in determining the SNIa zero point is not in doubt."
    },
    "9603/astro-ph9603146_arXiv.txt": {
        "abstract": "This paper describes new results on the identification of the complex gravitational lens responsible for the double quasar Q2345+007. A gravitational shear field was detected recently about 45\\arcsec\\ away from the QSO, centered on an excess of faint blue galaxies. The redshift distribution is still unknown, so the mass and the photometric properties of the deflector are still a matter of debate. We present deep photometric data obtained in the near-IR (J and K'), which are used together with preexisting optical BRI photometry to build spectral energy distributions for all the objects of the field, and to derive a photometric redshift estimate by comparison with synthetic spectrophotometric data. We propose a statistical method to analyse the redshift distribution, based on the cumulative histogram of the redshift ranges allowed for the different objects. An excess of galaxies at a redshift of $z \\simeq 0.75$ is clearly detected in the field of Q2345+007, with a 2D distribution showing a maximum located at the center of the weak shear field. Besides, the redshift inferred for this cluster is also compatible with that found for an absorption system in the spectrum of the B component of the quasar. We interpret this overdensity of objects as a distant cluster of galaxies responsible for the gravitational shear field. Two other redshift concentrations are studied: $z = 0.28$ which corresponds to the spectroscopic redshift of three galaxies but for which no strong excess of objects is identified, and $z \\simeq 1.2$, where an excess of galaxies is also detected, but with a rather smooth 2D distribution over our field of view. We also discuss the existence of other possible excesses of galaxies at redshift planes compatible with the absorption systems detected in the spectra of the QSOs. Most cluster-member candidates at $z \\simeq 0.75$ are undergoing a star-formation process or are burst systems where the star formation stopped between 1 and 3 Gyr ago. ",
        "introduction": "Since its discovery (Weedman et al. 1982), the nature of the double quasar Q2345+007 has been a controversial matter. The spectra of both quasars and their redshifts are very similar ($z = 2.15$), therefore the gravitational lensing hypothesis was suggested at once (Foltz et al. 1984). The distance between the two quasars is $7.1\\arcsec$, a value so high to be explained by a single galaxy deflector that Subramanian \\& Chitre (1984) suggested a double lens model. The best model they found was that of a galaxy in the central part of a cluster at a redshift close to 1. Nevertheless, the first deep visible images of the field did not show any trace of a lens candidate around the two images of the QSO (Tyson et al. 1986). This negative result leaded Tyson et al. to conclude that the mass to light ratio for the lens might be at least 1000, a value which is compatible with the hypothesis of a massive halo of dark matter suggested by Narayan et al. (1984). Conversely, Steidel \\& Sargent (1990) and Weir \\& Djorgovski (1991) concluded that the double quasar is probably a physical pair. Nevertheless, after a careful analysis of both spectra, Steidel \\& Sargent (1991) showed that they are really similar in many respects (the redshift, the emission lines as well as the absorption lines) and they definitely favour the gravitational lens hypothesis. Besides, the metal absorption systems identified in both spectra, at redshifts ranging from $z = 0.754$ to $z =1.98$, allow to suspect the presence of lens candidates at these redshifts. The close neighbourhood of the quasar was also studied in the near IR, up to a magnitude $K' = 20$, in a recent paper by McLeod et al. (1994), but the detection of a lens candidate was still negative.  All the works reviewed before to search for the lens deflector assume implicitely that the main deflector must be located close to the two QSO images. In a different approach, Bonnet et al. (1993) reported the detection of a shear field due to the weak gravitational distortion of the background sources by a mass distribution compatible with that of a cluster of galaxies. Its center is located about $45\\arcsec$ away from the double QSO, close to two bright galaxies identified at a redshift of $z=0.28$. Moreover, several arclets candidates were also identified, reinforcing the hypothesis of a strong deflecting mass. Deeper optical imaging confirmed the suspected excess of faint blue  galaxies associated with the lens (Mellier et al. 1994, Fischer et al. 1994). All these authors argued that its redshift could be higher than 1 because of the low surface brightness and color indices, and even as large as 1.5 if the excess is associated with the absorbing systems of the QSOs. But a precise value for the redshift cannot be inferred from their observations. A secondary small clump of blue galaxies was found close to the double quasar. As there are absorption lines in the spectra of both quasars at $z = 1.49$, this value for the redshift was considered as the most convincing one.  The galaxies in excess in the field of Q2345+007 are so faint that they are hardly observable in spectroscopic mode, and the redshift estimate of the deflecting agents is still a matter of debate. In this paper, we propose a different approach by using deep near-IR J and K' photometry of the field in complement of the existing deep B,R and I photometry. We show that from this spectral information, which covers a wide wavelength range, and a careful analysis of the photometric errors, it is possible to infer a more constrained {\\it photometric redshift} for each object in the field. The observed spectral energy distribution (hereafter SED) is compared with the predictions for different redshifts and spectromorphological types of galaxies. If detectable, a cluster along the line-of-sight must appear as an excess of objects at a particular value of the redshift distribution with respect to an empty field. A discussion about the properties and limits of such a photometric-redshift technique can be found in previous papers by Couch et al. (1983) and Ellis et al. (1985). Recently, Belloni et al. (1995) have applied successfully a similar photometric method to study the population of galaxies in a distant cluster at $ z = 0.41$. Although the photometric redshift has limitations in the case of individual objects, we show that the method can be applied successfully to detect an excess of objects at a given redshift through a statistical analysis of the photometric redshift ranges allowed for the  different objects in the field. This is the basic principle of the method we apply on the field of Q2345+007.  The outline of the present paper is the following. In the second section, we briefly summarize the photometric data. The data reduction and the construction of the photometric catalogue are discussed in section 3. The method of photometric redshifts is introduced and discussed in section 4, together with two illustrative examples on well-known clusters of galaxies at different redshifts. In section 5, the population of galaxies responsible for the excess in number counts at $z \\sim 0.75$ is characterized and its spatial distribution is determined. We also discuss in section 6 on the existence of other concentrations at $z=0.28$ and $z \\simeq 1.2$, as well as at the redshifts identified in the absorption spectra of the QSO. The SEDs of the arclet-candidates are studied in section 7. The final discussion and conclusions are in section 8. We give the photometry of the double quasar in appendix.  \\begin{table*} \\caption[]{Journal of observations, characteristics of the photometric system (filters + detectors) and photometric properties} \\begin{flushleft} \\begin{tabular}{lllrrrrrrrrrrrr} \\hline\\noalign{\\smallskip} Date & filter & Detector & Field & Exp. & FWHM & $\\lambda_{eff} $ & $\\Delta\\lambda $ & $T_{max}$ & compl. & limiting & $\\mu_{\\lambda} (1 \\sigma$) \\\\ &  &  & $(\\arcmin)$ & (sec) & (\\arcsec) & (nm) & (nm) & &  mag. & mag. &  mag/$\\arcsec^2 $ \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 13-14-15/10/90 & B$_J$ & RCA2 & $2.1 \\times 3.4$& 14500 & 1.2 & 447.8 & 142.7 & 0.82 &  27.5 & 29.0 & 28.7\\\\ 16/10/90 & R & RCA2 & $2.1 \\times 3.4$& 3000 & 1.4 & 645.8 & 112.4 & 0.73 &  26.2 & 26.8 & 26.6 \\\\ 14-15-17/10/90 & I & RCA2 & $2.1 \\times 3.4$& 7200 & 1.2 & 812.8 & 126.0 & 0.45 &  25.8 & 26.2 & 26.2 \\\\ 5-7/10/93 & J & NICMOS3 & $2.1 \\times 2.1$& 9950 & 1.2 & 1237.0 & 203.1 & 0.94 &  24.5 & 25.5 & 25.4 \\\\ 6/10/93 & K' & NICMOS3 & $2.1 \\times 2.1$& 4295 & 1.1 & 2103.2 & 359.5 & 0.96 &  22.8 & 23.5 & 23.1 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} \\end{table*} ",
        "conclusions": "The present work suggests that Q2345+007 is probably the result of a very complex gravitational lens involving several mass concentrations at different redshifts. We confirm the presence of a group of galaxies rather than a cluster at $z=0.28$. The main result is the identification and characterization of a distant medium-to-rich cluster at $z=0.75 \\pm 0.08$, which appears as an excess in the redshift distribution with respect to a blank field. The center of the cluster is located close to the center of the shear-field. Bonnet et al. (1993) already analysed this field and they obtained a good fit of the shear pattern by a singular isothermal sphere (SIS) with $\\sigma = 1200 km s^{-1}$ at z=1. Assuming that the bulk of the mass in the lens is associated with the cluster, the velocity dispersion required in order to produce the same mean shear pattern should range from $620$ to $905 km s^{-1}$, the more probable value beeing $790km s^{-1}$. This value was estimated with the following assumptions: the mass-distribution follows a SIS, the redshift of the lens is $0.65 \\leq z \\leq 0.85$, and the shear field measured comes mainly from sources around $z \\simeq 1.2$ and $z \\sim 1.8$. The resulting K-corrected M/L ratios in the inner $300 h_{50}^{-1} kpc$ radius are $M/L_B \\simeq 250 M_{\\odot}/L_{\\odot}$ and $M/L_J \\simeq 40 M_{\\odot}/L_{\\odot}$.  Most cluster-member candidates at $z \\sim 0.75$ show SEDs with signatures of a recent episode of star-formation. About 50\\% of them are undergoing a star-forming process whereas the others are more likely burst-systems where the star-formation stopped between 1 and 3 Gyr ago. An old population of stars is present in about 50\\% of galaxies. These secondary-burst systems are similar to the population of galaxies that Barger et al. (1995) find in several medium-redshift clusters (about 30\\% of the population in such clusters) but, in our case, more than 40 \\% of the population within the completeness limit in B$_J$ shows bursts, including the 10 brightest galaxies. Only 2 galaxies from this sample show SEDs corresponding to a pure old-population of stars, such as a non-evolved E-type galaxy, and they are relatively faint. These results indicate that the evolutionary state of the bulk population of galaxies is different from what is observed for medium-rich clusters at lower redshifts. The existence of short bursts of star formation at redshifts $0.5 \\leq z \\leq 1$ involving low-mass galaxies has been suggested by different authors to explain the results on deep galaxy counts (for example, Cowie et al. 1991, Lilly 1993, Babul \\& Ferguson 1995). According to our results, these star-formation processes could also affect the bulk population of galaxies in clusters within this redshift range.  Two other excesses of galaxies appear in the field of the double QSO with respect to the blank field, which points out the existence of clustering at redshift-strips $z \\simeq 1.2$ and $z \\sim 1.8$. In both cases, the 2D distribution in projected number-density is quite smooth, and the objects tend to be located around or close to the cluster at $z \\sim 0.75$. These objects are faint, beyond the completeness magnitude in at least one filter in most cases, especially for the excess at $z \\sim 1.8$. But, even if the 2D distribution has to be taken with caution, 95\\% of objects compatible with  $z \\sim 1.8$ are located in the field 1. As they suffer the gravitational amplification induced by the foreground cluster, they are probably detected in the field 1 just because the magnification bias tend to enhance their total luminosities in this region compared to the blank field. A systematic effect of shear should appear on this population if this identification is correct. A more detailed study is needed to measure the shear on this pre-selected sample of high-redshift objects, using deep and well sampled images. In any case, we can already say that the typical magnitudes of objects belonging to the redshift-strips at $z \\sim 1.2$ and $z \\sim 1.8$ (B$_J$  $\\geq 27$) make them compatible with the population used by Bonnet et al. (1993) to determine the shear field. The detection of a so faint shear-effect was probably made easier (or even possible) by the existence of clustering along the line-of-sight. Concerning the double QSO, the existence of several excesses at different redshift-strips indicates that we are probably dealing with a complex lens. The high separation of the 2 images is the combined effect of these lens-planes, including the faint galaxy in the close neighbourhood of the fainter QSO reported by Fischer et al. (1994). As a final remark, we can mention that none of the relatively blue objects detected $10\\arcsec$ around the double QSO has a photometric redshift assigned, because their SEDs are quite different from the models adopted in this paper. Concerning the complexity of the lens, a similar result was found by Angonin-Willaime et al (1994) in the field of the double QSO 0957+561, where the main lens is the sum of a bright gE galaxy and of the cluster of galaxies associated with it. The field of QSO 0957+561 is also contaminated by a background group of galaxies identified spectroscopically, and many rather bright foreground galaxies. Their effects on the high separation of the 2 images are not clear and probably not dominant, but it should affect strongly the shear field of the faint background galaxies.  Among the arclets identified by Bonnet et al. (1993) and by Mellier et al. (1994), A1 and A3 are more likely elongated cluster members rather than gravitational arclets, whereas A4 seems to be a foreground object, only marginally compatible with the cluster redshift. No other elongated arclet-candidates are visible. The only exception is A2, for which a high redshift solution exists, at $z \\sim 3.6-3.8$, although a more reliable one identifies it as a member of the group at $z \\sim 0.3$. It is important for the lens modeling to confirm the high redshift hypothesis for this object, and it can be done because $Ly\\alpha$ is expected at $\\lambda \\sim 5800 \\AA$.  One of the priorities for the near future is to complete an extended spectroscopic survey on this field to confirm the predicted distributions in redshift. Recently, some additional spectroscopic data have been obtained at the CFHT, and the reduction is going on. The presence of strong emission-lines is expected in the visible range for those brightest cluster-member candidates which are undergoing an active process of star formation, making possible the measure of the redshift. It is worth noting that the study of the SEDs from B$_J$ to K' allows not only an estimate of the photometric redshift but also a determination of the optimal wavelength and feasibility for a successful measure of the spectroscopic redshift. In a more general context, we expect to check the reliability of the statistical method at high-redshift. Combining this technique of multicolor analysis with the systematic studies of weak-shear fields, and lensed QSOs in particular, is a very promising way to characterize the visible counterparts of large scale inhomogeneities."
    },
    "9603/astro-ph9603021_arXiv.txt": {
        "abstract": "Since cosmology is no longer ``the data-starved science\", the problem of how to best analyze large data sets has recently received considerable attention, and Karhunen-Lo\\`eve eigenvalue methods have been applied to both galaxy redshift surveys and Cosmic Microwave Background (CMB) maps. We present a comprehensive discussion of methods for estimating cosmological parameters from large data sets, which includes the previously published techniques as special cases. We show that both the problem of estimating several parameters jointly and the problem of not knowing the parameters a priori can be readily solved by adding an extra singular value decomposition step.  It has recently been argued that the information content in a sky map from a next generation CMB satellite is sufficient to measure key cosmological parameters ($h$, $\\Omega$, $\\Lambda$, {\\etc}) to an accuracy of a few percent or better --- in principle. In practice, the data set is so large that both a brute force likelihood analysis and a direct expansion in signal-to-noise eigenmodes will be computationally unfeasible. We argue that it is likely that a Karhunen-Lo\\`eve approach can nonetheless measure the parameters with close to maximal accuracy, if preceded by an appropriate form of quadratic ``pre-compression''.  We also discuss practical issues regarding parameter estimation from present and future galaxy redshift surveys, and illustrate this with a generalized eigenmode analysis of the IRAS 1.2 Jy survey optimized for measuring $\\beta\\equiv\\Omega^{0.6}/b$ using redshift space distortions. ",
        "introduction": "\\label{IntroSec}  The problem of analysis of large data sets is one that, until recently, has not been a major concern of cosmologists. Indeed, in some areas no data existed to be analyzed. In the last few years, this situation has rapidly changed. A highlight in this transition has been the discovery of fluctuations in the cosmic microwave background (CMB) by the Cosmic Background Explorer (COBE) satellite (Smoot {\\etal} 1992). In its short lifetime, COBE produced such a large data set that a number of sophisticated data-analysis methods were developed specifically to tackle it. In addition, the advent of large galaxy redshift surveys has created a field where the data sets increase by an order of magnitude in size in each generation. For instance, the surveys where the object selection was based on the Infrared Astronomy Satellite (IRAS) contain several thousand galaxies: $\\sim 2,000$ (QDOT; Lawrence {\\etal} 1996) $\\sim 5,000$ (Berkeley 1.2Jy; Fisher {\\etal} 1995) and $\\sim 15,000$  (PSC-z; Saunders {\\etal} 1996). The proposed next-generation surveys will have much larger numbers of objects --- around 250,000 in the Anglo-Australian Telescope 2 degree Field galaxy redshift survey (Taylor 1995) and $\\sim 10^6$ for the Sloan Digital Sky Survey (Gunn \\& Weinberg 1995). Similarly plate measuring machines, such as the APM at Cambridge and SuperCOSMOS at the Royal Observatory, Edinburgh, can produce very large catalogues of objects, and numerical simulations of galaxy clustering are even now capable of producing so much data that the analysis and storage of the information is in itself a challenge.  A standard technique for estimating parameters from data is the brute force maximum likelihood method, which illustrates why people have been driven towards developing more sophisticated methods. For $n$ data items ({\\eg}, pixels in a CMB map, or Fourier amplitudes from a transformed galaxy distribution), the maximum likelihood method requires inversion of an $n\\times n$ matrix for each set of parameter values considered --- and this is for the simplest possible case where the probability distribution is Gaussian. Since a next-generation CMB satellite might produce a high resolution sky map with $\\sim 10^7$ pixels, and the CPU time required for an inversion scales as $n^3$, a brute force likelihood analysis of this type of data set will hardly be feasible in the near future.  Fortunately, it is often possible to greatly accelerate a likelihood analysis by first performing some appropriate form of data compression, by which the data set is substantially reduced in size while nonetheless retaining virtually all the relevant cosmological information. In this spirit, a large number of data-compression methods have been applied in the analysis of both CMB maps ({\\eg} Seljak \\& Bertschinger 1993; G\\'orski 1994) and galaxy redshift surveys ({\\eg} Davis \\& Peebles 1983; Feldman {\\etal} 1994; Heavens \\& Taylor 1995). A powerful prescription for how to do this optimally is the Karhunen-Lo\\`eve eigenvalue method (Karhunen 1947), which has recently been applied to both CMB maps (Bond 1994; Bunn 1995; Bunn \\& Sugiyama 1995) and redshift surveys (Vogeley 1995; Vogeley \\& Szalay 1996). The goal of this paper is to review the more general framework in which these treatments belong, and to present some important generalizations that will facilitate the analysis of the next generation of cosmological data sets.   The rest of this paper is organized as follows. In Section~\\ref{FisherSec}, we review some useful information-theoretical results that tell us how well parameters can be estimated, and how to determine whether a given analysis method is good or bad. In Section~\\ref{CompressionSec}, we review the Karhunen-Lo\\`eve data compression method and present some useful generalizations. In Section~\\ref{ApplicationsSec} we illustrate the theory with various cosmology applications, including the special case of the signal-to-noise eigenmode method. In Section~\\ref{PreCompressionSec} we discuss limitations of method and possible ways of extending it to make the analysis feasible for huge data sets such as a $10^7$ pixel future CMB map. Finally, our results are summarized and discussed in Section~\\ref{ConclusionsSec}. ",
        "conclusions": ""
    },
    "9603/astro-ph9603092_arXiv.txt": {
        "abstract": "The age of three of the oldest clusters -- M15, M68, M92 -- has been redetermined. We use the latest EOS and opacity data available for calculating both isochrones and zero age horizontal branches and employ the brightness difference between turn-off and horizontal branch to determine the cluster age. Our best ages for all three clusters are about 13 Gyr, and even smaller ages are possible. Our results help to reconcile cluster ages with recent results on the age of the universe determined from the Hubble constant. ",
        "introduction": "If the age of the universe is obtained from the cosmological expansion  by determining the Hubble constant $H_0$, it turns out to be less than 15 Gyr for $H_0\\, \\ga\\, 50 {\\rm km/(s\\, Mpc)}$ for values of the density parameter $\\Omega_0\\, \\ga\\, 0.4$. The most recent investigations into $H_0$ (for a review, see van den Bergh 1994) yield values between 60 and 90, which limit the age of the universe to less than 13 Gyr for any reasonable $\\Omega_0$, even if a cosmological constant is allowed. This number even reduces to 10 Gyr, if $H_0 \\approx 75$ (from Cepheids or SNIa). Clearly, even if one stretches the limits, the universe appears to be younger than $\\approx 15$ Gyr.  On the other hand, the classical method of determining the age of the oldest known stellar objects, the isochrone fitting to metal-poor globular clusters, consistently yields ages above $\\approx 14$ Gyr for the bulk of metal-poor halo clusters, with the majority of the clusters being around $16$ Gyr and the oldest ones up to $18$ Gyr old. An additional Gyr has to be added to this for the time between the genesis of the universe and the creation of the first clusters. This ``Conflict over the age of the Universe'' has been brought to the point by Bolte \\& Hogan (1995) by the example of M92, probably the best observed very old cluster, whose age they give as $15.8 \\pm 2.1$ Gyr (this range was used by Kennicutt, Freedman \\& Mould~(1995) to illustrate the relation -- and conflict -- with $H_0$).  It therefore is necessary to reconsider age determinations of globular clusters and to investigate the errors more carefully. In this {\\sl Letter} we report about new age determinations of three of the oldest clusters, including M92, using standard approaches, but latest equation of state and opacities for the stellar models. Cluster ages were determined by using the difference in visual magnitude ($\\Delta(V)$) between main-sequence turn-off and horizontal branch (HB).  Our work is similar to Chaboyer \\& Kim (1995; CK95), who used the same method and already found  a reduction of cluster ages of about 7\\% (or 1.2 Gyr for the oldest clusters). However, they had determined $\\Delta(V)$ from the difference between the theoretical turn-off luminosity $V_{TO}$ and an empirical horizontal branch brightness $V_{HB}$ (based on RR~Lyrae luminosities). In contrast, we calculate theoretical ZAHBs as well and fit both isochrones and ZAHB to the observations. This way, the influence of the new input physics is taken into account twice.  In the next section we will shortly describe our model calculations; Sect.~3 contains the comparison with the globular clusters M15, M68 and M92. We have concentrated on these three very old clusters, because the intention of this {\\sl Letter} is to reconcile cluster and cosmological ages. In a forthcoming paper we will discuss a larger set of clusters in a broader context. The conclusions will follow in the last section, as usual. ",
        "conclusions": "In this paper we have reexamined the age of three of the oldest globular clusters (M15, M68, M92). Changes to earlier work included the use of the latest OPAL EOS and opacity tables, supplemented by low-temperature opacity tables for {\\em exactly} the same compositions. In particular, $\\alpha$-enhancement within the metals could be taken into account. For appropriate chemical compositions both isochrones {\\em and} ZAHB models were calculated and cluster ages were derived from the $\\Delta(V)$ difference between TO and ZAHB. We find that the age of the clusters is about 13 Gyr, and that the three clusters are practically coeval. These ages are lower by up to 4--5 Gyr as compared to earlier results, e.g. Salaris et al.\\ (1993). They made use of the bolometric corrections of Vandenberg \\& Bell (1985) for the TO, while we switched to BK92 and BK78. This, as we have checked, lowers the ages already by almost 2 Gyrs. The additional age reduction results from the new OPAL EOS. Our results confirm and extend those of CK95 and Mazzitelli et al.\\ (1995). The further reduction as compared to CK95 is due to our additional ZAHB calculations.  The derived ages {\\em can further be reduced} by the following means: a) the use of a higher helium content, justified by Big Bang Nucleosynthesis results on the best-fit predicted primordial helium content of 0.247 (Hata et al.\\ 1996) or by observations including a large systematic error (Olive \\& Steigman 1995); b) the use of the new Kurucz (1992) bolometric corrections; we did not use them because of the problems encountered with RGB colors; c) the inclusion of diffusion (Chaboyer et al.\\ 1992). Each of these factors reduces the ages by another 0.5 Gyr at least, such that an age of the oldest globular clusters of 12 Gyr seems to be in reach.  Interestingly, the mean age obtained for the three clusters is almost coincident with the age of the Disk (10--12 Gyr) obtained by Hernanz et al. (1994) by means of the luminosity function of the white dwarfs in the solar neighbourood, thus implying that the Galactic Disk began to form without time delay with respect to the halo.  In a forthcoming paper we will present extended results for a large sample of clusters. The bottom line of the present {\\sl Letter} is that with updated physics the oldest globular clusters are only 13 Gyr (or less) old. This reduces the ``Age Conflict'' drastically. Stellar evolution theorists have done the first step. It is now for the cosmologists to confirm that $H_0 \\approx 50$."
    },
    "9603/gr-qc9603008_arXiv.txt": {
        "abstract": "In this essay, we introduce a new effect of gravitationally induced quantum mechanical phases in neutrino oscillations. These phases arise from an hitherto unexplored interplay of gravitation and the principle of the linear superposition of quantum mechanics. In the neighborhood of a $1.4$ solar--mass neutron star, gravitationally induced quantum mechanical phases are roughly $20 \\%$ of their kinematical counterparts. When this information is coupled with the mass square differences implied by the existing neutrino--oscillation data we find that the new effect may have profound consequences for  type-II supernova evolution. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603050_arXiv.txt": {
        "abstract": "We report on the detection of dark matter in the cluster of galaxies Abell~2163 using the weak gravitational distortion of background galaxies, and an analysis of the X-ray emission {}from the cluster. We find that while the qualitative distributions of the cluster light and the dark matter are similar --- shallow and extended, with significant substructure --- the X-ray morphology shows a more regular overall appearance. We interpret the joint lensing and X-ray observations as a signature of a merger event in the cluster. We present new ROSAT/HRI data and reanalyze ROSAT/PSPC data, accounting for the effect of a varying background to determine the best fit parameters in the $\\beta$-model formalism. We combine the surface brightness fits with two determinations of the radial temperature profile to determine the total mass. Although there are slight variations in the total mass determinations introduced by the uncertainties in the $\\beta$-fit, the main contributor to the error arises {}from the uncertainties in the temperature determinations.  Even though the morphologies of the dark matter/light and X-ray gas are quite different, we find that the total mass determined {}from the X-ray and weak lensing estimates are consistent with each other within the $2\\sigma$ error bars, with the X-ray inferred mass a factor of $\\simeq 2$ larger. However, as the lensing mass estimates are differential (the surface density at any point is determined relative to the mean in a control annulus), the shallow, extended nature of the mass profile biases the lensing inferred mass downwards. We estimate the correction for this effect and find very good agreement between the corrected lensing and X-ray results. We determine the gas mass fraction in this cluster and find $f_g \\simeq 0.07 \\;h^{-3/2}$ at all radii and a constant mass-to-light ratio of $M/L_V = (300 \\pm 100) \\;h M/L_{\\odot V}$. ",
        "introduction": "Determining the masses of the visible and the dark components of the galaxy clusters and mapping out their relative distributions is the key towards understanding  both the dynamical state of the clusters and their evolutionary history. As the largest clearly defined objects in the Universe, galaxy clusters are important cosmological probes. They are used, for example, to measure the matter composition of the Universe - e.g., the luminous baryon to dark matter ratio (\\cite{white93}) or the mass-to-light ratio (e.g., \\cite{blumenthal84}) on large scales. They are also important tracers of the large scale structure as measured through the mass function or the spatial correlation function (e.g., \\cite{efstathiou95}).  Ideally, one would like to study X-ray bright clusters whose galaxies population has been spectroscopically studied and that also induce measurable, even if only weak, gravitational distortions in the images of faint background galaxies. The X-ray emission of the hot intracluster plasma provides an attractive method to determine cluster masses (\\cite{sarazin86}; \\cite{hughes89}). The derivation of the gravitational mass using X-ray data rests on the assumption, however, that the gas is in hydrostatic equilibrium in the cluster potential. Since the dynamical time scale for the formation of clusters is comparable to the Hubble time and, as many clusters are found to have substructure implying that the are dynamically young and unrelaxed,  the X-ray mass derivations are not without uncertainties.  The discovery of gravitational lensing effects, both strong and weak, in galaxy clusters opened a new way to probe the gravitational potential of clusters, which is free of any assumption of the cluster dynamical state (e.g., \\cite{fort94}). Strong lensing effects resulting in obviously distorted images of background galaxies -- observed as arcs -- were first used to obtain cluster masses by modeling the lens effects. However, such analyses are highly model dependent and are limited to the central regions of the cluster.  The analyses of the weak gravitationally induced distortions in the images of faint, background galaxies, on the other hand,  offer a unique opportunity to directly probe the total mass distribution (\\cite{tyson90}; \\cite{ks93}; see also \\cite{squires96} and references therein). The weak lensing effects can be measured and inverted to derive the mass distribution in the cluster in a model independent way and free of any assumption of the cluster symmetry and dynamical state. This permits both a non-parametric determination of the total mass, and a 2D map of the total mass distribution in the cluster.  Taken together, the weak lensing analysis and an analysis of X-ray observations offers a unique possibility to probe the relative distributions of the gas and the dark matter and study the dynamical relationship between the two. It is worth noting that, on a case-by-case basis, discrepancies in the masses determined by the two methods can naturally arise (for example, if the cluster is strongly elongated cluster along the line of sight). Thus, to be able to make general conclusions about the cluster population, the analyses need to be compared for several clusters for which both good optical and X-ray data are available.  Miralda-Escud\\'e \\& Babul (1995) (hereafter MB) performed the first joint X-ray and lensing study of the clusters A2218, A1689 and A2163. In the former two clusters they found that the mass in the central region implied by the observed giant arcs is a factor of 2-2.5 larger than the mass derived {}from X-ray data if the intracluster plasma is assumed to be isothermal at the observed temperature and in hydrostatic equilibrium in the cluster potential. A similar discrepancy, but not quite as strong, was confirmed by Kneib \\etal (1995) also for A2218 in a more detailed effort to model of the strong lensing effect. In contrast, the mass determined from the strong lensing in the regular cluster PKS0745-191 is in agreement with the mass derived from X-ray data, using a multiphase cooling flow model. If it is a common occurrence, the X-ray/lensing mass discrepancy could have important implications for quantities, such as the cluster gas fraction $M_{\\rm gas}/M_{\\rm tot}$, derived solely {}from X-ray data.  Typically, the cluster gas fraction is estimated to be $M_{\\rm gas}/M_{\\rm tot} \\geq 0.05 \\;h^{-3/2}$ (\\cite{white95}; \\cite{david95}) a result that has been a source of much discussion (\\cite{white92}; \\cite{babul93}; \\cite{white93}). It is worth noting, however, that the lensing mass determinations based on the strong lensing features are, however, restricted to probing the central regions of clusters.  More recently, Squires \\etal (1996) used the weak gravitational lensing distortions to map the mass distribution in A2218 out to a radius of $\\sim 0.5 \\Mpc$, and jointly analyzed the X-ray, optical and weak lensing results.  Under the assumption that the cluster mass distribution extends out to $\\sim 1 \\Mpc$, as indicated both by the lensing mass profile and the X-ray surface brightness profile, the results suggested that the lensing/X-ray mass discrepancy of kind found by MB may extend systematically beyond cluster core (although the combination of the uncertainties in the lensing and X-ray mass estimates, the latter largely due to poorly determined cluster temperature profile, did not preclude agreement between the two estimates). This discrepancy was recently confirmed by Loewenstein (1996) for both A2218 and A1689 using new ASCA determinations for the temperature profiles. Conversely, in a similar study of the clusters MS1455 and MS0016, Smail \\etal (1995) found that the lensing and X-ray masses are in agreement (although we note that the X-ray analyses in these latter cases were based on very sparse X-ray spectroscopic data and thus the uncertainties on the X-ray mass determinations were larger).  In this paper, we present a study of the galaxy, gas, and gravitational mass distribution in the cluster of galaxies A2163 ($z=0.201$). This cluster is the hottest cluster and one of the two most massive galaxy clusters known so far (\\cite{arnaud92}). Based on GINGA satellite measurements, A2163 has an X-ray temperature of $\\sim$ 14~keV and a X-ray luminosity of $6 \\times 10^{45}$~erg~s$^{-1}$ (\\cite{arnaud92}). It is one of the rare clusters for which the X-ray emission has been traced out to a distance similar to the virial radius; the emission was detected significantly with the ROSAT/PSPC up to $2.3 \\Mpc$ or 15 core radii (\\cite{elbaz95} -- hereafter EAB). While the total mass within that radius, derived {}from the GINGA and PSPC data, is exceptionally high (2.6 times greater than the total mass of Coma), the corresponding gas mass fraction, $\\sim 0.1 \\; h^{-3/2}$, is typical of other rich clusters.  A quick drop of the temperature at large radii ($\\sim$ 4~keV beyond $6$ core radii) was observed recently with ASCA, strongly constraining the total mass profile, assumed to follow a simple parametric law (\\cite{MMIYFT95} --- hereafter referred to as MMIYFT). The temperature distribution in A2163 also shows evidence for complex non-azimuthally symmetric temperature variations in the central regions (Markevitch \\etal 1994). A2163 exhibits the Sunyaev-Zel'dovich effect (\\cite{wilbanks94}) and is remarkable in the radio, having the most luminous and extended  halo yet detected (\\cite{herbig95}).  The spatially resolved  measurements of the the Sunyaev-Zel'dovich effect, combined with  the PSPC data, also confirm the decrease in the temperature in the outer part of the cluster (\\cite{holzapfel96}).  In spite of its exceptional X-ray and radio properties, the {\\em dynamical} state of A2163 is puzzling. The high mean X-ray temperature and high X-ray luminosity would suggest a massive cluster with a very deep potential well that ought to contain a plethora of strong lensing features such as arcs and arclets.  In fact, only two arcs have been observed and they lie at a relatively small distance {}from a brightest cluster galaxy. This is very much in contrast to the lensing features detected in other hot clusters (e.g., A2218).  The optical properties of A2163 are also quite unassuming in comparison to compact clusters such as A2218 and A1689.  It is classified as an Abell richness class 2 cluster, its central galaxy is not a cD galaxy, and the cluster galaxy distribution is irregular and extended.  On the other hand, A2163 has a very high velocity dispersion ($\\sigma=1680$~km/s) and a flat, very extended velocity histogram (Soucail \\etal 1996; Arnaud \\etal~1994). The optical data, together with the detailed X-ray morphology, have been interpreted as signature of a recent or ongoing merger of two large clusters (EAB; Soucail \\etal 1996).  For the present study, we acquired optical images in V- and I-bands using the Canada-France-Hawaii telescope. The total mass distribution of A2163 is determined {}from the distorted images of the faint background galaxies  using the algorithm of  Kaiser \\& Squires (1993) and amendments (\\cite{sk95}). We also derive the cluster galaxy light and surface number density distribution. We reanalyze the ROSAT/PSPC data of EAB, allowing for statistical uncertainties in the background, and utilize ASCA temperature information provided by MMIYFT. Furthermore we present new ROSAT/HRI data, which resolves the center of the cluster to a much higher accuracy than the PSPC data. We reinvestigate the total mass and gas mass fraction  determination based on X-ray data using the method developed by Neumann and B\\\"ohringer (1995) which is free of any assumption on the gravitational potential shape, in contrast with parametric methods used in previous studies of this cluster. We make qualitative comparisons of the morphology of the gas, galaxy and mass distribution and give new insight into the dynamical state of the cluster. We compare the X-ray and lensing total mass estimates, determine the baryon fraction as well as the mass-to-light ratio profile. Finally we consider the cosmological implications of our results. ",
        "conclusions": "\\label {sec:summary}  We have studied optical, gas and dark matter distributions in the cluster of galaxies Abell~2163. We have traced the cluster galaxy light distribution over a $7' \\simeq 1 \\Mpc$ field centered on the dominant central galaxy. Using the weak gravitational distortion of background galaxies and correcting for systematic effects that can bias the galaxy shapes, we have mapped the dark matter distribution in the cluster over the same scale. Combining the observed X-ray surface brightness profile {}from ROSAT/PSPC observations with the spectroscopically determined temperature profiles, we have estimated the projected total mass and gas mass distribution to a radius of $\\simeq 2 \\Mpc$. In the present case, we find agreement at the $2\\sigma$ level between the raw weak lensing and X-ray inferred masses,  with the X-ray mass determinations being a factor  of $\\simeq 2$ higher. Correcting for matter in the control annulus, as determined {}from the X-ray observations, yields a much better agreement between the two results. The extended nature of the mass distribution is consistent relatively weak shear profile that was measured in this cluster. The extended and clumpy nature of the cluster galaxy light distribution, the broad cluster galaxy velocity histogram, the high value of $\\beta_T$, the extended luminous radio halo, and the irregular mass distribution inferred {}from the weak lensing analysis, suggests that A2163 is a cluster in the process of formation.  The observations presented here have been a useful check of the consistency between the two mass determination methods and have enabled us to speculate on the dynamical state of the cluster. The somewhat unsatisfactory aspect of this is that, at least for the lensing analysis, the mass determinations were confined to a relatively small physical radius {}from the cluster center. Clearly, any conclusions about the universality of the gas mass fraction and mass-to-light ratio are subject to interpretations regarding the physics of clusters centers, and how representative a remarkable cluster like A2163 is of the Universe as a whole. A more pleasing comparison should, and now could, be done at over larger scales. Indeed, the technology now exists for such large field lensing observations with the MOCAM 14$^\\prime$ and the UH $\\simeq 0.5^\\circ$ cameras at CFHT. One exciting example is that, in hierarchical clustering scenarios, clusters of galaxies tend to form at intersections of filaments.  In clusters such as A2163, which appear to be very much in formation, it may be possible to use large field observations of the weak lensing to detect filaments along which mass is flowing into the cluster. With the types of observations possible with the current generation of instruments, a wide range of outstanding issues can be probed regarding the dark matter content and distribution in the Universe."
