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+{
+    "9602/astro-ph9602028_arXiv.txt": {
+        "abstract": "We consider stochastic particle acceleration in plasmas around stellar mass black holes to explain the emissions above 1 MeV from Galactic black hole candidates. We show that for certain parameter regimes, electrons can overcome Coulomb losses and be accelerated beyond the thermal distribution to form a new population, whose distribution is broad and usually not a power law; the peak energy of the distribution is determined by the balance between acceleration and cooling, with particles piling up around it. Radiation by inverse Compton scattering off the thermal (from background) and non-thermal (produced by acceleration) particles can in principle explain the hard X-ray to gamma-ray emissions from black hole candidates. We present model fits of Cyg X-1 and GRO J0422 in 50 keV -- 5 MeV region observed with OSSE and COMPTEL. ",
+        "introduction": "Most of Galactic black hole candidates (GBHCs) show X-ray spectra well fit by a thermal Comptonization model with temperatures $\\sim$ 50 -- 100 keV and Thomson depths of a few (e.g., \\cite{har94}; \\cite{lia93}). OSSE and COMPTEL experiments on {\\em Compton Gamma-Ray Observatory} (CGRO), however, have recently revealed that persistent gamma-rays $>$ MeVs are being produced in some GBHCs, in particular Cyg X-1 (\\cite{joh93}; \\cite{mcc94}) and GRO J0422+32 (\\cite{dij95}). These gamma-ray tails are hard to fit with a single component Sunyaev \\& Titarchuk (1980) thermal model (or a more recent model by Titarchuk 1994). Hence non-thermal processes are strongly hinted by these observations.  Models of non-thermal $e^\\pm$ pairs have been studied with the emphasis on the power-law X-ray emission from active galactic nuclei (AGNs) (e.g., \\cite{fab86}; \\cite{lz87}; \\cite{sv87}; \\cite{cop92}), and they can be also applied to plasmas around stellar mass black holes. Those models assume that either mono-energetic or power-law leptons with a large Lorentz factor ($\\gamma \\gg 1$) are injected and they initiate cascade processes such as $e^{\\pm}$ production and Compton scattering (see \\cite{sv94} for a recent review). However, those models did not specify any acceleration mechanism, thereby lacking the self-consistency in determining the particle distributions and photon spectra.  Motivated by the MeV photons from GBHCs, we study the role of stochastic particle acceleration in accreting plasmas near GBHCs. The mechanisms of wave-particle resonant interactions that lead to particle acceleration have been directly observed in solar-wind (e.g., \\cite{mar91}) and extensively studied in the context of solar flares (e.g., \\cite{mel74}; \\cite{ram79}; \\cite{hp92}; \\cite{mr95}). Stochastic acceleration has been applied to diffusive shock acceleration and cosmic rays (\\cite{sch94}), the lobes of radio galaxies (e.g., \\cite{lac77}; \\cite{ach79}; \\cite{eil84}), and recently, the central regions of AGNs (\\cite{dml95}). Here in this {\\em Letter}, we couple the particle energization and radiative processes to determine self-consistently the steady state particle and photon distributions in GBHC environments. We then present fits to gamma-ray data of GBHCs observed with OSSE and COMPTEL. ",
+        "conclusions": "\\label{dis-sec}  We have examined stochastic particle acceleration via wave-particle resonant interactions near Galactic black holes. We find that under certain conditions, stochastic electron acceleration can overcome both Coulomb and Compton losses, resulting in a suprathermal population. Preliminary model spectra show good fits to the recent OSSE and COMPTEL observations of Cyg X-1 and GRO J0422. We find that Cyg X-1 has a much higher component in gamma-rays than J0422, as it is also evident in the data. This can be understood in terms of the higher $\\ell_{\\rm nt} / \\ell$ and lower compactness of Cyg X-1, compared to J0422, if they have the same black hole mass. On the other hand, the high energy emissions from J0422 can be fitted by different masses of the putative black holes. But these fittings are different from the standard thermal (i.e. e-p plasma only) models because a large amount of pairs are produced and thermalized due to the high compactness. Since both sources are highly variable, especially in gamma-ray emissions, we expect the non-thermal energy content $\\ell_{\\rm nt} / \\ell$ might vary but the detailed microphysics of how the system partitions the energy flows is presently unclear.  The non-thermal population we obtain here is relatively soft with few pairs, not capable of explaining either the transient MeV bump discovered by Ling et al. (1987) or the COMPTEL power-law tail if it extends to much higher energies. This softness is due to our assumptions of a copious soft photon source permeating the entire emission region and {\\em the coexistence of thermal and non-thermal particles in a single homogeneous volume}. Coulomb interaction {\\em efficiently} transfers the non-thermal energy to the thermal reservoir which strongly limits the existence and energy content of non-thermal particle population. To achieve much harder non-thermal populations and emissions we need to postulate the existence of physically distinct regions for the thermal and non-thermal components, either in the form of a thermal disk plus non-thermal corona or jets, or thermal outer disk and non-thermal inner torus. We plan to undertake such ventures in future work. But it is highly encouraging that a simple-minded single-component spherical model with few parameters is already capable of explaining from first principles the quiescent (low-hard state) emissions of Cyg X-1 and GRO J0422 detected by CGRO.  We note, of course, that there are many other alternative interpretations of the hard tail, including the pion decay model (\\cite{jr94}) and pair-dominated hot cloud model invoking quenching of the soft photon source (\\cite{ld88}). It is possible that particles accelerated by shocks may also account for the $>$ MeV emissions."