    },
    "9603/astro-ph9603099_arXiv.txt": {
        "abstract": "Photoionization models dictate that many prominent quasar emission lines are sensitive to both the luminosity and shape of the quasars' high energy continuum - primarily the extreme ultraviolet (EUV) and soft X-ray continuum.  Unfortunately, the EUV band is severely obscured by Galactic absorption.  Using data from the adjacent UV and soft X-ray bandpasses, we initiate the first large-scale, multi-line investigation of correlations between the QSO soft X-ray continuum and line emission in a sample of QSOs observed by \\ein~ and \\IUE.  We present a new error analysis for objective, automated line measurements, which enables us to include the information contained in weak or undetected lines.  We tabulate more than 300 UV emission line equivalent widths from \\IUE~ spectra of 85 QSOs in the atlas of Lanzetta et al. (1993), then characterize the distributions of line equivalent and velocity widths (\\ew\\, and FWHM).  We then compare these line parameters to the QSO continuum spectral energy distributions from optical through soft X-ray wavelengths, using survival analysis to incorporate any non-detections for X-ray flux and/or UV emission lines.  Several correlations noted in previous studies are {\\em not} reproduced here. However, we illustrate that the exclusion of undetected lines from such studies may spuriously enhance apparent correlations.  We find significant correlations between \\ew\\, and UV luminosity (e.g., the well-studied Baldwin effect) for \\lya, \\civ, \\heii, and \\ciii.  \\wciii\\, and \\wheii\\, also show previously unreported correlations with X-ray luminosity which, for \\ciii, appears to be primary. The line ratios \\rciii~ and \\rheii\\, both show strongest dependence on \\lx.  \\wlya\\, correlates strongly with spectral slopes \\auv\\, and \\aox\\, (between 2500\\AA\\, and 2~keV), but {\\em not} with X-ray luminosity.  Using these results, we argue that one simple geometrical interpretation of the Baldwin effect (BEff) as a result of a distribution of disk inclinations is not plausible.  We also provide evidence that the BEff weakens or disappears when the line emission is correctly compared to the luminosity in the continuum bandpass relevant to its production.  We thus support the interpretation of the BEff as a change in spectral energy distribution with luminosity, and we predict that no BEff relative to X-ray luminosity should be found for Fe\\,II or Mg\\,II emission lines.  Extensions of our method to samples of a wider redshift/luminosity range will be presented in a later paper, which will test these predictions. ",
        "introduction": "\\label{intro}   \\subsection{The Ionizing Continuum of Quasars} \\label{seds}  The majority of the nearly 8000 quasars known to date were discovered either via their prominent optical and ultraviolet (OUV) emission lines, or from their distinct colors in these bandpasses. The production of emission lines in QSO spectra is widely attributed to photoionization and heating of the emitting gas by the UV to X-ray continuum (e.g., Ferland \\& Shields 1985, Krolik \\& Kallman 1988).  Individual emission lines from a given ion are particularly sensitive to photons of energy above the corresponding ionization threshold.  As an example, the continuum flux relevant to the production of Ly$\\alpha$ emission is above 13.6eV, while He\\,II$\\lambda$1640 is produced by photons above the 54eV ionization edge of He$^{+}$, which at 228\\AA\\, is in the EUV.  Note however that the production of many emission lines may be sensitive to continuum energy ranges both softer and harder than the ionization potential of the species in question because such photons may ionize from excited states and also heat the gas via free-free and H$^-$ absorption. Many important lines respond to the extreme ultraviolet (EUV) or soft X-ray continuum.  Unfortunately, the EUV band is severely obscured by Galactic absorption.  However, constraints on the EUV ionizing continuum are available both through analysis of the emission lines, and through the adjacent UV and soft X-ray windows.  Both radio-loud (RL) and radio-quiet (RQ) quasars are seen to have soft ($\\lapprox$1~keV) X-ray emission that exceeds the extrapolation from the power-law continuum observed at higher energies (e.g., Turner \\& Pounds 1989, Masnou et al. 1992). This X-ray `soft excess' has often been interpreted as the high energy continuation of the big blue bump (BBB), possibly thermal emission from the surface of an accretion disk (although see Barvainis 1993).  From the optical/UV side, the bump is an upturn in emission toward shorter wavelengths commonly observed in quasar spectral energy distributions (SEDs; e.g., Elvis et al. 1994). Somewhere in the EUV band, the SEDs must peak and turn down again to meet the observed X-ray emission.   \\subsection{ Emission Lines as Continuum Diagnostics} \\label{balqsos}  There are pressing reasons to investigate the relationship between available measurements of their high energy continuum and the OUV emission lines in QSOs.  First is to investigate observational constraints on photoionization models for the broad line region (BLR) of active galactic nuclei (AGN): do spectral energy distributions (SEDs) directly determine emission line strengths or line profile parameters?  Conversely, do similar emission line parameters in QSOs provide empirical testimony for similar high energy SEDs?  The overall similarity of QSO emission line spectra had been taken as evidence of fairly uniform, robust physical conditions in the BELR, which encouraged the assumption that clouds in the BELR inhabit a narrow swath of parameter space (in density, size, and ionization parameter).  Early photoionization pioneers such as Mushotsky \\& Ferland (1984) ran models on a single cloud. Refinements using cloud ensembles showed a reduced dependence of total line emission on intrinsic QSO SEDs (Binette et al. 1989).  Details of individual clouds or even clouds in a single ``zone'' can be lost in the mix, and correlations between continuum shape and observed line parameters diluted.  Baldwin et al.  (1996) reiterate that averaging of emission from clouds with a wide variety of properties (but uniformly large columns) results in QSO line spectra robustly consistent with those observed.  Correlations of \\ew\\, with SED would provide evidence against such models.  Perhaps emission lines can be used to infer the strength and shape of the high energy SED (Krolik \\& Kallman 1988, Zheng 1991), even in the presence of extrinsic effects such as absorption along the line of sight.  As an example, radio-quiet QSOs with broad UV absorption lines (BALs) are now known to exhibit markedly weak X-ray emission as a class (Green et al. 1995, Green \\& Mathur 1996).  The similarity of emission-line properties in BAL and non-BAL QSOs (Weymann et al. 1991) has been cited as evidence that orientation is the cause of the BAL phenomenon (i.e., {\\em all} radio-quiet QSOs have BAL clouds).  If similar emission lines indeed vouch for similar intrinsic high energy SEDs, then the large observed \\aox\\, values for BAL QSOs are likely to be caused by strong absorption along the line-of-sight rather than by differences in their intrinsic SEDs.  However, the UV and X-ray absorbers have yet to be positively identified as one (e.g., see the techniques of Mathur 1994).  Since BAL QSOs {\\em may} be heavily absorbed, the question of whether similar emission lines are testimony for similar intrinsic SEDs must be answered through study of line/continuum correlations in unabsorbed QSOs.  The simple question of whether line equivalent width \\ew\\, correlates with \\aox, for example, remains to be explored across a range of emission lines and for QSO samples spanning a range of luminosities.  \\subsection{\\bf The Baldwin Effect and Changes in Continuum Shape with Luminosity}  If the proportionality between line and continuum strength were linear, then diagnostics such as line ratios and equivalent widths would be independent of continuum luminosity.  Baldwin (1977) first noticed that in high redshift quasars, the equivalent width (hereafter, \\ew) of the CIV $\\lambda1550$\\AA\\, emission line in quasars decreases with increasing UV ($1450$\\AA) luminosity.  The Baldwin effect (BEff) was also found to be strong for ions such as OVI, NV, He\\,II, CIII], Mg\\,II, and Ly$\\alpha$ (e.g., Tytler \\& Fan 1992, Zamorani et al. 1992). Several possible explanations for the BEff have been offered, one being a dependence of blue bump strength on luminosity.  The shape of the continuum (i.e., the SED) of quasars does appear to correlate with luminosity.  In the UV regime, Zheng \\& Malkan (1993) found that the UV continuum increases in strength relative to the optical toward higher luminosities, and that the strength of the BEff decreases once the effect of the increasing UV (BBB) continuum is removed. In the X-ray bandpass, the largest, most uniform study -- ROSAT All-Sky Survey (RASS) observations of 908 QSOs in the Large Bright Quasar Survey (the LBQS/RASS; Green et al. 1995) -- confirmed earlier reports (e.g., Wilkes et al. 1994, Tananbaum et al.  1986) that the hypothetical power-law index between UV (rest $\\lambda2500$\\AA) and soft X-ray regimes, \\aox\\, increases significantly with luminosity. The increase in \\aox\\, is equivalent to a {\\em decrease} in soft X-ray relative to UV emission with increasing luminosity.  There are also hints (e.g., Schartel et al. 1996) that the soft X-ray spectral index \\ax\\, of QSOs may decrease with luminosity and/or redshift.  The decrease in \\ax\\, could mean that the soft X-ray spectrum {\\em hardens} with increasing redshift and/or luminosity \\footnote[1]{Redshift and luminosity dependence can be hard to disentangle in magnitude-limited surveys, but several recent results, e.g., Wilkes et al. 1994, confirm the primacy of (optical) luminosity in the correlation with \\aox.}. Alternatively, a soft excess may shift out of the \\ros\\, passband toward higher redshift, and/or move toward lower energies in higher luminosity sources.  Any or all of these trends of continuum shape could strongly influence the efficiency of the ionizing continuum, and should affect observed emission line strengths and ratios. Investigations of the relationship between emission line and continuum strengths abound in the literature, but only a handful of small samples have been studied relating the shape and strength of the {\\em high energy} QSO continuum to emission lines.   Zheng, Kriss, \\& Davidsen (1995; hereafter ZKD) find a strong anti-correlation between the \\ew\\, of OVI$\\lambda1034$ and \\aox. Although some models predict such behavior for other UV emission lines, no other such trends have been observed.  The increase in \\aox\\, with luminosity in QSOs when combined with the observed BEff is not nearly sufficient to explain the trend in \\wovi\\, with \\aox.  What might be responsible?  Higher luminosity QSOs may undergo spectral evolution such that fewer photons from a soft X-ray excess/BBB component are available for ionization.  The intriguing results of ZKD are based on a variety of published X-ray fluxes, and a heterogeneous compilation of rest-frame UV spectra (2 from HUT, 16 from \\IUE, 14 from $HST$, and 29 ground-based), excluding all non-detections.  Thus, although they may well prove robust, such results are open to challenge on the basis of the strong, diverse selection effects inherent in such a sample.  On the other hand, even in complete, flux-limited samples of QSOs (which often constitute a large fraction of other more heterogeneous samples) there is a strong correlation between redshift and luminosity.  At a {\\em given} redshift, the more luminous objects will usually have higher signal to noise (S/N) spectra.  As a result, most of the weak-lined QSOs remaining in a sample that ignores non-detections will be luminous (a Malmquist bias).  In addition, noise that randomly enhances the apparent line strength will bump low luminosity objects into the sample with spuriously high line \\ew\\, (an Eddington bias). Thus the apparent statistical significance of line/continuum correlations may spuriously enhanced by a {\\em combination} of selection effects if undetected lines are left out of the sample. Although some general selection effects in line/continuum studies of the BEff have been considered in the literature (e.g., Zamorani et al. 1992), few studies can be found incorporating line error estimates and upper limits to line \\ew, both essential to unbiased line/continuum studies.  Here we initiate a line/continuum investigation of wide scope, using (1) large, homogeneous samples (2) uniform data and analysis, and (3) a wider range of lines and (consequently) ionization potentials.  We outline new error analysis for a simple automated line measurement technique (\\S~\\ref{meas}), and include limits in all analyses.  To facilitate further study, we tabulate these data for individual QSOs. Via correlation tests (see \\S~\\ref{results} for details), we seek to determine which of \\luv, \\auv, \\lx, or \\aox\\, dominates emission line formation, or at least which parameter most reliably predicts measured line parameters. In combination with other (e.g., higher redshift/luminosity) samples, these data, techniques, and results should prove useful for further studies of the effect of QSO SEDs on the broad emission line region (BELR).  One such followup study is now underway, using LBQS and RASS data (Green et al. 1996). ",
        "conclusions": "\\label{conclude}  Complex activity is likely to be associated with the nuclear environment in QSOs.  Rapid star formation and evolution, supernovae, accretion/merging of galaxies or protogalactic fragments, and a supermassive accreting  black hole may all contribute.  Although QSO spectra are surprisingly homogeneous given such a flamboyant cast, the simplest geometric and photoionization models do not succeed in explaining the relationship of QSO emission lines to the observed continuum.  This may partly be alleviated if the line emission can be directly compared to its principal ionizing/heating continuum.  Objective, automated line measurements including line upper limits are crucial to avoid spurious enhancement of apparent line/continuum correlations.  We find significant correlations between \\ew\\, and UV luminosity (e.g., the well-studied Baldwin effect) for \\lya, \\civ, \\heii, and \\ciii.  \\wciii\\, and \\wheii\\, also show previously unreported correlations with X-ray luminosity which, for \\ciii, appears to be primary. The line ratios \\rciii\\, and \\rheii\\, both show strongest dependence on \\lx.  \\wlya\\, correlates strongly with spectral slopes \\auv\\, and \\aox\\, (between 2500\\AA\\, and 2~keV), but {\\em not} with X-ray luminosity.  Using these results, we argue that one simple geometrical interpretation of the BEff that assumes ionizing X-ray emission to be more isotropic than UV continuum emission is not plausible.  If indeed the BEff were a result of a distribution of disk inclinations in this case, weak anti-correlations of line \\ew\\, with X-ray luminosity would be expected at best.  The significant anti-correlations of \\ciii\\, and \\heii\\, emission with \\lx\\, thus render the simplest geometrical model unlikely.  When we are able to compare the line \\ew\\, most directly to a portion of its PIHC using extant soft X-ray observations (for \\lya\\, and \\civ) the {\\em X-ray} BEff is not significant.  For \\heii\\, and \\ciii\\, lines, where the entire PIHC is softer than the \\ein\\, bandpass (KK88), a significant BEff persists relative to soft X-ray luminosities.  We thus argue that the BEff weakens or disappears when the line emission is compared to the luminosity in the bandpass of its principle photoionizing continuum.   This supports an interpretation of the BEff as a change in spectral energy distribution with luminosity.  We predict that no BEff relative to soft X-ray luminosity should be found for Fe\\,II or Mg\\,II emission lines. Extensions of our method to samples of a wider redshift/luminosity range would test these predictions.  Now that we have outlined a technique for the efficient measurement of large numbers of comparatively low S/N QSO spectra, we will apply it to the largest, most uniformly-selected such sample to date, the LBQS. Optical spectra and X-ray fluxes or upper limits are available for 908 QSOs in the LBQS from the {\\em ROSAT} All-Sky Survey. Analysis of that database will be combined with the results presented here to offer a truly wide luminosity baseline for further study of the interdependence of QSO continuum and emission line properties.  Of course, we would have preferred for each line an accurate measure of its rest-frame principle ionizing/heating continuum.  For \\ovi, this means the the `He\\,I' continuum ($24.5-54.4$eV). For \\lya\\, \\ciii\\, and \\civ\\, this means the ranges $13.6-24.5$eV (the `Lyman continuum') dominates (KK88). \\ciii\\, and \\civ\\, should depend also on ionizing photons from $0.3-0.4$keV.  Since these include the EUV range, which is observationally inaccessible for all but a handful of nearby AGN (Marshall et al. 1995), we have attempted an indirect examination of the relevant continuum via the adjacent UV and X-ray luminosities, and through \\aox.  The true strength of the EUV is best estimated using spectral index and normalization in the adjacent UV and soft X-ray bands together. Slopes are available in the UV data set, but the great majority of X-ray data provide only net counts in the \\ein\\, bandpass ($\\sim 0.16-3.5$keV).  Unfortunately, there are too few QSOs in the \\IUE~ sample with published \\ax\\, (about 18, judging from Elvis et al. 1994) to permit any convincing statistical tests.  The 2keV monochromatic fluxes used in \\aox\\, are thus derived from these \\ein\\, counts assuming a single power-law slope, and absorption due to Galactic $N_H$ only.  Although there may be absorption (either warm or cold; see e.g., Netzer 1993) intrinsic to the QSOs, this effect is unlikely to be strong.  A small soft excess above the power-law, which could be important to the ionizing continuum is, however, expected in many of the QSOs (Fiore et al.  1994).  Estimates of X-ray spectral slopes in the \\ros\\, band ($0.1-2.4$keV) for more than 100 bright QSOs should be available from ROSAT observations within the next few years (e.g., Bade et al.  1995).  These data should prove a valuable addition to the studies initiated here.  Hearty thanks to Ken Lanzetta for providing his \\IUE~ spectral atlas, including error spectra and continuum fits, in digital form.  Paul Eskridge helped simplify the task of PSR analysis. Craig Foltz provided a copy of the LBQS composite spectrum. I gratefully acknowledge Avi Loeb for our discussions of the emission line error analysis.  \\clearpage  \\appendix"
    },
    "9603/astro-ph9603159_arXiv.txt": {
        "abstract": "We examine the speed of inward traveling cooling fronts in accretion disks.  We show that their speed is determined by the rarefaction wave that precedes them and is approximately $\\alpha_F c_{F} (H/r)^q$, where $\\alpha_F$ is the dimensionless viscosity, $c_{F}$ is the sound speed, $r$ is the radial coordinate, $H$ is the disk thickness, and all quantities are evaluated at the cooling front.  The scaling exponent $q$ lies in the interval $[0,1]$, depending on the slope of the $(T,\\Sigma)$ relation in the hot state. For a Kramer's law opacity and $\\alpha\\propto (H/r)^n$, where $n$ is of order unity, we find that $q\\sim 1/2$.  This supports the numerical work of Cannizzo, Chen and Livio (1995) and their conclusion that $n\\approx3/2$ is necessary to reproduce the exponential decay of luminosity in black hole X-ray binary systems.  Our results are insensitive to the structure of the disk outside of the radius where rapid cooling sets in.  In particular, the width of the rapid cooling zone is a consequence of the cooling front speed rather than its cause.  We conclude that the exponential luminosity decay of cooling disks is probably compatible with the wave-driven dynamo model.  It is not compatible with models with separate, constant values of $\\alpha$ for the hot and cold states. ",
        "introduction": "The rate of mass transfer in accretion disks depends on the rate at which angular momentum can be transferred outward.  This is normally expressed in terms of a dimensionless viscosity $\\alpha$, which is defined as \\begin{equation} \\alpha\\equiv {\\nu\\over c_s H}, \\end{equation} where $H$ is the disk half-thickness, $c_s$ is the local sound speed, and $\\nu$ is the local effective viscosity (\\cite{SS73}). Initially $\\alpha$ was assumed to be constant. There are now strong grounds, both empirical and theoretical, for concluding that $\\alpha$ must be a variable. The task is to combine empirical evidence with theoretical guidance to construct a self-consistent theory of angular momentum transport in accretion disks that accounts for the wealth of observations.  A successful theory is likely to reveal that angular momentum transport is non-local, so that the concept of a local viscosity is itself of limited value.  Communication between observational studies of accretion disks and theories of angular momentum transport is facilitated by models of time-dependent accretion disks.  The time dependence is critical since the emissivity of steady-state disks is independent of the viscosity. An excellent review of the history of this area is given by Cannizzo (1993a; see also Cannizzo 1993b).  The first major constraint on $\\alpha$ came from comparison of limit cycle disk instability models with observations of dwarf novae.  The models provide a very credible basic interpretation of the dwarf nova phenomenon, but only if $\\alpha$ is not a constant (Smak 1984). This constraint does not determine the functional form of $\\alpha$.  Models in which $\\alpha$ has one, radially constant value, $\\alpha_{hot}$ in outburst and another lower, but also radially constant value $\\alpha_{cold}$ in quiescence work about as well as a model in which $\\alpha=\\alpha_0(H/r)^{n}$ which would apply if $\\alpha$ were a function of the sound speed and hence the temperature of the disk.  Another perspective on the behavior of $\\alpha$ can be obtained by comparing the dwarf novae with soft X-ray transients.  In the latter case, the only quantitative work has been done on those that are black hole candidates, but those are especially interesting laboratories because of the suspicion that the compact star lacks a hard surface and an associated magnetosphere and boundary layer.  Several of the black hole candidates have outbursts with rapid rise and subsequent slower decline that are in reasonable agreement with the same limit cycle disk instability models that account for dwarf novae (\\cite{MW89}).  The quantitative and even qualitative behavior of the black hole models depends on the prescription for $\\alpha$. In the case of a double valued, but radially constant prescription, the outburst will tend to occur in the inner disk, giving rise to somewhat slower rise, more symmetric outbursts.  A prescription in which $\\alpha=\\alpha_{0}(H/r)^{n}$ will give very small values of $\\alpha$ in quiescence where H/r is found to decrease inward. This will yield a very long viscous time in the inner disk and promote outbursts that begin in the outer disk and propagate inward. This yields model outbursts with rapid rise and slower decline, in accord with the observations for the optical and soft X-ray light curves of the X-ray novae.  One of the interesting features of the black hole X-ray novae is the tendency to show an exponential decline. Simple models in which one quickly reduces the transfer rate to a hot disk with with constant $\\alpha$ generate geometrically declining, not exponential, light curves.  Even models in which the decline is driven by the cooling wave of the disk instability tend to have geometrically declining light curves with constant $\\alpha$.  Mineshige et al. (1993) have argued that to produce an exponential decline, the angular momentum of the inner disk must be removed at a rate proportional to the angular momentum.  They note that this tends to be the behavior of disk instability models with $\\alpha=\\alpha_{0}(H/r)^{n}$ with n $\\sim$ 1 - 2.  Cannizzo (1994) has also addressed this argument by noting that both dwarf novae and the black hole transients have exponential declines.  Cannizzo concluded that to reproduce the exponential one needs $\\alpha\\propto r^{\\epsilon}$, with $\\epsilon\\sim~0.3 - 0.4$, which is consistent with Mineshige et al.  Cannizzo carried the argument one step further, however, by making the case that the precise value of $\\epsilon$ that leads to exponential decline is itself a function of other parameters of the problem such as the transfer rate and inner disk radius.  From this he concluded that exponential decline requires some form of feedback to operate in the disk to give just this behavior. This may hint that the angular momentum transport process is non-local, as the theories where internal waves play a critical role imply.  These arguments have been extended significantly by Cannizzo, Chen, and Livio (1995).  Cannizzo et al. used well-resolved numerical studies to show that the width of the cooling front can be approximated very closely by w = $\\sqrt{Hr}$ and that for such a behavior, exponential decay of the light curve during the cooling wave phase is obtained only for a prescription of the form $\\alpha=\\alpha_0(H/r)^n$ with n very close to 3/2.  The critical point in their argument is that angular momentum is removed from the hot part of the disk by the advance of the cooling front and if the cooling front velocity is proportional to $r$, as it is for $n\\sim 1.5$, then this loss of angular momentum is proportional to the total angular momentum in the disk. This preferred value of $n$ is consistent with the theory of angular momentum transport by an internal wave-generated dynamo driven by tidal instabilities at the outer edge of the disk (\\cite{VJD90}, \\cite{VD92}).  Nevertheless, the physical underpinnings of this behavior of the cooling wave were not clear.  In particular, it is not clear why the cooling front should have this width.  Given this width, it is possible to argue that the cooling front should have the velocity characteristic of torque induced mass flows with a radial length scale of $w$, i.e. \\begin{equation} V_r\\sim {\\alpha c_s^2\\over w\\Omega}\\sim \\alpha c_s \\left({H\\over r}\\right)^{1/2}. \\label{eq:cwidth} \\end{equation} Since $c_s$ near the cooling front is approximately constant, and since $H\\sim c_s/\\Omega\\propto r^{3/2}$, this gives a cooling front velocity which is proportional to $r$ when $n=3/2$. In what follows we will argue that although this expression for the cooling front velocity is approximately correct, the direction of causality has been reversed. The cooling front width is a consequence of the cooling front velocity.  In this paper we present an analysis of the behavior of the cooling wave and show that its propagation depends on the viscous flow in the hot state and is nearly independent of the actual cooling process and of the state of the disk in the cool, quiescent material that accumulates in the wake of the inward-propagating cooling wave.  We argue, in agreement with Cannizzo, Chen, and Livio, that the exponential decay gives strong evidence for the presence of the cooling wave, and hence of the disk instability phenomenon in general, and a powerful constraint on the physical nature of the local viscosity. The mechanisms that control the propagation of the cooling front are discussed in \\S 2.  Section 3 presents constraints on the opacity and other functional forms of $\\alpha$. The relation of these results to the internal wave-driven dynamo are presented in \\S 4.  Summary and conclusions are given in \\S 5. ",
        "conclusions": "We have constructed a simple model for the propagation of cooling fronts in accretion disks which reproduces the numerical results of Cannizzo, Chen, \\& Livio (1995).  In this model the cooling front speed is determined by the rarefaction wave that lowers the disk temperature to the point where rapid cooling can set in.  We find that the speed of the cooling front scales as \\begin{equation} v_{cF}\\sim \\alpha_F c_F \\left({c_F\\over r_F\\Omega(r_F)}\\right)^q, \\end{equation} where the subscript $F$ refers to the radius where the disk falls out of thermal equilibrium and begins rapid cooling.  The coefficient $q$ is given in equation (\\ref{eq:expq}) and depends on both the opacity law and the functional form of $\\alpha$.  However, $q$ will be close to $1/2$ for most models of the disk hot state.  Somewhat surprisingly, a local reduction of $\\alpha$ near the cooling front has only a modest effect on this result.  One striking aspect of our derivation is that we do not appeal to any aspect of the structure of the disk at radii greater than $r_F$, which marks the onset of rapid cooling.  This stands in contrast to the suggestion by CCL that the cooling front velocity is determined by its width, measured from the onset of rapid cooling to its finish. We have not discussed the structure of this region here, but we note that given a cooling front velocity determined by the structure of the disk inward from the cooling front, the width of the cooling front itself can be estimated by inverting equation (\\ref{eq:cwidth}). In other words, the scaling of the cooling front width is a consequence of the speed of the cooling front, not its cause.  Our success in modeling the propagation of cooling fronts as rarefaction waves suggests a similar effort could be made to model heating fronts as compressional waves.  We have not yet done this, but expect to examine this problem in a future paper.  Aside from the internal wave driven dynamo model, we have not discussed models for $\\alpha$ which are consistent with the results of this paper.  While this is largely from a lack of suitable candidates, there is one other prediction of $\\alpha$ with the required form (\\cite{MM83}).  However, this estimate is based on using large scale buoyant cells driven by magnetic buoyancy via the Parker instability (cf. \\cite{P75}).  Zweibel \\& Kulsrud (1975) showed that sufficiently strong turbulence in a shearing environment would suppress the Parker instability. Vishniac \\& Diamond (1992) pointed out that the Balbus-Hawley instability (\\cite{V59}, \\cite{C61}, \\cite{BH91}) always leads to a level of turbulence which is sufficiently strong by this criterion.  In fact, the linearly unstable modes of an azimuthal magnetic field suffer turbulent mixing at a rate roughly equal to the local orbital frequency. We conclude that the magnetic buoyancy driven model is not consistent, in its original form, with the dynamics of magnetic fields in accretion disks.  There are several conclusions to be drawn from this work.  {}First, we have provided a simple analytic derivation which supports the conclusion of CCL that the exponential decay of the luminosity of black hole disk systems following outbursts is consistent with a local law for the dimensionless disk viscosity $\\alpha\\propto (H/r)^n$ if, and only if, $n$ is approximately $3/2$.  Second, given this scaling for $\\alpha$ we find that disk systems in general should exhibit approximately exponential luminosity decay from peak luminosity whenever the hot state opacity follows a simple power law.  Exceptions will involve hot state $(\\Sigma, T)$ relations which are either unusually close to thermal instability, in which case the cooling front velocity can approach $\\alpha_F c_F$, or in which $T$ is extremely insensitive to $\\Sigma$, in which case the cooling front velocity will approach the accretion velocity in the inner disk.  Third, this result implies that any $\\alpha$ scaling for which $\\alpha$ is constant in the hot state is in conflict with current observations.  This includes models in which $\\alpha$ is given by $[\\alpha_{hot},\\alpha_{cold}]$, where $\\alpha_{hot}$ is a constant.  {}Fourth, since the cooling front speed depends only on the hot state, other models for $\\alpha$ can also give exponential decays, although they may fail on other grounds.  For example, if $\\alpha\\propto r^{2/3}$, then we can obtain a roughly exponential luminosity decay in spite of the fact that this law is insensitive to the local temperature.  {}Fifth, this result is apparently compatible with the internal wave driven dynamo model for disk viscosity.  This does not follow trivially from the prediction that $\\alpha\\propto (H/r)^n$, where $n$ is approximately $3/2$ in a stationary disk. The waves reach the cooling front after traveling through the cold part of the disk.  Consequently, the $\\alpha_F$ induced by the internal wave driven dynamo is greatly reduced.  Here we have relied on an independent mechanism, the incoherent dynamo, to give a minimal value for $\\alpha_F$. The final scaling law obtained in this way lies within observational limits.  Clearly further work on these dynamo mechanisms, and on the nature of the cold disk state, would be helpful for providing a definitive answer to this question.  Still,  this is the only internally consistent model for $\\alpha$ which is constructed from first principles and which satisfies the cooling wave constraint.  We note that CCL have shown that the value of $\\alpha_0$ can be estimated from the luminosity decay rate.  This value has not been calculated for the internal wave driven dynamo model, but when it is the existence of an observationally motivated estimate will provide another critical test of the model.  {}Finally we note that this whole analysis is predicated on the assumption that a cooling wave exists in the decline of the light curve of transient black hole candidates and related systems.  While the evidence is indirect, one can thus regard the exponential decline as a strong argument that a cooling wave is the fundamental mechanism of the decline of these transients. Furthermore, it adds to the evidence that the accretion disk ionization instability is the underlying physical cause of the transient outburst phenomenon."
    },
    "9603/astro-ph9603123_arXiv.txt": {
        "abstract": "Starting in 1991 we have performed a regular monitoring of EGRET sources with the Effelsberg 100-m telescope and the IRAM 30-m at Pico Veleta. In comparison with data on a sample of flat-spectrum quasars which have not been seen by EGRET we search for correlations of any kind in the behaviour of FSRQ's in the radio and gamma-ray regime. While there is no radio-to-gamma-ray luminosity correlation for these sources, it seems that gamma-ray high states coincide with increased activity in the radio regime with a strong tendency that gamma-ray outbursts precede radio outbursts. The gamma-ray spectra seem to harden with increasing flux level. ",
        "introduction": "We will first discuss PKS 0528+134, the most luminous $\\gamma$-ray source known besides $\\gamma$-ray bursts. 0528+134 exhibits superluminal motion with $\\beta_{app}\\simeq 4.4$ and indications for even higher values (Pohl et al. 1995). Superluminal motion was expected since given the variability time scales of a day, the redshift of z=2.07 and the $\\gamma$-ray flux in the OSSE range, strong Doppler boosting is required to satisfy the compactness limit and the Elliot-Shapiro relation (McNaron-Brown et al. 1995). It has been noted that the expulsion of new VLBI components may coincide with $\\gamma$-ray outbursts (Pohl et al. 1996). \\begin{figure} \\noindent \\plotone{pohlm1.eps} \\caption{Here we show the radio light curve of 0528+134 based on Effelsberg data and observations at 86\\,GHz taken with the IRAM 30-m. The histograms give the two-weeks-averages of the NRL-GBI data at 2.25 GHz (solid line) and 8.3 GHz (dotted line). Observational uncertainties are below 5\\% except for the data at 32 GHz and 86 GHz which have uncertainty levels around 10\\%. The unusual depression in July 1993 is most likely an extreme scattering event and not intrinsic to the source. For comparison the state in the EGRET range is indicated by empty circles for low state ($S< 5\\cdot 10^{-7}\\ {\\rm ph.\\,cm^{-2}\\,sec^{-1}}$ above 100 MeV), medium level ($S< 10^{-6}$), and high state ($S> 10^{-6}$). A question mark indicates protected data. It is tempting to relate the $\\gamma$-ray outbursts in 1991 and 1993 to the mm outbursts a few month later. However, no strong $\\gamma$-ray outburst has been reported yet for 1995, a few months before the brightest ever-recorded mm outburst. Or is there a time lag of two-and-a-half years, relating the 1991 $\\gamma$-ray flare to the 1993 radio outburst, respectively the 1993 $\\gamma$-ray flare to the 1995 radio outburst? Given the superluminal motion we would expect a corresponding VLBI knot to have a core separation of 0.3 mas after that time. In 1992.85 more than 80\\% of the 22 GHz flux was confined within 0.1 mas to the core (Pohl et al. 1995).} \\end{figure}  \\noindent In Fig.1 we show the radio light curve of 0528+134 in comparison to its $\\gamma$-ray state for the period 1991 to 1995. The source was quiet in the radio regime between 1985 and 1991 (Zhang et al. 1994). The following conclusion can be drawn: 0528+134 has been very bright in $\\gamma$-rays either when it was weak in radio or a few months before a mm outburst. It was at medium $\\gamma$-ray level at the time of the brightest ever-recorded radio outburst, at the end of 1995. Though the $\\gamma$-ray high states seem to precede the radio outbursts, which is also supported by backextrapolation of the position of VLBI knots, there is no simple one-to-one relation. The analysis of individual sources is mainly hampered by the limited coverage of the $\\gamma$-ray light curve and the problem that many sources -- like 0528+134 -- exhibit $\\gamma$-ray variability on time scales of days, less than the standard integration time in EGRET observations. ",
        "conclusions": ""
    },
    "9603/astro-ph9603137_arXiv.txt": {
        "abstract": "We have examined consequences of strong tidal encounters between a neutron star and a normal star using SPH as a possible formation mechanism of isolated recycled pulsars in globular clusters. We have made a number of SPH simulations for close encounters between a main-sequence star of mass ranging from 0.2 to 0.7 $\\msun$ represented by an n=3/2 polytrope and a neutron star represented by a point mass. The outcomes of the first encounters are found to be dependent only on the dimensionless parameter $\\eta^{\\prime} \\equiv (m/(m+M))^{1/2} (r_{\\min}/R_{MS})^{3/2}(m/M)^{(1/6)}$, where m and M are the mass of the main-sequence star and the neutron star, respectively, $r_{\\min}$ the minimum separation between two stars, and $R_{MS}$ the size of the main-sequence star. The material from the (at least partially) disrupted star forms a disk around the neutron star.  If all material in the disk is to be acctreted onto the neutron star's surface, the mass of the disk is enough to spin up the neutron star to spin period of 1 ms. ",
        "introduction": "The cores of globular clusters have high stellar densities. Recent studies showed that the physical interactions between stars are main driving force in determining the dynamical evolution after the core collapse (e.g., Goodman 1988). The physical interactions include direct collisions between normal stars, tidal captures between a normal star and a compact star (neutron star or white dwarf), and formation of binaries via three-body processes.  Such interactions can lead to the formation of objects that are not common in low stellar environments. Indeed globular clusters reveal high abundance of X-ray binaries and short period pulsars. These systems are thought to be related to the physical interactions between stars. X-ray binaries can be formed either by tidal capture between a neutron star and a main-sequence star, or by the evolution of primordial binary.  Another observational evidence for the high abundance of the physical interactions in clusters is the enhancement of short-period pulsars. A large number of short period pulsars have been detected in globular clusters. Unlike many pulsars in the solar neighborhood, cluster pulsars are known to be predominantly short period ones. Neutron stars are believed to be born with short spin periods ($\\lsim$ 0.1 sec) and strong magnetic fields ($10^{12} - 10^{13}$ G). As a neutron star evolves, its rotational energy is released through magnetic dipole radiation which causes deceleration of the spin. Since the spin-down time scales for pulsars with strong magnetic fields are of order $10^8$ years while the production of neutron stars in globular cluster must have ceased long time ago, present day population of short period pulsars in clusters must be `recycled' pulsars, which has been spun up by accretion of mass.  Aside from the fact that the abundance of recycled pulsars in globular clusters relative to ordinary stars is far greater than that for solar neighborhood, the cluster pulsars are preferably isolated (as opposed to in a binary system). According to compilation of 558 pulsars by Taylor, Manchester, \\& Lyne (1993), 19 out of 22 in clusters are {\\it isolated} pulsars while the ratio is only 2 out of 11 in field pulsars. Therefore any theory for the recycled pulsars in globular clusters should explain both the high abundance and high ratio of isolated to binary pulsars.  Proposed mechanisms for producing isolated recycled pulsars in dense stellar regions include close encounters between a neutron star and a main-sequence star leading to a complete or partial disruption of the main-sequence star (Krolik 1984, Lee 1992), and the three-body interactions which can liberate a neutron star in the binary through various ways (see Rappaport,Putney, \\& Verbunt 1989 for summary).  In the present paper, as a start of the study on the first mechanism, we examine the consequences of close encounters between a neutron start and a main-sequence star.  During a close encounter involving a neutron star, the orbital energy of an incident normal star can be dissipated and stored in its interior by tidal force exerted by a neutron star.  Depending on the amount of energy deposited to the stellar envelope, one of the following outcomes are possible: tidal capture of a normal star, or total or partial disruption of the normal star. Based on SPH simulations, Davies, Benz, \\& Hills (1992) find that encounters involving a main-sequence star of $0.8 \\msun$ can leave a single object for $r_{\\min} \\lsim 1.75 \\, R_{MS}$ and detached binaries for $1.7 \\, R_{MS} \\lsim r_{\\min} \\lsim 3.5 \\, R_{MS}$, while encounters involving a red giant star leave a compact neutron star-white dwarf binary for $r_{\\min} \\lsim 1.8 \\, R_{RG}$ and a detatched binary for $1.8 \\, R_{RG} \\lsim r_{\\min} \\lsim 2.5 \\, R_{RG}$, where, $r_{\\min}$ is the separation at periastron passage. In the present study, while Davies \\etal concentrated on only one main-sequence star mass, we perform SPH simulations of encounters between a $1.4 \\msun$ neutron star and a main sequence star with a mass in the range between $0.2$ and 0.7 $\\msun$ for various $r_{\\min}$. This extension will bring a full comparison between simulation and theory, and will eventually provide more realistic information to the study of dynamical evolution of globular clusters.  This paper is organized as follows.  In \\S 2, we describe our simulations of stellar encounters, and we discuss the consequences of the encounters based on the simulations in \\S 3.  Final section summarizes our major findings. ",
        "conclusions": "We have investigated the consequences of close encounters between a neutron star and a main-sequence star using the SPH method.  We have found that the mass fraction stripped from the main-sequence star during the encounter, $\\Delta m$, is determined only by a single parameter $\\eta^{\\prime}$ in the ranges of $m$ and $v_\\infty$ used in our simulations.  The orbital energy deposited in the stellar envelope per unit mass, $\\Delta E/m$ seems responsible for the mass stripping because it also can be well approximated to be a function of $\\eta^{\\prime}$.  From simulations with relatively large $\\eta^{\\prime}$, the core contraction of the main-sequence star remnant has been witnessed as well as the envelope expansion.  The radiation from the contracted core will help the envelope expansion and the envelop may not be able to restore itself to the original place.  The mass of the disk formed around a neutron star after an encounter is as much as of order 0.1 $\\msun$.  If all of this mass could be accreted onto the neutron star's surface, the accretion could accelerate the spin of the neutorn star down to 1 ms.  However, this kind of heavy accretion disks are subject to high radiation pressure and considerable fraction of the disk is believed to be blown away, which is against the spin-up of the neutron star. The fate of these heavy accretion disks are uncertain and remains to be studied intensively, to explain the high abundance of isolated, recycled pulsars in globular clusters."
    },
    "9603/astro-ph9603071_arXiv.txt": {
        "abstract": "Two dimensional concordance plots involving the baryon-to-photon ratio, $\\eta$, and an effective number of light neutrinos, $N_{\\nu}$, are used to discuss the overall consistency of standard big-bang nucleosynthesis in light of recent determinations of the primordial deuterium ($^2$H) abundance. Observations of high-redshift Ly-$\\alpha$ clouds have provided discordant $^2$H/H determinations: one cloud with $^2$H/H high compared with the previously accepted upper limit on ($^2$H + $^3$He)/H; and one system with a significantly lower upper bound on $^2$H/H than those previously obtained. The high value of $^2$H/H agrees well with the current observationally-inferred primordial abundances of $^4$He and $^7$Li for $N_{\\nu}=3$. The low value of $^2$H/H does not fit well with the current observationally-inferred primordial abundance of $^4$He for $N_{\\nu}=3$. In addition, if the low value of $^2$H/H is indicative of the primordial deuterium abundance, then significant depletion of $^7$Li in old, hot Pop II halo stars is probably required to obtain a concordant range of $\\eta$, for {\\em any} effective number of neutrino flavors. The use of conservative ranges for the primordial abundances of $^2$H, $^4$He, and $^7$Li allow success of the standard picture for $N_{\\nu}=3$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603094_arXiv.txt": {
        "abstract": "We present a method for solving the two-dimensional linearized collisionless Boltzmann equation using Fourier expansion along the orbits. It resembles very much solutions present in the literature, but it differs by the fact that everything is performed in coordinate space instead of using action-angle variables. We show that this approach, though less elegant, is both feasible and straightforward.  This approach is then incorporated in a matrix method in order to calculate self-consistent modes, using a set of potential-density pairs which is obtained numerically. We investigated the stability of some unperturbed disks having an almost flat rotation curve, an exponential disk and a non-zero velocity dispersion. The influence of the velocity dispersion, halo mass and anisotropy on the stability is further discussed. ",
        "introduction": "Before presenting the method of the mode calculation in more detail, we first introduce the unperturbed galaxy models for which the calculations were performed.  All our models have the same unperturbed potential and mass density, but feature different distribution functions, with varying velocity dispersion, streaming velocity and anisotropy.  \\subsection{The unperturbed potential} The potential of the unperturbed galaxy was constructed as a sum of two Kuzmin-Toomre disk potentials with different core radii. In the plane of the disk, it is given by \\begin{equation} V_0(r)={1 \\over \\sqrt{1+r^2}} + {1 \\over \\sqrt{1+(r/4.4)^2}}. \\label{pot0} \\end{equation} Note that potentials are defined as binding energies, with a positive sign. This potential produces a rotation curve which is much flatter than a single component Kuzmin-Toomre potential (see fig. \\ref{figrot}). The ratio of the flat part to the rising part is about $6/1$.  \\begin{figure} \\vspace{5.5cm} \\special{ hscale=80 vscale=80 hoffset=-24 voffset= -15 hsize=480 vsize=500 angle=0 psfile=4583F1.ps } \\caption{Rotation curve of the unperturbed \\label{figrot} galaxy potential} \\end{figure}  Although the rotation curve tends to rise somewhat too slowly near the centre, it has, at least qualitatively, a realistic behaviour.  \\subsection{The mass density profile} It is well-known that strongly flattened galaxies usually have a substantial amount of dark halo mass, extending much further than the visible component. It is sufficient for our purpose to model it by a spherical and pressure-supported galaxy component.  Therefore it is reasonable to assume that this halo component only influences the stability behaviour of the disk by its contribution to the global potential. The same roughly holds for the central bulge, which is hot and has a three-dimensional structure as well. Thus both the halo and the bulge are ``inert'' and the corresponding potential and mass density are taken to be spherical and denoted by $V_{0,H}(r)$ and $\\rho_{0,H}(r)$.  The disk itself is the only component which is supposed to consist of ``active mass'', sensitive to instabilities. The unperturbed disk is supposed to be two-dimensional and axisymmetric, with a surface density $\\rho_{0,D}$ and a potential $V_{0,D}$.  We chose an exponential mass profile with a central core: \\begin{equation} \\rho_{0,D}=\\alpha e^{-1.3 \\sqrt{0.2+r^2}} \\label{dens0}. \\end{equation} Although this mass density extends up to infinity, the actual models have an outer limit at $r=6$, and the mass density reaches zero at that point. This we achieve by fitting the distributions (\\ref{dens0}) with several finite components. The relative error made at the outer edge is only of the order of 0.1\\%. The parameter $\\alpha$ determines the total disk mass.  Since the total system should be self-consistent, the total potential in the plane of the galaxy is given by the sum of the disk and halo component: \\begin{equation} V_0(r)=V_{0,D}(r)+V_{0,H}(r). \\end{equation} Since the total potential $V_0$ has a fixed form (\\ref{pot0}) and the potential of the disk is determined by the surface mass density (\\ref{dens0}), we can calculate $V_{0,H}=V_0-V_{0,D}$.  Thus follows the halo mass density \\begin{equation} \\rho_{0,H}= -{1 \\over 4 \\pi G} {1 \\over r^2} {d \\over dr} \\left( r^2 {d V_{0,H} \\over dr} \\right), \\end{equation} and the total halo mass within a radius $r$ \\begin{equation} M_{0,H}(r)= -{1 \\over G} r^2 {d V_{0,H} \\over dr}. \\end{equation}  \\begin{figure} \\vspace{5.5cm} \\special{ hscale=80 vscale=80 hoffset=-62 voffset=-30 hsize=480 vsize=500 angle=0 psfile=4583F2.ps } \\caption{Total mass inside $r$ of the disk (full curves) and the halo (dashed curves) for $H/D$=$2.5$, $5$ and $+\\infty$.} \\end{figure}  Of course, the halo mass density should everywhere be positive. This puts an upper limit on the value of $\\alpha$ in (\\ref{dens0}). The relative contribution of the disk and halo components are quantified using the halo-to-disk factor, $H/D$, which gives the proportion of the total mass inside the radius $r_{\\rm max}$ (=6) for the halo and the active disk. Self-consistency with a non-negative spherical halo puts a lower limit on $H/D$ of about $2.5$. Note that the assumption of a spherical halo is not a crucial one. If one would assume an oblate halo, the only effect would be a somewhat lower minimum $H/D$.  \\subsection{The distribution function} We will examine the stability behaviour for different stellar distributions. According to Jeans' theorem, the unperturbed part of the distribution is a function of two integrals of motion, the binding energy $E$ and the angular momentum $J$, defined by \\begin{equation} E=V_0(r)-{1 \\over 2} ( v_r^2+v_\\theta^2) {\\rm \\ \\ \\  and \\ \\ \\ } J=r v_\\theta. \\end{equation} In order to generate a variety of finite disks, based on the potential (\\ref{pot0}), the distribution function is written as a linear combination of basic distributions: \\begin{equation} \\df_0(E,J)=\\sum_{t=1}^{n_t} c_t \\df_{0,t}(E,J). \\end{equation} All components only have a (everywhere positive) contribution for orbits lying completely inside $r_{\\rm max}$, i.e. for the region where (see also fig \\ref{lind1}) \\begin{equation} E \\ge V_0(r_{\\rm max}) - {J^2 \\over 2 r_{\\rm max}^2}, \\end{equation} \\begin{equation} E \\ge V_0(r_{\\rm max})-{1 \\over 2} v_{circ}^2(r_{\\rm max}),\\label{excl2} \\end{equation} with $r_{\\rm max}$ the radius of the edge of the disk. Equation (\\ref{excl2}) is required in order to exclude the region $r_- \\ge r_{\\rm max}$. In addition, the distribution goes to zero at this edge in a smooth way, so that the first derivative remains finite everywhere.  The explicit form of the components is listed in the Appendix.  The expansion coefficients $c_t$ are determined by a least square fit of the corresponding mass density to the proposed exponential form (\\ref{dens0}) (the coefficients are forced to be positive, in order to avoid negative distributions). By choosing an appropriate set of components $\\df_{0,t}$, we were able to create the desired orbital densities. In all the models, the error on the fit to the mass density never exceeds 1\\% of the central value.  We constructed 4 models, labeled I to IV (the explicit form of the distribution functions is listed in the appendix). Along the sequence, the models become more and more rotation-supported, having an increasing streaming velocity and decreasing temperature. Fig. \\ref{figvelocI} shows the streaming velocities and dispersions for the coldest and hottest case. The dispersions all go to zero at the edge of the disk. Note that model I is perfectly isotropic and has a linearly increasing mean velocity curve. For all disks, Toomre's local axisymmetric instability criterion (\\cite{Toomre64}) \\begin{equation} Q= { \\sigma_r \\kappa \\over 3.36 G \\rho_{0,D} } \\label{defQ} \\end{equation} (with $\\sigma_r$ the radial velocity dispersion and $\\kappa$ the epicyclic frequency) is everywhere higher than $1$.  \\begin{figure} \\vspace{8.0cm} \\special{ hscale=80 vscale=64 hoffset=-22 voffset= -17 hsize=480 vsize=500 angle=0 psfile=4583F3.ps } \\caption{Streaming \\label{figvelocI} velocity (full line), azimuthal velocity dispersion (long dashed line) and radial velocity dispersion (short dashed line) for Model IV (top panel) and Model I (bottom panel).} \\end{figure}  In fig. \\ref{figdf0}, the distribution function of model III is shown in turning point space. The turning points of an orbit are the largest (resp. smallest) distance from the centre that an orbit can reach, and are called apocentre $r_+$, or pericentre $r_-$.  By convention, we take $r_+$ always positive, while $r_-$ has the same sign as $J$. Evidently, these quantities are integrals of the motion and $r_+ \\ge | r_- |$ (on circular orbits, $r_+=r_-$).  For a given pair $(r_+,r_-)$, the energy and angular momentum follow immediately from \\begin{equation} E=V_0(r_{+,-})-{1 \\over 2} { J^2 \\over r_{+,-}^2} \\label{Eext}, \\end{equation} which leads to \\begin{equation} E={r_+^2 V_0(r_+)    -     r_-^2 V_0(r_-) \\over r_+^2 - r_-^2} \\end{equation} and \\begin{equation} J=\\sqrt{2}r_+ r_- \\sqrt{V_0(r_+)-V_0(r_-) \\over r_-^2 - r_+^2}. \\end{equation} Inversely, $r_+$ and $r_-$ are found as the roots for $r$ of (\\ref{Eext}) for a particular $E$ and $J$.  We preferred these variables over the normal $(E,J)$ space, not only because they have an easy physical interpretation, but also because this representation is more related to our method for solving the linearized Boltzmann equation, which employs a grid interpolation in the turning point space (see following section).  We have chosen finite disks in order to compactify phase space. However, as can be seen from the distribution function (fig. \\ref{figdf0}), the disk reaches this limit at $r_{\\rm max}$ in a very smooth way. This is important since it has been proven that a sharp edge or, more generally, a sharp feature in $(E,J)$ space, can introduce additional instabilities which might not always be physical (\\cite{Toomre64}, \\cite{Sellwood_Kahn}).  \\begin{figure} \\vspace{5.0cm} \\special{ hscale=70 vscale=70 hoffset=-17 voffset= -10 hsize=480 vsize=500 angle=0 psfile=4583F4.ps } \\caption{Turning point \\label{figdf0} representation of the distribution function of model III.} \\end{figure} ",
        "conclusions": "A big part of this article was devoted to the description of a method for finding linear modes in stellar disks. The proposed strategy heavily relies on existing techniques, such as the matrix method and Fourier expansion along the unperturbed orbit, but differs from previous approaches by the fact that everything is calculated in ordinary coordinate and velocity space and that a numerical set of potential-density pairs is used. With the proposed scheme, the full perturbed distribution function is obtained with no extra calculation costs. In this way we have shown that calculations in coordinate space, although much less suited for theoretical considerations, can offer a fast and flexible alternative to action-angle variables when it comes to numerical computation of normal modes.  This method was applied to a set of unperturbed disk models, having a more or less realistic potential and an exponential mass density. This disks are embedded in a spherical, inert halo in order to obtain a self-consistent model. In agreement with various other studies (e.g. \\cite{Ath:Sellwood}; \\cite{Vau_Dej}), the calculations showed that these disks can be stabilised by increasing the velocity dispersion and/or the halo mass.  For almost isotropic velocity dispersions, a $Q_{cent}$ value of $2.0-2.5$ turned out to be a reasonable stability limit.  Comparison of the stability of the present models with the behaviour of disks embedded in quadratic potentials (\\cite{Vau_Dej}) shows a striking resemblance. Qualitatively, the stability behaviour of those simple uniformly rotating systems shares all the features that we have found for the more sophisticated models discussed in this paper. And, to a certain degree, there is even a quantitative agreement. It seems that, for some important aspects of the stability behaviour, the structure of the unperturbed distribution might be more inportant than the nature of the unperturbed potential."
    },
    "9603/astro-ph9603161_arXiv.txt": {
        "abstract": "Using recent high--resolution ($<$0.1\") radio observations of a large sample of Seyfert galaxies (Roy et al., 1994), we analyze the relations between their compact radio core emission and several nuclear and host galaxy properties (galaxy morphology, optical, infrared, X-ray, extended radio emissions, interaction parameters, and some emission line properties). We apply survival analysis techniques in order to exploit the information contained in the numerous \"censored\" data (upper limits on fluxes).  We find that Seyfert nuclei hosted in early--type galaxies are, on average, characterized by stronger radio core emission than the norm for Seyfert galaxies.  Galaxies with a nearby companion display enhanced radio core emission with respect to objects without companions. Furthermore, we confirm that Seyfert types 2 host more powerful compact radio cores than types 1.  Remarkably, radio core emission appears to be unrelated to optical, near--infrared, and far--infrared radiations, but shows some correlation with total radio emission and with tracers of nuclear activity such as the IRAS 12 and 25 $\\mu$m band, hard X-ray and narrow--line emissions. This favours the view that Seyfert radio cores are typically powered by AGN rather than by radio supernovae.  A link between radio core strength and bolometric luminosity is suggested, in analogy to what is observed in the cores of radio--quiet QSOs. ",
        "introduction": "Very high resolution radio observations of active galactic nuclei (AGN) can be profitably used to probe the central engine, since they are principally sensitive to nuclear non-thermal emission at high brightness temperature from AGN activity and help us to discriminate against low brightness temperature extended radio emission expected from the presence of circumnuclear HII regions. For instance, compact radio cores are fairly frequented detected in samples of Seyfert galaxies, but rarely detected in optically selected starburst galaxies (e.g., Norris et al., 1990).  Consistently, the deeper, very long baseline interferometry (VLBI) survey of ultraluminous infrared (IR) galaxies by Lonsdale, Smith \\& Lonsdale (1993), which did not reveal any correlation between the detection of compact radio cores and the optical classification of galactic nuclei, has been used to suggest the likely existence of AGNs embedded in dust within their nuclei. The observations of compact radio cores in Seyfert galaxies have recently been also employed as a new test of the canonical unification schemes of AGNs, according to which Seyfert type 1 objects (S1) would contain the same central engines as Seyfert types 2 (S2), but would have the broad--line region (BLR) and the strong optical, UV, and X-ray continua of the central source concealed from our view by a disk or torus of dusty molecular clouds (see, e.g., the review by Antonucci, 1993). If the torus surrounding the nucleus is optically thin at radio wavelengths, orientation effects should have no effect on the radio appearance of Seyfert galaxies. As a matter of fact, contrarily to old contentions (see, e. g., the review by Lawrence, 1987) based on poor statistics, recent low--resolution radio observations show no significant differences in the major radio properties (central and total radio powers, radio size, radio spectral index) of S1 and S2 objects (e.g., Edelson, 1987; Ulvestad \\& Wilson, 1989; Giuricin et al., 1990), as is expected from standard unified schemes.   But, surprisingly, in the high--resolution radio survey by Roy et al. (1994) compact radio cores were found more commonly in S2 than in S1 galaxies, although the disk radio powers, [OIII]$\\lambda 5007$\\AA $\\!$ emission line luminosities, and far--infrared (FIR) luminosities of the two classes of objects are similar (Roy et al., 1994, 1996). In order to reconcile their findings with the unified view of AGNs, the authors offered some models which resort to free--free absorption of S1 radio cores by the narrow--line region (NLR), if the radio cores lie in the BLR, or by individual NLR clouds, if the radio cores are located in the NLR.  In this paper we use Roy et al.'s (1994) radio core data set in order to examine the relations (as yet unexplored) between radio core emission and several properties of Seyfert galaxies and nuclei. Specifically, we examine the 2.3 GHz radio core fluxes $F_c$ of the optical+infrared--selected sample of Seyfert galaxies observed by Roy et al. (1994). The authors observed 157 Seyfert galaxies with the 275 km Parkes--Tidbinbilla Interferometer (PTI) at 1.7 or 2.3 GHz. They combined fluxes (or 5$\\sigma$ upper limits on fluxes) from different observations, converting 1.7 GHz fluxes to 2.3 GHz with the adoption of a spectral index $n=-0.7 ~ (F_{\\nu}\\propto\\nu^{n})$. Their high--resolution survey, characterized by uniform sensitivity and resolution, is blind to Kpc--scale emission of low brightness temperature, whereas it is sensitive only to structures with brightness temperature greater than $10^{5}$ K and sizes less than 0\".1, which correspond to 20--200 $pc$ over the redshift range $0.01<z<0.1$, typical of the sample.  Radio core data of 22 Seyfert galaxies have recently been published by Sadler et al. (1995) as part of their multifrequency (at 1.7 GHz, 2.3 GHz, 8.4 GHz) PTI survey of radio cores in a sample of nearby spiral galaxies. Their radio core fluxes appear to be in substantial agreement with Roy et al.'s (1994) for the 10 objects in common.  In this paper we restrict ourselves to Roy et al.'s (1994) 149 Seyfert galaxies having $z<0.1$, since distant galaxies might represent a different class of objects. The radio data set considered comprises about 35\\% of detections and 65\\% of upper limits (25\\% and 47\\% of detections for S1 and S2, respectively). The presence of many upper limits on the radio fluxes dictates the use of the techniques of Survival Analysis suitable for the treatment of \"censored\" data (see, e. g., the review by Feigelson, 1992). In practice we use the software package ASURV (Rev 1.2, La Valley, Isobe \\& Feigelson, 1992) in which the methods presented in Feigelson \\& Nelson (1985) and in Isobe, Feigelson \\& Nelson (1986) are implemented. We pay particular attention to removing distance effects when we study luminosity--luminosity correlations and when we compare luminosities of subsets of objects which lie at different average distances.   Basically, we search for correlations (as yet unexplored) between the radio core emission and several properties of Seyfert nuclei and host galaxies. These include galaxy morphology, optical, infrared, X-ray, extended radio emissions, interaction parameters and some emission--line nuclear properties. Several of these properties are indicators of nuclear activity.  Our study will also ensure that the above-mentioned difference between the radio core emissions of S1 and S2 objects is not simply due to a systematic difference of the two samples in some properties, to which radio core emission might be strongly related.   In \\S 2 we describe our statistical analysis. In some tables or plots we present the most outstanding results for the interesting cases. We shall not bother with extremely weak correlations (at the $<$90\\% significance level); correlations at a level falling between 90\\% and 95\\% will be referred to as marginal. All data compiled and results obtained are available on request. \\S 3 summarizes our results and contains our conclusions. ",
        "conclusions": "We may summarize the principal results of our paper as follows:  1) The compact radio cores of Seyfert galaxies are characterized by a median power of $\\log P_c$= 20.9 (W/Hz) together with a median core-to-total power ratio of 0.17. They are roughly comparable in power with the radio cores observed in the nearest (z$<$0.1) objects of the Bright Quasar Sample, whilst they appear to be typically weaker than the cores (frequently) detected in early--type galaxies and stronger than the analogous radio structures (rarely) detected in non--Seyfert spiral galaxies.  The radio core power correlates with the total radio power, ($P_c \\propto P_T^{1.0\\pm0.2}$), and the extended radio power ($P_T-P_c$), but it appears to be unrelated to the structural characteristics of extended radio emission (on an arcsec angular scale), which is not dominated by compact radio cores.  2) Seyfert nuclei hosted in early--type galaxies show a somewhat stronger radio core emission (typically by a factor of $\\sim$2) than the norm for Seyfert objects.  This appears to reflect the general strong association of radio--loud objects with early--type galaxies. Moreover, early--type Seyferts tend to show a slightly flatter $P_c$--$P_T$ relation ($P_c \\propto P_T^{0.8\\pm0.2}$) than the average for the whole sample. Our results indicate that also within the Seyfert class a connection between radio core activity and galaxy morphology is present (see, e.g., Wilson \\& Colbert, 1995; Fabian \\& Rees, 1995 for recent, different interpretations of this connection).  3) Galaxies with a nearby companion (especially non--early--type ones) display enhanced radio core emission with respect to objects without one.  This effect, which is in line with the behaviour of extended radio emission, lends further support to the idea that Seyfert activity is stimulated by interactions (e.g. Osterbrok's (1993) review), especially in early--type spirals (e.g., Monaco et al., 1994).  4) Seyfert 2 galaxies are confirmed to have more powerful radio cores (typically by a factor of $\\sim$1.5) than Seyfert 1. We have verified that this is not due to a different proportion of early types or objects with nearby companions in the two Seyfert classes.  5) Radio core emission is unrelated to optical, near--infrared, far--infrared emissions and ratios of some prominent narrow emission lines. In this respect, the behaviour of the radio core emission appears to be markedly different from that of total and extended radio emission, which well correlate with the optical and far--infrared emissions of the host galaxy.  6) Radio core emission shows some relation with prominent (albeit dissimilar in nature) signatures of AGN activity, such as hard X-ray, mid--infrared (in the IRAS $\\sim12\\mu$m and $\\sim25\\mu$m bands), and narrow--line ([OIII]$\\lambda5007$\\AA, $H\\beta$) emissions. In this case, the behaviour of radio cores preserves some similarity with that of extended radio power, which is known to be related to the same emissions.  The last results (point 6) suggest that Seyfert radio cores are typically powered by AGNs rather than by radio supernovae, although the power of radio supernovae (e.g., Weiler et al., 1986; Colina \\& P\\'erez--Olea, 1992) is comparable with that of the weakest detected radio cores in Seyferts.  Our results (points 1) and 2)) emphasize the difference (Sadler et al., 1995) between Seyfert radio cores and cores of early--type galaxies. On the other hand, a comparison with the compact radio cores observed in nearby QSOs is hardly feasible.  Although in many respects the latter structures do not seem identical to Seyfert cores, the discrepancies could be entirely accounted for by differences in spatial resolution and radio survey sensitivity (see, e.g., the end of \\S 2.4 for a discussion of the radio core--optical relation in Seyferts and QSOs).  In this respect, we deem it interesting that radio core strength appears to correlate with mid--infrared emission, since mid--infrared, being an approximate constant fraction of the bolometric flux, is believed to provide the best indication of the bolometric luminosity of Seyferts (e.g., Spinoglio \\& Malkan, 1989; Spinoglio et al., 1995).  Therefore, we suggest that the radio core power is likely to be linked to the bolometric luminosity of Seyfert nuclei, despite the fact that it represents a negligible fraction of the power output of the central engine. This link would assimilate Seyfert radio cores to radio--quiet QSO radio cores, which have recently been claimed to be reliable tracers of the QSO bolometric power output (Lonsdale et al., 1995), which in QSOs, unlike Seyferts (precisely Seyferts 2), is typically dominated by blue and UV radiation.  Besides high--resolution observations at higher frequencies (less affected by free--free absorption), there is a strong need for radio measurements on an arcsec/subarsec scale for many Seyfert objects. These measurements would probe the connection (if any) between compact radio cores and inner radio structures, for which limited and inhomogeneous information is now available."
    },
    "9603/astro-ph9603027_arXiv.txt": {
        "abstract": "A model of Lyman limit QSO absorption systems is investigated where they are produced in gaseous galactic halos with a two-phase structure: a hot phase at the halo virial temperature in approximate pressure equilibrium, and a cold, photoionized phase in the form of clouds confined by the pressure of the hot medium, and falling through it to the halo center, probably to accrete on a galactic disk. We show that the masses of these clouds must be over a relatively narrow range so that they are stable against gravitational collapse and evaporation. These masses, near $10^6 \\msun$, also lead to a covering factor near unity. The hot phase is required to have a core radius such that its cooling time in the core is equal to the age of the halo, and the mass in the cold phase is determined by the rate at which the hot gas cools. We calculate the number of Lyman limit systems arising in halos of different masses in a CDM model and their impact parameters implied by these assumptions, and find them to be in reasonable agreement with observations. The evolution with redshift is correctly reproduced at $z\\ls 2$, while at higher redshift we predict fewer absorbers than observed. We argue that this implies that high redshift Lyman limit systems arise in more extended infall regions. The observed low ionization systems such as MgII are well reproduced as arising from the photoionized phase, and are mostly in halos around massive galaxies, while CIV selected systems are predicted to be more commonly associated with lower mass galaxies and larger impact parameters. Some CIV can also arise from regions at intermediate temperature at the boundaries of the clouds. The hot phase may give rise to detectable absorption lines in OVI, while the column densities predicted for other highly ionized species are low and difficult to observe. ",
        "introduction": "In theories of hierarchical gravitational collapse for galaxy formation, it is usually assumed that luminous galaxies form at the center of dark matter halos as the gas collapses, is shocked, and cools (see White 1995 for a review). At any given epoch, halos can be parameterized by their circular velocity $\\vcir$ at the virial radius $\\rvir$. Here, we shall use as the virial radius that which contains an average density equal to $18\\pi^2 \\rho_{\\rm {crit}} = 178 \\rho_{\\rm {crit}}$, where $\\rho_{\\rm {crit}}$ is the critical density, and we assume $\\Omega=1$. This is the density obtained for a spherical perturbation that virializes at half the radius of maximum expansion (e.g., Peebles 1980). The mass of the virialized halo is then $M_{\\rm halo} = 24\\pi^3\\, \\rho_{\\rm {crit}}\\, \\rvir ^3$; using also $\\vcir = (GM_{\\rm halo}/\\rvir)^{1/2}$, we have $$\\rvir = {\\vcir\\, t\\over 2\\pi} , \\eqno(\\new) $$ and $$M_{\\rm halo}={\\vcir ^3 t\\over 2\\pi G} , \\eqno(\\new) $$ where $t$ is the age of the universe.  According to these theories, halos should continuously be merging and gravitationally accreting gas from the intergalactic medium. The accreted new gas should be shock-heated to the virial temperature. In halos of low $\\vcir$, the cooling time is shorter than the dynamical time, and much of the gas can cool before it has time to move to the center of the halo (e.g., White \\& Rees 1978). However, there must always be some hot gas left in the halo after rapid cooling, because once a certain fraction of the gas has cooled the density of the remaining gas will be sufficiently low for the cooling time to be longer than the dynamical time, and pressure equilibrium should be established in the hot halo (Fall \\& Rees 1985). The cooled gas should form clouds in pressure equilibrium with the hot gas and should then fall through the galactic halo. In the presence of an external ionizing background (as well as the X-ray emission from the halo and other ionizing radiation from local stars in the central galaxy), the clouds will remain photoionized at an equilibrium temperature near $10^4 \\kelvin$.  In a massive system where the cooling time of the hot gas is longer than the dynamical time of the halo, most of the gas does not cool rapidly. Only in the inner region, where the density is higher, the gas will radiate fast enough to cool. The subsequent evolution of the system is then related to the problem of cooling flows. The standard cooling flow model used to explain the observed structure of X-ray halos in elliptical galaxies and clusters of galaxies (see Fabian 1994 for a review) involves the formation of a multiphase medium, where the hot gas is separated in phases at different temperatures cooling independently of each other and forming clouds as they move to the center. The main difficulty of these models is that they need to assume that all the phases are comoving, when in fact the densest phases should rapidly fall to the center (e.g., Balbus \\& Soker 1989; Loewenstein 1989; Malagoli \\etal 1990). Although a thermal instability may exist in the presence of magnetic fields (Loewenstein 1990; Balbus 1991), it is not clear if relative motions can be prevented when the overdensities of different phases become large, and if clouds can ever form in cooling flows. Once formed, clouds should either turn into stars as they move through the halo, or otherwise coalesce in a gaseous disk where they are supported by angular momentum. The latter process should lead to a spiral galaxy, but it clearly does not take place in the X-ray halos where cooling is observed, which are known observationally to contain very gas-poor, giant ellipticals. The only possibility is then that clouds turn into stars in the halo, but the stellar mass function must be very different from that in normal galaxies: massive stars must be suppressed in order to avoid a very luminous young population of stars, which is not observed in the giant ellipticals in the center of X-ray halos with large cooling rates.  An alternative possibility is that there is a central adiabatic core in X-ray halos where energy is injected by a central source. Tabor \\& Binney (1993) have proposed this model for elliptical galaxies, where the energy may arise from an active galactic nucleus. In disk galaxies, the energy may arise from supernova explosions. Observations of ``superbubbles'' above the disk of our galaxy indicate that some of the energy from supernovae is indeed released into the halo (Tenorio-Tagle et al. 1990). Simple calculations show that the total energy released in supernova explosions in normal spiral galaxies is similar to the thermal energy of all the gas accreted: for a galaxy like the Milky Way, a total of $10^8$ supernovae can produce $\\sim 10^{59}$ ergs of energy, while an accreted baryon mass of $10^{11}\\msun$ with a velocity dispersion $\\sim 200 \\kms$ has a thermal energy of $10^{59}$ ergs. When a sufficient amount of hot gas is produced near the center, the halo will be globally unstable to convection, and this will set up an adiabatic core extending to a radius that should increase with the amount of energy released. As the accreted gas reaches the adiabatic core, it may cool and form photoionized clouds that fall through the adiabatic substrate, eventually being deposited in a gaseous galactic disk, or evaporating into the hot gas. This could be a way to generate a hot component with a large core radius, while a cold component moving through accretes to the galaxy disk.  In the following, we consider a simple phenomenological model for these two components of gaseous galactic halos, which should hopefully capture the essence of their structure.  \\subsection {The hot phase }  Let us consider the profile of the hot gas in a halo after a time $t_M$ since the last big merger. In other words, we assume that a large amount of gas was accreted into the system at the time of the merger, and that during the time $t_M$ since then there has only been a much lower rate of accretion. We also assume that the feed-back effect from the central part of the halo does not alter the properties of gas at large radius. We consider first the case where the cooling time at the virial radius $\\rvir$ is longer than $t_M$, and $t_M$ is longer than the dynamical time. In this case, the gas at large radius should preserve its initial density after virialization, since cooling is not very important. Numerical simulations of the formation of halos show that the gas profile will approximately follow the mass profile, with a slope close to isothermal (e.g., Evrard 1990, Navarro, Frenk, \\& White 1995); thus, we assume that the density of hot gas is $\\rho_h \\propto r^{-2}$ at large radius. There will then be a radius $r_c$ ($<\\rvir$) where the cooling time is equal to $t_M$, which we call the cooling radius.  At $r < r_c$, the density profile must be shallower than isothermal, because otherwise most of the cooling radiation would be emitted near the center. In cooling flow models, and when the gravitational potential is important, the core radius of the gas profile is similar to the cooling radius (Waxman \\& Miralda-Escud\\'e 1995). In general, the detailed density profile of the hot gas depends on the initial conditions of the assumed multiphase gas when it starts cooling (Nulsen 1986). A narrow initial temperature distribution will result in a steep gas density profile ($\\rho_h \\propto r^{-3/2}$ for an isothermal potential well). This does not result in a good model for galactic halos, since it produces an X-ray surface brightness much higher than observed. A more extended profile for a cooling flow model can be obtained when the hot gas forms a multiphase medium with a wide temperature distribution. The flattest possible density profile occurs when the temperature profile is adiabatic. An adiabatic temperature profile can also be produced by the presence of a central energy source, as described previously. For simplicity we will adopt a model where the adiabatic halo extends all the way to the cooling radius $r_c$. This gives the flattest core that is consistent with the convective stability of the hot halo, if the profile outside $r_c$ is isothermal.  In such an adiabatic halo, the pressure of the hot gas $P_h$ is related to $\\rho_h$ by $P_h\\propto \\rho_h^{5/3}$. From the equation of hydrostatic equilibrium in an isothermal sphere it then follows that the density and temperature profiles are given by $$ \\rho_h(r)=\\rho_h(r_c)\\left\\lbrack 1-{2\\over 5} {\\mu \\vcir ^2\\over kT_h(r_c)} {\\rm ln} {r\\over r_c} \\right\\rbrack ^{3/2} ; \\eqno(\\new ) $$ $$ T_h(r)=T_h(r_c)\\left\\lbrack 1-{2\\over 5} {\\mu \\vcir ^2\\over kT_h(r_c)} {\\rm ln} {r\\over r_c} \\right\\rbrack , \\eqno(\\new ) $$ where $\\mu$ is the average mass per particle. We assume that the temperature of the hot gas at the cooling radius $r_c$ is equal to the virial temperature, so that $T_h(r_c)=T_v\\equiv \\mu \\vcir^2/2k$. The density of hot gas at $r_c$, $\\rho_h(r_c)$, is obtained by requiring the cooling time at $r_c$ be equal to $t_M$: \\eqnam{\\rhoh} $$ \\rho_h (r_c)= {5 \\mu\\, kT_v\\over 2 \\Lambda (T_v)t_M } , \\eqno(\\new ) $$ where $\\Lambda$ is the cooling rate. To determine $r_c$, we assume $$ \\rho_h(r_c)={f_g\\vcir ^2 \\over 4\\pi G r_c^2} , \\eqno(\\new ) $$ so that the density of hot gas at this radius is a fraction of $f_g$ of the total density of the halo.  When $r_c>\\rvir$, the cooling time at the virial radius is already short compared to $t_M$, and the fraction of gas in the hot phase is small. In this case, the cooling radius $r_c$ defined by equations (5) and (6) is simply a parameter rather than a physical cooling radius. We assume that the residual hot gas at $\\rvir$ is still at the virial temperature and has a density such that its cooling time is equal to $t_M$, so that $T_h(\\rvir)=T_v$ and $\\rho_h(\\rvir)=(5 \\mu\\, kT_v)/[2\\Lambda (T_v)t_M]$. For simplicity we also assume that the hot gas within $\\rvir$ has density and temperature profiles similar to those described by equations (3) and (4), but with $r_c$ replaced by $\\rvir$. Outside $\\rvir$, the gas has typically not been shocked and not much hot gas should exist there.  The time $t_M$ in equation (5) corresponds to the interval between major mergers, since the gas is then heated to a stage from which it starts cooling. This time depends on the halo mass and the power spectrum; we ignore this complication here, and assume $t_M = t/(1+ \\Omega)$, where $t$ is the age of the universe. This reproduces an earlier formation of halos for low $\\Omega$ (see also Fig.8 of Lacey \\& Cole 1993 for $\\Omega=1$). The result obtained for the cooling radius in terms of the halo circular velocity is shown in Figure 1 at several redshifts, assuming $f_g=0.05$ and a metallicity $Z = 0.3 \\zsun$ for the cooling function adopted from Sutherland \\& Dopita (1993). For typical $L_*$ galaxies ($\\vcir \\simeq 250 \\kms$) at $z\\sim 0$, the cooling radius is near $150 \\kpc$, and it stays almost constant as we increase the halo velocity, up to scales of rich clusters of galaxies. For smaller galaxies, the cooling radius is larger and eventually reaches the virial radius; the value of $\\vcir$ at which this occurs is shown by the thick vertical ticks in Figure 1.  The total surface brightness emitted in soft X-rays by these hot halos can be computed given the emissivity per unit volume, $\\epsilon_x(r) = \\epsilon_0 \\rho^2_h(r) /\\mu_e^2$, where $\\mu_e$ is the gas mass per electron. The result is: \\eqnam{\\sxprof} $$ S_x(R) = {2\\epsilon_0\\, f_g^2\\, \\vcir ^4\\over 64\\pi^3 G^2\\, \\mu_e ^2 r_c^3}\\, F(R/r_c) $$ $$= 1.5\\times 10^{-8} \\left({f_g\\over 0.05} \\right)^2 \\left( {\\vcir\\over 200\\kms } \\right)^4 \\left( { 100\\kpc \\over r_c } \\right)^3 F\\left({R\\over r_c}\\right) \\erg\\,\\cm^{-2}\\sec^{-1}\\sr^{-1}, \\eqno(\\new) $$ where $R$ is the radius in projection and, for the density profile given by equation (3), the dimensionless function $F$ is $ F(x) = \\int_0^{1} dz \\left[1- (2/5) {\\rm ln}\\, (z^2 + x^2)\\right]^{2}$, where we have neglected the minor contribution from the gas outside the adiabatic core. Taking the characteristic values in equation (\\sxprof), and $\\epsilon_0=2\\times 10^{-23} \\erg\\cm^{3}\\sec^{-1}$, (adequate for temperatures $\\sim$ few $10^6 \\kelvin$ and $Z \\sim 0.1 \\zsun$), we find that for halos with $r_c=100\\kpc$ the surface brightness at $R= 100 (20) \\kpc$ is $S_x = 1.5\\times 10^{-8} (4\\times 10^{-8}) \\erg\\cm^{-2}\\sec^{-1}\\sr^{-1}$, and the total luminosity within $r_c$ is $L_x \\sim 7 \\times 10^{40} \\erg\\sec^{-1}$.  The value for the surface brightness obtained at radii $R\\sim 20 \\kpc$ is consistent with the observational upper limits. For example, McCammon \\& Sanders (1984) find that the X-ray surface brightness within $30 \\kpc$ around the galaxy M101 (where the disk circular velocity is $\\vcir \\simeq 200 \\kms$ ; Dean \\& Davies 1975) is smaller than $\\sim 3\\times 10^{-8} \\erg\\cm^{-2}\\sec^{-1}\\sr^{-1}$, corresponding to $L_x < 10^{40} \\erg \\sec^{-1}$ within the same radius. This upper limit was determined by subtracting the emission between $30$ and $50 \\kpc$ from the emission within $30 \\kpc$, so for a very extended X-ray halo the central surface brightness could probably be larger. For our galaxy, the emission measure that is derived from the hot gas at radius $\\sim 100\\kpc$ is $\\sim 2\\times 10^{-3}\\,{\\rm pc\\, cm}^{-6}$, which is consistent with the soft X-ray background (e.g. Wang \\& McCray 1993). Despite of this, it is clear that the density profile of the hot gas cannot be much steeper than that assumed in our model. The hot phase in our model (and the large total X-ray luminosity it produces) is still consistent with the X-ray observations of normal galaxies only because we have assumed a large core radius for the hot gas.  It is possible that $f_g$ is substantially smaller than $0.05$, and a steeper density profile is then allowed. However, the value of $f_g$ in our model is also constrained by the absorption line systems, as will be discussed later in the paper.  \\subsection {The cold phase}  The accreted gas that does not stay in the hot phase and cools must fall through the hot halo in the form of photoionized clouds. We refer to the gas in these clouds as the ``cold phase''. As discussed before, the rate at which cold gas accumulates in a halo is determined by both gas infall and gas cooling. The accumulation rate is in general not uniform in time: it may be substantially higher than the average during a big merger when large amount of gas is accreted and can cool. We shall, however, ignore this in our simple model.  In small halos, where the cooling time of the hot gas within the virial radius is much shorter than the dynamical time, most of the accreted gas will have cooled. Since the total gas mass accreted in a halo with circular velocity $\\vcir$ is $f_g \\vcir ^2 \\rvir /G$, the total mass of gas that has been in the cold phase can be written as $$M ={f_g\\vcir ^2 \\rvir \\over G} -\\int_0^{\\rvir} 4\\pi x^2\\rho_h(x)\\,dx , \\eqno(\\new) $$ where $\\rho_h(x)$ is the density of gas in the hot phase, as discussed in \\S2.1. For massive halos where the cooling time in their outer region is longer than the dynamical time, we assume that the infalling gas is initially shock-heated to the virial temperature, and cold gas results from subsequent cooling. In this case, the total mass of gas that has been in the cold phase can approximately be written as $$M ={f_g\\vcir ^2 r_c \\over G} -\\int_0^{r_c} 4\\pi x^2\\rho_h(x)\\,dx , \\eqno(\\new) $$ where we have assumed that the amount of cold gas outside $r_c$ is negligible.  We now assume that most of the cold gas has been formed near a radius $r_{min} \\equiv {\\rm min}(r_v, r_c)$, and that a constant mass inflow rate of the cold gas is present within $r_{min}$, and no cold gas is present outside this radius. This approximation is adequate when there is heated gas in an adiabatic core, since the accreting gas would then start forming clouds at $r_c$ when it mixes with the heated gas, and in a cooling flow model where, for a wide initial temperature distribution, most of the gas is deposited near $r_c$. The cold gas should form clouds that will fall through the halo, and it will therefore stay in the cold phase only for a short time, before forming stars or merging with the disk being formed in the halo center. We assume the mass flow rate to be ${\\dot {M}} =M/t_M$, and that the clouds move to the halo center with a constant velocity $v_c$. Assuming also spherical symmetry for the gas distribution, we can write the density of the cold gas as a function of the distance $r$ to the halo center as $$\\rho_c (r) = {\\dot M \\over 4\\pi r^2 v_c } = { f_g\\vcir ^2 \\rmin \\over 4\\pi G\\,t_M\\, r^2\\, v_c} \\left\\lbrack 1-{\\rmin ^2\\over r_c^2} \\int _0^1 x^2 \\left(1-{4\\over 5}\\ln x\\right)^{3/2} \\right\\rbrack  ~. \\eqno(\\new) $$ In the absence of a hot phase, the cold gas is in free fall and $v_c$ must be of the order of the virial velocity $\\vcir$. The friction of hot gas may cause cold clouds to move at a terminal velocity which is smaller than $\\vcir$. However, since halos on galactic scales have cooling times that are similar to the dynamical times, the inflow velocity should not be much smaller than the virial velocity (see \\S 3). The inflow velocity could be much smaller than $\\vcir$ in massive systems, and also if relative motions of different phases could be prevented by a magnetic field. For simplicity, we will assume $v_c$ to be a constant throughout a halo and to have a value of the order of $\\vcir$. We will examine the effect of changing the value of $v_c$. ",
        "conclusions": ""
    },
    "9603/astro-ph9603033_arXiv.txt": {
        "abstract": "We present a new method for calculating linear cosmic microwave background (CMB) anisotropy spectra based on integration over sources along the photon past light cone. In this approach the temperature anisotropy is written as a time integral over the product of a geometrical term and a source term. The geometrical term is given by radial eigenfunctions which do not depend on the particular cosmological model. The source term can be expressed in terms of photon, baryon and metric perturbations, all of which can be calculated using a small number of differential equations. This split clearly separates between the dynamical and geometrical effects on the CMB anisotropies. More importantly, it allows to significantly reduce the computational time compared to standard methods. This is achieved because the source term, which depends on the model and is generally the most time consuming part of calculation, is a slowly varying function of wavelength and needs to be evaluated only in a small number of points. The geometrical term, which oscillates much more rapidly than the source term, does not depend on the particular model and can be precomputed in advance. Standard methods that do not separate the two terms and require a much higher number of evaluations. The new method leads to about two orders of magnitude reduction in CPU time when compared to standard methods and typically requires a few minutes on a workstation for a single model. The method should be especially useful for accurate determinations of cosmological parameters from CMB anisotropy and polarization measurements that will become possible with the next generation of experiments. A programm implementing this method can be obtained from the authors. ",
        "introduction": "The field of cosmic microwave background (CMB) anisotropies has seen a rapid development since its first detection by the COBE satellite only a few years ago. There are now several reported experimental results that are detecting anisotropies on degree angular scales (see \\cite{Scott95} and \\cite{Bond96} for a recent review), which together with a few upper limits on smaller angular scales already give interesting limits on cosmological models. With the development of the new generation of experiments now being proposed one hopes to accurately map the CMB sky from arcminute scales to several degree scales. The amount of data thus provided would allow for an unprecedented accuracy in the determination of cosmological parameters. As theoretical modelling shows (\\cite{Bond94,Hu95c,Jungman95,Seljak94}), CMB anisotropies are sensitive to most of the cosmological parameters and have a distinctive advantage over other cosmological observations in that they probe the universe in the linear regime. This avoids the complications caused by physical processes in the nonlinear regime and allows to use powerful statistical techniques to search over the parameter space for the best cosmological model (see e.g. \\cite{Jungman95}). A large stumbling block in this program at present is the speed of theoretical model calculations, which are still too slow to allow for a rapid search over the parameter space. This limitation was partially removed by the development of appoximation methods (\\cite{Hu95a},b; \\cite{Seljak94}), which can give fast predictions of CMB anisotropy with a 10\\% accuracy. However, these approximations are not accurate enough to exploit the complete amount of information that will be present in the future CMB observations. This is especially true for some of the more extreme cosmological models, where simple approximations break down and lead to systematic inaccuracies in the results. Obviously, it would be useful to have a fast method that would not be based on any approximations and would lead to accurate results for any cosmological model. The purpose of this paper is to present a new method of CMB calculation that satisfies these requirements.  Theoretical calculations of the CMB anisotropies are based on linear theory of cosmological perturbations, developed first by Lifshitz (1946) and applied to the CMB anisotropies by Peebles and Yu (1970). In this early calculation only photons and baryons were included, but later workers extended the calculations to include dark matter (Bond \\& Efstathiou 1984, 1987; Vittorio \\& Silk 1984), curvature (Wilson \\& Silk 1981; Sugiyama \\& Gouda 1992; White \\& Scott 1995), gravity waves or tensor modes (Crittenden et al. 1993) and massive neutrinos (Bond \\& Szalay 1983; Ma \\& Bertschinger 1995; Dodelson, Gates \\& Stebbins 1995). Most of these and more recent calculations (e.g. Holtzmann 1989; Stompor 1994; Sugiyama 1995) solve for each Fourier mode of temperature anisotropy $\\Delta_T(\\vec k)$ by expanding it in Legendre series up to some desired $l_{max}$ and then numerically evolve this system of equations in time from the radiation dominated epoch until today. Typically this means evolving a system of several thousand coupled differential equations in time, a slow process even for the present day computers. In addition, because each multipole moment is a rapidly oscillating function one has to densely sample it in values of $k$ with typical number of evaluations of the order of $l_{max}$. Even the fastest codes at present require several hours of CPU time for each theoretical model (Sugiyama 1995), while some numerically more accurate ones (e.g. \\cite{BB95}) require more like tens or hundreds of hours.  In this paper we explore a different approach to compute CMB anisotropies based on integration of the sources over the photon past light cone. The method is a generalization of approximate method developed by one of the authors (Seljak 1994). It differs from it in that it is exact, in the sense that it can achieve arbitrary precision within the limits of linear perturbation theory. By rewriting the system of equations in the integral form one can separate between the geometrical and dynamical contributions to the anisotropies. The former do not depend on the model and need to be computed only once, while the latter contain all the information on the particular model and can be computed with a small system of equations. Solving for CMB anisotropies using this integral form greatly reduces the required computational time. The outline of the paper is as follows: in \\S 2 we present the basic system of equations that needs to be solved both in the standard and in the integral method. In \\S 3 we present in some detail a practical implementation of the integral method, highlighting the computational differences between it and the standard Boltzmann method. We conclude in \\S 4 by discussing possible applications where the new method can be particularly useful. ",
        "conclusions": "In this paper we presented a new method for accurate calculations of CMB anisotropy and polarization spectra. The method is not based on any approximations and is an alternative to the standard Boltzmann calculations, which are based on solving large numbers of differential equations. The approach proposed here uses a hybrid integro-differential approach in solving the same system of equations. By rewriting the Boltzmann equations in the integral form the solution for the photon anisotropy spectrum can be written as an integral over a source and a geometrical term. The first is determined by a small number of contributors to the photon equations of motion and the second is given by the radial eigenfunctions, which do not depend on the particular cosmological model, but only on the geometry of space.  One advantage of the split between geometrical and dynamical terms is that it clarifies their different contributions to the final spectrum. A good example of this is the temperature anisotropy in the non-flat universe, which can be be written using a similar decomposition, except that spherical Bessel functions have to be replaced with their appropriate generalization (\\cite{abbott86}). This will be discussed in more detail in a future publication, here we simply remark that replacing radial eigenfunctions in a non-flat space with their flat space counterpart (keeping comoving angular distance to the LSS unchanged) is only approximate and does not become exact even in the large $l$ (small angle) limit. The geometry of the universe leaves its signature in the CMB spectra in a rather nontrivial way and does not lead only to a simple rescaling of the spectrum by $\\Omega_{\\rm matter}^{-1/2}$ (\\cite{Jungman95}).  The main advantage of our line of sight integration method is its speed and accuracy. For a given set of parameters it is two orders of magnitude faster than the standard Boltzmann methods, while preserving the same accuracy. We compared our results with the results by Sugiyama (1995) and by Bode \\& Bertschinger (1995) and in both cases the agreement was better than 1\\% up to a very high $l$ for all of the models we compared to.  The method is useful for fast and accurate normalizations of density power spectra from CMB measurements, which for a given model require the CMB anisotropy spectrum and matter transfer function, both of which are provided by the output of the method. Speed and accuracy are even more important for accurate determination of cosmological parameters from CMB measurements. In such applications one wants to perform a search over a large parameter space, which typically requires calculating the spectra of a several thousand models (e.g. \\cite{Jungman95}). One feasible way to do so is to use approximation methods mentioned in the introduction. These can be made extremely fast, but at a cost of sacrificing the accuracy. While several percent accuracy is sufficient for analyzing the present day experiments, it will not satisfy the requirements for the future all-sky surveys of microwave sky. Provided that foreground contributions can be succesfully filtered out (see Tegmark \\& Efstathiou 1995 for a recent discussion) one can hope for accuracies on the spectrum close to the cosmic variance limit, which for a broad band averages can indeed reach below 1\\% at $l>100$. It is at this stage that fast and accurate CMB calculations such as the one presented in this paper will become crucial and might enable one to determine many cosmological parameters with an unprecedented accuracy."
    },
    "9603/astro-ph9603012_arXiv.txt": {
        "abstract": "We have obtained very deep near-infrared images in the fields of 10 QSOs whose spectra contain damped Lyman-$\\alpha$ absorption (DLA) systems with $1.7<z_{\\rm abs} <2.5$.  The main aim of our investigation is to provide new constraints on the properties of distant galaxies responsible for the DLA absorption. After subtracting the point spread function associated with the QSO light, we have detected galaxies very close to the QSO line of sight (projected distance $1.2$--$1.3\\,$arcsec) in two fields.  These sources therefore represent promising candidate galaxies responsible for the DLA absorption. Placed at the absorber's redshift, the impact parameter is $10h_{50}^{-1}\\,$kpc and the luminosity is close to $L_K^{\\ast}$. Such parameters are consistent with the hypothesis, verified for metallic systems at lower redshift, that slowly-evolving massive galaxies produce at least some of  the absorption systems of high column density in QSO spectra out to beyond $z\\simeq2$.  In addition to detecting these candidate DLA galaxies, the radio-loud QSOs in our sample show a significant excess of  sources on larger scales ($\\theta \\simeq 7\\,$arcsec); this excess is not  present in the radio-quiet QSO sightlines. We calculate that such an excess could be produced by luminous galaxies in the cores of clusters associated with radio-loud QSOs. Both results confirm that deep imaging of selected QSOs can be a powerful method of finding samples of likely $z\\simeq2$ galaxies.  Follow-up near-infrared spectroscopy is required to secure galaxy redshifts and star formation rates, while deep HST imaging can determine sizes and morphologies, providing valuable information on galaxy properties at large look-back times. ",
        "introduction": "Considerable progress has been made in recent years in understanding the evolutionary behaviour of starlight from distant field galaxies. Systematic redshift surveys (Ellis et al. 1995, Lilly et al. 1995) have provided large datasets from which the evolution of the luminosity function of galaxies has been directly determined to redshifts $z\\simeq1$. Such studies are complemented by smaller, but independent surveys based on the absorbing properties galaxies present to background luminous QSOs (Bahcall \\& Peebles 1969).  Following the pioneering work of Bergeron and Boiss\\`e (1990), Steidel \\& Dickinson (1995) have identified those galaxies responsible for the MgII absorption seen in the sightlines to 68 QSOs and have thus derived luminosity functions and rest-frame colours for the absorbing field galaxies at a mean redshift of $z\\simeq0.7$.  It is gratifying that both techniques, viz surveys that select galaxies by their {\\it emission\\/} and those via their {\\it absorption}, give consistent results for the evolutionary behaviour of massive galaxies to $z\\simeq1$ (Steidel \\& Dickinson 1995, Ellis 1995). Surprisingly little evolution is seen in both the colours and luminosities of such galaxies indicating they were already in place at large look-back times. The principal difference between the techniques lies in understanding the origin and fate of the numerous star-forming dwarfs seen in the redshift surveys (Ellis 1995); such galaxies do not appear to be prominent in the metallic absorber-selected samples.  The slow evolution of massive galaxies to $z\\simeq1$ has important implications for the likely detection of similar systems at higher redshift, particularly for early-type disk galaxies whose star formation rates are expected to increase slightly with redshift. Limited field spectroscopy has been performed beyond $z=1$ (Cowie, Hu \\& Songaila 1995) but QSO absorption samples are available in abundance. Indeed, some constraints on the nature of galaxies responsible for high $z$ metallic systems are available from statistical studies of a few MgII systems (Steidel et al. 1995, private communication) and a sample of more distant CIV absorbers (Arag\\'on-Salamanca et al. 1994).  Although no spectroscopic identifications have yet been made for these systems, the implied luminosities and impact parameters are consistent with continuation of the trends found at lower redshift.  Of particular interest in this regard is the nature of the material which produces the damped Lyman alpha absorption (DLA, Wolfe et al. 1986).  Conventional wisdom holds that such systems arise infrequently when the disk of an absorbing galaxy presents a small impact parameter to the QSO.  Statistical analyses indicate the abundance of MgII, CIV and DLA systems is consistent with a single population of massive galaxies whose volume density approximates that seen today, the principal difference arising in the cross-section. If this is the case, deep imaging in good seeing will be required to isolate the galaxies responsible for the DLA systems. Although this explanation is not accepted universally (see, e.g., Yanny \\& York 1992), any constraints on the properties of galaxies close to the QSO sightline where a DLA absorber may be present will be of interest.  This paper begins to explore the possibilities observationally. In Section 2 we present new infrared imaging observations of 10 QSO sightlines each chosen to contain a damped Lyman alpha absorption. Section 3 analyses the distribution of faint sources found as a function of projected distance from the QSO, and discusses two promising cases where the DLA absorber may have been found. The implications are briefly discussed and Section 4 presents our conclusions. ",
        "conclusions": "We have obtained very deep ($K_{\\rm lim}\\simeq 21.3$--$21.7$) images in good seeing conditions of the fields of 10 QSOs with damped Lyman-$\\alpha$ absorption systems in the $1.7<z_{\\rm abs} <2.5$ range. The main aim is to constrain the properties of  those galaxies responsible for the DLA absorption and, more generally, to find samples of high redshift galaxies for evolutionary studies. From the analysis of those images, several results emerge:  1. For two of the QSOs, point spread function subtraction techniques reveal the existence of one galaxy in each case very close to the QSO line of sight (projected distance $1.2$--$1.3\\,$arcsec, or $\\simeq10\\,$kpc at the absorber redshift). We argue that, given the small probability of a chance alignment and the implied impact parameter at $z_{\\rm abs}$,  these two objects are promising candidates for the DLA absorption.  Both would have luminosities close to $L_K^{\\ast}$ and their properties would be consistent with trends identified in metal line absorber galaxies found at lower redshifts.  2. The radial distribution of faint sources around the QSOs shows good agreement with that expected from published $K$-band number counts for projected distances $\\gtrsim7\\,$arcsec.  At $\\theta\\lesssim7\\,$arcsec, radio quiet QSOs show no sign of excess in the number of objects, but the radio loud QSOs show an excess of $2.8$ sources per line of sight significant at the $99.5$\\% level.  3. Recognising that the statistical significance of the excess is not very high given the number of lines of sight sampled,  the difference in behaviour detected for the radio quiet and radio loud QSOs is nonetheless suggestive.  Since the DLA systems should not be affected by the radio properties of the QSOs ($z_{\\rm abs}\\ll z_{\\rm em}$),  we suggest  the excess may be associated with the QSOs themselves. If that is the case, radio loud QSOs at $z\\simeq2.4$ could well reside in groups or clusters of galaxies much in the same way as closer counterparts, and we are just seeing the 2--3 brightest galaxies there.  4. The average luminosity of the excess galaxies, if placed at $z_{\\rm em}$, is $\\simeq 2L_K^{\\ast}$, and the magnitude distribution reaches $\\simeq 2\\,$mag brighter than $L_K^{\\ast}$. If these galaxies are the high redshift counterparts of present day brightest cluster galaxies, it implies luminosity evolution in the $K$-band of $\\sim1\\,$mag from $z\\simeq2.4$ to the present, in the sense that high-$z$ galaxies were brighter than today.  Imaging high redshift QSO sightlines is an efficient method to identify likely high-$z$ galaxies for evolutionary studies.  To the precision where evolutionary comparison with lower redshift galaxies are concerned, the small difference in look-back-time makes it immaterial whether the galaxies are at $z_{\\rm abs}$ or at $z_{\\rm em}$.  A larger sample of radio loud and radio quiet QSOs, with and without DLA systems is needed in order to add weight to our statistical conclusions.  Once the excess is firmly established,  HST imaging would provide morphology and colours and be useful in searching for objects even closer to the QSO line of sight.  Obtaining redshifts can resolve the ambiguities and this may be practical shortly via the use of  near-infrared spectrographs offering OH-suppression where faint emission lines may be seen against the contrast of the strong infrared background ."