+    },
+    "9602/astro-ph9602046_arXiv.txt": {
+        "abstract": "In the light of the recent finding that nebular emission lines are commonly found in the inner regions of early-type galaxies, we evaluate their effect to stellar absorption-line indices. We derive analytical expressions for changes induced by the presence of nebular emission lines, both for atomic and molecular absorption-line indices. We find that the \\NI\\ emission-line doublet at 5199 \\AA\\ can significantly affect \\mgb\\ line-strength measurements. For typical equivalent widths of the \\NI\\ doublet in nuclei of early-type galaxies featuring nebular emission, the \\mgb\\ index is artificially enhanced by 0.4\\,--\\,2 \\AA, which represents a significant fraction of the  typical equivalent width of the \\mgb\\ line in early-type galaxies.  We illustrate this effect in the case of NGC~2974.  \\keywords galaxies: elliptical and lenticular -- elliptical galaxies: line strengths, emission lines, kinematical substructure ",
+        "introduction": "\\label{intro}  The \\mgb\\ index at 5177\\,\\AA\\ is one of the most prominent features in the optical spectra of old stars. It is one of the 11 well-known strong spectral line-strength indices that are widely used to study stellar populations in old stellar systems such as globular clusters and early-type galaxies (e.g., Burstein \\etal \\cite{bfgk84}; Faber \\etal \\cite{fab+85}; Gorgas, Efstathiou \\& Arag\\'on-Salamanca \\cite{gea90}; Davies, Sadler \\& Peletier \\cite{dsp93}; Gonz\\'alez \\cite{gonz93}; Carollo \\& Danziger \\cite{cardan94};  Sansom, Peace \\& Dodd \\cite{sams+94}; Fisher, Franx \\& Illingworth \\cite{fish+95}; Surma \\& Bender \\cite{surben95}). Using a library of standard stars, these indices have been empirically calibrated as a function of stellar colour, surface gravity and metallicity (Gorgas \\etal \\cite{gorg+93}), allowing the construction of semi-empirical population synthesis models (e.g., Worthey \\cite{guy94}). The behaviour of these indices in old stellar systems has provided useful insights on the star formation history of elliptical galaxies (e.g., Davies \\cite{dav95} and references therein).  These line strengths are a measure of the depth of a given absorption feature, based on the flux in a central bandpass covering the feature of interest versus a local continuum level, which is linearly interpolated from a pair of bracketing bandpasses at either side of the feature. The continuum bandpasses have been defined in regions close to the feature, having no strong absorption lines; the bandpasses have been chosen wide enough to enable a calibration of index dilution due to velocity dispersions in galaxy nuclei. The indices have been defined using old Galactic stars and globular clusters, free of emission lines. However, it is now known that a significant fraction (at least 50\\%) of `normal' elliptical galaxies contains ionized gas (Goudfrooij \\etal \\cite{paul+94a} (hereafter GHJN) and references therein). In this {\\it Letter}, we study the effect of nebular emission lines to measured line strengths. Particular attention will be devoted to the effect of the \\NI$\\,\\lambda\\lambda 5197.9, 5200.4$ emission-line doublet (often, and hereafter, refered to as \\NI\\,5199), since it is situated right in the continuum bandpass redward of the \\mgb\\ lines, and is thus a potential source of error. ",
+        "conclusions": "\\label{concl}  In this {\\it Letter}, we have discussed the effect of nebular emission lines on the derivation of absorption-line strengths in the central region of early-type galaxies. We derived analytical expressions for changes induced by the presence of nebular emission lines, both for atomic and molecular absorption-line indices. In particular, we emphasize the fact that the measurement of the well-known \\mgb\\ index can be strongly perturbed by the presence of the \\NI\\,5199 doublet, with a difference in equivalent width $\\Delta \\mbox{\\mgb}$ up to 2~\\AA. We have illustrated this effect using spectrographic observations of NGC~2974. We argue that this effect should be considered when studying line-strength gradients in the inner regions of early-type galaxies."