    },
    "9603/astro-ph9603154_arXiv.txt": {
        "abstract": "We have examined the power spectral behavior of \\cyg\\ using approximately 1100 days of BATSE data.  These data have been searched for correlations between different power spectral features, and for correlations between power spectral features and energy spectral parameters derived from BATSE occultation analysis. ",
        "introduction": "A data base has been created combining 1100 days of \\cyg\\ observations with BATSE, from JD 2\\,448\\,387 to 2\\,449\\,480 (\\mbox{May 11, 1991} to \\mbox{May 8, 1994}). It includes observations of fast variability (0.01--0.488 Hz) as well as flux measurements (\\mbox{45--140 keV}) and energy spectral fits. The goal of this work is to look for correlations between various parameters derived from power spectral, flux, and energy spectral measurements.  This work is part of a larger on-going project to systematically study the power spectral behavior of all black-hole candidate sources observed with BATSE.  In Section 2 we describe the methods used for creating the power spectra and the results of the occultation analysis.  In Section 3 we show that correlations occur between several quantities derived from the energy spectral and power spectral fits, and briefly discuss these results. ",
        "conclusions": ""
    },
    "9603/hep-ph9603307_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  There is a wide consensus that, although very successful wherever it has been tested, our present Standard Model leaves open too many fundamental issues in particle physics to be an ultimate theory of nature. These lectures will focus on two of these that have to do with the understanding of what lies behind the mechanism of mass generation in general, and with the masses and the properties of neutrinos in particular.  A basic assumption of the Standard Model is the Higgs mechanism, which is introduced as a way to generate the masses of all the fundamental particles. This mechanism implies the existence of a fundamental scalar boson \\cite{HIGGS}. If such an elementary boson exists it is widely believed that some stabilising principle - e.g. supersymmetry (SUSY) - should be operative at the electroweak scale in order to explain the stability of its mass scale against quantum corrections associated with physics at super-high energies. The observed joining of the three \\gau coupling strengths as they are evolved from the presently accessible energies up to a common scale of $\\sim 10^{16}$ GeV provides circumstantial evidence that supersymmetry seems to set in somewhere around $M_{SUSY} \\sim 10^3$ GeV. Probing the details of this rich structure constitutes one of the main goals in the agenda of the next generation of elementary particle colliders.  Another fundamental issue in the Standard Model refers to the properties of neutrinos. Apart from being a theoretical puzzle, in the sense that there is no principle that dictates that neutrinos are massless, as postulated in the Standard Model, nonzero masses may in fact be required in order to account for the data on solar and on atmospheric neutrinos, as well as for an explanation of the dark matter in the universe. The possible detection of nonzero neutrino masses would have far-reaching implications for the understanding of fundamental issues in particle physics, astrophysics, as well as for the large-scale structure of our universe.  The above two different types of extensions of the Standard Model may be inter-connected. An example is provided by supersymmetric models with spontaneously broken R parity, which necessarily imply non-vanishing neutrino masses. As a result in some of these models there are novel processes that could be observed at high energy colliders whose magnitude would be correlated to the mass of the neutrinos. One interesting aspect of these models is that they may affect the physics of the electroweak sector in such a remarkable way, that can be probed in various present and future experiments, both at accelerator as well as underground installations, as we will describe.  \\subsection{Standard Model Basics} \\label{321}  The Standard Model is a Yang-Mills theory based on the \\321 gauge group, and described by the field representations in table 1, where all fermions are left-handed. \\bet \\caption{Multiplets of the Standard Model} \\begin{center} \\begin{math} \\begin{array}{|c|c|} \\hline & \\ \\ \\ {\\mbox SU(3)\\otimes SU(2)\\otimes U(1)} \\\\ \\hline W     & (1,3,0)\\\\ B     & (1,1,0)\\\\ \\hline \\phi  & (1,2,1)\\\\ \\hline \\ell_a  & (1,2,-1)\\\\ e_a^c   & (1,1,2)\\\\ Q_a     & (3,2,1/3)\\\\ u_a^c   & (\\bar{3},1,-4/3)\\\\ d_a^c   & (\\bar{3},1,2/3)\\\\ \\hline \\end{array} \\end{math} \\end{center} \\label{SM} \\eet The fundamental constituents of matter - quarks and leptons - interact mainly due to the exchange of the gauge bosons. The Lagrangian describing the Standard Model can be found in many textbooks and reviews \\cite{nubooks}. In order to comply with the fact that the weak interaction is mediated by massive vector bosons, the $W$ and the $Z$, the gauge symmetry has to be broken. The best way to do this in a way consistent with renormalizability and unitarity is through the nonzero vacuum expectation value (VEV) \\beq v/\\sqrt{2} = \\VEV{\\phi^0} \\label{WEAKSCALE} \\eeq of the neutral component of a complex Higgs scalar doublet $\\phi$ \\beq \\phi = \\left(\\begin{array}{ccccc} \\phi^+\\\\ \\phi^0 \\end{array}\\right). \\label{DOUBLET} \\eeq Through the term $|D \\phi |^2$, where $D$ denotes the covariant derivative that follows from the quantum numbers in table 1, this gives mass to all gauge bosons except one. This corresponds to the surviving $U(1)_{em}$ gauge symmetry corresponding to the unbroken QED electro-magnetic gauge invariance, whose generator is the electric charge, identified as \\beq \\label{Q} Q = T_3 + Y/2 \\eeq where $T_3$ and $Y$ denote the diagonal isospin and hyper-charge generators. Diagonalizing the neutral gauge boson mass matrix leads to a mixing between the gauge bosons $B_{\\mu}$ and ${W_3}_{\\mu}$ associated to the generators $Y$ and $T_3$, \\beq \\begin{array}{ccccc} Z_{\\mu}  = \\cos \\theta_W {W_3}_{\\mu} + \\sin \\theta_W B_{\\mu}\\\\ A_{\\mu}  = -\\sin \\theta_W {W_3}_{\\mu} + \\cos \\theta_W B_{\\mu} \\label{Weinberg1} \\end{array} \\eeq where \\beq \\theta_W = \\arctan \\frac {g'}{g} \\label{Weinberg2} \\eeq is the electroweak mixing angle, $g$ and $g'$ are the $SU(2)$ and $U(1)$ gauge coupling strengths, respectively. The massless vector boson $A_{\\mu}$ is the photon, while $Z_{\\mu}$ has a mass \\beq m_Z = \\frac {g v}{2 \\cos \\theta_W} \\label{Z} \\eeq while the charged gauge bosons $W^{\\pm} = \\frac{1}{\\sqrt2} (W_1 \\mp iW_2)$ have a mass \\beq m_W = \\frac {g v}{2}. \\label{W} \\eeq The $W$ and the $Z$ gauge bosons have been discovered in 1983 by the UA1 and UA2 experiments at CERN. The properties of the $Z$ have now been very precisely determined by the LEP experiments \\cite{Martinez}, while those of the $W$ mostly come from UA2 and the Fermilab CDF and D0 experiments \\cite{top}. The measured \\gau boson mass values agree well with the electroweak theory predictions, once radiative corrections are included.  \\subsection{The Fermion Sector}  The Standard Model contains three generations of quarks and leptons ($a=1,2,3$), also given in table 1. Different generations are identical in their gauge properties, but differ in mass. The fermion assignment is clearly not invariant under parity, since only the left handed fermions (not the anti-fermions) transform as $SU(2)$ doublets. As a result the fermions can not be given gauge invariant masses. Fermion masses arise from Yukawa-type interactions with the Higgs doublet scalar, \\bea \\begin{array}{ccc} {h_u}_{ab} \\: u_a^c \\: Q_b \\: \\tau_2 \\phi^* + {h_d}_{ab}\\: d_a^c \\: Q_b \\: \\phi + {h_e}_{ab} \\: e_a^c \\: \\ell_b \\: \\phi \\: . \\label{YUKAWA} \\end{array} \\eea which generate masses to all charged fermions $f=u,d,e$ \\bea h_f v / \\sqrt2 = M_f \\eea after electroweak breaking. Thus in the Standard Model the weak gauge boson as well as the charged fermion masses are generated spontaneously, by the Higgs mechanism. Note that the coupling matrices $h_u$, $h_d$ and $h_e$ are arbitrary complex non-symmetric matrices in generation space. This gives rise to a vast proliferation of parameters in the model. The understanding the pattern of fermion masses and mixing is one of the main challenges of the theory and clearly lies beyond the Standard Model.  The currents describing the interaction of quarks and leptons with the standard electroweak \\gau fields can be written simply from the gauge covariant derivatives that follow from table 1. For example the terms in the Lagrangian describing the interaction of quarks and leptons with the charged gauge bosons $W^{\\pm}$ are \\beq \\label{CC} \\frac{ig}{\\sqrt{2}} W_{\\mu}^+ \\sum_{a=1}^{n} \\bar{u}_{La} \\ \\gamma_{\\mu} \\ d_{La} + \\frac{ig}{\\sqrt{2}} W_{\\mu}^+ \\sum_{a=1}^{n} \\bar{\\nu}_{La} \\ \\gamma_{\\mu} \\ e_{La} + H.C. \\eeq where $u_a$ and $d_a$ are weak-eigenstate quarks ($u,c,t$ and $d,s,b$), and $\\nu_a$ and $e_a$ denote the three \\neus and the three charged leptons, $e,\\mu ,\\tau $. The contact with the old four-Fermi weak interaction theory is established through \\beq \\frac {G_F}{\\sqrt{2}} = \\frac {g^2}{8 m_W^2}\\:. \\label{fermi} \\eeq where $G_F$ is the Fermi constant, well determined from $\\mu$ decay. In addition to the charged current weak interactions, quarks and leptons also have neutral current weak interactions, described by \\beq \\label{NC} \\frac{ig'}{\\sin \\theta_W}\\ Z_{\\mu} \\ \\sum_A \\bar{\\Psi}_{A} \\ \\ \\gamma_{\\mu} \\ (T_3 - \\mbox{sin}^2 \\theta_W Q) \\ \\Psi_{A} \\eeq where $g \\sin \\theta_W = g' \\cos \\theta_W = e = \\sqrt{4\\pi\\alpha}\\:$ and $\\Psi_A$ denotes an arbitrary helicity fermion belonging to any of the multiplets in table 1. By including electroweak radiative corrections in the charged and neutral currents one obtains a very good description of all known weak interaction processes. In particular, all known neutral current phenomena are well described by a single parameter, the electroweak mixing angle $\\theta_W$. The success of the LEP experiments in the precise determination of the electroweak parameters has been so remarkable that just the internal consistency of the various measurements is sufficient to provide a very good determination of the mass of the top quark $m_t = 180 \\pm 14$ GeV, with the error largely due to the lack of knowledge of the Higgs boson mass \\cite{Martinez}. This is in excellent agreement with the direct measurement at Fermilab \\cite{top}.  Similarly, the total $Z$ decay width, as well as its partial widths have been precisely measured by the LEP collaborations, leaving little room for new physics. Of special interest to us is the measurement of the invisible $Z$ width \\cite{PDG95} \\beq \\label{NNU} \\Gamma_{inv}^{Z} = 499.9 \\pm 2.5 \\: \\mbox{MeV} \\eeq which can be translated into a measurement of the effective number of Standard Model neutrino generations. This places a very stringent constraint on models of neutrino mass where lepton number is a global symmetry spontaneously broken at low energies (see section \\ref{majorons} below).  Let us now turn to a discussion of the flavour structure of the charged current weak interactions. First we consider the case of quarks. In order to determine the form of the charged current we first diagonalize separately the quark mass matrices $M_u$ and $M_d$ \\bea \\label{DIAG} U_{Ru}^\\dagger M_u U_{Lu} = diag(m_u,m_c,m_t)\\\\ \\nonumber U_{Rd}^\\dagger M_d U_{Ld} = diag(m_d,m_s,m_b) \\eea Since these matrices can not in general be made diagonal in the same basis, it follows that the charged current quark weak interactions, rewritten in terms of the mass-eigenstate quarks $u_i$ and $d_i$, take the following form \\beq \\frac{ig}{\\sqrt{2}} W_{\\mu}^+ \\ \\bar{u}_{Li} \\ \\gamma_{\\mu} \\ V_{ij} \\ d_{Lj} + H.C. \\label{CKM} \\eeq where $V = U_{Lu}^\\dagger U_{Ld}$. Some of the parameters in the unitary matrix $V$ are not physical, since they can be eliminated by appropriate quark redefinitions. For $n$ generations $V_{ij}$ involves a set of \\beq n(n-1)/2 \\label{ANGLES1} \\eeq mixing angles $\\theta_{ij}$ and \\beq n(n-1)/2-(n-1) \\label{PHASES1} \\eeq CP violating phases $\\phi_{ij}$. They can be brought to the following useful form \\cite{2227,OLD} \\beq V = \\omega_0 (\\alpha) \\prod_{i<j}^{n} \\omega_{ij} (\\eta_{ij})\\: \\omega_0^\\dagger (\\alpha) \\label{CKM2} \\eeq where $\\theta_{ij} = \\mid \\eta_{ij} \\mid$ and $\\omega_0 = diag (e^{i \\alpha_i})$ is a diagonal matrix of phases, for example, \\beq \\omega_{12} (\\eta_{12}) = \\left(\\begin{array}{ccccc} c_{12} & e^{i \\phi_{12}} s_{12} & 0  & 0 & ...\\\\ -e^{-i \\phi_{12}} s_{12} & c_{12} & 0 & 0 & ...\\\\ 0  & 0 & 1 & 0 & ...\\\\ 0  & 0 & 0 & 1 & ...\\\\ ... & ... & ... & ... & ... \\end{array}\\right) \\label{CKM3} \\eeq For the simplest 2 generation case one sees that the $2 \\times 2$ mixing matrix $V$ is completely determined by one parameter, the Cabibbo angle. In this case CP is necessarily conserved. As first noted by Kobayashi and Maskawa \\cite{KM}, for three generations there are 3 mixing angles and a single phase that describes CP violation in the Standard Model. The first parametrizations of the three generation mixing matrix $V$ were given in ref. \\cite{KM} and \\cite{MAIANI}. A useful truncated form was introduced by Wolfenstein \\cite{WOL}. So far this scheme has been proven sufficient to account for the observed pattern of quark mixing and CP violation in the Standard Model \\cite{PDG95}. However, one lacks a fundamental theory of mass and flavour and, as a result, the parameters in $V$ are all undetermined and must be extracted from experiment. Note that the CP phases in \\eq{CKM2} are not directly observable: only special combinations (as many as given by our counting) which are invariant under quark rephasings have physical meaning. For 3 generations this CP violating combination is unique.  Finally, note that the weak neutral current of the mass-eigenstate quarks has a trivial flavour structure, i.e. it is given by the identity matrix both in the weak and in the mass-eigenstate basis. This important property is usually referred to as the Glashow-Iliopoulos-Maiani (GIM) mechanism \\cite{GIM}. It accounts for the absence of flavour changing neutral current mediated processes so far observed. Any deviations from this rule would indicate the presence of physics beyond the Standard Model. ",
        "conclusions": ""
    },
    "9603/astro-ph9603140_arXiv.txt": {
        "abstract": "Recent observations have allowed the geometry and kinematics of the M87 jet to be tightly constrained.  We combine these constraints with historical Very Long Baseline Interferometry (VLBI) results and the theory of synchrotron self-absorbed radio cores in order to investigate the physical properties of the jet.  Our results strongly suggest the jet to be dominated by an electron-positron (pair) plasma. Although our conservative constraints cannot conclusively dismiss an electron-proton plasma, the viability of this solution is extremely vulnerable to further tightening of VLBI surface brightness limits. The arguments presented, coupled with future high-resolution multi-frequency VLBI studies of the jet core, will be able to firmly distinguish these two possibilities. ",
        "introduction": "Jets are intimately linked with accretion processes.  They can be found in almost all astrophysical situations in which matter is believed to be undergoing disk accretion onto some central collapsed object.  Scales range from the slow ($\\sim 100$\\kmps) sub-parsec jets associated with forming stars to those found originating from powerful active galactic nuclei (AGN) which are initially relativistic and can propagate for hundreds of kpc before disrupting.  Despite intense study, many basic questions remain open.  The underlying formation mechanism is still uncertain (although hydromagnetic processes associated with the accretion disk seem a promising candidate mechanism: Blandford \\& Payne 1982; Lynden-Bell 1995 and references therein).  Moreover, whereas jets from forming stars are believed to be molecular material entrained by a faster outflow of atomic material, the basic nature of AGN jets has not been unambiguously determined.  Possibilities include an electron-proton plasma (i.e. normal material), an electron-positron plasma (i.e. a pair plasma), Poynting flux or combinations of these.  This is an observationally difficult issue to address: even the handful of well studied jets only show power-law continuum radiation (presumably of synchrotron origin) with no sign of any spectral features from material moving with the bulk flow.  This is to be contrasted with the Galactic source SS~433 in which the bulk flow of the jet is a known source of atomic line emission.  Celotti \\& Fabian (1993) have addressed the issue of the matter content of jets in a sample of radio-loud quasars and radio galaxies. Combining the synchrotron self-Compton constraints with indications of the kinetic luminosity suggest that, for the sample as a whole, the jets are either cold electron-positron flows or electron-proton flows with an electron low-energy cutoff of $\\sim 50\\MeV$.  They argue for the electron-proton case on the basis of annihilation constraints under the assumption that the pairs originate from the inner regions of the accretion flow.  However, the issue of the matter content of any one given extragalactic jet (as opposed to a large, possibly heterogeneous, sample) has not been addressed in detail.  An (observationally based) determination of the matter content of an individual well studied jet would be an important step in the study of jet formation, propagation and emission.  M87 (NGC~4486) is a giant elliptical galaxy near the centre of the Virgo cluster of galaxies.  It is associated with the radio source (Virgo-A, 3C~274) and classified as a FR-I radio galaxy (Fanaroff \\& Riley 1974) on the basis of its edge-darkened morphology and low radio luminosity ($P_{\\rm 178MHz}\\sim 1\\times 10^{32}$\\ergps\\Hz$^{-1}$). This galaxy contains the most spectacular example of an extragalactic jet in the northern sky.  It was the first extragalactic jet to be discovered (Curtis 1918) and, since then, it has been subjected to intense observational studies at all wavelengths (see Biretta 1993 for a recent review).  The complex knotty structure of the jet has been resolved at wavelengths from the radio through to X-rays (Biretta, Stern \\& Harris 1991).  The spectrum and proper motions (Biretta, Zhou \\& Owen 1995) of the knots yield direct constraints on the physical processes operating and the geometry of the system.  Hubble Space Telescope (HST) studies of the base of the jet reveal a rotating gas disk apparently lying normal to the jet direction (Ford et al. 1994; Harms et al. 1994).  The projected appearance and velocity of the disk suggests it to be at an inclination of $42\\pm 5\\degmark$ and to enclose a mass of $(2.4\\pm 0.7)\\times 10^{9}\\Msun$.  Also, Very Long Baseline Interferometry (VLBI) has resolved structure within the core of M87 on spatial scales of $\\sim 0.01$\\pc.  The VLBI studies will be discussed in more detail in Section 3.1.  The synchrotron spectrum of this jet has been recently examined by Meisenheimer, R{\\\"o}ser \\& Schl{\\\"o}telburg (1996).  This wealth of data makes M87 the most promising candidate for any physical study of extragalactic jets.  In this paper we combine these observational constraints of the M87 jet with the theory of synchrotron radiation in order to constrain the physical parameters of the jet.  In particular, we address the particle density, magnetic flux density and the matter content of the jet.  In order to keep this paper focussed, our treatment is as model independent as possible.  The arguments we use depend only upon an observationally well determined kinematic model and the standard synchrotron jet model of Blandford \\& K{\\\"o}nigl (1979).  Our arguments are independent of any models (of which there are a plethora) for the initial formation, collimination and acceleration of the jet.  Section 2 briefly reviews the relevant observations and the kinematic model to which they lead (Biretta 1993; Biretta, Zhou \\& Owen 1995).  Section 3 discusses the synchrotron self-absorption model for the radio emission from the core and the subsequent constraints on the particle density and the magnetic flux density.  Section 4 places further constraints on the particle density of the jet using the large scale limits on the total kinetic luminosity.  The resulting constraints suggest the jet to be electron-positron dominated rather than electron-proton dominated.  Section 5 discusses the robustness of the result and some implications.  It is found that an electron-proton jet possessing an electron population with a low-energy cutoff $\\sim 10m_{\\rm e}c^2$ also satisfies our constraints although this solution is extremely vulnerable to further tightening of the VLBI surface brightness lower limit.  Section 6 summarizes our conclusions.  Stellar surface brightness fluctuations give a distance to M87 of $15.9\\pm 0.9\\Mpc$ (Tonry 1991), independently of Hubble's constant. In the rest of this paper we assume a distance to M87 of $16\\Mpc$. ",
        "conclusions": "The M87 jet has the best determined geometry and kinematics of any extragalactic jet in the sky.  M87 is also one of the very few galaxies in which the mass of the central black hole ($2.4\\pm 0.7\\times 10^9\\Msun$) is well constrained (by HST observations).  This makes M87 an obvious system in which to study the physical properties of extragalactic jets.  We have utilized the standard theory of synchrotron self-absorbed cores in order to constrain the magnetic field, $B$, and the relativistic (proper) particle density, $n$, of the jet.  The 5\\,GHz data implies $n\\gamma_{\\rm min}\\approxgt 10\\pcmcu$.  For this jet to carry the kinetic liminosity inferred from global energetic arguments, the density is either $n\\sim 1\\pcmcu$ (electron-proton jet) or $n\\sim 10^2\\pcmcu$ (electron-positron jet).  We cite this as evidence for an electron-positron dominated jet rather than an electron-proton dominated jet.  An electron-proton jet with $10<\\gamma_{\\rm min}<100$ is consistent with the present (conservative) constraints but is extremely vulnerable to further tightening of the VLBI limits.  If the core seen in 100\\,GHz VLBI data is self-absorbed, the electron-proton jet is ruled out.  Some theoretical arguments are also presented detailing the problems of the high-$\\gamma_{\\rm min}$ scenario.  To make further observational progress, the self-absorbed core has to be identified and resolved.  This requires high-resolution multi-frequency VLBI observations.  The sub-arcsec X-ray imaging capability of {\\it AXAF} will separate the core X-ray emission from other components (e.g. inner knot emission) and thus allow stricter limits to be placed on the SSC radiation of the self-absorbed core. Such data, together with arguments similar to those presented here, should allow the matter content of this archetypal extragalactic jet to be unambiguously determined."
    },
    "9603/astro-ph9603006_arXiv.txt": {
        "abstract": "We model the surface brightness profiles of a sample of 119 Abell Brightest Cluster Galaxies (BCG), finding a generalised deVaucouleurs R$^{1/n}$ law, where $n$ is a free parameter, to be appropriate.  Departures from the $R^{1/4}$ law are shown to be a real feature of galaxy profiles, not due to observational errors or coupling of $n$ with the other model parameters. BCG typically have values of $n$ greater than 4. The shape parameter $n$ is shown to correlate with effective half-light radius, such that the larger BCG have larger values of $n$.   This continues a trend noticed amongst ordinary elliptical galaxies and dwarf ellipticals, such that the fainter galaxies have smaller values of $n$. ",
        "introduction": "The structure of the first-ranked members of Abell clusters, or brightest cluster galaxies (BCG), has long been of interest given that these galaxies can be used as standard candles for exploring the large-scale structure of the universe (\\cite{HMS56}; \\cite{AS72a}, \\cite{AS72b}; \\cite{GaO75}; Lauer \\& Postman 1994). \\cite{Hos80}, for example, developed the $L-\\alpha$ relationship, which relates BCG structure to their luminosities through fixed physical apertures, thus reducing the cosmic scatter in the aperture luminosities. Understanding the $L-\\alpha$ relationship has been of particular interest recently, given its use by Lauer \\& Postman (\\cite{LP92}, 1994) to probe the linearity of the Hubble flow and to detect large scale bulk flows out to the 15,000 km/s scale.  The structure of BCG may also reflect galactic cannibalism (\\cite{OaT75}; \\cite{HaO78}; Hoessel 1980; \\cite{SGH83}), where other cluster members may be captured through dynamical friction, changing both the luminosity and the shape of the dominant BCG.  While to first order BCG appear to be ordinary, if highly luminous, elliptical galaxies, in many ways they are a special class of objects. \\cite{TaR77} showed, for example, that their luminosity function was not consistent with them being drawn simply as the brightest member of a standard \\cite{Sch76} luminosity function. \\cite{Oem76} showed that as a class, BCG differed structurally from other giant ellipticals of similar luminosity, having more extended limiting radii at a given total luminosity than other giant elliptical galaxies. \\cite{Sbt86} conducted an extensive survey of BCG brightness profiles, finding them to be shallower than those for more ordinary elliptical galaxies, such that the BCG were always more extended at a given surface brightness level than would be suggested by simply scaling by total luminosity.  \\cite{HOS87} also concurred with this, showing BCG to follow a shallower (and tighter) relationship between effective radii, $R_e$, and effective surface brightness, $I_e$, than do ordinary elliptical galaxies (\\cite{OaH91}; Graham 1996). This distinction in form was additionally interesting as \\cite{Sbt86} showed it to arise in the relatively brighter mid-portions of the BCG profile, regardless of whether or not the BCG had additional envelopes at faint surface brightness making them cD galaxies as well. A cD galaxy is a giant elliptical that has a separate extended low surface-brightness envelope, which is evident as an inflection in the brightness profile (Oemler 1976) --- typically at $\\mu_V\\sim24$ or greater (\\cite{KaD89}).  \\cite{Sbt86} emphasized that the classic $r^{1/4}$ law, introduced by de Vaucouleurs (\\cite{1948}, \\cite{1953}) to describe elliptical galaxies at all luminosities, was often a poor match to the BCG profiles, in many cases fitting the profiles only over a restricted range of surface brightness. Looking at the profiles presented by \\cite{Sbt86}, it becomes apparent that many of the BCG would be better fitted by power laws rather than $r^{1/4}$ laws. This is also found to be the case in Ledlow \\& Owen's (1995) work with galaxies in rich clusters.  This appears to have little to do with whether or not a BCG is also a cD galaxy. An additional caveat is that because a constant power law will rise above an $r^{1/4}$ law at large radii, a cD envelope may be erroneously detected as a separate component in $r^{1/4}$ plots, even though a single power law could describe the BCG completely. In short, classification of a BCG as cD or not is problematic; our present investigation, however, is not affected by this issue as this apparently has little to do in any case with the distinct over all structural properties of BCG from those of giant ellipticals.  If some BCG are better described by power laws, and others by $r^{1/4}$ laws, an inspection of the profiles presented by \\cite{Sbt86} shows that BCG fall along a continuum between the two forms. We are thus motivated to advance the \\cite{Ser68} form, which includes both laws as a better description and generalization of BCG brightness profiles. In the Sersic form, $I(r)\\propto\\exp(-(r/R_e)^{1/n}),$ where $n$ is a free parameter. As $n\\rightarrow1$ from the $n=4$ de Vaucouleurs form, the Sersic form becomes increasingly exponential, while as $n\\rightarrow\\infty,$ the Sersic form becomes a pure power law.  \\cite{YaC94} fitted the Sersic form to a sample of dwarf elliptical (dE) galaxies, showing a correlation between galaxy absolute magnitude and the shape parameter $n$, such that the fainter galaxies had the smaller values of $n$. The same pattern in profile shape had also been noted by \\cite{Dav88}. Working with a sample of Fornax low surface brightness (LSB) diffuse dwarf elliptical and spheroidal galaxies, they showed a trend between $n$ and the logarithm of the scale-radius $R_{e}$ such that the scale radius increased as $n$ increased.  Caon, Capaccioli, and D'Onofrio (1993) showed that this trend extended into the domain of ordinary elliptical galaxies.  It has also recently been shown to hold for the bulges of spiral galaxies by \\cite{APB95} and \\cite{CJB96}.  We are now in an excellent position to explore the systematics of the BCG structural properties as \\cite{PL95} present a complete and volume-limited sample of BCG out to 15,000 km/s. The data was obtained with large-area CCDs and reduced with photometric uniformity in mind, given the goal to use this data set to detect the subtle effect of bulk flows on the $L-\\alpha$ relationship residuals (Lauer \\& Postman 1994). Here we show that the correlation between $n$ and scale-radius (Caon, Capaccioli, and D'Onofrio 1993) extends to BCG, where $n$ is typically greater than 4, whereas for the dE it is seen to be less than 4. The existence of this relation across such a large range of galaxy sizes must be telling us something fundamental about the structure and formation processes of elliptical galaxies in general.  In the next section we present some of the theory behind the $R^{1/n}$ model. Section 3 presents the BCG surface brightness data and the parameterised model fits to it. In Section 4 we discuss the results of the model fitting and compare this in the context of work on other galaxies in Section 5.  We investigate how $n$ relates to other measures of galaxy structure, both local, Section 6, and global, Section 7.  Our conclusions are given in Section 8. ",
        "conclusions": "The nature of the $R^{1/n}$ law is such that as $n$ increases the galaxy surface brightness profile flattens, and for $n>15$ it is very well represented by a power law. BCG typically have values of $n$ greater than 4 - that is their light profiles are less curved than the classic $R^{1/4}$ law; 40$\\%$ of our sample is well described by a power law. The range in profile shapes is real and not due to noise in the galaxy profiles or due to a coupling of the three parameters in the model. There is a trend between $n$ and the half-light radius, such that the larger galaxies have larger values of $n$. This trend appears to be a continuation of that noticed for dE galaxies, LSB galaxies, through to normal E and S0 galaxies and on to BCG, suggesting some common physical processes must be at play in the formation of all of these galaxies.  This global shape parameter, $n$, is shown to be independent of Richness Class and Bautz-Morgan type, suggesting that the galaxy environment (in so far as RC and BM type represent this) is not responsible for the shaping of the bulk distribution of stars in the galaxy.  We note that our analysis excludes the outer envelopes of the cD galaxies.   Whilst $n$ is a global measure of the galaxies light profile, $\\alpha $ is a measure the galaxies structure at a point.  $\\alpha $ is shown to be related to $n$ and $R_{e}$ in an opposing manner.   As one moves to larger galaxies, $n$ increases causing $\\alpha $ to decrease but at the same time $R_{e}$ increases causing $\\alpha $ to increase.  The dominating factor depends on which part of the profile one samples the value of $\\alpha $ at."
    },
    "9603/astro-ph9603076_arXiv.txt": {
        "abstract": "{\\it HST} has so far provided Cepheid distances to nine galaxies. Although not sufficient yet to determine the distance of the {\\it extended} Virgo cluster, they are decisive for the distance scale in two ways. (1) Seven of the galaxies contribute to a much improved calibration of the Tully-Fisher relation.  Applying this to a {\\it complete} sample of Virgo spirals one obtains a cluster distance of $(m-M)=31.79\\pm 0.09$. Other distance indicators support this value. The adopted linear distance of $22.0\\pm 0.8$~Mpc combined with the cluster velocity of $1178\\pm 32~{\\rm km\\,s^{-1}}$ (in the CMB frame) gives $H_{0}=54\\pm 2$ (internal error). (2) Six of the galaxies have been the site of seven SNe Ia with well observed maxima. Their resulting calibration in absolute magnitudes gives $M_{\\rm B}({\\rm max})=-19.53\\pm 0.07$ and $M_{\\rm V}({\\rm max})=-19.49\\pm 0.07$ with negligible {\\it intrinsic} scatter.  If this calibration is used to determine the distances of {\\it all} distant SNe Ia with known maxima and with $1100 < v < 30\\,000\\ {\\rm km\\,s^{-1}}$, $H_{0}$ becomes $56\\pm 3$ (internal error). Systematic errors tend to make this an upper limit; in particular the case $H_{0}\\ge 70$ can be excluded. -- The conclusion is that the large-scale value of the Hubble constant is $H_{0}=55\\pm 10$ (external error). ",
        "introduction": "The influence of {\\it HST} on the determination of $H_{0}$ is already enormous and it can only grow. In Table~1 a compilation is given of distance determinations with {\\it HST} having some bearing on $H_{0}$.  The values scatter between 55 and 80 and the formal mean is $H_{0}=61\\pm 3$. This, however, is not the best value, because some values shown are mutually incompatible.  There are now two self-consistent routes to $H_{0}$ to which {\\it HST} has heavily contributed. The first route via the Virgo cluster is described in Section 2, the second route using supernovae of type Ia (SNe Ia) is discussed in Section 3.  \\begin{table}[thb] \\caption{$H_{0}$ determinations from {\\it HST}} \\begin{center} \\scriptsize \\begin{tabular}{llr} \\tableline \\noalign{\\smallskip} \\multicolumn{2}{l}{I.~Cepheids} & \\\\ \\noalign{\\smallskip} a) in M101 & Kelson (1995) & \\\\ & confirms Sandage \\& Tammann (1974), & \\\\ & from which follows                  & $55\\pm \\phantom{1}9$ \\\\ \\noalign{\\smallskip} b) in Leo Group & Tanvir et al.~(1995)           & $69\\pm \\phantom{1}8$ \\\\ & present paper                  & $57\\pm \\phantom{1}6$ \\\\ \\noalign{\\smallskip} c) in Virgo Cluster & & \\\\ \\quad NGC 4321 & Freedman et al.~(1994)     & $80\\pm 17$ \\\\ \\quad NGC 4639 & Sandage et al.~(1996)      & $47\\pm 10$ \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{3}{l}{II.~Tully-Fischer distance of Virgo cluster using 11 Calibrators}  \\\\ \\multicolumn{3}{l}{\\quad with Cepheid Distances (7 of which from HST)} \\\\ \\noalign{\\smallskip} & Federspiel et al.~(1996) & $52\\pm \\phantom{1}6$ \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{3}{l}{III.~SNe Ia calibrated through Cepheids} \\\\ \\noalign{\\smallskip} SN 1937C & Saha et al.~(1994)           & $52\\pm \\phantom{1}9$ \\\\ SN 1972E & Hamuy et al.~(1995)           & $65\\pm 10$ \\\\ SN 1972E & Riess et al.~(1995)           & $67\\pm \\phantom{1}7$ \\\\ SN 1895B & Schaefer (1995)               & $51\\pm \\phantom{1}7$ \\\\ 6 SNe Ia & Branch et al.~(1996)           & $57\\pm \\phantom{1}7$ \\\\ 7 SNe Ia & Sandage et al.~(1996)           & $58\\pm \\phantom{1}4$ \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{3}{l}{IV.~Globular Clusters} \\\\ in M87 & Whitmore et al.~(1995) & $78\\pm 11$ \\\\ & Sandage \\& Tammann (1996) & $62\\pm \\phantom{1}9$ \\\\ in Coma & Baum et al.~(1995) & $<65$ \\\\ \\noalign{\\smallskip} \\tableline \\tableline \\end{tabular} \\end{center} \\end{table} Much of what follows depends on Cepheid distances. A word on the reliability of their P-L relation is therefore in place. The P-L relation of Madore \\& Freedman (1991), adopted in the following, {\\it assumes} an LMC modulus of 18.50. Actual confirmation of this value to better than 0.10~mag comes from the P-L relation calibrated by Galactic Cepheids (Sandage \\& Tammann 1968; Feast \\& Walker 1987) and independently from the ring of SN~1987A (Panagia et al.~1991; Gould 1994); further support is given by RR~Lyr stars and other distance determinations (cf. Tammann 1996). The zero point of the adopted P-L relation is therefore secure to $<5\\%$ in linear distance.  The slope of the relation is well determined by the LMC Cepheids; it is of less importance as long as the Cepheids under consideration cover a sufficient period range, which is also needed to avoid selection effects (Sandage 1988). Metallicity effects are believed to be small (Freedman \\& Madore 1990; Chiosi et al.~1993). ",
        "conclusions": "The two independent routes towards the large-scale value of $H_{0}$, via the Virgo cluster and SNe Ia, give 54$\\pm 2$ and 56$\\pm 3$ (internal errors), respectively.  Their only inter\\-depend\\-ence is that they rely on Cepheids (predominantly observed with {\\it HST}), which are the least controversial distance indicators at present. Together they make a strong case for $H_{0}=55\\pm 10$ (external error). Values of $H_{0}<40$ are equally unlikely as values of $H_{0}>70$\\,.  The relatively low value of $H_{0}$ is supported by additional methods, e.g.~TF and other distances of field galaxies (Sandage 1994, 1996 and references therein), and the Zeldovich-Sunyaev effect (Lasenby \\& Hancock 1995, Rephaeli 1995).  Baum's et al.~(1995) {\\it HST} photometry of globular clusters in the Coma cluster requires $H_{0}<65$\\,. A gravitationally lensed quasar sets $H_{0}<70$ (Dahle et al.~1994). Models of SNe Ia could not be understood if $H_{0}$ was $\\geq$ 60 (Branch et al.~1996) or in no case $\\geq$70 (H\\\"oflich \\& Khokhlov 1996; Ruiz-Lapuente 1996).  We believe that literature values significantly larger than $H_{0}=65$ are explained by an unwarrantedly high Virgo velocity, the unrealistic hope to fathom the depth of the Virgo cluster with only a single galaxy, the myth of a sharp, dispersionless cutoff of the luminosity function of planetary nebula shells, the reliance on the suspicious surface brightness fluctuation method, and/or simply by Malmquist bias which always artificially increases the value of $H_{0}$\\,."
    },
    "9603/astro-ph9603062_arXiv.txt": {
        "abstract": "We present {\\it Hubble Space Telescope} images of star-forming galaxies at redshifts $z>3$. These galaxies have been selected using ground--based images and color criteria sensitive to the presence of a Lyman discontinuity in the otherwise flat (in $f_{\\nu}$ units) UV spectral energy distribution of unreddened star formation. The spectroscopic confirmation of these $z > 3$ galaxies is reported in a companion paper (Steidel et al. 1996). The {\\it HST} images, which probe the rest-frame UV between 1400 and 1900 \\AA, show that the morphologies of the $z>3$ galaxies are generally compact and exhibit a relatively high degree of spherical symmetry, although we find a few cases of more diffuse light profiles and several cases where the objects are comprised of multiple compact structures. Overall, the dispersion of morphological properties is relatively narrow, in contrast to the variety found in star-forming galaxies at intermediate redshifts ($z\\sim 1$). The galaxies with compact morphology are typically characterized by a small but resolved ``core'', approximately $\\simlt 0.7$ arcsec in radius, or about $5\\h50$ ($8.5\\h50$) kpc with $q_0=0.5$ (0.05), and half-light radii of 0.2--0.3 arcsec, or $1.4$--$2.1\\h50$ ($2.4$--$3.6\\h50$) kpc. These sizes and scale lengths are similar to those of present-day bulges or intermediate-luminosity spheroids. The ``cores''  are often surrounded by lower surface-brightness  nebulosities, generally asymmetrically distributed. The minority of more diffuse galaxies do not possess this core, and an exponential function provides a very good fit to their light profiles. In contrast to highly elongated or irregular structures, such as ``chain galaxies'', that are found at $z \\sim 1$, the $z>3$ galaxies are characterized by a relatively high degree of spherical symmetry. The morphological properties, space density, star-formation rates, masses, and early epoch of the star-formation phase all support the hypothesis that we have identified the progenitors of present-day luminous galaxies at the epoch when they were forming the stars of their spheroidal components. ",
        "introduction": "The physics of galaxy formation remains largely unconstrained from an empirical perspective. In the last few years, several deep redshift surveys (Lilly et al. 1995a and 1995b; Steidel et al. 1995; Glazebrook et al. 1995a; Cowie et al. 1994; Cowie, Hu \\& Songaila 1995a) and deep post-refurbishment {\\it HST} imaging (Driver et al. 1995a and 1995b; Glazebrook et al. 1995b; Schade et al. 1995; Cowie Hu \\& Songaila 1995b) have extensively probed the evolutionary state of galaxies at intermediate redshifts ($z<1.0$), corresponding to $<55$ ($<60$)\\% of the life of the Universe ($H_0=50$ \\hbl\\ and $q_0=0.5$ ($0.05$). This work seems to show that the evolution of the luminosity function has followed rather diverse tracks for galaxies of different luminosity and morphological type. A common conclusion is that the population of luminous galaxies, i.e., the systems currently identified as relatively massive ellipticals and spirals, are characterized by at most a modest amount of evolution in luminosity and/or number density since $z \\sim 1$. The available data therefore suggests that the epoch of formation of the most massive and oldest systems predates that probed by the current surveys. While there has evidently been substantial evolution of later type systems in number and/or luminosity in the relatively recent past, the population of galaxies possessing a substantial spheroidal component (i.e., ellipticals and early--type spirals) has been remarkably quiescent over the redshift range probed by the redshift surveys, suggesting that the important epoch for their formation lies far beyond $z \\sim 1$. The fact that spheroidal systems, which contain approximately half of the present-day stars (Schechter \\& Dressler 1987), had assembled rather early in the course of the evolution has long been the rationale for searches for ``primeval'' galaxies. Establishing the epoch and the formation mechanism of these systems will place significant constraints on theories of galaxy and structure formation.  In a companion paper (Steidel et al. 1996, S\\&96 hereafter) we report the discovery of a substantial population of star-forming, but otherwise normal (i.e. {\\it non-- AGN}), galaxies at redshifts $z>3$. These galaxies were found using color criteria sensitive to the presence of a Lyman discontinuity (due to a combination of a galaxy's opacity to its own UV continuum radiation, the intrinsic energy distribution of hot stars, and the opacity of intervening gas at high redshift) in the otherwise flat (in $f_{\\nu}$ units) and featureless spectral energy distribution (SED) of unreddened star formation (see Steidel, Pettini, \\& Hamilton 1995).  The color criteria are very efficient in selecting galaxies at high redshifts, and although we have been able only recently to confirm them spectroscopically, over the last few years we have been collecting a fairly large sample of $z>3$ galaxy candidates (Steidel \\& Hamilton 1992, 1993; Steidel et al. 1995; Giavalisco et al. 1996, in preparation) and we have been investigating their morphological properties with the {\\it Hubble Space Telescope} ({\\it HST}) (see, e.g. Giavalisco et al. 1995). At the time of this writing, only 23 galaxies out of about 100 have securely measured redshifts. However, as we have detailed in S\\&96, we expect that a very high fraction, probably $\\simgt 90$\\%, of the candidates not yet spectroscopically confirmed are also in our targeted redshift range of $3.0 \\le z \\le 3.5$. We have already presented a number of arguments linking the population we have identified with the luminous galaxies of the present epoch (S\\&96). In this paper, we present deep {\\it HST} images of 19 of the ``Lyman break'' galaxies, 6 of which have secure redshift measurements in the range $2.8 \\simlt z \\simlt 3.4$. The images probe the rest-frame UV spectrum in the range $1400$--$1900$ \\AA\\ and have resolution high enough that we can attempt a quantitative discussion of their morphological properties. Thus, for the first time, we can characterize in a statistically significant way the evolutionary status of galaxies at a time when the universe was $\\simlt 20$\\% of its current age. ",
        "conclusions": ""
    },
    "9603/astro-ph9603124_arXiv.txt": {
        "abstract": "A study of the distribution of barred and nonbarred disk galaxies in the Virgo cluster is presented in an attempt to use the frequency and spatial distribution of galaxies with specific morphological features to study the efficiency of various environmental effects on the evolution of disk galaxies in clusters. The velocity distribution of the barred spirals in the Virgo region is clearly different than that of the nonbarred spirals, suggesting that barred spirals are more common in the main condensation of the cluster. A sample cleansed of galaxies not belonging to the main cluster condensation using the subcluster assignments of Binggeli et al.\\markcite{members} (1993) bears this out, showing that the radial distribution of barred spirals is more centrally condensed than that of nonbarred spirals. In contrast to the spiral galaxies, the distribution of barred S0 galaxies is statistically indistinguishable from that of nonbarred S0's. Consideration of the level of tidal perturbation due to the cluster mass distribution as compared to that due to individual galaxies suggests that tidal triggering by the cluster mass distribution is the most likely source of the enhanced fraction of barred spirals in the cluster center. ",
        "introduction": "The galaxies in the cores of present day galaxy clusters are preferentially found to be elliptical and lenticular galaxies, rather than spiral galaxies which predominant in lower density regions of the universe (Gisler \\markcite{gisler} 1980; Dressler \\markcite{dress} 1980a). It seems unlikely that this is simply due to different galaxy types forming in different environments, since observations of clusters at redshifts of $z \\sim 0.3-0.4$ show a much higher percentage of spirals in these clusters than in present day clusters (Couch et al. \\markcite{couch} 1994; Dressler et al. \\markcite{dress2} \\markcite{dress3} 1994a-b). Apparently some environmental effect is responsible for altering cluster spirals seen at high redshift beyond recognition as spirals by the current day. Several possible mechanisms have been proposed; ram pressure sweeping of the interstellar medium by the gas responsible for the cluster x-ray emission (Gunn \\& Gott \\markcite{gunngott} 1972), the cumulative effects of galaxy-galaxy collisions in the dense cluster core (Richstone \\markcite{rich} 1975), or tidal effects due to the gravitational field of the cluster as a whole (Merritt \\markcite{merritt} 1983). Unfortunately to date, it has been difficult to distinguish observationally between the various possible mechanisms.  \\newpage  In a study of the Coma cluster, Thompson \\markcite{comabars} (1981) found that the percentage of barred galaxies within approximately $0.75 {\\rm Mpc}$ of the cluster center ( $H_{0} = 75 {\\rm km/sec/Mpc}$ and a Virgo cluster distance of 20 Mpc is used throughout this paper ) was significantly higher than in the outer parts of the cluster. Thompson noted that this either meant that the excess of barred galaxies represented a kinematically distinct component confined to the core of the cluster, or that bars were triggered by some mechanism as disk galaxies entered the cluster core. If the latter is true, the lifetime of the induced bars must be the less than or the order of the core crossing time; for the Coma cluster the core crossing time is approximately $10^{9} yr$, which is around 4--5 disk rotation times for a typical spiral galaxy. Simulations of galaxy-galaxy interactions (Noguchi \\markcite{noguchi2} \\markcite{noguchi} 1987, 1988) and interactions of a disk galaxy with a cluster gravitational field (Byrd \\& Valtonen \\markcite{byrdvalt} 1990) show that both types of interaction can stimulate the formation of bars in galaxies that would otherwise be stable against the development of a bar, and enhance the formation of bars in galaxies which are already unstable to bar formation (Gerin, Combes \\& Athanassoula 1990 \\markcite{gerin}). In his simulations, Noguchi\\markcite{noguchi} (1988) finds that the lifetime of the induced bars is in the range $5 \\times 10^{8}$--$1.5 \\times 10^{9}yr$.  Thus, the frequency of barred galaxies in a given cluster gives information about the interaction history of those galaxies in the recent past (i.e. over the last 3--6 disk revolutions or so). The purpose of this project is to examine the distribution of barred galaxies in the Virgo cluster, as an aid in discriminating between the relative importance of different environmental effects in cluster galaxies. ",
        "conclusions": "A study of the velocity distribution of disk galaxies in the Virgo cluster region has been carried out, with the result that velocity distribution of barred spirals is skewed to lower recession velocities when compared to weakly or unbarred spirals. Since galaxies with low velocities in the Virgo region are preferentially associated with the main condensation of the Virgo cluster, this suggests the difference in the velocity distribution for spirals of different bar strength is due to an increase in the fraction of barred spirals associated with the main cluster when compared to less dense structures in the region.  Isolation of galaxies which belong to Virgo's main condensation indicate that the radial distribution of SB galaxies is more centrally peaked than that of the SA and SAB galaxies. Furthermore, there is no indication that barred spirals have more nearby companions than do unbarred spirals. Although the lack of nearby companions by itself does not exclude the possibility that galaxy-galaxy interactions are the dominant bar triggering mechanism, the fact that the ratio of crossing to collision times is much less than unity for reasonable perturbation strengths suggests that galaxy collisions are not efficient enough to be responsible for the increased bar fraction. The inference from these results is that the central concentration of barred galaxies is due to triggering of bars by the cluster mass distribution, as suggested by the simulations of Byrd \\& Valtonen (1990)\\markcite{byrdvalt}.  Unlike the spiral galaxies, the S0 and SB0 galaxies in Virgo have velocity and spatial distributions that are indistinguishable from one another. It is speculated that the difference between the spiral and lenticular galaxies may be due to the cluster tidal forces being insufficient to over come the stabilizing effect of the S0's more massive central bulges. Taken together, the enhanced fraction of barred galaxies in the center of the Virgo and Coma clusters is compelling evidence for the importance of tidal interactions in the evolution of cluster galaxies."
    },
    "9603/astro-ph9603018_arXiv.txt": {
        "abstract": "We demonstrate the existence of an enhanced rate of angular momentum relaxation in nearly Keplerian star clusters, such as those found in the centers of galactic nuclei containing massive black holes.  The enhanced relaxation arises because the radial and azimuthal orbital frequencies in a Keplerian potential are equal, and hence may be termed {\\em resonant relaxation\\/}.  We explore the dynamics of resonant relaxation using both numerical simulations and order-of-magnitude analytic calculations. We find that the resonant angular momentum relaxation time is shorter than the non-resonant relaxation time by of order $M_\\star/M$, where $M_\\star$ is the mass in stars and $M$ is the mass of the central object. Resonance does not enhance the energy relaxation rate. We examine the effect of resonant relaxation on the rate of tidal disruption of stars by the central mass; we find that the flux of stars into the loss cone is enhanced when the loss cone is empty, but that the disruption rate averaged over the entire cluster is not strongly affected.  We show that relativistic precession can disable resonant relaxation near the main-sequence loss cone for black hole masses comparable to those in galactic nuclei. Resonant dynamical friction leads to growth or decay of the eccentricity of the orbit of a massive body, depending on whether the distribution function of the stars is predominantly radial or tangential. The accelerated relaxation implies that there are regions in nuclear star clusters that are relaxed in angular momentum but not in energy; unfortunately, these regions are not well-resolved in nearby galaxies by the Hubble Space Telescope. ",
        "introduction": "\\label{sec_analyt}  \\subsection{Introduction}  The force field $\\mbf{F}(\\mbf{r},t)$ in an equilibrium $N$-body stellar system can be divided into a mean force $\\Fbar({\\mbf r})\\equiv\\langle \\mbf{F}(\\mbf{r},t)\\rangle$ and a fluctuating force ${\\mbf f}(\\mbf{r},t)\\equiv \\mbf{F}(\\mbf{r},t)-\\Fbar(\\mbf{r})$, where $\\langle\\cdot\\rangle$ denotes time average. If $N\\gg1$ the fluctuating force is a Gaussian random field, and hence is completely described by the correlation function $C_{ij}(\\mbf{r}_1,\\mbf{r}_2,\\tau)\\equiv \\langle f_i(\\mbf{r}_1,t),f_j(\\mbf{r}_2,t+\\tau)\\rangle$.  This fluctuating force induces diffusion or relaxation of the deterministic orbits that stars would follow if only the mean force field $\\Fbar$ were present. The relaxation time, $\\trel$, is crudely defined so that the diffusion of integrals of motion such as energy $E$ and angular momentum $L$ (per unit mass) is given by $\\delE\\sim E\\,(t/\\trel)^{1/2}$ and $\\delL\\sim L\\,(t/\\trel)^{1/2}$.  The usual estimate of the relaxation time (Jeans 1913, 1916; Chandrasekhar 1942; Binney \\& Tremaine 1987) is based on an infinite homogeneous stellar system in which stars travel on straight-line orbits. This assumption is plausible because each octave in spatial scale between the system size $R$ and the much smaller scale $\\rmin\\sim R/N$ (the scale on which close encounters produce $\\sim 90\\arcdeg$ deflections) contributes equally to the relaxation rate, and for most of these octaves the approximations of homogeneity and straight-line orbits are legitimate (summing the contributions from different scales gives rise to the well-known Coulomb logarithm $\\ln\\Lambda\\simeq \\ln(R/r_{\\rm min})\\simeq \\ln(N)$ which appears in formulae for the relaxation rate). One consequence of the assumption of straight-line orbits is that the correlation function $C_{ij}(\\mbf{r}_1,\\mbf{r}_2,\\tau)$ decays rapidly to zero when $\\tau$ exceeds $|\\mbf{r}_1-\\mbf{r}_2|/V$, where $V$ is a typical velocity (for infinite homogeneous systems $C_{ij}\\to 1/\\tau$ as $\\tau\\to\\infty$ [e.g. Cohen 1975]).  Our focus here is on stellar systems in which the correlation function remains non-zero for much larger times, a condition that occurs if most of the stars are near resonance.  If motion in the mean force field $\\Fbar(\\mbf{r})$ is regular, then stellar orbits are quasi\\-periodic with three characteristic frequencies $\\Omega_i$. The orbits are resonant if there are linear combinations of the form $\\sum_{i=1}^3k_i\\Omega_i=0$ where the $k_i$ are small integers. The simplest important examples are (i) spherical potentials, in which one frequency is zero because all orbits remain in a fixed plane; (ii) Kepler potentials, in which one additional frequency is zero because the apsis does not precess; and (iii) the harmonic oscillator potential, in which the radial frequency is twice the azimuthal frequency, so that the orbit shape is a centered ellipse.  The possibility of enhanced relaxation in potentials that support many near-resonant orbits was discussed by J. Ostriker two decades ago, in lectures for a graduate course in stellar dynamics attended by one of us (Ostriker 1973).  \\subsection{Non-resonant Relaxation} \\label{sec_nrrel}  We begin by examining relaxation in a near-Kepler potential.  Consider a spherical volume of radius $R$ centered on a point mass $M$ and containing $N\\gg 1$ identical stars of mass $m$, where $M_\\star\\equiv Nm\\ll M$. We assume that the stellar orbits have random orientations and moderate eccentricities, and that the density of stars is approximately uniform within $R$. The typical stellar velocity is $V\\sim (GM/R)^{1/2}$ and the characteristic orbital period is $\\torb\\sim R/V$.  Since $M_\\star\\ll M$, each orbit is approximately a Kepler ellipse, which precesses slowly on a timescale $\\tprec$. If the precession is dominated by the mean field from the other stars (rather than, say, relativistic effects or an external tidal field), then \\be \\tprec\\sim {M\\over M_\\star}\\torb, \\label{eq:raucha} \\ee which is much longer than $\\torb$.  The usual (non-resonant) relaxation rate can be estimated by the following argument. Consider a volume of radius $r<R$, which typically contains $n\\sim N\\,(r/R)^3$ stars. The instantaneous number of stars in this volume fluctuates by an amount $n^{1/2}$; thus the fluctuating force on the length scale $r$ is $f\\sim Gmn^{1/2}/r^2$. The force fluctuates on a timescale $t_f\\sim r/V$, and the typical impulse that a star receives during this timescale is $\\delta v\\sim ft_f\\sim Gmn^{1/2}/(rV)$.  Since impulses in successive intervals of length $t_f$ are uncorrelated, the total impulse after time $t$ is described by a random walk, \\be (\\Delta v)^2\\sim(\\delta v)^2{t\\over t_f}\\sim{G^2m^2n\\over r^3V}t\\sim {G^2m^2N\\over R^3V}t\\sim V^2{m^2N\\over M^2}{t\\over\\torb}. \\ee Note that $\\Delta v$ is independent of the scale $r$, which confirms that each octave in scale contributes equally to the relaxation rate; thus the total diffusion rate must be multiplied by a factor $\\ln\\Lambda$ which represents the number of octaves that contribute to the relaxation.  We may define the non-resonant relaxation time by $(\\Delta v/V)^2=(t/\\trel^{\\rm nr})$, so that \\be \\trel^{\\rm nr}\\sim {M^2\\over m^2 N\\ln\\Lambda}\\,\\torb. \\label{eq:trnr} \\ee  The fluctuating force changes both energy and angular momentum, at rates given by \\begin{eqnarray} (\\delE/E)_{\\rm nr}& \\sim &(t/\\trel^{\\rm nr})^{1/2}\\sim\\alpha {mN^{1/2}\\over M}(t/\\torb)^{1/2},\\nonumber \\\\ (\\delL/\\Lmax)_{\\rm nr}& \\sim &(t/\\trel^{\\rm nr})^{1/2}\\sim\\eta_{\\rm s} {mN^{1/2}\\over M} (t/\\torb)^{1/2}, \\label{eq:alpred} \\end{eqnarray} where $E\\sim V^2\\sim GM/R$, $\\Lmax^2\\sim V^2R^2\\sim GMR$, and $\\alpha$ and $\\eta_{\\rm s}$ are dimensionless constants that equal the square root of the Coulomb logarithm to within a factor of order unity.  \\subsection{Resonant Relaxation in Near-Kepler Potentials} \\label{sec_resrel}  To estimate the resonant relaxation rate, we imagine averaging the stellar density over an intermediate timescale that is $\\gg\\torb$ but $\\ll\\tprec$.  On this timescale each star can be represented by a fixed wire whose mass is the stellar mass, whose shape is a Kepler ellipse, and whose linear density is inversely proportional to the local speed in the elliptical orbit.  The gravitational potential from these wires is stationary and hence does not lead to energy relaxation; thus \\be (\\delE/E)_{\\rm res}=0. \\label{eq:erell} \\ee However, the wires exert mutual torques which induce angular momentum relaxation. The typical specific torque on a wire is $T\\sim N^{1/2}Gm/R$, and fluctuates on a timescale $\\sim\\tprec$ as the wires precess in different directions. Thus the characteristic change in specific angular momentum $\\delL\\sim Tt$ over a timescale $t<\\tprec$ is given by \\be (\\delL/\\Lmax)_{\\rm res}\\sim \\beta_{\\rm s} {mN^{1/2}\\over M}(t/\\torb),\\qquad t\\ll\\tprec, \\label{eq:trlrq} \\ee where $\\beta_{\\rm s}$ is a dimensionless constant of order unity.  Over timescales $t\\gg\\tprec$, the change in angular momentum is described by a random walk with increments $\\delL/\\Lmax\\sim \\beta_{\\rm s} (mN^{1/2}/ M)(\\tprec/\\torb)$ over a characteristic time $\\tprec$; thus, \\be (\\delL/\\Lmax)_{\\rm res}\\sim\\beta_{\\rm s}{mN^{1/2}\\over M}\\left(\\tprec\\, t\\over \\torb^2\\right)^{1/2},\\qquad t\\gg\\tprec; \\label{eq:genres} \\ee if the precession is determined by the mean field of the other stars then equation~(\\ref{eq:raucha}) implies that \\be (\\delL/\\Lmax)_{\\rm res}\\sim \\beta_{\\rm s}\\left(m\\over M\\right)^{1/2}(t/\\torb)^{1/2},\\qquad t\\gg\\tprec. \\label{eq:bigrel} \\ee In this case the angular momentum relaxation time, defined by $(\\delL/\\Lmax)_{\\rm res}\\allowbreak\\sim(t/\\trel^{\\rm res})^{1/2}$, is given by \\be \\trel^{\\rm res}\\sim {M\\over m}\\,\\torb, \\label{eq:trelk} \\ee which, remarkably, is independent of the number of stars $N$. The resonant relaxation time is shorter than the non-resonant relaxation time (eq.~[\\ref{eq:trnr}]) by a factor $(mN/M)\\ln\\Lambda$. If there is a range of stellar masses, the factor $m$ in equation~(\\ref{eq:trelk}) should be replaced by $\\int m^2\\dd N(m)/\\int m\\dd N(m)$.  There is no analog to the Coulomb logarithm in resonant relaxation, since the relaxation is dominated by large-scale fluctuations. The absence of the Coulomb logarithm implies that the resonant diffusion rate depends on the overall structure of the stellar system and cannot be computed using the assumption of local homogeneity, as is done for the non-resonant relaxation rate.  On timescales longer than the resonant relaxation time but shorter than the non-resonant relaxation time, a stellar system in a Kepler potential should be in the maximum-entropy state consistent with its original total angular momentum $\\mbf{L}_{\\rm tot}$ and energy distribution $N(E)$ (the number of stars with energy $<E$, which is invariant on this timescale since there is no resonant energy relaxation). This state is described by the phase-space distribution function \\be f(\\mbf{r},\\mbf{v})=w(E)\\exp(-\\mbf{b}\\cdot\\mbf{L}), \\label{eq:equil} \\ee where $E={1\\over2}v^2-GM/r$, $\\mbf{L}=\\mbf{r}\\times\\mbf{v}$, and $\\mbf{b}$ and $w(E)$ are determined implicitly by the constraints \\begin{eqnarray} \\mbf{L}_{\\rm tot}&=&\\int f(\\mbf{r},\\mbf{v})\\mbf{L}\\, \\dd\\mbf{r}\\dd\\mbf{v},\\nonumber \\\\ \\noalign{\\vskip -0.5em} & & \\\\ \\noalign{\\vskip -0.5em} {\\dd N(E)\\over \\dd E}&=&\\int f(\\mbf{r}_1,{\\mbf v}_1)\\delta(E-E_1)\\dd \\mbf{r}_1\\dd \\mbf{v}_1.\\nonumber \\end{eqnarray} If $\\mbf{L}_{\\rm tot}=0$ then $\\mbf{b}=0$ and the distribution function is a function of energy alone (i.e., the cluster is isotropic). If there is a range of masses then (\\ref{eq:equil}) is replaced by \\be f(\\mbf{r},\\mbf{v},m)=w(E,m)\\exp(-m\\,\\mbf{b}\\cdot\\mbf{L}). \\ee  \\subsection{Resonant Relaxation in Near-spherical Potentials} \\label{sec_spherpot}  A more limited form of resonant relaxation is present in {\\em any} spherical potential. Consider a generic spherical potential and average the stellar density over a timescale $\\gg\\torb$. On this timescale each star is smeared into an axisymmetric annulus whose inner and outer radii are the pericenter and apo\\-center distances.  The gravitational potential from these annuli is stationary and hence does not lead to energy relaxation; thus, as in equation~(\\ref{eq:erell}), \\be (\\delE/E)_{\\rm res}=0. \\ee The annuli exert mutual torques which induce angular momentum relaxation; however, in contrast to the Kepler case the torques are perpendicular to the orbit normals (because the averaged orbits are axisymmetric annuli rather than eccentric wires). Thus the vector torque $\\mbf{T}_{ij}$ between two annuli with vector angular momenta $\\mbf{L}_i$ and $\\mbf{L}_j$ satisfies ${\\mbf T}_{ij}\\cdot\\mbf{L}_i=\\mbf{T}_{ij}\\cdot\\mbf{L}_j=0$; in other words the torques change the directions of the angular momentum vectors but not their magnitudes. Thus, in contrast to equation~(\\ref{eq:trlrq}), there is no resonant relaxation of the scalar angular momentum, \\be (\\delL/\\Lmax)_{\\rm res}=(\\Delta|\\mbf{L}|/\\Lmax)_{\\rm res}=0, \\ee but there {\\em is} resonant relaxation of the {\\em vector} angular momentum. The resonant relaxation rate may be estimated by analogy with equation~(\\ref{eq:trlrq}). The typical specific torque on an annulus is $T\\sim N^{1/2}Gm/R$, and fluctuates on a timescale $\\sim\\tprecL$. Here $\\tprecL$ is the precession time for the angular momentum vector ({\\em not} the apsis, as in \\S\\S~\\ref{sec_nrrel}\\ and \\ref{sec_resrel}), and is determined by the stochastic component of the potential, $\\tprecL\\sim N^{1/2}\\torb/\\mu$.  The characteristic change in vector angular momentum over a timescale $t<\\tprecL$ is given by $|\\Delta\\mbf{L}|\\sim Tt$.  Since we are considering general near-spherical potentials, we no longer require that the potential is dominated by a black hole, $M_\\star\\ll M$; thus the maximum angular momentum is given by $\\Lmax^2\\sim V^2R^2\\sim G(M+M_\\star)R$ and we have \\be (|\\Delta\\mbf{L}|/\\Lmax)_{\\rm res} \\sim \\mu{\\beta_{\\rm v}\\over N^{1/2}} (t/\\torb),\\qquad t\\ll\\tprecL, \\label{eq:sphtr} \\ee where $\\beta_{\\rm v}$ is a dimensionless constant of order unity, and \\be \\mu={M_\\star\\over M_\\star+M} \\label{eq:mudef} \\ee ($\\sim 1$ if the potential is dominated by the stars themselves).  Over timescales $t\\gg\\tprecL$, the change in angular momentum is described by a random walk. The increments $\\Delta\\mbf{L}$ of the random walk are determined by evaluating (\\ref{eq:sphtr}) at $t\\sim\\tprecL$, which yields $|\\Delta\\mbf{L}|/\\Lmax\\sim 1$. In other words the angular momentum vectors drift over the whole velocity sphere on a timescale of order $\\tprecL$, implying $\\trel^{\\rm res}\\sim \\tprecL\\sim N^{1/2}\\torb/\\mu$, which is shorter than the non-resonant relaxation time $\\trel^{\\rm nr}$ (eq.~[\\ref{eq:trnr}]) by a factor $(\\mu/N^{1/2})\\ln\\Lambda$.  A closely related form of resonant relaxation is present in general axisymmetric potentials that are nearly spherical. Once again there is no resonant relaxation of the scalar angular momentum, but there is resonant relaxation of the vector angular momentum. In this case only the $z$-component of the vector angular momentum is conserved by motion in the mean potential, so we focus on this component. By analogy with equation (\\ref{eq:sphtr}) we have \\be (|\\Delta L_z|/\\Lmax)_{\\rm res} \\sim \\mu{\\beta_{\\rm z}\\over N^{1/2}} (t/\\torb),\\qquad t\\ll\\tprecL, \\label{eq:sphtrz} \\ee where in this case $\\tprecL$, the precession time for the angular momentum vector, is determined by the non-spherical component of the mean potential. Over timescales $\\gg \\tprecL$, the change in $L_z$ is described by a random walk with increments $|\\Delta L_z|/\\Lmax\\sim (\\mu\\beta_{\\rm z}/N^{1/2})(\\tprecL/\\torb)$ over a characteristic time $\\tprecL$; thus, \\be (|\\Delta L_z|/\\Lmax)_{\\rm res} \\sim \\mu{\\beta_{\\rm z}\\over N^{1/2}} \\left(\\tprecL\\,t\\over\\torb^2\\right)^{1/2},\\qquad t\\gg\\tprecL. \\label{eq:vecrelz} \\ee The relaxation time, defined by $(|\\Delta L_z|/\\Lmax)_{\\rm res}\\sim (t/\\trel^{\\rm res})^{1/2}$, is then given by \\be \\trel^{\\rm res}\\sim {N\\over \\mu^2}{\\torb^2\\over\\tprecL}, \\ee which is shorter than $\\trel^{\\rm nr}$ by a factor $(\\torb/\\tprecL)\\ln\\Lambda$.  The vector angular momentum also diffuses from non-resonant relaxation, at a rate (cf. eq.~[\\ref{eq:alpred}]) \\be (|\\Delta\\mbf{L}|/\\Lmax)_{\\rm nr} \\sim \\mu{\\eta_{\\rm v}\\over N^{1/2}} (t/\\torb)^{1/2}, \\label{eq:nrvec} \\ee where $\\eta_{\\rm v}$ is a dimensionless constant that equals the square root of the Cou\\-lomb logarithm to within a factor of order unity.  \\subsection{Resonant Relaxation in Regular Potentials}  Motion in any time-independent, regular potential can be described by action-angle variables $(J_i,\\phi_i)$ and a Hamiltonian $\\overline H({\\mbf J)}$. For spherical potentials the actions $(J_1,J_2,J_3)$ can be chosen to be, respectively, the radial action, the total angular momentum, and the $z$-component of the angular momentum (e.g., Tremaine \\& Weinberg 1984). The motion is quasiperiodic with fundamental frequencies $\\dot\\phi_i=\\Omega_i({\\mbf J})=\\partial \\overline H/\\partial J_i$. In a spherical potential, $\\Omega_1$ is the radial frequency, $\\Omega_2$ is the azimuthal frequency, and $\\Omega_3=0$; in a Kepler potential $\\Omega_1=\\Omega_2$. Resonant relaxation can be important if the resonance condition $\\sum_{i=1}^3 k_i\\Omega_i= 0$ is approximately satisfied for most stars, where the $k_i$'s are small integers with no common factor (e.g., $k_1=k_2=0$ for a spherical potential; $k_1=-k_2$ for a Kepler potential).  To isolate the effects of resonant relaxation we perform a canonical transformation to new action-angle variables $(K_i,\\psi_i)$ defined by the generating function \\be S(K_i,\\phi_i)=K_1\\sum_{i=1}^3 k_i\\phi_i +K_2\\phi_2+K_3\\phi_3. \\ee Then \\be J_1=k_1K_1,\\quad J_2=k_2K_1+K_2,\\quad J_3=k_3K_1+K_3, \\ee and \\be \\psi_1=\\sum_{i=1}^3 k_i\\phi_i,\\quad \\psi_2=\\phi_2,\\quad \\psi_3=\\phi_3. \\ee The resonance condition implies that $\\psi_1$ is slowly varying; the effects of resonant relaxation are therefore described by a fluctuating Hamiltonian of the form $h(\\mbf{K},\\psi_1)$. The resonant Hamiltonian does not depend on the non-resonant angles $\\psi_2$ and $\\psi_3$ or on the time, so Hamilton's equations imply that the conjugate momenta $K_2$ and $K_3$ and the total energy $E=\\overline H+h$ are unaffected by resonant relaxation; this in turn implies that the changes in the energy and actions caused by resonant relaxation satisfy the constraints \\be \\Delta E=0,\\qquad \\Delta \\mbf{J}=C(t)\\mbf{k}, \\ee where $C(t)$ is scalar.  Resonant relaxation leads to diffusion of $C(t)$ which is described by analogs of equations (\\ref{eq:sphtr}) and (\\ref{eq:vecrelz}), \\begin{eqnarray} (\\Delta C/\\Lmax)_{\\rm res} &\\sim &\\mu {\\gamma\\over N^{1/2}} (t/\\torb), \\qquad t\\ll\\tpreck;\\nonumber \\\\ \\noalign{\\vskip -0.5em} & & \\\\ \\noalign{\\vskip -0.5em} (\\Delta C/\\Lmax)_{\\rm res} &\\sim &\\mu {\\gamma\\over N^{1/2}} \\left(\\tpreck\\, t\\over\\torb^2\\right)^{1/2}, \\qquad t\\gg\\tpreck,\\nonumber \\end{eqnarray} where $\\gamma$ is a dimensionless constant of order unity, $\\mu$ is defined in equation~(\\ref{eq:mudef}), $\\tpreck\\sim (\\sum k_i\\Omega_i)^{-1}$ is the time required for the slow angle $\\psi_1$ to change by one radian, and we have assumed that $|{\\mbf k}|$ is small.  \\subsection{Resonant Friction}  Chandrasekhar (1943) showed that the stochastic changes in velocity associated with relaxation are accompanied by a systematic drag which he named ``dynamical friction.'' For massive objects (MOs) orbiting in a stellar system, orbital evolution from dynamical friction is much faster than evolution from stochastic relaxation. Near-resonances can enhance dynamical friction just as they enhance stochastic relaxation, leading to an effect we may call ``resonant friction.''  To analyze resonant friction we use the formalism developed by Lynden-Bell \\& Kalnajs (1972); the $z$-component of the specific torque on the MO from dynamical friction is (Tremaine \\& Weinberg 1984, eq. 65) \\be T_z=4\\pi^4m_0\\sum_{\\mbf{\\scriptstyle k},k_3\\ge0}k_3 \\int \\dd \\mbf{J} \\sum_i k_i {\\partial f\\over\\partial J_i}\\left|\\Psi_{\\mbf{\\scriptstyle k}}\\right|^2 \\delta({\\mbf k}\\cdot{\\mbf{\\Omega}}-\\omega_{\\mbf{\\scriptstyle k}}). \\label{eq:wein} \\ee Here $m_0\\gg m$ is the mass of the MO, $f(\\mbf{J})$ is the phase-space density of bound stars, and the potential of the MO has been written in the form \\be U(\\mbf{r},t)=m_0\\,\\hbox{Re}\\left\\{\\sum_{\\mbf{\\scriptstyle k},k_3\\ge0} \\Psi_{\\mbf{\\scriptstyle k}}({\\mbf J})\\exp[i(\\mbf{k}\\cdot{\\mbf{\\phi}}-\\omega_{\\mbf{\\scriptstyle k}}t)]\\right\\}. \\ee Equation (\\ref{eq:wein}) contains both resonant and non-resonant contributions to dynamical friction.  Now let us assume that the stellar system is spherically symmetric and that the distribution function depends only on energy and angular momentum, $f=f(E,L)$, consistent with Jeans' theorem. Without loss of generality we may assume that $z=0$ is the orbital plane of the MO, so that the total specific torque on the MO is $T=T_z$. Using $\\partial E/\\partial J_i=\\partial\\overline H/\\partial J_i=\\Omega_i$ and $L=J_2$ we have \\be T=4\\pi^4m_0\\sum_{\\mbf{\\scriptstyle k},k_3\\ge0}k_3\\int \\dd \\mbf{J}\\left( \\omega_{\\mbf{\\scriptstyle k}}{\\partial f\\over\\partial E}+k_2{\\partial f\\over\\partial L}\\right)\\left| \\Psi_{\\mbf{\\scriptstyle k}}\\right|^2\\delta(\\mbf{k}\\cdot{\\mbf{\\Omega}}- \\omega_{\\mbf{\\scriptstyle k}}). \\label{eq:weine} \\ee Now suppose that there is a near-resonance for the triplet ${\\mbf k}_r$. Formally, we set $\\mbf{k}_r\\cdot{\\mbf{\\Omega}}=\\epsilon\\,{\\mbf k}_r\\cdot\\widetilde{\\mbf{\\Omega}}$ and $\\omega_{\\mbf{\\scriptstyle k}_r}=\\epsilon\\,\\widetilde\\omega_{\\mbf{\\scriptstyle k}_r}$ and let $\\epsilon\\to 0$. Keeping only the largest terms in (\\ref{eq:weine}) gives the contribution of the triplet $\\mbf{k}_r$ to the specific torque from resonant friction, \\be T={4\\pi^4m_0\\over \\epsilon}k_{3,r}k_{2,r}\\int \\dd \\mbf{J} {\\partial f\\over\\partial L}\\left|\\Psi_{\\mbf{\\scriptstyle k}_r}\\right|^2\\delta({\\mbf k}_r\\cdot\\widetilde{\\mbf{\\Omega}}-\\widetilde\\omega_{\\mbf{\\scriptstyle k}_r}), \\label{eq:weinee} \\ee which diverges as $\\epsilon\\to0$.  Evaluating expressions such as (\\ref{eq:weinee}) is arduous, mostly because of the complexity of the action-angle expansion $\\Psi_{\\mbf{\\scriptstyle k}}$ of the potential of a point mass (Weinberg 1986; see also Hernquist \\& Weinberg 1989). However, it is simple to derive some qualitative effects of resonant friction:  \\begin{itemize}  \\item Equation (\\ref{eq:weinee}) shows that there is no resonant friction when the distribution function is isotropic ($\\partial f/\\partial L=0$). The term $\\propto \\partial f/\\partial E$ in equation~(\\ref{eq:weine}) contributes only non-resonant friction.  \\item The resonant torque is zero if the stars are on circular orbits (in this case $\\Psi_{\\mbf{\\scriptstyle k}}=0$ unless $k_1=0$, and if $k_1=0$ then resonance requires $k_2=0$ since $\\Omega_2\\not=0$ and $\\Omega_3=0$); similarly, the resonant torque is zero if the MO is on a circular orbit.  \\item The orbits that contribute most to the friction are those with near-zero inclinations, since these remain closest to the zero-inclination orbit of the MO and precess in the same direction. Zero-inclination orbits have $\\Psi_{\\mbf{ \\scriptstyle k}}=0$ unless $k_2=k_3$ (Weinberg \\& Tremaine 1984). Together with equation (\\ref{eq:weinee}), this suggests that the dominant contribution to the resonant torque comes from terms with $k_{3,r}k_{2,r}>0$, which in turn implies that the sign of the torque on the MO is the sign of $\\partial f/\\partial L$.  In the common case where orbits are predominantly radial, $\\partial f/\\partial L <0$, resonant friction removes angular momentum from the MO orbit at constant energy, thereby increasing its eccentricity. This effect can dominate the eccentricity evolution of black-hole binaries in galactic nuclei and may promote the merger of binary black holes, since emission of gravitational radiation is more efficient for an eccentric binary (Begelman et al. 1980, Quinlan 1996).  \\item The rate of growth or decay of angular momentum of the MO through resonant friction is approximately given by \\be {1\\over\\Lmax}\\left({\\dd L\\over \\dd t}\\right)_{\\rm res}\\sim {T\\over \\Lmax} \\sim\\pm {m_0 M_\\star\\over M^2}{\\tpreck\\over\\torb^2}, \\ee where we have assumed that the orbit of the MO has moderate eccentricity and that $|\\partial f/\\partial L|\\sim f/\\Lmax$. In a near-Kepler potential, where ${\\mbf k}_r=(1,-1,0)$ and $\\tpreck$ is given by equation (\\ref{eq:raucha}), we have \\be {1\\over\\Lmax}\\left(\\dd L\\over \\dd t\\right)_{\\rm res}\\sim\\pm{m_0\\over M}{1\\over\\torb}. \\label{eq:fricc} \\ee In this case the frictional torque is---remarkably---independent of the number of stars, although equation~(\\ref{eq:fricc}) holds only when $m_0\\lesssim M_\\star$. When $m_0\\gtrsim M_\\star$ the precession time is $\\tprec\\sim(M/m_0)\\torb$, so that \\be {1\\over\\Lmax}\\left(\\dd L\\over \\dd t\\right)_{\\rm res}\\sim \\pm{M_\\star\\over M}{1\\over\\torb}. \\label{eq:friccc} \\ee The two expressions (\\ref{eq:fricc}) and (\\ref{eq:friccc}) can be combined, \\be {1\\over\\Lmax}\\left(\\dd L\\over \\dd t\\right)_{\\rm res}\\sim \\pm{\\hbox{min}\\,(M_\\star, m_0)\\over M}{1\\over\\torb}. \\ee For comparison, non-resonant dynamical friction removes energy and angular momentum at the slower rate \\be {1\\over E}\\left(\\dd E\\over \\dd t\\right)_{\\rm nr}\\sim {1\\over\\Lmax}\\left(\\dd L\\over \\dd t\\right)_{\\rm nr}\\sim -{m_0M_\\star\\over M^2}{1\\over \\torb}. \\ee  \\end{itemize}  Another manifestation of resonant friction occurs when the potential is nearly spherical but the mean rotation of the stars is non-zero (for example a stellar disk) and the orbit of the MO is inclined. In this case the distribution function $f=f(E,L,L_z)$ depends not just on $E$ and $L$ but also on the $z$-component of angular momentum $L_z=J_3$. Because the potential is nearly spherical, $\\Omega_3\\simeq0$ (i.e. inclined orbits precess slowly) so terms with $k_1=k_2=0$ are near-resonant. The $z$-component of the torque on an MO from the triplets $\\mbf{k}=(0,0,k_3)$ is \\be T_z=4\\pi^4m_0\\sum_{k_3\\ge0}k_3^2\\int \\dd \\mbf{J} {\\partial f\\over\\partial L_z}\\left|\\Psi_{0,0,k_3}\\right|^2\\delta(k_3\\Omega_3 -\\omega_{0,0,k_3}). \\ee If the stellar system rotates in a prograde direction, then $\\partial f/\\partial L_z$ is generally positive so $T_z$ is positive; thus the resonant friction erodes the inclination of the MO until it settles into the equatorial plane of the stellar system. The resonant torque is zero if the stars have precisely zero inclination (for zero-inclination stars $\\Psi_{0,0,k_3}=0$ unless $k_3=0$); the rate of change of the orbital inclination $I_0$ of the MO is given approximately by \\be {1\\over I_0}\\left(\\dd I_0\\over \\dd t\\right)_{\\rm res}\\sim \\pm {M_\\star m_0\\over M^2}{\\tprec\\over\\torb^2}\\langle I^2\\rangle, \\ee where $\\langle I^2\\rangle$ is the mean-square inclination in the disk and we have assumed $|\\partial f/\\partial L_z|\\sim f/\\Lmax$. ",
        "conclusions": "Resonant relaxation is the dominant source of angular momentum relaxation for stellar systems in near-Keplerian potentials, and thus plays an important role in determining the structure of stellar cusps around black holes in galactic nuclei or globular clusters. Resonant relaxation enhances the angular momentum relaxation rate by roughly the ratio of the mass of the black hole to the mass in stars but does not affect the energy relaxation rate (more precisely, the combination of actions $J_1-J_2$ is conserved, where $J_1$ is the radial action and $J_2$ is the angular momentum).  Resonant relaxation is also present in the harmonic potentials that characterize constant-density cores, and may enhance the rate of angular momentum relaxation in the cores of globular clusters. In constant-density cores, resonant relaxation preserves the combination of actions $J_1-2J_2$. This form of resonance is not likely to be important for elliptical galaxies, which do not generally have constant-density cores (e.g., Gebhardt et al. 1996).  One might speculate that generic potentials contain (high-order) resonances that are strong enough to support resonant relaxation. In this case the angular momentum relaxation time would be much shorter than the energy relaxation time throughout most of a galaxy. We suspect that this speculation is not correct, since our $N$-body simulations in the isochrone potential yield very similar relaxation rates for the energy and angular momentum. Nevertheless, the presence of resonant relaxation is a reminder that our understanding of relaxation in stellar systems is crude, and has not been numerically verified under conditions ($N\\sim 10^{11}$) found in real galaxies.  Resonant friction leads to growth or decay of the eccentricity of massive objects orbiting in near-Kepler potentials, depending on whether the star orbits are predominantly radial or tangential. Resonant friction can strongly influence the orbital evolution of a binary black hole (at least if the mass ratio of the binary is sufficiently far from unity). In radially biased star clusters the eccentricity of the binary will grow, at a rate faster than the decay of the orbital energy, at least if the friction is dominated by cluster stars rather than unbound stars. The binary eccentricity will grow until the resonant friction is quenched by relativistic precession, at which point gravitational radiation may erode the energy of the binary faster than non-resonant dynamical friction. The details of this evolution are relevant to the merger rate of black holes, the gravitational-wave background, the prevalence of binary black holes in AGNs, and the viability of massive black holes as dark matter candidates (see Quinlan 1996 for references).  Resonant friction can also erode the inclination of a massive object in a rotating stellar system. This process may be relevant to a star cluster in which there is a massive accretion disk (Ostriker 1983; Syer et al.\\ 1991); resonant friction could accelerate the evolution of massive stars into low-inclination orbits embedded in the accretion disk.  The analytical treatment of resonant relaxation that we have offered in \\S~\\ref{sec_analyt} is only approximate. Accurate expressions for the resonant and non-resonant relaxation rates in a given star cluster could be derived by expanding the potential from a stellar orbit in action-angle variables. So far this procedure has only been carried out for the dynamical friction component (Weinberg 1986).  The relative simplicity of the diffusion coefficients that describe non-resonant relaxation (e.g., Binney \\& Tremaine 1987) is illusory in near-Keplerian and other near-resonant potentials---except to describe energy relaxation---since resonant relaxation is stronger, and depends more sensitively on the structure of the stellar system. For order-of-magnitude estimates we have used the formulae in \\S~\\ref{sec_analyt}, with the dimensionless coefficients given in Table~\\ref{aveparas}.  The estimates of tidal disruption rates in \\S~\\ref{sec_newtloss} suffer from the absence of a consistent treatment of relativistic precession, which detunes the Kepler resonance near the loss cone in galactic nuclei. However, resonant relaxation is unlikely to increase substantially the tidal disruption rate, which is mostly determined by the location of the critical radius $r_{\\rm crit}$, set by the angular momentum changes in a single orbital period. For similar reasons, resonant relaxation will not greatly affect the disruption rate resulting when the wandering (``Brownian motion'') of the black hole from the center of the nucleus (Quinlan 1995) is taken into account.  Thus tidal disruption appears to remain incapable of powering typical AGNs; however, since the relativistic detuning is largely ineffective at hole masses $\\lesssim 10^7\\;\\Msun$ (eq.~\\ref{eq:grlimit}), resonant effects may offer modest improvements in the feasibility of disruption-dominated mass loss in Seyferts and other nuclei containing low-level AGN activity, for which the energy requirements are less severe. Similarly, resonant relaxation may modestly enhance the rate of flares from tidally disrupted stars in nearby galaxies with central black holes (Rees 1988).  The effectiveness of relativistic precession in disabling resonant relaxation illustrates that general relativity can have dramatic physical consequences even where the motion is predominantly Newtonian; in particular, the shape of the density cusp inside $r_{\\rm crit}$ can be strongly dependent on relativistic precession. Thus resonant relaxation might one day be used to show that the massive dark objects observed in galactic nuclei are indeed black holes (or at least behave as such on a scale of $\\sim 10^2$ Schwarzschild radii)---a conclusion which today must be reached by indirect (albeit compelling) arguments.  The discussion in \\S 3.1 also illustrates that the Fokker-Planck equation used to describe non-resonant relaxation (e.g. Binney and Tremaine 1987) is not always adequate to describe resonant relaxation. The Fokker-Planck equation assumes that the fluctuating forces at different times and locations are uncorrelated, i.e., that the correlation function of \\S 1.1 has the form $C_{ij}({\\mbf r}_1,{\\mbf r}_2,\\tau)=K_{ij}({\\mbf r}_1)\\delta({\\mbf r}_2- {\\mbf r}_1)\\delta(\\tau)$. This is a reasonable approximation for non-resonant relaxation, which is dominated by close encounters (cf. \\S 1.1).  In contrast, the resonant forces are correlated over large spatial scales and over times $\\sim\\tprec$. The inadequacy of the Fokker-Planck approximation (or the master equation, or the approximation that relaxation is a Markov process), is particularly acute in the diffusion limit (eq. \\ref{eq:lossres}), when the size of the loss cone in angular-momentum space $\\Lmin$ is much greater than the change in angular momentum per orbit $\\delLo$ but much less than the change in angular momentum per precession time.  There is an appealing analogy between relaxation of stars in angular-momentum space and models of stellar structure. Non-resonant relaxation is a random walk in $L$-space, as for the motion of ions in the radiative zone of a star. Resonant relaxation implies a large-scale drift in $L$ superimposed upon a small-scale random walk, analogous to ionic motion in a convective stellar envelope. The quenching of resonant relaxation by relativistic precession can produce a random-walk dominated ``core'' in $L$-space surrounded by a drift-dominated envelope---conceptually similar to the radiative core/convective envelope structure found in solar-type stars. There are fundamental differences: $L$-space has no analog to gravity, but there {\\em is\\/} a net flux of stars towards small $L$ due to removal of stars by tidal disruption.  The structure of a relaxed star cluster surrounding a black hole has been examined by several authors (Peebles 1972; Bahcall \\& Wolf 1976, 1977; Lightman \\& Shapiro 1977; Cohn \\& Kulsrud 1978). These analyses do not take resonant relaxation into account and therefore some of their conclusions may be suspect: we expect that including resonant relaxation will not strongly affect the structure of the star cluster outside the critical radius $r_{\\rm crit}$ (\\S~\\ref{sec_newtloss}) or the total flux of stars into the loss cone, but may substantially reduce the density of stars inside $r_{\\rm crit}$. This classic problem should be re-investigated.  Resonant relaxation implies that there may be regions near the centers of elliptical galaxies (typically $\\lesssim 1\\pc$ in radial extent; cf. Table~\\ref{galtab}) that are relaxed in angular momentum but not energy. If non-rotating, such regions will have isotropic distribution functions; if rotating, the mean rotation speed will depend on the stellar mass. Unfortunately, these regions are not accessible at Hubble Space Telescope resolutions in most nearby galaxies. A more fundamental problem regarding the possible observational detection of resonant relaxation is that the isotropy it produces will be undetectable unless the initial distribution function is significantly anisotropic; there is currently no evidence for this in observed nuclear star clusters."