+    },
+    "9602/astro-ph9602121_arXiv.txt": {
+        "abstract": "A protogalaxy candidate at $z=2.72$ has been discovered serendipitously by the CNOC cluster redshift survey. The candidate is an extremely luminous ($V=20.5$ mag, absolute mag --26) and well resolved (2$''$$\\times$3$''$) disk-like galaxy. The redshift is identified from  a dozen strong UV absorption lines, including lines with P-Cygni profiles, which are indicative of the presence of young O and B stars. No emission lines are found between 1000 and 2000\\AA~(rest), including Ly$\\alpha$. The surface brightness profile of the galaxy fits an exponential law with a scale length of $\\sim$3.5 kpc. The multi-color photometric data fit the spectral energy distributions of a stellar population from 400 million years to an arbitrary young age, dependent on the amount of dust extinction. However, the presence of a strong P-Cygni profile in \\CIV~indicates that a very substantial component of the stellar population must be younger than $\\sim$ 10 Myr. These models predict that this galaxy will evolve into a bright galaxy of several $L^*$ in brightness. We can interpret this object as an early-type galaxy observed within about 100 million years of the initial burst of star formation which created most of its stellar mass, producing the extremely high luminosity. Because of the resolved, regular, and smooth nature of the object, it is unlikely that the high luminosity is due to gravitational lensing. We estimate the sky density of this type of objects observable at any one time to be $10^{0\\pm1}$ per square degree. ",
+        "introduction": "The discovery of a primeval, or proto-, galaxy -- loosely defined as a galaxy in its initial star formation stage --  has long been one of the most sought-after observational goals of extragalactic astronomy. The importance of proto-galaxies (PGs) to our understanding of the formation of structures and evolution of galaxies is enormous. By observing galaxies in their early stages of evolution, we may be able to delineate the dynamical processes that lead to and accompany the formation of galaxies. Furthermore, because evolution is slow for galaxies containing mostly stars older than about 1 Gyr (Tinsley 1972), it is only by observing the properties of galaxies during the first Gyr or so of their lifetime that we will be able to place  strong constraints on the star formation and chemical evolution history of galaxies.   It is precisely because no definitive observational data for PG exist, that over the years PGs have been predicted to have properties covering the whole imaginable range of morphologies, colors, redshift of formation, and emission-line characteristics (e.g., see the excellent reviews by Koo 1986 and Pritchet 1994, and references therein). In response to the predictions arising from these diverse models of PGs, many different methods have been suggested and carried out in their search -- narrow-band imaging from the near-UV to the infrared searching for Ly$\\alpha$ emission (e.g., Pritchet \\& Hartwick 1990); broad-band imaging searches for the Lyman break (e.g., Steidel \\& Hamilton 1993, Djorgovski 1992); and searching around other known high-redshift objects discovered by their non-stellar emission, such as quasars, radio-galaxies, and damped Ly$\\alpha$ absorbers (e.g., Djorgovski 1985, McCarthy 1993, and Wolfe 1995). Most of these searches have one thing in common: the dependence of redshift identification on the existence of the Ly$\\alpha$ emission line.  Many different types of objects have been put forward as possible PG candidates, ranging from high-redshift radio galaxies to Ly$\\alpha$ emitters found around quasars and damped Ly$\\alpha$ systems. Most recently, two promising candidates were announced in the literature. One was the ultra-luminous IRAS galaxy 10214+4724 at $z=2.29$ (Rowan-Robinson et al.~1991). However, recent observations (Matthew et al.~1994, Eisenhardt et al.~1996) have clearly shown that gravitational lensing is the culprit behind the very large apparent luminosity; the object is consistent with being similar to low-redshift luminous IRAS galaxies. Another recent object of interest is Hawaii-167, an unresolved object at $z=2.33$ with strong UV absorption lines, discovered by Cowie et al.~(1994) from their $K$-band redshift survey. Egami et al.~(1996) interpret this object as a buried quasar with a host galaxy undergoing a strong burst of star formation. Cowie et al.~(1995) recently reported the discovery of many galaxies at $1<z<1.6$ with star formation rates in excess of 10 ${\\rm M}_\\odot$ yr$^{-1}$. However, these massive star forming galaxies are most likely to be galaxies with an on-going (if episodic) star formation history, and not in their first Gyr of their life time. Hence it appears that a PG candidate that everyone can agree on continues to elude searchers.  