    },
    "9603/astro-ph9603105_arXiv.txt": {
        "abstract": "Studies of limits on active-sterile neutrino mixing derived from big bang nucleosynthesis considerations are extended to consider the dependance of these constraints on the primordial deuterium abundance. This study is motivated by recent measurements of D/H in quasar absorption systems, which at present yield discordant results. Limits on active-sterile mixing are somewhat relaxed for high D/H. For low D/H ($\\approx 2 \\times 10^{-5}$), no active-sterile neutrino mixing is allowed by currently popular upper limits on the primordial $^4$He abundance $Y$. For such low primordial D/H values, the observational inference of active-sterile neutrino mixing by upcoming solar neutrino experiments would imply that $Y$ has been systematically underestimated, unless there is new physics not included in standard BBN. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603039_arXiv.txt": {
        "abstract": "We develop an analytic model for the hierarchical correlation amplitudes $S_{j,g}(R)\\equiv \\bxi_{j,g}(R)/\\bxi_{2,g}^{j-1}(R)$ [where $j=3$,4,5, and $\\bxi _{j,g}(R)$ is the $j$th order connected moment of counts in spheres of radius $R$] of density peaks and dark matter halos in the quasi-linear regime. The statistical distributions of density peaks and dark matter halos within the initial density field  (assumed Gaussian) are determined by the peak formalism of Bardeen et al. (1986) and by an extension of the Press-Schechter formalism, respectively. Modifications of these distributions caused by gravitationally induced motions are treated using a spherical collapse model. We test our model against results for $S_{3,g}(R)$ and $S_{4,g}(R)$ from a variety of N-body simulations. The model works well for peaks even on scales where the second moment of mass ($\\bxi_2$) is significantly greater than unity. The model also works successfully for halos that are identified  earlier than the time when the moments are calculated. Because halos are spatially exclusive at the time of their identification, our model is only qualitatively correct for halos identified at the same time as the moments are calculated.  For currently popular initial density spectra, the values of $S_{j,g}$ at $R\\sim 10\\mpch$ are significantly smaller for both halos and peaks than those for the mass, unless the linear bias parameter $b$ [defined by $b^2=\\bxi_{2,g} (R)/\\bxi_2 (R)$ for large $R$] is comparable to or less than unity. The $S_{j,g}$ depend only weakly on $b$ for large $b$ but increase rapidly with decreasing $b$ at $b\\sim 1$. Thus if galaxies are associated with peaks in the initial density field, or with dark halos formed at high redshifts, a measurement of $S_{j,g}$ in the quasi-linear regime should determine whether galaxies are significantly biased relative to the mass. We use our model to interpret the observed high order correlation functions of galaxies and clusters. We find that if the values of $S_{j,g}$ for galaxies are as high as those given by the APM survey, then APM galaxies should not be significantly biased. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603111_arXiv.txt": {
        "abstract": "We present high-sensitivity multi-frequency radio continuum observations of the starburst galaxy NGC~2146. We have fitted these data with a three-dimensional diffusion model. The model can describe the radio emission from the inner disk of NGC~2146 well, indicating that diffusion is the dominant mode of propagation in this region. Our results are indicating that NGC~2146 has recently undergone a starburst, the star forming activity being located in a central bar. The spatial variation of the radio emission and of the spectral index yield tight constraints on the diffusion coefficient $D_0$ and the energy dependence of the diffusion. Away from the central bar of the galaxy the radio emission becomes filamentary and the diffusion model was found to be a poor fit to the data in these regions; we attribute this to different transport processes being important in the halo of the galaxy. ",
        "introduction": "Multifrequency radio observations of nearby, face-on galaxies present a unique way to study the Cosmic Ray (CR) propagation within their disks. These studies provide constraints on the main parameters of this propagation, i.e. the value and energy dependence of the diffusion coefficient and the form of the electron energy distribution immediately after acceleration.  Furthermore, because of the relatively long radiative life-time of CR electrons (which is in normal, non-starburst galaxies up to several $10^7$ years) such observations also provide information about the temporal evolution of the sources of the CR's on the same time scale; in turn, since the sites of acceleration are supernova remnants the data are probing the history of star formation on time scales up to of order $ \\approx 10^8$ years.  To date, the parameters describing CR propagation have only been derived for our Galaxy; radio observations are the only way of obtaining information for external galaxies. In this paper we analyse high-sensitivity multi-frequency radio data for the nearby galaxy NGC~2146.  Our data have sufficient frequency coverage that we can include in our model fits both a thermal and diffuse synchrotron component for the radio emission.  NGC~2146 is a very luminous radio source ($\\pnu=5.9\\,10^{22}$~W/Hz) at a distance of 21.8 Mpc (Benvenuti, Capacioli \\& D'Odorico 1975, adjusted to $H_0$=50 km s$^{-1}$ Mpc). Most of its radio luminosity is due to a compact starburst in its centre. The extent of the radio emission (1.5 arcmin $\\approx$ 9.6 kpc) is much less than the extent of the optical emission ($\\approx$ 6 arcmin). This suggests that active star formation (both current and in the recent past) is confined to this central region of NGC~2146. In this paper we are only discussing this inner region of the galaxy ($\\approx$ 1.5 arcmin $\\times$ 0.7 arcmin) that is characterized by the strong radio emission. Optical images of NGC~2146 show it to be a peculiar spiral galaxy (e.g. Benvenuti et al. 1975). Hutchings et al. (1990) have presented a detailed study of NGC~2146 in several wavebands and proposed that the galaxy is in the late stages of a merger which has given rise to an intense burst of star formation in the nucleus. ",
        "conclusions": "\\subsection{Model uncertainties}  As illustrated in Fig. 5-8 the model describes the data well and allows us to constrain important parameters such as the diffusion coefficient and its energy dependence and infer the existence of a central bar. The error in the determination of these parameters depends however not only on the quality of the model fit, but also on a number of parameters that were determined independently such as $B, \\urad$, the inclination angle and distance of the NGC~2146.  In order to fit the spectral data a model that produces a break in the synchrotron spectrum is required -- this is evident from the data since the spectrum at the centre of NGC~2146 steepens between 5 and 8.4 GHz before flattening towards 15 GHz which we attribute to thermal emission. The simplest way to achieve such a break is to assume some temporal dependence of the CR injection and hence of the star-formation rate. The time-scale of the changes implied by the data is rather short, less than $10^6$ years, however there is a strong dependence on $B$ and $\\urad$: $t\\propto B^{1/2}/(\\urad+\\ub)$. Whereas our estimate for $\\urad$ does not suffer from large uncertainties, $B$ is very uncertain and an improved estimate would require an independent determination of the magnetic field.  The main sources of error for the diffusion coefficient, $D_0$, are the uncertainty in $\\urad$, $B$, the distance to NGC~2146 and the assumed structure of the emitting region. In principle it could be possible to increase/decrease the size of the CR source along the minor axis ($R_{y}$) and try to fit the data by decreasing/increasing $D_0$. If $R_y$ is very small, even with a large value of $D_0$ the model cannot produce the observed shape of $\\prad$. On the other hand, a larger assumed source size would make it impossible to fit the spectral variations unless we assumed that either the injection index of the CR population varied within the disc of the galaxy, and/or that different regions of the galaxy had undergone very different star-formation histories.  Within the range of acceptable $R_{y}$ the estimated values for $D_0$ all lie within our quoted error bounds. The uncertainty in the inclination angle of NGC~2146 does not affect our estimates of $D_0$ and $\\mu$ greatly and  it most strongly affects the assumed source size. Our conclusion regarding the bar-like shape of the star forming region  depends therefore sensitively on the assumed inclination angle. However, only for a severe underestimate of the true inclination angle of NGC~2146 (i.e. if $i\\ga 80^\\circ$) this conclusion cannot be maintained. In this case the radio emission along the minor axis would be almost entirely from the halo.  \\subsection{CR propagation in the central disk and halo}  The diffusion model is a good fit to the radio data for NGC~2146 along the major axis of the galaxy and along the minor axis for distances less than 10 arcsec (2.6~kpc) from the major axis. This suggests that diffusion is the dominant mode of propagation in this region which corresponds to the inner, most actively star forming disc of NGC~2146.  At distances greater than 10 arcsec from the centre of the galaxy along the minor axis the model is no longer an adequate fit to the data (Figs. 5 and 6). The shape of the 1.5~GHz emissivity is very flat suggesting that transport of CR electrons is more efficient than would be expected from diffusion alone. Furthermore, the spectral index flattens which indicates that (i) the energy losses have decreased rapidly, or (ii) the diffusion coefficient increases, or (iii) a much faster process than diffusion is responsible for the transport of the CR electrons. Upon examination of the radio emission in the region where the diffusion model begins to fail we note that the radio images (Figs. 1 and 3) show filamentary structure. The most likely reason for this behaviour is that away from the centre of the galaxy a large fraction of the radio emission is from a halo which is seen in projection.  The propagation of CR electrons in the halo might be  more complicated than we have considered for the propagation within the disc. Evidence for this comes from observations of the edge-on galaxies NGC~891 and NGC~4631 where extended radio continuum emission in the halo with a rather flat spectral index has been observed (Hummel 1991). Breitschwerdt (1994) has explained this flat spectral index in a galactic wind model including diffusive and convective transport of CR's. Siemieniec and Cesarsky (1991) modelled the spectral index in the halo of NGC 891 by a diffusion model with an outwardly increasing diffusion coefficient. Both galaxies show that in the halo the propagation of CR electrons is faster than that predicted by steady-state diffusion alone.  Due to the strong star formation activity in NGC~2146, an outflow from the disc triggered by correlated SN explosions  may be present, a mechanism which is variously described as a  ``galactic fountain'' (Shapiro \\& Field 1976) or ``chimney'' model (Ikeuchi 1988, Norman \\& Ikeuchi 1989). Such an outflow would transport CR's and magnetic field into the halo and could therefore account for a higher radio emission in the halo than expected from a pure diffusion model. The analysis of X-ray and optical data of NGC~2146 (Armus et al. 1995) has shown that there are indeed indications of the existence of such a starburst-driven superwind.   \\subsection{Star formation in NGC~2146}  Hutchings et al. (1990) proposed that NGC~2146 is in the late stage of a merger. The merger must have begun $\\geq 10^9$ years ago as the rotation curve does not appear to be disrupted in the outer regions of the galaxy (Young et al. 1988a, show the inner regions of the galaxy to have non-circular motions of the order of 100 km s$^{-1}$). In this scenario material has collapsed into the centre of the galaxy and a starburst has commenced in the nucleus.  The data and model presented in this paper further emphasize that active star formation is confined only to the centre of the galaxy. The radio images show that the star formation activity is located in a peculiar spiral {\\bf S} shape around the peak of the radio emission. The merger may have triggered the formation of a bar which feeds the nuclear starburst with material and is responsible for the unusual structure of the radio emission. Star formation is proceeding all along the central bar, and the nucleus is not prominent in the high-resolution radio images. Young et al. (1988a) suggest that if the galaxy is in the late stage of a merger the  nucleus of the merging companion could be obscured by the dust lane. The images of the radio emission presented here argue against this as we do not find two spatially distinct peaks of radio emission  which could be identified as the nuclei of the merging systems. Instead, the star formation is distributed throughout the bar-like central region.  The unresolved sources seen on our highest resolution images have sizes of less than 100~pc and may represented isolated pockets of star formation or could be individual high-luminosity supernovae as seen in M82 (Muxlow {\\em et al.} 1994). The presence of dynamical features such as the arm in the H$\\alpha$ emission (Young, Kleinmann \\& Allen 1988b) and the bar in the radio emission suggest that this starburst is dynamically driven, i.e. due to a merger or an interaction.  {\\bf Acknowledgements:} The VLA is operated by the National Radio Astronomy Observatory for Associated Universities Inc., under a cooperative agreement with the National Science Foundation. UL gratefully acknowledges the receipt of a postdoctoral fellowship of the Deutsche Forschungsgemeinschaft (DFG). We would like to thank the referee, Prof. Davies, for useful comments."
    },
    "9603/astro-ph9603057_arXiv.txt": {
        "abstract": "We obtained X-ray spectra of the Seyfert 1 galaxy NGC~3516 in March 1995 using the Japanese X-ray satellite ASCA. Simultaneous far-UV observations were obtained with the Hopkins Ultraviolet Telescope on the Astro-2 shuttle mission. The ASCA spectrum shows a lightly absorbed power law of energy index 0.78. The low energy absorbing column is significantly less than previously seen. Prominent {\\sc O~vii} and {\\sc O~viii} absorption edges are visible, but, consistent with the much lower total absorbing column, no Fe K absorption edge is detectable. A weak, narrow Fe~K$\\alpha$ emission line from cold material is present as well as a broad Fe~K$\\alpha$ line.  These features are similar to those reported in other Seyfert 1 galaxies. A single warm absorber model provides only an imperfect description of the low energy absorption. In addition to a highly ionized absorber with ionization parameter $U = 1.66$ and a total column density of $1.4 \\times 10^{22}~\\rm cm^{-2}$, adding a lower ionization absorber with $U = 0.32$ and a total column of $6.9 \\times 10^{21}~\\rm cm^{-2}$ significantly improves the fit. The contribution of resonant line scattering to our warm absorber models limits the Doppler parameter to $< 160~\\rm km~s^{-1}$ at 90\\% confidence. Turbulence at the sound speed of the photoionized gas provides the best fit. None of the warm absorber models fit to the X-ray spectrum can match the observed equivalent widths of all the UV absorption lines. Accounting for the X-ray and UV absorption simultaneously requires an absorbing region with a broad range of ionization parameters and column densities. ",
        "introduction": "Intrinsic absorption is a valuable tool for probing structures in active galactic nuclei (AGN). While absorption may in principle arise anywhere in the host galaxy, the most interesting absorbers are those that appear to be associated with the central engine.  These ``warm absorbers\" commonly appear in the X-ray spectra of AGN (\\cite{Turner93}; \\cite{NP94}), and they could be material in the broad-emission-line region (BELR) (e.g. \\cite{Netzer93}; \\cite{RF95}) or the X-ray heated wind which forms the reflecting region in type 2 AGN (\\cite{KK95}), or an entirely new component. If X-ray warm absorbers are related to associated UV absorption systems (\\cite{Mathur94}; \\cite{Mathur95}), then UV and X-ray observations together place powerful constraints on the ionization structure of the absorber. In the X-ray one can measure the column densities of highly ionized species (e.g. {\\sc O~vii} and {\\sc O~viii}) while simultaneously observing lower ionization relatives in the UV ({\\sc O~vi}, {\\sc N~v}, and {\\sc C~iv}). Objects with strong UV absorption lines and soft X-ray absorption are therefore good candidates for further tests of this hypothesis.  The Seyfert 1 galaxy NGC~3516 exhibits unusually strong, variable UV absorption lines (\\cite{UB83}; \\cite{Voit87}; \\cite{Walter90}; \\cite{Kolman93}; \\cite{Koratkar96}), and has a variable X-ray spectrum characteristic of the warm absorber phenomenon (\\cite{Halpern82}). Observations obtained with {\\it Ginga} (\\cite{Kolman93}; \\cite{NP94}) show a flat power law with energy index $\\sim 0.5$ over the 2--18 keV range, a highly ionized iron edge with a corresponding total column density of $\\sim 3 \\times 10^{23}~\\rm cm^{-2}$, and a cold fluorescent Fe K$\\alpha$ line with EW = 377 eV.  To measure simultaneously the X-ray and UV absorption in NGC~3516 we coordinated ASCA observations with the flight of the Astro-2 space shuttle mission in March 1995.  Far-ultraviolet spectra obtained with the Hopkins Ultraviolet Telescope (HUT) that allow us to measure the resonance doublets of {\\sc O~vi}, {\\sc N~v}, Si~{\\sc iv} and {\\sc C~iv} are discussed in a companion paper by Kriss et al. (1996)\\markcite{Kriss96a}. ",
        "conclusions": "The broad Fe K$\\alpha$ emission line in our spectrum of NGC~3516 resembles features seen in MCG--6-30-15 (\\cite{Tanaka95}; \\cite{Fabian95}), in NGC~5548 and IC~4329A (\\cite{Mushotzky95}) and in NGC~4151 (\\cite{Yaqoob95}). At the signal-to-noise ratio of our spectrum we are unable to place significant constraints on relativistic disk models. In fact, the width and luminosity of the line depends strongly on our model for the underlying continuum. In the empirical model, summarized in Table \\ref{empirical_table}, the line width is smaller largely because the underlying continuum is flatter and has less curvature than in the warm absorber models. In the warm absorber models the widths and equivalent widths are large, but given the uncertainties, they are compatible with the maximum equivalent widths of $\\sim 200$ eV expected in models of X-ray reflection from cold disks (e.g., \\cite{GF91}).  A simple, single-zone photoionized absorber is an imperfect description of the soft X-ray opacity in NGC~3516. At least two zones are required to give as good a fit to the data as our empirical model containing discrete absorption edges due to {\\sc O~vii} and {\\sc O~viii}. These two zones may be a simplification of a broad distribution of ionization parameters in the absorbing gas, or they may be indicative of two entirely different regions as suggested by Otani et al. (1996)\\markcite{Otani96} in their study of the variability of {\\sc O~vii} and {\\sc O~viii} opacity in MCG--6-30-15. The complexity of the absorbing gas in NGC~3516 increases even more when one considers the UV absorption lines observed simultaneously with HUT. In the companion to this paper Kriss et al. (1996)\\markcite{Kriss96a} find that neither a single photoionized absorber nor the multiple warm absorber models considered here are adequate to explain the UV absorption lines. This confirms the conclusion of Kolman et al. (1993)\\markcite{Kolman93} that there is probably not a direct connection between the warm X-ray and the UV absorbers in NGC~3516.  Having separate regions for the X-ray and UV absorption contrasts with the conclusions of Mathur et al. (1994,1995)\\markcite{Mathur94}\\markcite{Mathur95} in their studies of the warm absorbers in 3C~351 and in NGC~5548. For those objects they described a single absorbing zone that could account for both the X-ray warm absorber and the associated UV absorption lines. Not all X-ray and UV absorbers may be so simple. A clue to the additional complexity of the absorption in NGC~3516 is the presence of large columns of lower ionization species such as Si~{\\sc iv} and an optically thick Lyman limit. Low ionization species also present problems for single zone models in attempting to model UV and X-ray absorption in NGC~4151 (\\cite{Kriss95}). Since Si~{\\sc iv} absorption is strong in broad absorption line quasars, by analogy this may indicate that single zone models will also not suffice to explain both the UV and X-ray absorption in these objects, contrary to the suggestion of Mathur, Elvis, \\& Singh (1995).\\markcite{MES}  Although the absorbing medium in NGC~3516 appears to be highly stratified, the presently observed weak X-ray absorption coincident with an episode of weak UV absorption suggests that some underlying mechanism ties them together. At the epoch of our observation, NGC~3516 was brighter than usual by about a factor of two, but the low energy X-ray absorption decreased by more than can be accounted for simply by photoionization due to the increased luminosity. As suggested by Walter et al. (1990)\\markcite{Walter90}, these large changes in absorption column may be caused by different clouds moving across the line of sight. Rather than a single cloud, however, we require a whole population of clouds of differing column densities and ionization parameters moving into place. If the absorption arises in a wind driven from the accretion disk or the obscuring torus, large changes in opacity that are correlated in the UV and the X-ray may be linked by fluctuations in the mass supply to the outflow. X-ray and UV absorption in AGN may ultimately have a common origin, but the absorption probably occurs in distinctly different regions with a variety of physical conditions."
    },
    "9603/astro-ph9603147_arXiv.txt": {
        "abstract": "We have examined fifteen days of CGRO/BATSE data, obtained during the first outburst of the black-hole candidate source \\groj, to search for rapid variability of its X-ray flux. We find no evidence for significant variability of \\groj\\ during our observations, with a 2$\\sigma$ upper limit to the fractional r.m.s.\\ amplitude in the frequency range 0.03--0.488 Hz of $6.6\\%$. We cannot, on the basis of our observations, determine the source state (low, high, or very-high state) of \\groj. ",
        "introduction": "The X-ray transient source \\groj\\ (also known as X-ray Nova Scorpii 1994) was discovered with BATSE on the Compton Gamma Ray Observatory on 27 July 1994. The evolution of the hard (20--100 keV) X-ray intensity during the subsequent 5 months has been discussed by Harmon et al.\\ (1995). Three outbursts were observed, each characterized by a fast ($< 1$ day) rise to a level of 600--700 mCrab in the 20--100 keV range, with no well-defined single maximum during the outbursts (10--50 days). Typical peak flux levels were $0.30\\; \\rm photons \\; cm^{-2}\\, sec^{-1}$ in the energy range of the observations. There was no significant emission from the location of \\groj\\ for a 50 day period prior to 27 July.  \\groj\\ is a nearby (2--3 kpc) dynamical black-hole candidate with a reported mass function of $3.35 \\pm 0.14\\rm\\, M_\\odot$ (Bailyn et al.\\ 1995a,b; see Tanaka \\& Lewin 1995 for a recent review of black-hole candidates).  The source shows strong radio outbursts associated with superluminal expansion events (Tingay et al. 1995; Hjellming \\& Rupen 1995) that are correlated with the increase in hard X-ray flux. Bailyn et al.\\ (1995b) found that the optical light curve shows eclipses, at an orbital period of 2.6 days, and suggested that the orbital inclination of \\groj\\ is close to 90$^\\circ$. The energy spectrum can be well described by a power law out to at least 300 keV, with spectral (photon) index varying between 2.5 and 3.1 (Wilson et al. 1995).  Here we report on the analysis of the fast variability of the hard X-ray intensity of \\groj.  In section 2 we describe the observations and data analysis.  We discuss our results in section 3. Section 4 summarizes our conclusions. ",
        "conclusions": "During the first outburst of \\groj\\ we detected no (0.03--0.488 Hz) variability in its (20--100 keV) flux, with an upper limit to its r.m.s. amplitude of 6.6\\% ($2 \\sigma$). Several bright ($\\geq 1$ Crab, 20--100 keV)  black-hole candidate sources observed with BATSE (in particular, Cygnus X-1, \\frank, and GRO~J0422+32) have shown enhanced noise power in the 0.03--0.488 Hz range, and occasional quasi-periodic oscillations (Kouveliotou 1994; Van der Hooft et al. 1995). Fractional r.m.s.\\ values for this noise observed from Cygnus~X-1 (10--30\\%, Crary et al. 1995) and \\frank\\ ($\\sim \\rm 15\\%$, Van der Hooft et al. 1995) are characteristic of the low state (Van der Klis 1994), and it is likely that they were encountered in that state with BATSE.  Based on the BATSE data for the period July 28--August 11, 1994 alone we cannot interpret the observed low variability of \\groj\\ in terms of the source state scheme for black-hole candidates. For that a better understanding of the hard X-ray properties of black-hole candidates (e.g., their spectral slope) in the different source states is required. Investigations of this type will be facilitated by low-energy ($< 10$ keV) observations, performed concurrently with a BATSE observation.  \\vspace{.25in} We thank Dr. W. Lewin for his comments on this paper. This project was performed within NASA grant NAG5-2560 and supported in part by the Netherlands Organization for Scientific Research (NWO) under grant \\mbox{PGS 78-277.} This work was performed while DJC held a National Research Council-NASA Research Associateship. FvdH acknowledges support by the Netherlands Foundation for Research in Astronomy with financial aid from NWO under contract number \\mbox{782-376-011}, and the Leids Kerkhoven--Bosscha Fonds for a travel grant. JvP acknowledges support from NASA grant \\mbox{NAG5-2755}.  \\pagebreak"
    },
    "9603/astro-ph9603129_arXiv.txt": {
        "abstract": "We study star-formation-inducing mechanisms in galaxies through multi-wavelength measurements of a sample of dwarf galaxies in the Virgo cluster described in paper I. Our main goal is to test how star formation inducing mechanisms depend on several parameters of the galaxies, such as morphological type and hydrogen content. We derive the star formation rate and star formation histories of the galaxies, and check their dependence on other parameters.   Comparison of the sample galaxies with population synthesis models shows that these objects have significantly lower metallicity than the Solar value. The colors can generally be explained as a combination of two different stellar populations: a young (3--20 Myr) metal-poor population which represents the stars currently forming presumably in a starburst, and an older (0.1--1 Gyr) population of previous stellar generations. There is evidence that the older stellar population was also formed in a starburst. This is consistent with the explanation that star formation in this type of objects takes place in short bursts followed by long quiescent periods.  No significant correlation is found between the star formation properties of the sample galaxies and their hydrogen content. Apparently, when star formation occurs in bursts, other parameters influence the star formation properties more significantly than the amount of atomic hydrogen. No correlation is found between the projected Virgocentric distance and the rate of star formation in the galaxies, suggesting that tidal interactions are not significant in triggering star formation in cluster dwarf galaxies. ",
        "introduction": "\\label{sec_int}  Star formation is probably the most fundamental process in galaxies. Understanding its nature, and its dependence on galactic type and environment may contribute to our knowledge about development of galaxies, as well as about the development of the entire Universe.  The star formation process is characterized by two main parameters: the initial mass function (IMF) and the total star formation rate (SFR). First introduced by Salpeter (1955), the IMF is usually described as a power law in the range of 2--2.5 (the original value proposed by Salpeter is 2.35). Other characteristics of the IMF are the low and high mass limits, usually taken as 0.1M$_\\odot$ and 60--120M$_\\odot$. Several typical IMFs describe star formation in different types of galaxies and environments, such as the solar neighborhood, the Magellanic clouds, etc. Some IMFs (e.g, Miller and Scalo 1979, and Scalo 1986) were not published originally as power laws, but can be fitted as power laws with varying coefficients for different mass ranges. In Fig.~\\ref{fig_IMF} four different IMFs, used in various models, are shown.  \\begin{figure}[htbp] \\vspace{9.4cm} \\special{psfile=\"f-fig_IMF.ps\" angle=270 hscale=55 vscale=50 hoffset=0 voffset=290} \\caption{\\protect \\footnotesize{Different IMFs by several authors. The power law coefficients are indicated close to each segment of the relations. The tick marks on the abscissa are one unit apart, i.e., each tick mark corresponds to a factor of 10 in the number of stars formed per unit mass. The designations are as follows: S55 $\\equiv$ Salpeter (1955), K83 $\\equiv$ Kennicutt (1983), MS97 $\\equiv$ Miller \\& Scalo (1979), S86 $\\equiv$ Scalo (1986). The MS79 and S86 IMFs are shown as parameterized by Bruzual \\& Charlot (1993).}} \\label{fig_IMF} \\end{figure}  As for the SFR, its time dependence varies dramatically between different galactic types (e.g., Gallagher, Hunter \\& Tutukov 1984; Kennicutt {\\it et al.} 1994). In early type galaxies, the star formation usually decays smoothly with time. In late-type galaxies, on the other hand, the star formation is generally more intense, and the SFR is subject to significant changes on short timescales.  Despite extensive progress in understanding the star formation process, there are several issues still not fully understood. In paper I we introduced two open questions concerning the star formation processes in galaxies: (1) What are the mechanisms that govern the star formation process, and how do they depend on the galactic type and environment? (2) How do the SFR and IMF depend on various galactic properties, such as interstellar gas density, morphology of the interstellar gas, metallicity and the amount of dust in the interstellar medium.  The simplified assumption is that the SFR depends directly on the density of the interstellar gas (Schmidt law, Schmidt 1959). Actually, the observable quantity is the gas surface density.  Testing this parameter against the SFR in normal galaxies indicates that there is a threshold gas surface density, below which no star formation takes place (Kennicutt 1989). The value of this threshold surface density varies from galaxy to galaxy, thus it is not a global parameter, though it is believed to be of order a few times $10^{20} atoms/cm^2$.  Larson (1987) counts four major mechanisms that induce star formation in galaxies: large scale gravitational instabilities of gas clouds, compression of interstellar gas clouds due to the passage of a gravitational density wave, compression in a rotating galactic disk - shear forces acting on the clouds due to differential rotation and random collisions between clouds. In addition, star formation can be triggered by external influences, such as tidal interaction induced by another galaxy during a close encounter of two galaxies, or by interaction with interstellar matter, during a passage near a cluster core.  In order to investigate the above questions, we have constructed four samples of galaxies, described in paper I. These galaxies are all dwarfs in the Virgo cluster. The aim in selecting dwarf galaxies is to eliminate some of the mechanisms described above, i.e. the grand design density wave and shear differential rotation forces do not act in dwarf galaxies. Therefore the theoretical situation is simplified.  Our goal in this paper is to derive the star formation properties of the sample galaxies from the observational data gathered in paper I, and to check these againts various star formation scenarios. This will enable to test the feasibility of several star formation mechanisms.  Since all methods used to determine galactic star formation properties combine observational data with theoretical models, the result is model dependent. Realistic models which use conventional IMFs and reliable stellar evolutionary tracks would result in SFR values that may differ usually by at most $\\sim$50\\% from each other, given the same observational data. This is approximately the accuracy of our observed parameters from which the SFR is derived. As will be discussed below, the uncertainty is mainly due to effects of dust extinction and other observational biases.  In order to estimate the SFR in a sample of galaxies, one needs to know the IMF. This can be found by fitting a number of observed properties to a set of models of different stellar populations, with different IMFs (population synthesis).  Since the colors of a stellar population depend also on its age and history, a specific star formation history is usually assumed, and the observed properties can be tested against different IMFs. Naturally,  this approach can be used only for entire samples of galaxies and not for individual galaxies.  Once the IMF is determined (or assumed), the SFR can be calculated on the basis of models + observations. All the observational techniques basically measure the number of massive stars in the galaxy to trace the 'current' star formation. If the lifetime of stars with a certain mass $M$ is $T(\\!M\\!)$ and the star formation rate at this mass is $R(\\!M\\!)$, then the number of such stars currently observed is simply $N(\\!M\\!)=R(\\!M\\!)\\times T(\\!M\\!)$, where both $R(\\!M\\!)$ and $N(\\!M\\!)$ are given per unit stellar mass.  The total SFR is calculated by extrapolating from the massive stars to the entire mass range, using the IMF. The higher the mass of the stars we use, the larger the error of the total SFR that can be introduced due to the extrapolation. On the other hand, concentrating on higher mass stars leads to a more 'up to date' star formation result. Therefore, there is a tradeoff between how `current' the derived SFR is, and the accuracy of the total SFR. In paper I we described the observations of UV and H$\\alpha$ line radiation from the sample galaxies, which we use here for determination of their SFR.  The relation between the number of observed LyC photons and the SFR can be reduced to the knowledge of the dependence of $R(\\!M\\!)$ on $M$, and the number of LyC photons emitted by a star with mass $M$ during its entire lifetime $P(\\!M\\!)$. With these two parameters, the number of photons emitted per unit time by an $M$ mass star (per unit stellar mass) is $F(\\!M\\!)=P(\\!M\\!)\\times R(\\!M\\!)$, and the total flux of LyC photons is given by:  \\begin{equation} \\label{e_Nly} N_c = \\int_{M_{min}}^{M_{max}} F(\\!M\\!)\\; dM = \\int_{M_{min}}^{M_{max}} P(\\!M\\!)\\cdot R(\\!M\\!) \\;dM \\end{equation}  where $M_{min}$ and $M_{max}$ are the low and high mass limits of the IMF, and the dependence of $R(\\!M\\!)$ on $M$ is the IMF.  We can now express the number of LyC photons as: \\begin{equation} \\label{e_NsR} N_c = S\\!F\\!R \\; \\frac{{\\displaystyle \\int_{M_{min}}^{M_{max}} P(\\!M\\!) \\cdot I\\!M\\!F(\\!M\\!)\\; dM}} {{\\displaystyle \\int_{M_{min}}^{M_{max}} M \\cdot I\\!M\\!F(\\!M\\!)\\; dM}} \\end{equation}  where $S\\!F\\!R$ is in M$_\\odot$/yr, and the integrals in the equation originate from the theoretical models. We therefore obtain a direct relation between the SFR and the number of LyC photons.  This treatment assumes a constant IMF. However, some studies indicate a dependence of the IMF on the metallicity of the galaxy (e.g., Terlevich \\& Melnick 1983). The metallicity of a galaxy gradually increase wirh time, thus, the IMF will also depend on the galactic age. In this case the SFR and IMF would be coupled and will be changing in time.  The LyC flux is derived here through its influence on the ambient hydrogen, e.g., by measurement of the resulting Balmer lines and assuming case B recombination. For the calculation of the SFR we use the H$\\alpha$ line. This line has been used by many (Kennicutt 1983, hereafter K83, Kennicutt \\& Kent 1983; Gallagher, Hunter \\& Tutukov 1984; Pogge and Eskridge 1987; Kennicutt {\\it et al.} 1994), mainly because of its high intensity. A few percent of the total ionizing flux are reemitted as H$\\alpha$ (Kennicutt 1989), so it is convenient for tracing the star formation in faint galaxies. The internal dust extinction influences the line intensity significantly, and will be discussed in more detail in the following section. ",
        "conclusions": "\\label{sec_conc}   The main goal of this study is to investigate mechanisms that govern star formation processes in galaxies and their dependence on various galactic parameters. The intention in focusing on our sample of late-type dwarfs in the Virgo cluster was to exclude some of these mechanisms thought to be responsible for star formation in large galaxies. In addition, we concentrate on cluster members to test for effects of the cluster environment on the star formation properties of the galaxies.  The observational data used to evaluate the star formation parameters, such as the SFR, IMF, and star formation history, are affected primarily by the internal dust extinction in the galaxies. This depends on the amount and distribution of dust in the galaxy. We adopted general correction parameters for the entire sample to account for the effects of dust.  Using a data base consisting of a number of broad band colors and H$\\alpha$ observations it is possible to track the ongoing star formation process, as well as the star formation history of the sample. This is done for both low HI and high HI subsamples, which enables one to check the dependence of these parameters on the neutral hydrogen content of the galaxies. In all cases, no significant dependence of the star formation properties on the HI content was found. This may be explained by the following argument: the differences among the various galaxies in our sample lie in the relative weight of the flux originating from their `current' and `previous' star formation episodes. Since the hydrogen is depleted during each such burst of star formation, its current amount depends on the number of bursts that occurred in the past, as well as on the original amount present. The strength of the current starburst apparently does not depend on this star formation history and, thus, cannot be correlated with the neutral hydrogen content.  Considering the entire set of observations, together with various population synthesis models, a star formation scenario can be sketched for the sample galaxies. The process of star formation in late-type dwarf galaxies takes place in short bursts, presumably much shorter than 1Gyr (10--100 Myr). The past burst(s) probably occurred within the last $\\sim$Gyr. Redder galaxies can be fitted with a single, longer burst with this decay time (1 Gyr), but the observational data of the entire sample can only be explained in terms of a series of short bursts. In this picture, the difference among galaxies lies in the relative weight of the starburst and the older population. This depends on the age and size of the current burst and of former bursts, therefore the dwarf galaxies may be seen as a single type, which are observed at different epochs of their evolution.  In addition, the galaxies appear to have low metal abundance, mainly from their blue R--I color. This is in agreement with previous data for irregular dwarf galaxies, known to have typically low metallicities (e.g., Kunth \\& Sargent 1986). A low metallicity of a galaxy indicates its young age, since the amount of heavy elements produced by massive stars increases during the lifetime of the galaxy, which implies that some of our sample galaxies are genuinely young. This also is not a new finding (see Gondhalekar {\\it et al.} 1984). A development trend can be sketched, in which the galaxies undergo a series of starbursts, and their metallicity increases from one burst to the other. Since the HII regions in many of the galaxies occupy a large fraction of the galactic volume, it is possible that each burst changes significantly the total galactic metal abundance. The relation between the fraction of the galactic surface covered by H II regions and star formation properties will be discussed in a subsequent paper.  We have tested the cluster influence on the sample galaxies' star formation properties. No correlation was found between the Virgocentric distande of the galaxies and any of their star formation properties. This may manifest the low significance of tidal forces when acting on dwarf galaxies. It can be understood as these galaxies are small in size and, thus, a gradient in an external gravitational field induced by another galaxy may not cause a great difference from side to side relative to the galaxy's own gravitational well. This finding is, therefore, not surprising.  A remarkable galaxy in our sample is VCC144. It has strong SFR and an exceptional SFR/area, compared with the other galaxies. It is also the only object in which the burst population is believed to dominate in luminosity and in mass over older populations, if any. However, it does not show any special behavior in other parameters such as the HI flux, velocity dispersion, infrared flux, or distance from cluster center, in which it seems a `normal' object in the sample. It is the most condensed object of our sample, in terms of star formation, but with no other peculiarities. It is possible that this object is experiencing its very first burst of star formation, which indicates that galaxies are still being formed these days in the Virgo cluster.  To conclude - the star formation in late-type dwarf galaxies in the Virgo cluster occurs probably in bursts. The bursts do not seem to depend on galactic history or on cluster environment. The details of the IMF of the burst population cannot be determined from the data collected in this study, but a general scenario of the galactic evolution is sketched. During their lifetime, the galaxies evolve from late-type to earlier type, their metallicity increases, and they become redder objects in the optical."
    },
    "9603/astro-ph9603108_arXiv.txt": {
        "abstract": "We present new FUV/UV observations of the DA white dwarf Wolf~1346 obtained with the Hopkins Ultraviolet Telescope. The atmospheric parameters of this object are estimated from a fit of model atmospheres to several optical spectra to be \\Teff~ = 20000~K, \\logg~ = 7.90. From the optical spectrum this star is a normal DA without any indications for chemical elements other than hydrogen. The hydrogen line \\Lbeta, however, shows a very unusual shape, with a steep red wing and two absorption features on this wing. The shape is reminiscent of the effects of quasi-molecular line broadening, as observed in \\Lalpha\\ in cooler DA white dwarfs. We show that this is indeed the correct explanation, by identifying 4 quasi-molecular satellites caused through perturbations by the H$^+$ ion (H$_2^+$ quasi-molecule). The steep red wing is caused by the exponential decline of the line profile beyond the satellite most distant from the line center at 1078~\\AA. ",
        "introduction": "Throughout the recent history of astronomy it has happened repeatedly that the opening of a new observing window into space has lead to discoveries that were not expected.  Just in the field of white dwarfs the systematic exploration of the UV window has brought unexpected surprises. We mention just a few out of a large number: the discovery of metal features in hot DA, very strong carbon lines in DC white dwarfs, which show no features at all in the optical, and quasi-molecular satellites in cool DA, which were expected to show just a simple Stark broadened \\Lalpha\\ line.  One of the least explored spectral windows remains the FUV region, roughly defined as 912 - 1215 \\AA, the region from the Lyman edge to \\Lalpha. While the interstellar matter is quite transparent at these wavelengths, as opposed to the region below the Lyman edge, progress has been slow due to technological difficulties in the production of mirrors and gratings. This situation is improving recently with projects like ORFEUS (Orbiting Retrievable Far and Extreme Ultraviolet Spectrograph; Grewing et al. 1991; Hurwitz \\& Bowyer 1991) and HUT (Hopkins Ultraviolet Telescope, Davidsen et al. 1992).  In this paper we report on an observation of the bright DA white dwarf Wolf~1346 (WD2032+248), obtained with the HUT instrument on a flight in 1995. The observation was part of an Astro-2 Guest Investigator program (Finley, Kimble, and Koester) aimed at studying the Stark broadening of the higher Lyman lines in DA. While several hotter DA show the whole Lyman spectrum compatible with symmetrical Stark broadened profiles without any unexplained features, Wolf~1346 at 20000~K has a \\Lbeta\\ line with a strong asymmetry, a very steep red wing, and absorption features in the wing near 1060 and 1078~\\AA. We demonstrate below that all these features are due to quasi-molecular H$_2^+$ absorption (or broadening of \\Lbeta\\ by protons as perturbers), very similar in nature to the 1400~\\AA\\ feature observed at lower temperatures in the red wing of \\Lalpha. ",
        "conclusions": "We have not made an effort to find the best fitting UV model, and the fit is clearly not perfect, especially in the region between \\Lalpha\\ and \\Lbeta. We have made some experiments and our conclusion is that far wing absorption of \\Lgamma\\ and higher Lyman lines, which are still calculated with standard Stark broadening, are part of the problem. Fig.\\ref{hutsat}, however, clearly proves, by the coincidence of position and shape, that the two observed features near 1060 and 1078 \\AA\\ are indeed satellite features of \\Lbeta, and that the steep rise of the wing is caused by the exponential decline of the line profile beyond the last satellite. Further detailed studies of the line profiles of all Lyman lines will hopefully improve the quantitative agreement in the future, whereas new observations of hotter and cooler objects should establish the range where these features are observable, and provide a challenge to experimental physicists to measure these line profiles in laboratory plasmas."
    },
    "9603/astro-ph9603087_arXiv.txt": {
        "abstract": "Constraints on big bang nucleosynthesis (BBN) and on cosmological parameters from conflicting deuterium observations in different high red-shift QSO systems are discussed.  The high deuterium observations by Carswell {\\it et al}., Songaila {\\it et al}., and Rugers \\& Hogan is consistent with $^4$He and $^7$Li observations and Standard BBN ($N_\\nu$ =3) and allows $N_\\nu \\leq 3.6$ at 95\\% C.L., but is inconsistent with local observations of D and $^3$He in the context of conventional theories of stellar and Galactic evolution.  In contrast, the low deuterium observations by Tytler, Fan \\& Burles and Burles \\& Tytler are consistent with the constraints from local Galactic observations, but require $N_\\nu = 1.9 \\pm 0.3$ at 68\\% C.L., excluding Standard BBN at 99.9\\% C.L., unless the systematic uncertainties in the $^4$He observations have been underestimated by a large amount.  The high and low primordial deuterium abundances imply, respectively, $\\Omega_{\\rm B}h^2 = 0.005 - 0.01$ and $\\Omega_{\\rm B}h^2 = 0.02 - 0.03$ at 95\\% C.L.  When combined with the high baryon fraction inferred from x-ray observations of rich clusters, the corresponding total mass densities (for $50 \\le H_0 \\le 90$) are $\\Omega_{\\rm M} = 0.05 - 0.20$ and $\\Omega_{\\rm M} = 0.2 - 0.7$, respectively (95\\% C.L.)  The range of $\\Omega_{\\rm M}$ corresponding to high D is in conflict with dynamical constraints ($\\Omega_{\\rm M} \\ge 0.2 - 0.3$) and with the shape parameter constraint ($\\Gamma = \\Omega_{\\rm M}h = 0.25 \\pm 0.05$) from large scale structure formation in CDM and $\\Lambda$CDM models. ",
        "introduction": "Among the light nuclides synthesized during the early evolution of the universe, deuterium is unique in its sensitivity to the universal density of baryons and in the simplicity of its galactic evolution. As gas is incorporated into stars and the heavy elements are synthesized, D is only destroyed \\cite{Epstein-Lattimer-Schramm} so that any D abundance inferred from observations provides a {\\em lower} bound to its primordial value. \\footnote{ $X_{2{\\rm P}} > X_{2{\\rm OBS}}$, where the D mass fraction is $X_2 = 2 X\\; n_{\\rm D}/n_{\\rm H}$; $X$ is the hydrogen mass fraction, and $n_x$ is the number density for nuclide $x$; the subscript P is for the primordial abundance.  As an estimate of the primordial value $X_{\\rm P} = 1 - Y_{\\rm P}$, we will adopt $X_{\\rm P} = 0.76 \\pm 0.01$.  In this paper we quote 1$\\sigma$ uncertainties unless otherwise indicated. } Unfortunately, an {\\em upper} bound to the primordial D abundance is more uncertain, depending on the evolutionary history of the matter being observed.  Thus, although estimates of the D abundance in the presolar nebula \\cite{Geiss,Steigman-Tosi-95} and in the local interstellar medium (ISM) \\cite{McCullegh,Linsky-etal} provide interesting lower bounds to primordial D [$X_{2{\\rm P}} \\ge X_{2\\odot}= (3.6 \\pm 1.3) \\times 10^{-5}$, $X_{2{\\rm P}} \\ge X_{2{\\rm ISM}}= (2.2 \\pm 0.3) \\times 10^{-5}$, where the H mass fraction has been taken to be $X_\\odot = X_{\\rm ISM} = 0.70 \\pm 0.01$], upper bounds are more model dependent (see, for example, Ref.~\\cite{Vangioni-Flam-Audouze,Tosi,Steigman-Tosi-92,Steigman-Tosi-95,% D-paper}).  For this reason, observations of D in (nearly) unevolved systems (high red-shift, low metallicity QSO absorbers) have been eagerly anticipated.  If, indeed, $X_{2{\\rm P}} \\sim X_{2{\\rm QSO}}$, then because of the sensitivity of the D abundance to the nucleon abundance ($\\eta = n_{\\rm B}/n_\\gamma$; the ratio of the present baryon density to the critical density is $\\Omega_{\\rm B}h^2 = 0.0037 \\eta_{10}$, where the Hubble parameter is $H_0 = 100 h$ km/s/Mpc and $\\eta_{10} = 10^{10}\\eta$), a measurement of (D/H)$_{\\rm QSO}$ to $\\sim$30\\% accuracy will lead to a determination of $\\eta$ to $\\sim$ 20\\% accuracy.  Armed with $\\eta$, reasonably accurate predictions of the primordial abundances of $^3$He, $^4$He, and $^7$Li will follow (see, for example, Ref.~\\cite{D-paper}).  For example, for $1.5 < \\eta_{10} < 10$, a 20\\% uncertainty in $\\eta_{10}$ will lead to an uncertainty in the predicted $^4$He mass fraction which ranges from $\\sim 0.003$ (at low $\\eta_{10}$) to $\\sim 0.002$ (at high $\\eta_{10}$).  Deuterium is the ideal baryometer.  In the last two years, observations of D in high red-shift, low metallicity QSO absorbers have begun to appear in the published literature \\cite{Carswell-etal,Songaila-etal,Rugers-Hogan,Tytler-Fan-Burles,% Burles-Tytler}. The first observations of D in absorption against Q0014+813 \\cite{Carswell-etal,Songaila-etal,Rugers-Hogan} suggested a surprisingly high abundance [for our quantitative comparisons we will adopt the recent reanalysis by Rugers and Hogan: D/H = $(1.9 \\pm 0.4) \\times 10^{-4}$, $X_2 = (2.9 \\pm 0.6) \\times 10^{-4}$], roughly an order of magnitude larger than the presolar or ISM values ($X_{2{\\rm QSO}}/X_{2{\\rm ISM}} \\sim 8 \\pm 3$, $X_{2{\\rm QSO}}/X_{2\\odot} \\sim 13 \\pm 3$).  As such efficient D destruction in the Galaxy is not expected \\cite{Vangioni-Flam-Audouze,Tosi,Steigman-Tosi-92,Steigman-Tosi-95,Edmunds} it has been suggested that the feature identified as D in Q0014+813 might be a hydrogen interloper \\cite{Steigman-94}.  However, the Rugers-Hogan reanalysis argues against this possibility.  Further, recent papers \\cite{Carswell-etal-1996,Wampler-etal} present evidence for D absorption in front of two other QSOs (Q0420-388 and BR1202-0725, respectively) which, if the identifications are correct, suggest ${\\rm D/H} \\ge 2 \\times 10^{-5}$ and ${\\rm D/H} \\le 1.5 \\times 10^{-4}$, respectively.  Although puzzling from the point of view of chemical evolution in the Galaxy, the high D abundance points towards a low baryon density ($\\eta_{10} \\sim 2$) which is consistent with the predicted and observed (inferred) primordial abundances of $^4$He and $^7$Li \\cite{Steigman-94,crisis-paper}.  As we shall see, however, this low baryon density ($\\Omega_{\\rm B} h^2 \\sim 0.007$) is in conflict with determinations of the total mass density and the baryon fraction inferred from x-ray observations of rich clusters.  In contrast, from recent observations, Tytler, Fan, and Burles \\cite{Tytler-Fan-Burles} and Burles and Tytler \\cite{Burles-Tytler} derive a low D abundance: $({\\rm D/H}) = [2.3 \\pm 0.3$ (stat) $\\pm 0.3$ (sys) $]\\times 10^{-5}$ towards the QSO1937-1009 \\cite{Tytler-Fan-Burles} and $({\\rm D/H}) = [2.5^{+0.5}_{-0.4}$ (stat) $^{+0.4}_{-0.3}$ (sys)] $\\times 10^{-5}$ towards the QSO1009+2956 \\cite{Burles-Tytler}.  We have combined their two results to obtain (D/H)$_{\\rm QSO} = (2.4 \\pm 0.5) \\times 10^{-5}$; $X_{2{\\rm QSO}} = (3.6 \\pm 0.8) \\times 10^{-5}$.  Although marginally larger than ISM deuterium ($X_{2{\\rm QSO}}/X_{2{\\rm ISM}} = 1.6 \\pm 0.4$), the low D abundance is not very different from the presolar value ($X_{2{\\rm QSO}}/X_{2\\odot} = 1.0 \\pm 0.4$), suggesting that even though the absorbers are at high redshift ($z_{\\rm abs}$ = 3.572 and 2.504) and have very low metallicity ($\\sim 10^{-3}$ solar), some D may have already been destroyed ($X_{2{\\rm P}} \\ge X_{2{\\rm QSO}}$).  If indeed $X_{2{\\rm P}} \\sim X_{2{\\rm QSO}}$ (no significant D destruction), then the problems for BBN identified by Hata {\\it et al}.\\ \\cite{crisis-paper}, which were based on $X_{2{\\rm P}}$ inferred from solar system observations of D and $^3$He, persist.  The higher baryon density suggested by the low D result is, however, in agreement with the x-ray cluster data (but still supports a low density universe).  It is hoped that future observations of D in other high red-shift, low metallicity QSO absorbers will resolve the current dichotomy between the high D result for Q0014+813 on the one hand \\cite{Carswell-etal,Songaila-etal,Rugers-Hogan} and the low D results for Q1937-1009 and Q1009+2956 on the other \\cite{Tytler-Fan-Burles,Burles-Tytler}.  Here, we explore the implications for cosmology (the baryon density), for the primordial abundances of the other light nuclides ($^4$He and $^7$Li), and for particle physics (bounds to the effective number of equivalent light neutrinos, $N_\\nu$), of the high D abundance and contrast them with those for the low D abundance. ",
        "conclusions": "A determination of the deuterium abundance in a nearly uncontaminated environment such as that provided by high redshift, low metallicity QSO absorption clouds could be a key to testing the consistency of primordial nucleosynthesis in the standard, hot, big bang cosmology, to pinning down the universal density of baryons, and to constraining physics beyond the standard model of particle physics.  Such data is beginning to be acquired but, at present, the observational situation is in conflict.  On the one hand there is evidence in favor of high D \\cite{Songaila-etal,Carswell-etal,Rugers-Hogan,Wampler-etal,% Carswell-etal-1996}: (D/H)$ \\sim 2 \\times 10^{-4}$.  In contrast, Tytler, Fan, \\& Burles \\cite{Tytler-Fan-Burles} and Burles \\& Tytler \\cite{Burles-Tytler} find evidence for low D: (D/H)$ \\sim 2 \\times 10^{-5}$.  If the former, high-D values are correct, it is surprising that Tytler, Fan, \\& Burles and Burles \\& Tytler fail to find such a large abundance in their high redshift ($z$ = 3.57 and 2.50), very low metallicity ($Z/Z_\\odot \\sim 10^{-3}$) absorbers; high $z$ and low $Z$ argue against an order of magnitude destruction of primordial D.  If, instead, the low D result is correct, such weak D-absorption might often go unnoticed and the high-D cases might be accidental interlopers.  Based on velocity information Rugers and Hogan \\cite{Rugers-Hogan} argue against this possibility which, if more high-D cases are found, will become increasingly unlikely. Presumably, the present confused situation will be clarified by the acquisition of more data.  Here, we have considered separately the consequences for cosmology and particle physics of the high-D and low-D data.  For the high-D case, SBBN ($N_\\nu = 3$) is consistent with the inferred primordial abundances of D, $^4$He, and $^7$Li provided that the baryon density is small (see Table~\\ref{tab:constraints} and Figs.~\\ref{fig:bbn_comb_qsos}, \\ref{fig:eta-P_qsos}, \\ref{fig:eta-Y_qsos}, \\ref{fig:eta-y7_qsos}, \\ref{fig:nnu-P_comb}, and \\ref{fig:dY-nnu_qsos}). However, for consistency with the solar system and/or present interstellar D abundances, such a large primordial D abundance requires very efficient D destruction.  The low baryon density which corresponds to high-D still leaves room for dark baryons and reinforces the case for non-baryonic dark matter (see Fig.~\\ref{fig:H-Omega_B}).  However, when combined with the x-ray cluster data, low $\\Omega_{\\rm B}$ and high $f_{\\rm HG}$ suggest a very low density universe ($\\Omega_{\\rm M} \\lesssim 0.21$ for $H_0 \\ge 50$ km/s/Mpc; see Fig.~\\ref{fig:H-Omega_M}).  The conflict between the upper bound on $\\Omega_{\\rm M}$ and the evidence for a lower bound $\\Omega_{\\rm DYN} \\gtrsim 0.2 - 0.3$ argue against high-D and low $\\eta_{10}$ (see Figs.~\\ref{fig:H-Omega_M_comb} and \\ref{fig:H-Omega_ML_comb}) .  Such a low value for $\\Omega_{\\rm M}$ is also in conflict with the constraint from the shape parameter $\\Gamma$ (see Figs.~\\ref{fig:H-Omega_M_comb} and \\ref{fig:H-Omega_ML_comb}).  These problems persist even allowing for a non-vanishing cosmological constant (which could resolve the age-expansion rate problem).  See Figs.~\\ref{fig:H-Omega_M_comb} and \\ref{fig:H-Omega_ML_comb}.  In contrast, the low-D case leads to severe tension between the SBBN prediction and the inferred primordial abundance of $^4$He (see Figs.~\\ref{fig:bbn_comb_qsos}, \\ref{fig:eta-Y_qsos}, and \\ref{fig:eta-y7_qsos}).  This stress on SBBN can be relieved if the primordial helium mass fraction, derived from observations of low metallicity HII regions, is in error --- due, perhaps, to unaccounted systematic effects --- by an amount $\\Delta Y_{\\rm sys} \\ge 0.011$ (see Fig.~\\ref{fig:dY-nnu_qsos}).  Alternatively, this conflict could be evidence of ``new physics'' \\cite{Langacker-TASI} ($N_\\nu \\ne 3$; see Figs.~\\ref{fig:nnu-P_comb} and \\ref{fig:dY-nnu_qsos}).  The best fit between predictions and observations with low D is for $N_\\nu = 1.9 \\pm 0.3$.  One way to alter standard BBN is to change the physics of the neutrino sector. For example, many models predict the existence of sterile neutrinos, which interact only by mixing with the ordinary neutrinos.  Such sterile neutrinos would not contribute significantly to the number of effective neutrinos ($2.991 \\pm 0.016$) inferred from the Z line-shape \\cite{LEP-Nnu-limit}, but could be produced cosmologically for a wide range of masses and mixings \\cite{Langacker-nu_s}.  However, they only increase $N_\\nu$, exacerbating the discrepancy.  Another possibility arises if $\\nu_\\tau$ has a mass in the range $10 \\ {\\rm MeV}\\lesssim M_{\\nu_\\tau} \\le 24\\ {\\rm MeV}$ (the upper limit is the recent result from ALEPH \\cite{ALEPH}).  In this case BBN production of $^4$He can be either increased or decreased (relative to the standard case), depending on whether $\\nu_\\tau$ is stable or unstable on nucleosynthesis time scales ($\\sim 1$ sec).  An effectively stable $\\nu_\\tau$ ($\\tau \\ge 10$ sec) in this mass range always increases $Y$ relative to the standard case \\cite{nu-tau} (but, see \\cite{Hannestad-Madsen}) and would thus make for a worse fit with the data.  However, if $\\nu_\\tau$ has a lifetime $\\lesssim 10$ sec and decays into $\\nu_\\mu +\\phi$ (where $\\phi$ is a `majoron-like' scalar), \\footnote{% Decays with $\\nu_e$ in the final state can directly alter the neutron-to-proton ratio and thus affect $Y_{\\rm P}$ somewhat differently \\cite{Gyuk-Turner}.}  it is possible to decrease the predicted $Y$ relative to the standard case (see figures 3 and 7 of Ref.~\\cite{Kawasaki-etal}).  Such an unstable $\\nu_\\tau$ contributes less than a massless neutrino species at the epoch of BBN, thereby reducing the yield of $^4$He.  For example, a $\\nu_\\tau$ with mass $20 - 30$ MeV which decays with a lifetime of $\\sim 0.1$ sec reduces $N_\\nu$ by $\\sim 0.5 - 1$ (and $Y$ by $\\sim 0.006 - 0.013$, respectively), thus helping to resolve the apparent conflict between theory and observation.  It is also possible to alter the yield of BBN $^4$He by allowing $\\nu_e$ to be degenerate \\cite{degeneracy}.  If there are more $\\nu_e$ than $\\bar{\\nu}_e$, $Y$ is reduced relative to the standard (no degeneracy) case as the extra $\\nu_e$'s drive the neutron-to-proton ratio to smaller values at freeze-out.  A reduction of $Y$ of $\\sim 0.01$ can be accomplished with a $\\nu_e$ chemical potential of $\\mu_e/T_\\nu \\sim 0.03$, corresponding to a net lepton-to-photon ratio of 0.005.  This is to be compared to the net baryon asymmetry which is smaller by $\\sim$7 orders of magnitude.  Nevertheless, scenarios for a large lepton asymmetry are possible \\cite{Langacker-Segre-Soni}.  Lastly, one can relax the assumption that baryons are homogeneously distributed.  However, inhomogeneous BBN typically results in higher $Y_{\\rm P}$, and therefore does not naturally resolve the $^4$He-D conflict \\cite{inhomogeneous-BBN}.  Provided that the high-$Y$, low-D challenge can be resolved (by $\\Delta Y_{\\rm sys} \\ge 0.011$ and/or $N_\\nu < 3$), low-D is consistent with the Pop II $^7$Li abundance if there has been a modest amount of lithium destruction/dilution in the oldest stars (see Figs.~\\ref{fig:bbn_comb_qsos} and \\ref{fig:eta-y7_qsos}).  The higher baryon density for low-D strengthens the case for dark baryons (see Fig.~\\ref{fig:H-Omega_B}), although that for non-baryonic dark matter, while still very strong, is somewhat weakened.  When folded with the hot gas bound on the x-ray cluster baryon fraction, a ``cluster baryon crisis'' persists, arguing for $\\Omega_{\\rm M} < 1$ (see Figs.~\\ref{fig:H-Omega_M}--\\ref{fig:H-Omega_ML_comb})."
    },
    "9603/astro-ph9603034_arXiv.txt": {
        "abstract": "Interferometric gravitational wave detectors could measure the frequency sweep of a binary inspiral [characterized by its chirp mass] to high accuracy. The observed chirp mass is the intrinsic chirp mass of the binary source multiplied by $(1+z)$, where $z$ is the redshift of the source. Assuming a non-zero cosmological constant, we compute the expected redshift distribution of observed events for an advanced LIGO detector. We find that the redshift distribution has a robust and sizable dependence on the cosmological constant; the data from advanced LIGO detectors could provide an independent measurement of the cosmological constant. ",
        "introduction": "A non-zero cosmological constant may help solve some of the current observational puzzles, most notably, the conflict between the age of globular clusters and the apparent high value of the Hubble constant [which suggests a younger Universe]. \\cite{Lambda} A sizable cosmological constant can make the Universe old in spite of a high Hubble constant, although a non-zero cosmological constant is ugly from the theoretical viewpoint. Whatever our aesthetic preferences, the value of the cosmological constant should ultimately be determined by observational measurements.  Advanced LIGO detectors can expect to observe approximately 50 neutron star binary inspiral events per year, from distances up to 2000$\\,$Mpc, the accuracy in the measurement of the signal strength can be better than 10\\%, and the accuracy in the measurement of the chirp mass [which characterizes the frequency sweep of a binary inspiral] can be better than 0.1\\%. \\cite{FinnCher93}\\cite{CutFla94}. The cosmological implications of gravitational wave observations of binary inspiral have been discussed by several authors \\cite{Schutz86,CherFinn93,Marko93}. Most recently, Finn pointed out that the observations of binary inspirals in an interferometric gravitational wave detector, in terms of the distribution of observed events with signal strength and chirp mass, can be quite sensitive to cosmology \\cite{Finn96}.   Previous discussions of the cosmological implications of gravitational wave observations have considered the measurements of the Hubble constant $H_0$ [$H_0=100h\\,$km$\\,$sec$^{-1}$Mpc$^{-1}$, $0.5 \\leq h <1$] and the deceleration parameter $q_0$, assuming the cosmological constant to be zero. Even though Markovi\\'{c} discussed the measurement of the cosmological constant, he assumed its true value to be {\\em zero}. \\cite{Marko93} In this paper, we consider the measurement of a {\\it non-zero} cosmological constant $\\Omega_{\\Lambda}$, and find that the $\\Omega_{\\Lambda}$ dependence of the observed chirp mass spectrum is more robust than its $H_0$ dependence.  Markovi\\'{c} proposed measuring the cosmological parameters using the observed distribution of measured luminosity distances and (estimated) redshifts \\cite{Marko93}, while Chernoff and Finn suggested an alternative method without using the measured luminosity distances \\cite{CherFinn93}. We have chosen to follow the method and notations of Refs.{\\cite{CherFinn93,Finn96}}. ",
        "conclusions": "In summary, we have calculated the expected maximum source redshift $z_{\\rm max}$, the source redshift distribution $P(z,>\\rho_0)$, the signal-to-noise ratio distribution $P(\\rho,>\\rho_0)$, and the total number of events per year $\\dot{N}(\\rho>\\rho_0)$, for advanced LIGO detectors in a Universe with nonzero cosmological constant. $z_{\\rm max}$, $P(z,>\\rho_0)$, and $P(\\rho,>\\rho_0)$ all depend on $\\Omega_{\\Lambda}$ and $\\Omega_0$ in a fundamental way through the angular diameter distance, and they all depend on $h$ through the combination $h\\,A(r_0,\\rho_0, {\\cal M}_0)$. $\\dot{N}(\\rho>\\rho_0)$ is very sensitive to the local binary merger rate $\\dot{n}_0$ through $\\dot{n}_0\\,\\left(cH_0^{-1}\\right)^3$, the value of which is quite uncertain at this time.  The expected redshift distribution of observed events in an advanced LIGO detector has a robust and sizable dependence on the cosmological constant. Although the redshift distribution has an apparent dependence on $h$ which is more dominant, this dependence on $h$ is superficial in the sense that it always appears in the combination of $h\\, A(r_0,\\rho_0, {\\cal M}_0)$ [$A$ is given by Eq.(\\ref{eq:A})]; increasing $h$ has exactly the same effect on the redshift distribution as increasing $r_0$ or ${\\cal M}_0^{5/6}$, or decreasing $\\rho_0$, by the same amount. On the other hand, the redshift distribution depends on $\\Omega_{\\Lambda}$ and $\\Omega_0$ in a fundamental way, this dependence is detector and source independent.  If we live in a flat Universe, then the cosmological constant can be determined quite accurately from the expected redshift distribution of observed events with a cut on the signal-to-noise ratio. Assuming arbitrary geometry of the Universe, the expected redshift distribution of observed events with a cut on the signal-to-noise ratio may correspond to a family of related cosmological models with different values of $\\Omega_{\\Lambda}$ and $\\Omega_0$ [the combination $\\alpha \\equiv \\Omega_0 (1+z_{max})^2 - \\Omega_{\\Lambda} (z_{max}+2)$ is easily and accurately measured]; this degeneracy can be lifted by either accurately measuring the distribution of observed events in signal-to-noise ratio, or by increasing the detector sensitivity. However, it may well prove more practical to use other astrophysical constraints on $\\Omega_{\\Lambda}$ and $\\Omega_0$ to resolve this degeneracy. The data from advanced LIGO detectors should provide an independent and robust measurement of the cosmological constant."
    },
    "9603/hep-ph9603392_arXiv.txt": {
        "abstract": "We consider the possibility that neutrinos are coupled very weakly to an extremely light scalar boson.  We first analyze the simple problem of one generation of neutrino and show that, for ranges of parameters that are allowed by existing data, such a system can have serious consequences for the evolution of stars and could impact precision laboratory measurements.  We discuss the extension to more generations and show that the general conclusion remains viable.  Finally, we note that, should such a scalar field be present, experiments give information about effective masses, not the masses that arise in unified field theories. ",
        "introduction": "In the last few years, it has been suggested that neutrinos might interact weakly among themselves through the exchange of a very light scalar particle~\\cite{MKY,TRE}, with possible consequences for the evolution of the Universe and for the propagation of neutrinos from distant events.  In many of these discussions, one assumes that neutrinos are distributed according to the usual Big Bang scenario and Standard Model physics, the effects of scalar exchange being treated as a perturbation.  In this paper we examine that assumption.  This problem is a special case of the general problem of relativistic fermions interacting through the exchange of scalar and vector bosons, and the general formalism has been worked out under the name Quantum Hadrodynamics (QHD) and applied to the study of nuclear physics~\\cite{QHD}.  We shall show that, for a wide range of parameters, neutrinos will tend to cluster, that these clusters could be of a size to affect stellar formation and dynamics and that there may be observable consequences for physics within the solar system.  In fact, if the clustering is strong enough so that the density of neutrinos within a cloud is sufficiently large, terrestrial laboratory experiments can be affected.  Over the last decade, many groups studying the endpoint of the Tritium beta ray spectrum for signs of neutrino mass have reported a best fit value for the square of the anti-neutrino mass less than zero~\\cite{TRIT1,TRIT2,TRIT3,TRIT4,TRIT5,TRIT6}.  Robertson et.al~\\cite{TRIT1} point out that this result could be obtained by assuming a very small branch due to the absorption of relic neutrinos, provided the density of such neutrinos were some $10^{13}$ higher than usual cosmological values.  In choosing parameter ranges for the examples in this paper, we have kept this idea in mind.  We must emphasize, however, that the general feature of cloud formation and its consequence for the evolution of structure in the history of the Universe is quite robust, and must be considered whatever the eventual resolution of the anomaly in Tritium beta decay.  In this work, we only consider the effects of light scalar exchange. The effects of the known vector exchange ($Z_0$) are far too small (really, too short ranged) to affect these results.  The exchange of a light vector particle is severely constrained by data.  To avoid problems with the axial anomaly (the neutrinos do couple to the $Z_0$) one must either invent many new fermions or demand that the light boson couple to known leptons or quarks, which quickly leads to conflict with experiments designed to test for a fifth force~\\cite{EA}.  Furthermore, the self shielding of a vector exchange, while it might allow for the development of a neutrino-antineutrino plasma, will not allow for the coherent action required to drive cloud formation.  Thus, we only treat scalar exchange.  Even so, there remains the rich possibilities of different couplings to different generations. Aside from a few comments, we leave that to further work, concentrating our discussion on the simpler system of one flavor of neutrino.  The paper is organized as follows.  In section 2 we review the treatment of infinite matter in QHD and apply that to the problem at hand.  In section 3 we do the same for finite clouds of neutrinos.  In section 4, we describe the consequences such clouds would have on the the evolution of structures in the Universe and, in section 5, we confront this picture with what data exists, extracting limits on the parameters.  In section 6 we discuss the changes in the analysis of Tritium beta decay experiments within such neutrino clouds, and comment on the effects of such clouds on stellar dynamics in section 7.  In section 8 we discuss some aspects of the extension to include more than one generation, illustrating the remarks with special cases applied to two generations.  We offer our conclusions in section 9. ",
        "conclusions": "We have applied the techniques of Quantum Hadrodynamics to the study of a system of neutrinos interacting through a light, weakly coupled scalar boson.  We have shown that, for a wide range of parameters, neutrinos will tend to condense into clouds, with dimensions the scale of the  inverse boson mass.  In fact, for parameters which cause no conflict with laboratory measurements, such clouds could easily be the right size and density to affect experiments on and around the earth.  We have shown that it is likely that any such condensation would have occurred before recombination and that the formation of neutrino clouds could form a natural seeding mechanism for the formation of hadronic objects on the scale of stars.  Neutrino cloud formation, being a phase change, occurs very quickly, so these seeds are available at the earliest possible epoch for star formation.  The extension of this work to more than one neutrino flavor depends on the mass hierarchies of the neutrinos and the scalars as well as the details of the coupling of each scalar to different generations.  For the case of two generations of neutrinos and one scalar field, we have looked at two simple choices for the couplings.  While different in detail, both generate concentric spherical (assumed) distributions with the lighter neutrinos (as determined by vacuum mass) extending farther out.  We would expect this general feature to survive for three generations, raising the possibility of the heaviest species being essentially within the star, the other two occupying different regions of space out to a distance, depending on the detailed history of the system, of a fraction of a parsec.  If the density of the electron component of the neutrinos and antineutrinos around the Sun is high enough, there could be observable effects on very sensitive experiments such as the study of Tritium beta decay to search for antineutrino mass effects or double beta decay measurements seeking evidence that neutrinos are Majorana particles.  One consequence of the existence of such an interaction would be that all such measurements would have to be interpreted in terms of effective masses, rather than the vacuum masses that are relevant to model building.  Whether terrestrial effects are observed or not, evidence for or against the existence of such a scalar interaction is most likely to come from astronomy and astrophysics.  The implications of the existence of neutrino clouds, with respect to both the time scale and the distribution of sizes, should be amenable to testing through modelling and observation.  The gravitational effects within our own Solar system, while subtle, could be observable in very high accuracy satellite tracking data.  Depending on the precise model for several generations, one may be able to observe the modifications of oscillation and propagation in the extremely dense neutrino fluxes associated with supernovae.  Whatever the experimental outcome of such tests may be, this problem remains as a fascinating extension of the theoretical techniques of QHD, developed for the study of atomic nuclei, to vastly different regions of parameter space.  This work was supported in part by the US Department of Energy, the National Science Foundation and the Australian ARC.  Some of the work was carried out during visits to the Institute for Nuclear Theory at the University of Washington; their hospitality is gratefully acknowledged.  One of us (GJS) particularly wants to thank C.  Horowitz for several very useful conversations.  \\newpage"
    },
    "9603/astro-ph9603020_arXiv.txt": {
        "abstract": "The Hubble Deep Field\\footnote{Based on observations with the NASA/ESA \\HST, obtained at the Space Telescope Science Institute, which is operated by AURA, under NASA contract NAS 5-26555} (HDF) offers the best view to date of the optical sky at faint magnitudes and small angular scales.  Early reports suggested that faint source counts continue to rise to the completeness limit of the data, implying a very large number of galaxies.  In this {\\it letter,} we use the two-point angular correlation function and number-magnitude relation of sources within the HDF in order to assess their nature.  We find that the correlation peaks between $0.25\\arcsec$ and $0.4\\arcsec$ with amplitude of 2 or greater, and much more for the smallest objects.  This angular scale corresponds to physical scales of order $1 \\kpc$ for redshifts $z \\ga 1$.  The correlation must therefore derive from objects with subgalaxian separations. At faint magnitudes, the counts satisfy the relation $\\mbox{Number} \\propto 1/\\mbox{flux}$, expected for images which are subdivisions of larger ones.  Several explanations for these observed correlations are possible, but a conservative explanation can suffice to produce our results.  Since high redshift space ($z \\ga 0.5$) dominates the volume of the sample, observational redshift effects are important.  Rest-frame ultraviolet radiation appears in the HDF's visible and near-UV bands, and surface brightness dimming enhances the relative brightness of unresolved objects versus resolved objects.  Both work to increase the prominence of compact star-forming regions over diffuse stellar populations.  Thus, a ``normal'' gas-rich galaxy at high redshift can appear clumpy and asymmetric in the visible bands.  For sufficiently faint and distant objects, the compact star-forming regions in such galaxies peak above undetectable diffuse stellar backgrounds.  Our results do not exclude asymmetric formation or fragmentation scenarios. ",
        "introduction": "The Hubble Deep Field (Williams \\etal 1996) affords us an unprecedented view of the optical sky at small angular scales and faint flux levels.  It thus allows us to study faint (and presumably high redshift) galaxies without complicating field crowding effects caused by comparatively poor seeing in ground-based faint galaxy studies (\\cf Tyson 1995).  Preliminary results (Giavalisco \\etal 1996) show that source counts in the HDF continue to rise as a power law to the completeness limit of the data.  Such an effect may be due to ever larger numbers of galaxies at fainter flux levels. However, it may also be due to the increasingly clumpy appearance of galaxies at high redshift, which can confuse source detection algorithms into counting parts of each physically distinct galaxy as several faint sources.  To test this possibility, we consider how redshift effects can lead to over-counting of whole sources in a deep field like the HDF.  $K$-correction and surface-brightness dimming tend to enhance the relative prominence UV bright and compact objects, such as active star-forming regions (O'Connell \\& Marcum 1996).  If the enhancement is sufficient, several star-forming regions occurring in a single galaxy will produce the appearance of several small sources separated by an angular scale comparable to the size of a normal galaxy.  In the later sections of the paper, we use two different statistics to test the extent to which HDF source counts reflect the subdivision of galaxies.  These tests exploit the weak dependence of angular size on redshift $z$ at $z \\ga 1$. Since galaxies of present day size ($10\\kpc$) remain resolved at all redshifts in the HDF (Peebles 1993), we can compare the physical separations and sizes of objects in the HDF to those of galaxies in the low-redshift universe.  In section~4, we discuss the two-point angular correlation function $w(\\theta)$ of HDF sources.  Considerable correlations may be expected for physical scales $\\la 10 \\kpc$ if many galaxies in the field break up into multiple giant H\\ii regions in the source catalogs.  In section~5, we present number-magnitude relations derived from our source catalogs, which show a smooth increase to the completeness limit, with a flatter faint-end slope than in deep ground-based images, and a rough relation $N \\propto 1/\\mbox{flux}$, consistent with the hypothesis that many of the faintest images are parts of larger objects. ",
        "conclusions": "We have cataloged objects in the Hubble Deep Field in a way that is less prone to spurious detections than were previous efforts.  From the catalog, we have drawn the angular correlation for high and low color-redshift subsets using two different cuts.  We have found similar correlations down to $0.5\\arcsec$ for both subsets.  Since the signal is ostensibly dominated by H\\ii regions in the lower redshift subset, we surmise that that signal is also dominated by subgalaxian structure in the higher redshift subset.  We have compared our results to the correlation derived from an independent catalog (Couch 1996) which appears not to over-count nearby objects. As expected, there is less correlation in the lower redshift subset of this catalog than in ours. However, at higher redshift, the correlations of the catalogs agree rather well, so that both of our catalogs include as distinct objects what are likely subgalaxian structure at high redshift.  A dramatic increase in sub-arcsecond correlation occurs in the subset of objects with the smallest angular sizes, in agreement that the correlated objects are small, subgalaxian objects.  The qualitative difference between deep space- and ground-based optical data is due to a conspiracy of scales.  Because the characteristic angular sizes of galaxies at redshifts $1 \\la z \\la 5$ correspond to the $\\sim 1''$ angular resolution of ground-based data, deep optical counts from the ground will see galaxy-sized objects as single peaks.  At the higher resolution available from space, substructure becomes detectable in galaxies at any redshift, and overcounting becomes a possibility.  We have also computed the magnitude-radius relation, which shows that a large fraction of the objects have characteristic sizes around $0.15\\arcsec$, corresponding to scale lengths $\\sim 1 \\kpc$, typical of both high redshift galaxian scale-lengths and diameters of giant star-forming regions.  The peak at $0.15\\arcsec$ allows several objects to fit into a single galaxy, as one requires for the subgalaxian structure scenario.  The number-magnitude relations for our catalogs show convergent flux in all bands with $N \\propto 1/\\mbox{flux}$ as expected for images broken into fragments.  This physically reassuring result differs from the na\\\"\\i{}ve extrapolation of ground-based number-magnitude relations for U, B, and possibly R bands.  We find that after we smoothed the data, the counts drop dramatically at the faint end.  This illustrates how seeing reduces the faint counts in ground-based work, diluting isolated faint objects below detection thresholds while blurring substructures in brighter galaxies together to form single peaks.  The statistical tests presented herein suggest that the most distant objects in the HDF must be some combination of galaxies and star-forming fragments, a distinction increasingly hard to draw in deep fields.  This supports our hypothesis that ultraviolet bright and compact star-forming regions contribute substantially to the flux, and increasingly to the number counts, we receive from high redshift samples."
    },
    "9603/astro-ph9603093_arXiv.txt": {
        "abstract": "We describe a technique that uses radial color gradients in disk galaxies to detect the presence of bulk non-circular motion or elliptical orbits. In a disk galaxy with both a radial color gradient and non-circular motion, isochromes or iso-color contours should follow the shape of closed stellar orbits, and the ellipticity of the isophotes should vary as a function of wavelength. A difference in the ellipticity of isochromes and the isophotes can be used to detect the presence of non-circular motion. A model galaxy is constructed which demonstrates this phenomenon. The difference between isochrome and isophote ellipticity is directly related to the ellipticity of the potential. This provides a new way to measure the ellipticity of the dark matter in the outer parts of galaxies.  As an example, we apply this technique to two dwarf galaxies NGC 1800 and NGC 7764. We detect a bar in NGC 1800 which has only previously been suggested from the HI velocity field. In NGC 7764 there is no color gradient along its bar so we cannot detect non-circular motion in this region; however ellipticities observed in a star forming ring at the end of the bar are consistent with this ring being located near the corotation resonance. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603103_arXiv.txt": {
        "abstract": "We present our analysis of {\\sl{Hubble Space Telescope}} Wide Field Planetary Camera 2 observations in F555W (broadband $V$) and F450W (broadband $B$) of the globular cluster Hodge 11 in the Large Magellanic Cloud galaxy. The resulting $V$ vs.\\ $\\bmv$ color-magnitude diagram reaches $2.4$ mag below the main-sequence turnoff (which is at $V_{\\rm{TO}}=22.65\\pm0.10$ mag or $M_V^{\\rm TO}=4.00\\pm0.16$ mag). Comparing the fiducial sequence of Hodge 11 with that of the Galactic globular cluster M92, we conclude that, within the accuracy of our photometry, the age of Hodge 11 is identical to that of M92 with a relative age-difference uncertainty ranging from 10\\% to 21\\%. Provided that Hodge 11 has always been a part of the Large Magellanic Cloud and was not stripped from the halo of the Milky Way or absorbed from a cannibalized dwarf spheroidal galaxy, then the oldest stars in the Large Magellanic Clouds and the Milky Way appear to have the same age. ",
        "introduction": "The oldest stars for which reliable ages can be determined are found in globular clusters.  Of the various methods used to infer ages in these systems, the absolute magnitude of the main-sequence turnoff is considered to be the most secure measurement. The main-sequence turnoff point for old stars occurs at $M_V\\approx +4.0$ mag which makes it difficult to observe the old turnoff population in relatively nearby galaxies such as Andromeda (M31) even with the refurbished {\\em{Hubble Space Telescope}}. The ages of the oldest stars in a cluster of galaxies can be used to establish the chronology of galaxy formation within that cluster. The age-spread of the oldest stars provides an important cosmological probe for the investigation of synchronized galaxy formation. Current technology, unfortunately, allows us to conduct this experiment only with the Milky Way and its relatively nearby companions.  Efforts have been made to infer the ages of star clusters in the Magellanic Clouds from their integrated photometry. The Large Magellanic Cloud (LMC) star cluster Hodge 11 (Hodge 1960) was classified by Searle, Wilkinson, \\& Bagnuolo (1980) as being SWB class VII and suggested that Hodge 11 is similar to the old metal-poor Galactic halo globular clusters. This was confirmed by Elson \\& Fall (1988) and Girardi \\et (1995) whose new UBV cluster photometry places Hodge 11 clearly among the oldest Galactic globular clusters. The LMC globular cluster NGC 2257 was recently thought to be similar to Hodge 11, however Testa \\et (1995) has determined that NGC 2257 is 2--3 Gyr younger than the oldest Galactic globular clusters. Therefore, Hodge 11 is especially interesting because it might be the oldest globular cluster in the Magellanic Clouds.  This paper presents the first results of our {\\em{Hubble Space Telescope}} snapshot observation program of star clusters in the Magellanic Clouds using the Wide Field Planetary Camera 2 instrument. Our sample of star clusters was chosen to cover the full age range available in the Clouds and we have surveyed 46 star clusters using $\\sim$15 hours of spacecraft time. While the principal aim of our observational program was to investigate the global properties of star clusters in the Magellanic Clouds, we now turn to our results on the Large Magellanic Cloud globular cluster Hodge 11. ",
        "conclusions": ""
    },
    "9603/astro-ph9603098_arXiv.txt": {
        "abstract": "Quasi-coherent oscillations have been detected in the extreme ultraviolet flux of the dwarf nova SS~Cygni during observations with the {\\it Extreme Ultraviolet Explorer\\/} satellite of the rise and plateau phases of an anomalous outburst in 1993 August and a normal outburst in 1994 June/July. On both occasions, the oscillation turned on during the rise to outburst and persisted throughout the observation. During the 1993 outburst, the period of the oscillation fell from 9.3~s to 7.5~s over an interval of 4.4 days; during the 1994 outburst, the period fell from 8.9~s to 7.19~s (the shortest period ever observed in SS~Cyg, or any other dwarf nova) within less than a day, and then rose to 8.0~s over an interval of 8.0 days. For both outbursts, the period $P$ of the oscillation was observed to correlate with the 75--120~\\AA \\ count rate $I_{\\rm EUV}$ according to $P\\propto I_{\\rm EUV}^{-0.094}$. A magnetospheric model is considered to reproduce this variation. It is found that an effective high-order multipole field is required, and that the field strength at the surface of the white dwarf is 0.1--1~MG. Such a field strength is at the lower extreme of those measured or inferred for bona~fide magnetic cataclysmic variables. ",
        "introduction": "Rapid periodic oscillations are observed in the optical flux of high accretion rate (``high-$\\Mdot $'') cataclysmic variables (CVs; nova-like variables and dwarf novae in outburst) (\\cite{pat81}; \\cite{war95}; \\cite{war96}) These oscillations have high coherence ($Q\\approx 10^4$--$10^6$), periods $P\\approx 10$--30~s, amplitudes of less than 0.5\\%, and are sinusoidal to within the limits of measurement. They are referred to as ``dwarf nova oscillations'' (DNOs) to distinguish them from the apparently distinct longer period, low coherence ($Q\\approx 1$--$10$) quasi-periodic oscillations (QPOs) of high-$\\Mdot$ CVs, and the longer period, high coherence ($Q\\approx 10^{10}$--$10^{12}$) oscillations of DQ~Her stars. DNOs have never been detected in dwarf novae in quiescence, despite extensive searches; they appear on the rising branch of the dwarf nova outburst, typically persist through maximum, and disappear on the declining branch of the outburst. The period of the oscillation decreases on the rising branch and increases on the declining branch, but because the period reaches a minimum about one day after maximum optical flux, dwarf novae describe a loop in a plot of period versus optical flux.  The dwarf nova SS~Cygni routinely exhibits DNOs during outburst. Optical oscillations have been detected at various times with periods ranging from 8.2~s to 10.9~s (\\cite{pat78}; \\cite{hor80}; \\cite{pat81}). During one outburst, the period was observed over a interval of $\\approx 6$ days to fall from 7.5~s to 7.3~s and then rise to 8.5~s (\\cite{hil81}). At soft X-ray energies ($E\\approx 0.1$--0.5 keV), oscillations have been detected in \\hbox{{\\it HEAO 1\\/}} LED~1 data at periods of $\\approx 9$~s and 11~s (\\cite{cor80}; \\cite{cor84}) and in {\\it EXOSAT\\/} LE data at periods between 7.4~s and 10.4~s (\\cite{jon92}). In this Letter, we describe observations with the {\\it Extreme Ultraviolet Explorer\\/} satellite ({\\it EUVE\\/}; \\cite{bow91}; \\cite{bow94}) of the EUV oscillations of SS~Cyg. ",
        "conclusions": "If a magnetospheric model applies to SS~Cyg, it appears likely that the surface magnetic field strength of its white dwarf is 0.1--1~MG. This range of values is orders of magnitude lower than the field strengths of AM~Her stars ($B \\approx 10$--80~MG; \\cite{cro90}; \\cite{beu96}), but may overlap with the field strengths of DQ~Her stars (very uncertain, by $B\\sim 0.1$--10~MG; \\cite{pat94}). With such a field strength, it is a challenge for SS~Cyg in quiescence not to manifest the photometric variations associated with DQ~Her stars; the limit on such variations in the optical is 0.001 mag ($\\approx 0.01\\%$) near 0.1~Hz (\\cite{pat95}). In outburst, the accretion flow should be channeled down to the footpoints of the magnetic field and produce hard and soft X-rays in the manner of AM~Her and DQ~Her stars. However, neither the eponymous DQ~Her nor V533~Her are hard X-ray sources (\\cite{cor81}), demonstrating that magnetic accretion can take place without the production of hard X-rays. While SS~Cyg in outburst is a {\\it known\\/} hard X-ray source (\\cite{jon92}; \\cite{nou94}; \\cite{pon95}), it is {\\it not\\/} known to oscillate in hard X-rays; the limit is 6\\% when soft X-ray oscillations were observed by {\\it HEAO 1\\/} (\\cite{swa79}). To produce a more stringent upper limit, {\\it XTE\\/} observations are required. Planned simultaneous observations of SS~Cyg in outburst with {\\it XTE\\/} and {\\it EUVE\\/} will determine the shapes of the hard and soft X-ray light curves, the correlations between the soft and hard X-ray fluxes, and the extent to which the oscillations in the soft X-ray flux are manifest in the hard X-ray flux."
    },
    "9603/astro-ph9603117_arXiv.txt": {
        "abstract": "We study the  time evolution and gravitational wave emission properties of a black hole orbiting {\\it inside} an accretion disk surrounding a massive black hole. We  simultaneously solve the structure equations of the accretion disk in the presence of heating, cooling, and viscosity as well as the equations governing the companion orbit. The deviation from the Keplerian distribution of the angular momentum of the disk due to pressure and advection effects causes a significant exchange of angular momentum between the disk and the companion. This significantly affects the gravitational wave emission properties from the binary system. We show that when the companion is light, the effect is extremely important and must be taken into account while interpreting gravitational wave signals from such systems. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603070_arXiv.txt": {
        "abstract": "Standard big bang nucleosynthesis (BBNS) promises accurate predictions of the primordial abundances of deuterium, helium-3, helium-4 and lithium-7 as a function of a single parameter.  Previous measurements have nearly always been interpreted as confirmation of the model (\\cite{cop95}). Here we present a measurement of the deuterium to hydrogen ratio (D/H) in a newly discovered high redshift metal-poor gas cloud at redshift $z=2.504$. This confirms our earlier measurement of D/H (\\cite{tyt96}), and together they give the first accurate measurement of the primordial D abundance, and a ten-fold improvement in the accuracy of the cosmological density of ordinary matter. This is a high density, with most ordinary matter unaccounted or dark, which is too high to agree with measurements of the primordial abundances of helium-4 and lithium-7. Since the D/H measurement is apparently simple, direct, accurate and highly sensitive, we propose that helium requires a systematic correction, and that population II stars have less than the primordial abundance of $^7$Li. Alternatively, there is no concordance between the light element abundances, and the simple model of the big bang must be incomplete and lacking physics, or wrong. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9603/astro-ph9603136_arXiv.txt": {
        "abstract": "We have used a laser heterodyne spectrometer to resolve the emission line profile of the 63 \\micron\\ \\threepone\\ fine-structure transition of \\oi at two locations in M42. Comparison of the peak antenna temperature with that of the 158 \\micron\\ \\cii fine-structure line shows that the gas kinetic temperature in the photodissociation region near \\tonec\\ is 175 - 220 K, the density is greater than ${\\rm 2 \\times 10 ^{\\rm 5}}$\\ \\cmthree , and the hydrogen column density is about ${\\rm 1.5 \\times 10 ^{\\rm 22}}$\\ \\cmtwo . A somewhat lower temperature and column density are found in the IRc2 region, most likely reflecting the smaller UV flux. The observed width of the \\oi line is 6.8 \\kms\\ (FWHM) at \\tonec , which is slightly broadened over the intrinsic linewidth by optical depth effects. No significant other differences between the \\oi and \\cii line profiles are seen, which shows that the narrow emission from both neutral atomic oxygen and ionized carbon comes from the PDR.  The \\oi data do not rule out the possibility of weak broad-velocity emission from shock-excited gas at IRc2, but the \\cii data show no such effect, as expected from non-ionizing shock models. ",
        "introduction": "Photodissociation regions (PDR) occur at the interface between \\ion{H}{2}\\ regions and cooler molecular material.  They are most easily studied via the fine-structure line emission from C$^{\\rm +}$ and O and the rotational lines of the abundant CO molecule. Of particular interest are the fine-structure lines of \\oi at 63 \\micron\\ and \\cii at 158 \\micron , since these are believed to be the major cooling lines, with \\oi emission usually dominant at densities typical of dense molecular clouds (${\\rm n \\, > \\, 10 ^{\\rm 4}}$ \\cmthree ) and \\cii emission more significant at lower densities (\\cite{th85a}). Since the ionization potential of O is very similar to that of H, essentially all the oxygen in a PDR will be neutral.  Similarly, since the first ionization potential of C is similar to the dissociation energy of CO while its second ionization potential is close to that of He, carbon throughout the PDR will be singly ionized.  Observations of both major fine-structure cooling lines from the same region provide information on the temperature, density, and relative abundance of the species within the PDR.  Both the 63 \\micron\\ \\threepone\\ \\oi and the 158 \\micron\\ $^{2}P_{\\case{3}{2}}$\\, -\\,$^{2}P_{\\case{1}{2}}$ \\cii lines have been previously observed and mapped in Orion (e.g. \\cite{mgh79,swt79,wer84,cra86,sta93}). While comparison of integrated intensities and an assumed linewidth for both transitions allowed these authors to make the first estimates of physical conditions, a definitive determination was hampered by the inability to resolve line profiles.  The 5 \\kms\\ (FWHM) linewidth of the \\cii line near \\tonec\\ in M42 was first measured by Boreiko, Betz, and Zmuidzinas (1988), who examined the dynamics of the photoionized gas. This Letter describes observations of the 63 \\micron\\ \\threepone\\ line of \\oi in the Orion region with 0.2 \\kms\\ resolution, which gives the first fully resolved line profiles for M42 at this wavelength. Comparison with 158 \\micron\\ \\cii line profiles at similarly high resolution allows the determination of the optical depth in both these lines, and hence the temperature and limits to the density within the PDR. ",
        "conclusions": ""
    },
    "9603/astro-ph9603158_arXiv.txt": {
        "abstract": "\\rightskip = 0.0in plus 1em We show that Earth mass planets orbiting stars in the Galactic disk and bulge can be detected by monitoring microlensed stars in the Galactic bulge. The star and its planet act as a binary lens which generates a lightcurve which can differ substantially from the lightcurve due only to the star itself. We show that the planetary signal remains detectable for planetary masses as small as an Earth mass when realistic source star sizes are included in the lightcurve calculation. These planets are detectable if they reside in the ``lensing zone\" which is centered between 1 and 4 AU from the lensing star and spans about a factor of 2 in distance.  If we require a minimum deviation of 4\\% from the standard point-lens microlensing lightcurve, then we find that more than 2\\% of all $\\mearth$ planets and 10\\% of all $10\\mearth$ in the lensing zone can be detected. If a third of all lenses have no planets, a third have $1\\mearth$ planets and the remaining third have $10\\mearth$ planets then we estimate that an aggressive ground based microlensing planet search program could find one earth mass planet and half a dozen $10\\mearth$ planets per year. ",
        "introduction": "\\label{sec-intro} The recent discovery of several giant planets (\\cite{swiss51peg}, \\cite{marbut}) has confirmed the existence of planets orbiting main sequence stars other than the Sun. Two of these first 3 giant planets have orbits that were unexpected, and this together with the surprising discovery of planets in a pulsar system (\\cite{wol}) demonstrates the importance of observational studies of extra-solar planetary systems. Indirect ground based techniques which detect the reflex motion of the parent star through accurate radial velocity measurements or astrometry are likely to have sensitivity that extends down to the mass of Saturn ($\\sim 100\\mearth$) (\\cite{butetal}) or even down to $10\\mearth$ (\\cite{shao}) with interferometry from Keck or the VLT. There is great interest in searching for planets with a masses similar to that of the Earth, and NASA's new ExNPS program (\\cite{ExNPS}) seeks to build a spacecraft capable of imaging nearby Earth mass planets in the infrared. In order to ensure the success of such a mission, we will need to have at least a rough idea of how prevalent planets with masses close to that of the Earth really are.  A ground based gravitational microlensing survey system sensitive to planets down to $1\\,\\mearth$ has been proposed by Tytler (1995). This project would involve both a microlensing survey telescope to detect microlensing events in progress and a world-wide network of follow-up telescopes that would monitor the microlensing lightcurves on a $\\sim\\,$hourly timescale in search of deviations due to planets. Existing microlensing surveys (\\cite{macho-nat}, \\cite{eros-nat}, \\cite{ogle1}, and \\cite{duo}) have recently demonstrated real time microlensing detection capability (\\cite{macho-alert}, \\cite{ogle-ews}), and two world-wide microlensing follow-up collaborations (\\cite{planet} and \\cite{macho-gman}) are now in operation, but to detect Earth mass planets, more capable survey and follow-up systems will be required.  In this paper, we provide the theoretical basis for this enterprise by calculating realistic microlensing lightcurves and detection probabilities for planets as small as $1\\,\\mearth$. Previous authors (\\cite{mao-pac}, \\cite{gould-loeb}, and \\cite{bolatto}) have considered the deviations from the single lens lightcurve due to planets using the point source approximation. This is a poor approximation for planets in the 1-10$\\,\\mearth$ mass range, so we have calculated planetary-binary lensing event lightcurves for realistic finite size source stars, and we show that planets in the 1-10$\\,\\mearth$ mass range can cause deviations from the standard single lens lightcurve with amplitudes larger than 10\\% which last for a couple hours or more. We calculate planetary detection probabilities based upon a set of assumed event detection criteria and a simple planetary system model loosely based upon the solar system. ",
        "conclusions": "\\label{sec-char}  In the previous section, we have shown that a significant number of planets with masses down to $1\\mearth$ can be detected via gravitational microlensing if microlensing events towards the Galactic bulge are monitored $\\sim$ hourly with photometric precision of 0.5-1.0\\% which is readily achievable in crowded stellar images.  We can also use the results of our probability calculations to help determine the optimal planetary search strategy. For example, given a large number of events to monitor for planetary deviations and a limited amount of observing time, how long should we follow each event? The probabilities given in table~\\ref{tab-prob} assume that each event is followed from event detection at $A=1.58$ until $A$ drops to $1.13$, but if we stop the follow-up observations when $A>1.34$, then we will only be sensitive to planetary deviations from planets in the interval $0.62 < \\ell < 1.62$. Integration over the curves in figure~\\ref{fig-prob} indicates that this will reduce the chance of detecting a planet by 5-10\\% (for the ``factor-of-2\" model), but the total number of observations required drops by 27\\%. Thus, if the capacity of the follow-up system is saturated, it is best to concentrate follow-up observations on events with $A>1.34$. This effect is basically geometric: planets that are outside the lensing zone ($\\ell >1.6$) tend to give rise to ``isolated\" events that aren't associated with a stellar lensing event detected by the survey system\\rlap.\\footnote[3]{Isolated planetary lensing events might be detected by microlensing surveys, but the detection efficiency and variable star background rejection would be quite poor.} It is optimal to search for planets at $\\ell < 1.6$ where they would ``modulate\" a detectable stellar lensing lightcurve.  Our results also suggest that it will be easier to detect Earth mass planets by monitoring turn-off star lensing events than giant star events. (Gould and Welch (1995) have shown that combined infrared and optical observations may allow the detection of earth mass planets in giant star lensing events, however.)  We've established that low mass planets can be detected, but we should also address what can we learn about each planet that is discovered through microlensing. Planetary lightcurve deviations would be detected in real time so that observations can be repeated every few minutes during the planetary deviation. The lens parameters $\\ell$ (the separation perpendicular to the line of sight in units of $R_E$) and the mass ratio $\\epsilon$ can generally be determined from gross features of the lightcurve. $\\ell$ is easily determined (up to a 2-fold ambiguity) from the amplification that the unperturbed lightcurve would have in the deviation region, and the mass ratio $\\epsilon$ can be determined from the timescale of the planetary deviation. The 2-fold ambiguity in $\\ell$ is also easily resolved in most cases by the shape of the lightcurve deviation as can be seen in figures~\\ref{fig-lcs} and \\ref{plate}. For $\\ell < 1$, the the deviation region consists of positive deviation regions surrounding the two caustic curves with a long trench of negative deviations in between. This leads to light curves with regions of large negative perturbations surrounded by regions of smaller positive perturbations. For $\\ell > 1$, the situation is reversed and the dominant perturbation is a central positive one which has regions of small negative perturbations on either side of it.  Another parameter that may be measured is the angular Einstein ring radius of the planet itself. This comes about because the ratio of the this radius to the angular radius of the star is the parameter which describes the finite source effects. For planets of Earth mass, the finite source effects are almost always important, so in principle, this parameter may be measurable in most events.  In summary, we have calculated realistic lightcurves for microlensing events where the lens star has a low mass planetary companion, and we have shown that planets with masses as small as $1 \\mearth$ can be detected via gravitational microlensing. Thus, gravitational microlensing is the only ground based method that has been shown to be sensitive to Earth mass planets."
    },
    "9603/astro-ph9603122_arXiv.txt": {
        "abstract": "A new solution is presented for the puzzling, observed universality of the exponential luminosity profiles, perpendicular to the disk plane of spiral and lenticular galaxies. It is shown that such exponential $z$-profiles result naturally from gaseous protodisks which settle into isothermal equilibrium prior to star formation. Subsequent cooling leads to a gravitational contraction of the gas towards the equatorial plane and to a stellar exponential $z$-profile if the star formation rate is assumed to be comparable to the cooling rate of the gas. The final stellar scale height depends only on the initial gas temperature and local surface density. This model therefore provides a new method to investigate the early energetic state of galactic protodisks with measured scale heights and surface densities along the disk plane. ",
        "introduction": "Galaxies of similar morphological type exhibit similarities in their structural properties, independent of whether they are isolated in general fields or agglomerated in galaxy clusters.  These similarities under different environments may emerge due to some {\\it internal} regularities that have superseeded external, probabilistic disturbances such as interactions and/or mergers.  There is a well-known evidence from optical observations that many edge-on spiral and lenticular galaxies have a universal luminosity profile perpendicular to their disk plane, or a universal $z$-distribution of stellar mass density if a constant mass-to-luminosity ratio is assumed (Tsikoudi 1977; Burstein 1979; van der Kruit \\& Searle 1981).  It was once claimed that the $z$-distribution is best fitted by a self-gravitating isothermal model like sech$^2(z/2h_z)$ using a scale parameter $h_z$ (van der Kruit \\& Searle 1981).  Near-infrared observations of edge-on spiral galaxies however uncovered an excess over the isothermal model at small $|z|$ where the optical photometry is hindered by dust absorption (Wainscoat, Freeman, \\& Hyland 1989; Aoki et al. 1991; van Dokkum et al. 1994). These infrared observations showed that an exponential distribution $\\exp(-|z|/h_z)$ provides a superior fit to the observed $z$-profiles.  Furthermore it is known from analyses of star counts that the vertical structure of the Galactic disk in the solar neighbourhood is well fitted by an exponential distribution rather than an isothermal model (Bahcall \\& Soneira 1980; Prichet 1983).  An exponential $z$-distribution can be constructed by adding up several stellar disk components with different vertical velocity dispersions, but this is only possible if the contribution from stars with larger velocity dispersion is fine-tuned to dominate progressively at larger distances from the disk plane.  A mechanism that enables this tuning has not been found to date.  The vertical disk structure has been considered as a result of dynamical evolution of the stellar component through encounters with massive clouds and spiral structures (Villumsen 1983; Lacey 1984; Carlberg 1987). However, although collisions with clouds might be important in increasing the velocity dispersions of young stars this effect cannot account for the high velocity dispersion of old stars (Lacey 1984). In addition, the resulting disks are isothermal and not exponential.  We report that the hydrodynamical equations associated with a simple star formation law possess a remarkable solution which naturally produces an exponential stellar $z$-distribution in the final stage of gravitational settling of galactic protodisks.  This result might provide new insight into the formation of galactic disks and their early evolutionary phases. ",
        "conclusions": "The presented idea that the star formation timescale $t_*$ is comparable to the cooling timescale $t_c$ for the origin of exponential stellar $z$-profiles works when the protodisk has nearly reached an equilibrium prior to star formation and gas cooling.  Under a realistic situation the dynamical timescale $t_{dyn}$ of the protodisk may initially not be in balance with the cooling timescale $t_c$. If $t_{dyn}<t_c$ in the hot and tenuous gas, the protodisk achieves an equilibrium state and undergoes a quasi-static contraction until the gas density becomes sufficiently large so that $t_c$ falls below $t_{dyn}$.  If $t_{dyn}>t_c$ otherwise, the gas cools rapidly and condenses into cool and dense clouds.  In this stage, cloud collisions dissipate kinetic energy and enhance the formation of stars. As demonstrated by Burkert et al. (1992), the increase in energy input through supernova explosions will compensate energy dissipation and cooling, eventually halting the collapse and achieving a quasi-equilibrium state. In either cases, after reaching its equilibrium, the protodisk evolves due to the feedback effects of star formation, independent of its initial formation history and of variations in global physical conditions (e.g., Lin \\& Murray 1992).  Gas cooling or energy dissipation induces a slow gravitational settling of the protodisk while, at the same time, leading to the formation of stars in the disk.  This coupling yields a continuous change of various stellar characteristics (age, metallicity, color, velocity dispersion, etc) as a function of the $z$-distance from the disk plane.  Although our model is not sufficiently detailed to allow an extensive test against such data, the power of producing an exponential stellar $z$-profile from a vast range of physical conditions is very appealing and should be considered as a strong candidate for understanding the universality of exponential $z$-profiles.  Burkert et al. (1992) have performed more detailed simulations of the dynamical settling of hot protogalactic gas into the  equatorial plane, taking into account star formation and heating and cooling processes in a multiphase interstellar medium.  They achieved a good agreement with the observed vertical density structure of the Galactic disk in the solar neighborhood. The complexity of their model leads however to a very large parameter space which cannot be explored in detail due to computational limitations. Among many different processes involved in their physical model, the process of crucial importance is that the rate of star formation is adjusted sooner or later to balance with the local cooling rate by means of the self-regulated star formation mechanism (Cox 1983; Franco \\& Cox 1983). It is not clear from their models, which process is most important in leading to exponential $z$-profiles. Our calculations demonstrate that an ideal condition for such a profile is  $t_*\\sim t_c$, independent of the details of heating and cooling.  \\subsection*"
    },
    "9603/astro-ph9603064_arXiv.txt": {
        "abstract": "We investigate the effects of finite sky coverage on the spectral resolution $\\Delta\\ell$ in the estimation of the CMBR angular power spectrum $C^{\\ell}$. A method is developed for obtaining quasi-independent estimates of the power spectrum, and the cosmic/sample variance of these estimates is calculated. The effect of instrumental noise is also considered for prototype interferometer and single-dish experiments. By proposing a statistic for the detection of secondary (Doppler) peaks in the CMBR power spectrum, we then compute the significance level at which such peaks may be detected for a large range of model CMBR experiments. In particular, we investigate experimental design features required to distinguish between competing cosmological theories, such as cosmic strings and inflation,  by establishing whether or not secondary peaks are present in the CMBR power spectrum. ",
        "introduction": "The cosmic microwave background radiation (CMBR) is one of the most promising links between astronomical observation and cosmological theory. A significant amount of experimental data already exists and further observational progress is expected in the near future. The experimental success in this field  has prompted a large amount of theoretical effort on two closely related fronts. Firstly, theorists strive to assess the impact current  observations have on theories of the early Universe. Secondly, the design of future experiments is guided by what are believed to be crucial tests of our theoretical prejudices.  A particularly active field of research in CMBR physics are the so-called Doppler peaks (Hu \\& Sugiyama 1995a,b). These consist of a series of oscillations in the angular power spectrum of CMBR fluctuations $C^{\\ell}$ predicted for most inflationary models. They are predicted in the multipole range $100\\la \\ell \\la 1500$, corresponding to angular scales $0.05 \\la \\theta\\la 1$ degrees. Experimental measurement of the Doppler peaks' positions and heights would fix at least some combinations of cosmological parameters (e.g. $H_0$, $\\Omega_0$ etc.) which are left free in inflationary models (Jungmann 1995). Future experimental projects are usually designed with these goals in mind. In particular, models with $\\Omega_0 =1$ are of special interest.  In this paper we shall consider another theoretical context. For a long time inflation  (Steinhardt 1995), and topological defects (Kibble 1976; Vilenkin \\& Shellard 1994), have stood as conceptually opposing alternative scenarios for structure formation in the early Universe. It was shown by Albrecht et al (1996) and Magueijo et al (1996), that qualitative aspects of Doppler peaks should reflect this conceptual opposition. In particular the absence of secondary Doppler peaks was proved to be a robust prediction for a large section of defect theories. Interestingly enough this statement relies only on the role played by causality and randomness in these theories. Cosmic strings were unambiguosly shown to fall in this category. However, this may or may not be the case for textures (Crittenden \\& Turok 1995; and Durrer et al 1995). Therefore it appears that determining whether or not there are secondary Doppler peaks in the CMBR power spectrum offers an important alternative motivation for experimental design and data analysis. Answering this question is a far less demanding experimental challenge, which nevertheless will have a dramatic impact on our understanding of the Universe (Magueijo \\& Hobson 1996, Albrecht \\& Wandelt 1996).  We address this issue by proposing in Section~\\ref{peakstat} a statistic for detecting secondary oscillations, and studying how it fares for signals coming from various models, when measured by using different experimental strategies. The result is encoded in a detection function $\\Sigma$ telling us to within how many sigmas we can claim a detection of secondary oscillations, given a particular model and experiment. We consider signals coming from standard CDM (sCDM), and an open CDM model (henceforth called stCDM) which is tuned to confuse inflation and cosmic strings in all but the issue of secondary oscillations. In Secs.~\\ref{sfobs}, \\ref{cv}, and \\ref{noise} we set up a framework for computing errors in power spectra estimates. We consider errors  resulting from spectral resolution limitations due to finite sky coverage, cosmic/sample variance, instrumental noise and foreground subtraction. We consider a large parameter space of experiments including single-dish experiments and interferometers. For single-dish experiments we allow the beam size, sky coverage, and detector noise to vary. For interferometers we take as free parameters the primary beam, number of fields, and detector noise.  We then use this framework to compute the detection function $\\Sigma$ for sCDM and stCDM in this large class of experiments.  The results obtained are given in Sections~\\ref{peakstat} and \\ref{cs}, and provide experimental guidance in two different ways. Firstly they allow the choice of an ideal scanning strategy (choice of resolution and sky coverage) given a constraint such as finite funding (albeit disguised in a more mathematical form). Secondly, one may compute the expected value of the detection, assuming ideal scanning, as a function of available funding. This provides lower bounds on experimental conditions for a meaningful detection as well as an estimate of how fast detections will improve thereafter. We summarize the main results in the Section~\\ref{conc}. ",
        "conclusions": "\\label{conc} We have studied the significance of secondary peak detection for sCDM and stCDM in a large parameter space of experiments, including interferometers and single-dish telescopes, and have adopted a broad minded attitude towards sky coverage. If point source subtraction is to be done in parallel with a CMBR experiment (so as to account for variable point-sources), however, a large sky-coverage may never be possible (see Section~\\ref{point}). This detail, often overlooked in satellite proposals, could then radically undermine a large number of estimates assisting experimental design. We will however not dwell on this awkward possibility, but while keeping it in mind, shall consider unrestricted sky coverage.  The results obtained are reported in Sections~\\ref{peakstat} and \\ref{cs} and stress the contradictions of an all-purpose experiment. If the low-$l$ plateau of the spectrum is a theoretical target then one needs all-sky coverage, and satellite single-dish  experiments are to be favoured. As shown in Section~\\ref{peakstat} even if one wishes to study Doppler peak features for sCDM all-sky coverage might still be preferable. Depending on the noise levels, a large sky coverage might be desirable, even for a resolution of about $\\theta_b=0.4-0.5^{\\circ}$.  Our work shows how such a design relies heavily on the assumption that the signal is in the vicinity of sCDM. If instead one is to test the high-$l$ opposition between low $\\Omega$ CDM and cosmic strings, then we have seen that single-dish experiments are required to have rather high resolutions. Interferometers appear to be less constrained, providing 2-3 sigma detections under very unassuming conditions, with rapid improvements following further improvement in experimental conditions. Furthermore, in this context, all-sky scanning is not only unnecessary, but in fact undesirable. The best scanning is normally achieved with deep small patches. These two features contradict sharply the ideal experimental design motivated by the standard theoretical gospel.  Overall the parameter space of successful stCDM detections seems to increase for interferometers and shrink for single-dish when compared with sCDM. We believe that a variety of contrasting experimental techniques may equally well find their niche as regards important theoretical implications."
    },
    "9603/astro-ph9603032_arXiv.txt": {
        "abstract": "The structure of stationary photodissociation fronts is revisited. $\\HH$ self-shielding is discussed, including the effects of line overlap. We find that line overlap is important for $N(\\HH)\\gtsim10^{20}\\cm^{-2}$, with a factor-of-two suppression of pumping rates at column densities $N(\\HH)\\approx3\\times10^{20}\\cm^{-2}$. We compute multiline UV pumping models, and compare these with simple analytic approximations for the effects of self-shielding.  The overall fluorescent efficiency of the photodissociation front is obtained for different ratios of $\\chi/\\nH$ (where $\\chi$ characterizes the intensity of the illuminating ultraviolet radiation) and different dust extinction laws. The dust optical depth $\\tau_{pdr}$ to the point where 50\\% of the H is molecular is found to be a simple function of a dimensionless quantity $\\phi_0$ depending on $\\chi/\\nH$, the rate coefficient $R(T)$ for $\\HH$ formation on grains, and the UV dust opacity. The fluorescent efficiency of the PDR also depends primarily on $\\phi_0$ for $\\chi\\ltsim3000$ and $\\nH\\ltsim10^4\\cm^{-3}$; for stronger radiation fields and higher densities radiative and collisional depopulation of vibrationally-excited levels interferes with the radiative cascade. We show that the emission spectrum from the PDR is essentially independent of the color temperature $T_{color}$ of the illuminating radiation for $10^4\\K\\ltsim T_{color}$, but shows some sensitivity to the rotation-vibration distribution of newly-formed $\\HH$. The 1--0S(1)/2--1S(1) and 2--1S(1)/6--4Q(1) intensity ratios, the ortho/para ratio, and the rotational temperature in the $v$=1 and $v$=2 levels are computed as functions of the temperature and density, for different values of $\\chi/\\nH$.  We apply our models to the reflection nebula NGC 2023. Apparent inconsistencies between published K-band and far-red spectroscopy of this object are discussed; we adjust the two sets of observations for consistency. We are able to approximately reproduce the (adjusted) observations with models having $\\chi=5000$, $\\nH=10^5\\cm^{-3}$, and a reasonable viewing angle. Further observations of NGC 2023 will be valuable to clarify the uncertain spatial structure of the emission. ",
        "introduction": "It is by now widely recognized that reprocessing of OB starlight in photodissociation fronts plays an important role in the overall energetics of star-forming galaxies, particularly extreme ``starburst'' galaxies. An important mechanism in this reprocessing is the absorption of ultraviolet photons by molecular hydrogen, followed by ultraviolet and infrared fluorescence, and, about 15\\% of the time, by dissociation. The vibrational fluorescence process, first noted by Gould \\& Harwit (1963), has been investigated by a number of authors (e.g.: Black \\& Dalgarno 1976; Shull 1978; Black, Porter \\& Dalgarno 1981; van Dishoeck \\& Black 1986; Black \\& van Dishoeck 1987; Sternberg 1986, 1988; Sternberg \\& Dalgarno 1989).  A number of theoretical investigations have advanced our understanding of photodissociation regions, or ``PDRs''.\\footnote{ We use ``PDR'' for ``photodissociation region''.  Note, however, that some authors use these initials for ``photon-dominated region''. } Solomon (1965: see Field, Somerville, \\& Dressler 1966) apparently was the first to propose that destruction of interstellar $\\HH$ might be dominated by dissociating transitions to the vibrational continuum following permitted transitions to the $^1B\\Sigma_u^+$ state; the resulting photodissociation rate was first estimated by Stecher \\& Williams (1967), who called attention to the potential importance of $\\HH$ self-shielding. The $\\HH$ self-shielding process has been studied by Shull (1978), Federman, Glassgold \\& Kwan (1979), and de Jong, Dalgarno, \\& Boland (1980).  Tielens \\& Hollenbach (1985a) discussed the overall thermal and chemical structure of photodissociation fronts, and applied their theoretical model to explain observed properties of the photodissociation region in Orion (Tielens \\& Hollenbach 1985b). Black \\& van Dishoeck (1987) examined in detail the fluorescent excitation of $\\HH$, and the resulting infrared line emission. Recently Abgrall \\etal (1992) presented new data relating to the photodissociation of $\\HH$, and carried out a more accurate treatment of the self-shielding process in order to reexamine the structure of the H/$\\HH$ transition zone. Recent detailed self-shielding calculations for a cloud of moderate density subject to strong UV illumination (Le Bourlot \\etal 1992) and for diffuse cloud conditions (Heck \\etal 1992) show that in some regions line overlap significantly affected the $\\HH$ photodissociation rate.  In a previous paper (Bertoldi \\& Draine 1996) we discussed the structure of coupled ionization-dissociation fronts, and concluded that they were in many cases expected to be propagating at significant velocities, calling into question the interpretation of some regions, including Orion, which have been based on models of stationary photodissociation fronts. In the course of our investigation we have reexamined the self-shielding of $\\HH$.  In this paper we revisit the problem of $\\HH$ self-shielding. We confirm that line overlap can often be important, and we identify the region in parameter space where this occurs. We develop an approximate method to allow for the effects of line overlap in a statistical fashion. We apply this method to compute models for photodissociation fronts including the effects of line overlap; we also improve somewhat upon previous treatments by applying the results of recent studies of the wavelength dependence of dust extinction. We find a simple analytic description of the $\\HH$ self-shielding function which appears to give a good approximation to the results of detailed self-shielding calculations, including line overlap.  We construct models of stationary photodissociation fronts for different densities, temperatures, and intensities of incident FUV radiation. We examine how observable properties, such as the 1--0S(1) surface brightness, and the 1--0S(1)/2--1S(1) and 2--1S(1)/6--4Q(1) line ratios, depend on the model parameters. Two indicators of the excitation mechanism (fluorescent vs. shock) and the gas density are the ratio of the ortho-- (odd $J$) to para--H$_2$ (even $J$) level populations and the rotational temperature within a vibrational level. We compute these quantities as a function of gas density and FUV illumination.  The famous reflection nebula NGC 2023 is a well-studied example of fluorescent emission by $\\HH$. We attempt to reconcile published observations of this object. The available observations of NGC 2023 appear to be consistent with a photodissociation front with $\\nH\\approx10^5\\cm^{-3}$ irradiated by FUV radiation with intensity (relative to the Habing field) $\\chi\\approx5000$. ",
        "conclusions": "The principal results of this paper are as follows: \\begin{enumerate} \\item An approximate treatment of self-shielding [eq.\\ (\\ref{eq:withoverlap})] is developed which treats line overlap in a statistical fashion. This treatment allows individual lines to be treated using the single-line equivalent-width approximation, but allows statistically for overall suppression of the continuum due to line overlap. Comparison with exact calculation shows this approximation to be accurate. \\item Two simple approximations are provided for the self-shielding function for $\\HH$.  The first, a simple power-law [eq.\\ (\\ref{eq:powfit})], provides a fairly good fit to self-shielding over the range $10^{15}< N_2 < 10^{21}\\cm^{-2}$. The second [eq.\\ (\\ref{eq:goodfit})] is an analytic function which provides a very good approximation to the self-shielding including the effects of line overlap. These functions are recommended for use in future studies of photodissociation regions. \\item The effects of line overlap become important for column densities $N(\\HH)\\gtsim10^{20}\\cm^{-2}$, suppressing the UV pumping rate by a factor of 2 at $N(\\HH)\\approx3\\times10^{20}\\cm^{-2}$ (cf. Fig.~\\ref{fig:eqwidths}). \\item For dense cloud dust properties, the effects of line overlap become important while the dust is still optically thin for $\\chi/\\nH\\ltsim 0.05$ (cf. Fig.~\\ref{fig:NH2vsNH,dense}). \\item The dust optical depth $\\tau_{pdr}$ at the point where $2n(\\HH)=n({\\rm H})$ is primarily a function of a single parameter $\\phi_0 \\propto (\\chi/\\nH R)\\sigma_{1000}^{3/4}$, defined in eq.\\ (\\ref{eq:phi_0}). The approximation (\\ref{eq:taupdr}) provides a good estimate for $\\tau_{pdr}$. \\item The ``efficiency'' of UV pumping of $\\HH$ in the photodissociation front, $\\epsilon_{pump}$, is also a function of the single parameter $\\phi_0$, and is approximately given by eq.\\ (\\ref{eq:eps_pump_approx}). \\item For $\\nH\\ltsim10^5\\cm^{-3}$ the efficiency $\\epsilon_{1-0S(1)}$ for 1--0S(1) emission is given as a function of $\\phi_0$ and $\\chi$ by eq.\\ (\\ref{eq:epsilon10s1fit}). \\item The emission from PDRs is quite insensitive to the color temperature of the illuminating radiation for color temperatures $T_{color}\\gtsim10^4\\K$, or stars with spectral type A0 or earlier. \\item The emission spectrum from low $J$ levels [e.g., 1--0S(1), or 3--2S(1)] is insensitive to the distribution function $\\hhini(v,J)$ of newly-formed $\\HH$, but the emission from higher $J$ levels [e.g., 1--0S(9)] is increased (by a factor $\\sim2$) when the newly-formed $\\HH$ is rotationally-``hot''. \\item Observable properties of photodissociation fronts -- including the 1--0S(1)/2--1S(1) line ratio (Fig.~\\ref{fig:10S1/21S1}), the 2--1S(1)/6--4Q(1) line ratio (Fig.~\\ref{fig:21S1/64Q1}), the ortho/para ratio (Fig.~\\ref{fig:orthopara}), and the rotational temperatures (Fig.~\\ref{fig:trotv=1},\\ref{fig:trotv=2}) -- are computed for various values of $\\chi$, $\\nH$, and temperature. Complete $\\HH$ vibration-rotation spectra are available via anonymous ftp. \\item The reflection nebula NGC 2023 is considered. We correct the H87 intensities for assumed beam dilution, and adjust the B92 intensities for consistency. The (adjusted) observations are approximately reproduced by a model with $\\nH=10^5\\cm^{-3}$ and $\\chi=5000$, with $T_0\\approx900\\K$. The agreement is not perfect; possible explanations for the discrepancies are discussed. \\end{enumerate}"
    },
    "9603/astro-ph9603048_arXiv.txt": {
        "abstract": "We introduce the unbiased way statisticians look at the 2--point correlation function and study its relation to multifractal analysis. We apply this method to a simulation of the distribution of galaxy clusters in order to check the dependence of the correlation dimension on the cluster richness. ",
        "introduction": "The statistical description of the galaxy clustering is usually based on the two-point correlation function $\\xi(r)$. This function is, following the terminology used by statisticians working in point field statistics, a second-order characteristic of the point process (Diggle 1993; Stoyan \\& Stoyan 1994). The first-order characteristic is just the intensity measure $\\lambda(\\vec r)$ (Mart\\'{\\i}nez et al. 1993). Assuming the Cosmological principle, we accept that galaxies in large volumes represent a stationary and isotropic point process, having therefore constant intensity equal to the number density of galaxies per unit volume, denoted by $n$. ",
        "conclusions": ""
    },
    "9603/astro-ph9603160_arXiv.txt": {
        "abstract": "The standard method of modelling axisymmetric stellar systems begins from the assumption that mass follows light. The gravitational potential is then derived from the luminosity distribution, and the unique two-integral distribution function $f(E,L_z^2)$ that generates the stellar density in this potential is found. It is shown that the gravitational potential can instead be generated directly from the velocity data in a two-integral galaxy, thus allowing one to drop the assumption that mass follows light. The two-dimensional rotational velocity field can also be recovered in a model-independent way. Regularized algorithms for carrying out the inversions are presented and tested by application to pseudo-data from a family of oblate models. ",
        "introduction": "Modelling of elliptical galaxies has a long history, extending back to a time when it was taken for granted that elliptical galaxies were axisymmetric and that all of their mass could be accounted for in stars. We now know that dark matter is a commmon component of galaxies, both at large (Ashman 1992) and small (Kormendy \\& Richstone 1995) radii, and that triaxial configurations are possible ones for stellar systems (Schwarzschild 1979).  One goal of modern dynamical studies is accordingly to verify the presence of dark matter; another is to detect departures from axisymmetry. Because triaxial models are difficult to construct, one often begins by looking for simpler models that are consistent with the data. This ``model-building'' approach is popular since it requires a minimum of thought about the information content of the data: one simply builds a model and checks whether it reproduces the observations. But one can take a more sophisticated approach, and ask whether the data imply anything definite about the observed stellar system. For instance, one might attempt to falsify the axisymmetric hypothesis in a model-independent way. This ``inverse problem'' approach is more difficult but also more rewarding, since it leads to more secure conclusions about the dynamical state of the galaxy.  Beginning with Binney, Davies \\& Illingworth's (1990) pioneering study, a standard scheme has been developed for modelling elliptical galaxies. One assumes that the observed galaxy is axisymmetric and that the distribution of mass is known -- for instance, mass might be assumed to follow light. The gravitational potential is then derived from this mass distribution using Poisson's equation; the unique two-integral distribution function $f(E,L_z)$ that reproduces the stellar density in this potential (or more exactly, the even part of $f$) can also be found (Lynden-Bell 1962; Hunter 1975). The odd part of $f$, which determines the rotational velocity field, is usually represented via some ad hoc parametrization. One then projects this derived $f$ back into observable space to find the predicted kinematical variables and compares them with the observations. If there is agreement, one can claim to have found an axisymmetric model that is consistent with the data.  This technique has been used to reproduce the kinematical data for a few, well-observed elliptical galaxies (van der Marel et al. 1994; Dehnen 1995; Qian et al. 1995). But it is difficult to state precisely what has been learned from studies like these. Are the observed galaxies actually characterized by a constant $M/L$, axisymmetry, and a two-integral distribution function, as assumed? Or might there exist models with spatially-varying $M/L$'s or three-integral $f$'s that reproduce the data equally well? And if no model consistent with the data can be found, which of the various assumptions made in the model-building has been violated?  A useful comparison can be made here with spherical systems (\\S 2.1). If one assumes that the gravitational potential $\\Phi(r)$ of a spherical galaxy is known, then the observed velocity dispersion profile can be used to infer the unique dependence of anisotropy on radius (Binney \\& Mamon 1982). Alternatively, one can assume that the distribution of stellar velocities is everywhere isotropic, in which case the velocity dispersion profile implies a unique form for $\\Phi(r)$ (Merritt \\& Gebhardt 1994). Assuming {\\it both} isotropy {\\it and} a constant $M/L$ would clearly be an over-determination of the problem, since the two assumptions together imply a unique velocity dispersion profile, and this predicted profile would almost certainly be different from the observed profile -- indeed, the kinematical data would not have been used at all in the construction of the model. But this is precisely what is commonly done in modelling axisymmetric galaxies: one assumes both isotropy in the meridional plane (i.e. $f=f(E,L_z)$) as well as an ad hoc form for the potential, and these two assumptions together imply a unique dependence of the mean square velocity on position over the image of the galaxy. It is unlikely that a model constructed in this way would reproduce the two-dimensional velocity dispersion data; and if it did not, there would be no clear indication of which assumption had been violated.  It is clearly desirable to think more carefully about the information content of kinematical data in axisymmetric stellar systems. Unlike the spherical case, the functions to be derived depend on two spatial variables, but the kinematical data are likewise two-dimensional, depending on both projected radius and position angle. Many such data sets now exist, for stars in globular clusters (Meylan et al. 1995) and in galaxies (Bacon et al. 1994); emission-line objects around the Galactic bulge (te Linkel Hekkert et al. 1991) and around other galaxies (Hui et al. 1995); and galaxies in clusters (Colless \\& Dunn 1996). These data are most often in the form of discrete velocities, and in the best-studied systems, the measured velocities number in the hundreds or even thousands.  Here it is shown (\\S 2.2) that the availability of two-dimensional kinematical data allows one to approach the dynamical study of axisymmetric stellar systems as an inverse problem rather than as a model-building problem. From the single {\\it Ansatz} $f=f(E,L_z)$, one can infer the unique gravitational potential $\\Phi(\\varpi,z)$ that is consistent with the observed first and second velocity moments of $f$. The rotational velocity field can likewise be obtained in a model-independent way. By using the kinematical data in the construction of the model, rather than assuming that mass follows light, one thus arrives at the unique pair of functions $\\{f,\\Phi\\}$ that are consistent with the two-integral hypothesis. If this unique model can be falsified -- for instance, by using line-of-sight velocity distributions, proper motions, or some other data -- then the two-integral, axisymmetric hypothesis will have been convincingly ruled out. In the model-building approach, on the other hand, one can not discard the two-integral assumption until one has explored an effectively infinite number of possible forms for the potential (\\S 2.3).  The numerical techniques that are required for converting the velocity data into a map of the potential are rather more subtle than the ones that have been applied up to now in the study of axisymmetric galaxies. These techniques are accordingly described in some detail (\\S3.1) before applying them to pseudo-data generated from a family of oblate models (\\S3.2, 3.3).  Although this approach yields the unique potential consistent with the two-integral assumption for $f$, the true potential might be different if $f$ depends on a third integral. However it would be premature to investigate three-integral models until one had first used the algorithms described here, or equivalent ones, to rule out a two-integral model. Having done so, one could then search for the pair of functions $\\{\\Phi(\\varpi,z)$, $f(E,L_z,I_3)\\}$ that are most consistent with the data. This is a hard problem, and even in the spherical geometry there is no published algorithm that can extract $\\Phi(r)$ and $f(E,L^2)$ in a completely model-independent way from the data; the most sophisticated such algorithms are still based on parametrized forms for the potential (Saha \\& Merritt 1993; Merritt 1993b).  The focus here is on situations where the velocities are measured discretely, as would be the case for stars in a globular cluster. Most large, kinematical data sets are of this form.  It is also assumed throughout that the observer lies in the equatorial plane of the observed system. The reason for this unpleasant assumption is that deprojection of the luminosity density becomes underdetermined if the inclination angle is less than $\\pi/2$, even if this angle is assumed known (Rybicki 1986; Gerhard \\& Binney 1996). One can therefore probably not hope to uniquely infer the dynamical state of an axisymmetric galaxy that is not viewed edge-on. This fact is routinely ignored in the model-building studies but must be faced squarely if one wishes to solve the inverse problem.  Also presented here (\\S 3.4) is the first regularized algorithm capable of deriving $f(E,L_z)$ for an axisymmetric system from its density and potential (the Lynden-Bell -- Hunter problem) given imperfect or incomplete information about those functions. ",
        "conclusions": "Given measurements of the line-of-sight rotational velocity and the velocity dispersion over the two-dimensional image of an edge-on, axisymmetric galaxy, one can derive the unique functions $\\Phi(\\varpi,z)$ and $f(E,L_z)$ that reproduce those data, without the necessity of assuming that mass follows light. The validity of such a model can then be tested using other data, such as the full line-of-sight velocity distributions or the proper motions. Failure to reproduce the data using this method would imply that $f$ depends on a third integral of motion (or that one of the geometric assumptions, axisymmetry or zero inclination, was violated). This approach is a generalization of the one pioneered by Binney, Davies and Illingworth (1990) which assumes that mass follows light and which derives both $\\Phi$ and $f$ from the luminosity distribution alone, without using the kinematical data except in the normalization of $\\Phi$. We have presented numerical algorithms, suitable for use with noisy and incomplete data, for carrying out the required inversions and shown that they provide smooth and accurate estimates of $f$ and $\\Phi$ using simulated data sets derived from oblate models.  In a companion paper, the algorithms described here will be applied to velocity data from the globular cluster $\\omega$ Centauri.  The results described here suggest the following avenues for future work.  1. A mathematical -- as opposed to numerical -- demonstration of uniqueness (or non-uniqueness) of the solutions to the coupled equations (\\ref{magic}) and (\\ref{disp}) would be of fundamental importance.  2. The second-derivative penalty functions adopted here are almost certainly non-optimal and it would be worthwhile to investigate more general forms. In addition, it is possible that the solutions to the roughness-penalized inverse problems defined here can be found analytically, as in Wahba \\& Wendelberger (1980).  3. As shown most recently by Gerhard \\& Binney (1996), the axisymmetric inverse problem becomes under-determined when the stellar system is viewed from a direction that is not parallel to the symmetry plane. Additional information must therefore be added if one is to solve the inverse problem for an arbitrary orientation. (One would also like to {\\it infer} that orientation from the data.) There is at present no clear understanding of what form that extra information should take.  4. One can define an alternate inverse problem based on the {\\it Ansatz} that mass follows light, or more generally that the form of the potential $\\Phi(\\varpi,z)$ is known a priori; for instance, the potential might be derived from X-ray data, from the statistics of gravitational lensing, etc. This is of course the usual assumption that is made when modelling axisymmetric galaxies; however the goal would be to learn something unique about the degree of velocity anisotropy $\\sigma_{\\varpi}/\\sigma_z$ from the velocity moment data, as in the analogous spherical problem (Binney \\& Mamon 1982). Unfortunately one can show that the assumption of a known potential is {\\it not} sufficient in itself to uniquely constrain the anisotropy, even for an edge-on system; one must make additional assumptions, e.g. that the velocity ellipsoid is everywhere aligned with the coordinate axes. Nevertheless, more work along these lines would help to elucidate the more difficult inverse problem for a three-integral $f$.  \\bigskip  This work was supported by NSF grant AST 90-16515 and NASA grant NAG 5-2803. T. Fridman kindly assisted in the plotting of the figures. I thank W. Dehnen and O. Gerhard for lively discussions on the topic of axisymmetric modelling, and for pointing out some mathematical errors in an early version of the paper. I also thank H. Dejonghe for many useful insights on problems of this sort.  \\clearpage"
    },
    "9603/astro-ph9603013_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Perturbation theory already has a long and distinguished history in Cosmology, but for many years the attention focused mainly on the linear terms, \\eg for early universe calculations of the power spectra of light and matter at recombination. But for the latter evolution of large scale structures, perturbation theory went out of fashion in the era which saw the advent of massive computer simulations of their formation. These allowed investigating strongly non-linear scales, the only ones that could be carefully studied statistically in the galaxy surveys of the time. Another reason of this purgatory period of the theory may have been the perception that any early signature of the weakly non-linear phase would have been totally erased by the following strongly non-linear phase, \\ie that it was irrelevant. In any case, the situation has now dramatically changed, galaxy catalogs encompass ever increasing scales where density contrasts are weak; they display remarkable moments hierarchy on large scales, with a striking similarity to that observed on smaller scales. Meanwhile realistic calculations have become recently available (\\eg including the smoothing inevitable in any realistic measurement, redshift distortions, etc), and last but not least, comparisons with numerical simulations showed the predictive power of the theory. In a few years the beauty awakened, and this area of research is now striving. These notes just surveyed some of the simpler analyses, and new results are posted nearly every week. Of course, much remain to be done, and maybe the most important question to be answered is why does these perturbative approaches work so well?    \\newcommand{\\inprep}{{\\em in preparation}} \\newcommand{\\inpress}{{\\em in press}} \\newcommand{\\submit}{{\\em submitted to}} \\newcommand{\\preprint}{{\\em preprint}}  \\newcommand{\\mn}{{\\em Mon. Not. R. astr. Soc}} \\newcommand{\\apj}{{\\em Astrophys. J.}}  \\newcommand{\\apjsup}{{\\em Astrophys. J. Suppl.}} \\newcommand{\\aj}{{\\em Astron. J.}}  \\renewcommand{\\aa}{{\\em Astr. Astrophys.}} \\newcommand{\\ass}{{\\em Astrophys. Space Sci.}}  \\newcommand{\\nat}{{\\em Nature}} \\newcommand{\\et}{{\\em et al.}\\,\\,}"
    }
}