From a very simplistic view, an object that is at high redshift, looks like a galaxy, and is identified by spectral features consistent with arising entirely from a large number of young stars would certainly qualify as a PG candidate. In this paper, we describe the serendipitous discovery of an ultra-luminous galaxy at $z=2.72$, identified by its absorption-line spectrum which appears to have all of these attributes. The object was observed as part of the CNOC (The Canadian Network for Observational Cosmology) Cluster Redshift Survey (Carlberg et al. 1994; Yee, Ellingson, \\& Carlberg 1996, hereafter YEC). The most conservative interpretation of the spectrum is that it is primarily due to young  O, B, and possibly A stars. With its resolved morphology, galaxy-like surface profile, high redshift, and extremely high luminosity (about --26 mag at rest 1500\\AA), this object is consistent with being a $L^*$ galaxy observed within the first 100 Myr since its initial star burst. This object appears to fit every definition of a PG.  In Section 2 we describe the imaging and spectroscopic observations that led to this discovery. Section 3 presents the analysis of the spectra and images. We discuss the implications of the results in Section 4. Our conclusions are summarized in Section 5. In this paper, we use $H_0=75$ km s$^{-1}$ Mpc$^{-1}$, and $q_0=0.1$ throughout. ",
+        "conclusions": "\\subsection {cB58 as a Protogalaxy}  The extended nature, regular morphology, and the lack of a significant point source or emission lines rule out the possibility that cB58 is a quasar or a  quasar-like object. Furthermore, the lack of color gradients and the fact that the absorption lines appear across the whole spatial extent of the spectral images allow us to eliminate the possibility that we are observing the continuum of a background quasar with the absorption lines being produced by an extended (foreground) galaxy at $z=2.7$. Hence, both the continuum and the absorption lines must arise from a single object.  The strong continuum and absorption lines can be most conservatively be interpreted as arising from early type O, B, and  possibly A stars. The spectrum is best compared to the IUE star library in Fanelli et al.~(1992) and Kinney et al.~(1993), and the IUE spectral atlas of galaxies in Kinney et al. In particular, the spectrum bears a striking resemblance to IUE spectra of starbursting galaxies such as NGC7552 and others in Kinney et al. However, the ultra-violet continuum of cB58 is many hundred times brighter than those of nearby starburst galaxies.  The absorption features in the spectrum can be attributed to both stars and the interstellar medium (e.g., see Kinney et al.~1993; and Leitherer, Robert, \\& Heckman 1995, here after LRH, for a detailed discussion). Strong stellar lines include \\SiIV, \\CIV, and N$\\,${\\sb IV}$\\,\\lambda$1720 from O stars, and \\AlII~and \\AlIII~from A stars. Dominant interstellar lines include \\SiIIa, \\OISiII, \\CII, \\SiIIb, \\CIV, \\FeII, and \\AlII. Massive hot stars develop strong stellar winds with velocities up to 2000 to 3000 km s$^{-1}$. Hence, it is expected that lines predominantly arising from hot stars are broader than those from the interstellar medium. Although our spectral resolution is relatively poor ($\\sim 900$ km/sec), the widths of the strong stellar lines are consistent with being produced by massive hot stars. For example, \\SiIV~and \\CIV, strong lines associated with O stars,  are significantly broader than \\CII~and \\SiIIb, lines that are attributed primarily to the interstellar medium. In both spectra both the line shape and the shifted centroid of the \\CIV~line are strong evidence for the classic P-Cygni profile expected from early O stars. The lack of an emission line at Ly$\\alpha$ is consistent with starburst galaxies and HII regions seen locally (Hartman et al.~1988).  The morphology and the exponential profile of the optical image indicate that the object is galaxy-like, with a linear size of $\\sim$25 kpc. The scale length of $\\sim$3.5 kpc is entirely consistent with typical $L^*$ galaxies. The smooth light distribution, the lack of measurable color gradients or significant morphological distortion, and the extended nature of the absorption line in the long-slit spectrum all indicate that the young stars are distributed relatively uniformly over the whole object, and that the starburst activity is not localized.  The high luminosity and global nature of the starburst activity argue that this galaxy most likely has very recently undergone a starburst which is massive enough to be the initial burst associated with the formation of this galaxy. In the following sections, we examine in more quantitative detail this interpretation and its implications.  \\subsection {Models }   Constructing detailed and definitive models of the star formation history of a galaxy is difficult, even at such a young age. There are many uncertainties in input parameters for this type of modeling; e.g., initial mass function (IMF), metalicities, and the nature of dust. These uncertainties are greatly accentuated because the PG is being observed at a time when the universe is expected to be very different from the present. Nevertheless, the current data, though limited, allow us to set interesting limits on some of the parameters.  For any star forming galaxy, three basic components are usually required for modeling: an existing population of stars, on-going star formation, and dust. Each of these components has many parameters which are not well constrained. And often, depending on the data on hand, some important parameters are degenerate. For instance, it is often not possible to separate the effects of stellar age and dust using colors.  For first order modeling of the stellar population of the galaxy, we compare the broadband photometry with the spectral energy distributions (SEDs) from  GISSEL models of Bruzual \\& Charlot (1993). The range of models are then investigated in more detail in conjunction with the spectroscopic data. It should be noted that such modeling can only be considered as approximate, as it is certain that young galaxies at these redshifts will not have the near-solar metal abundance on which these models are based.  We have not used the continuum of the SIS spectrum for SED model fitting because of uncertainties in the fluxing. However, we note that the  SIS spectrum agrees with the $g$, $V$, and $r$ photometry very well, although it differs by about 25\\% (too bright) at $I$. The discrepancy in $I$ is  probably the result of light loss in the calibration standard star due to the narrow slit used; such a loss may have been much less severe for the resolved galaxy. The accuracy of the photometry is verified by comparing data from 3 cluster galaxies (at $z=0.373$) to the SEDs of evolved galaxies from Bruzual \\& Charlot (1993).  For dust extinction, we use a slightly modified version of the average LMC extinction curve from Fitzpatrick (1986). We have interpolated over the 2175\\AA~hump in the LMC extinction, based on the fact that in red end of the SIS spectrum, which are equivalent to 2200\\AA~rest, we see no evidence for the expected prominent dip due to this feature. We note that many local star burst galaxies do not show this bump in the extinction in their IUE spectra also (Kinney et al.~1993).  Figure 7 shows a comparison of single starburst models from Charlot \\& Bruzual (1993) and the broad-band colors, which have been corrected for differing amounts of extinction. Unfortunately, extinction and stellar age are degenerate in modeling the colors. Assuming that there is no dust, the optical photometry indicates a single-burst model at an age of about 400 Myr. This can then be considered as the upper limit of the age of the star formation, and will be referred to as the {\\it 400-Myr model}. We note that these results are based on the ``standard'' GISSEL default models with an upper mass cut-off of 125 \\MSUN; changing this cut-off does not affect any of the conclusions based on the SED.  As more extinction is added, younger ages fit the data better. The spectral energy distribution can be fitted by an arbitrarily young model up to an extinction of about $E(B-V)\\sim0.4$, beyond which the SED is definitely too steep even for the hottest stars. Figure 7 also shows two models with extinction corrections. For $E(B-V)=0.18$, the best fitting age is 200 Myr. As an example of a very young model, using $E(B-V)=0.3$, the data fit an age of 5 Myr or less.  Within our observational error bars, models for constant star formation essentially look identical to the very young single-burst model at these wavelengths. Figure 8 shows the comparison of constant star forming models and the data with an extinction correction of $E(B-V)$=0.3. In the observed optical bands, the SED is not able to distinguish models of ages between zero and 500 Myr. Observations in the near-IR bands should be able to set an upper limit to the age of the constant star formation models. We note that a SED with less extinction will not fit any of the constant star formation models. We shall refer to the dusty young extremes of these models as the {\\it single-burst young model} and the {\\it constant formation young model}, or generically as the {\\it young models}. We note that for all the models discussed so far, an additional much older population, if it exists, does not contribute significantly to the flux at the observed optical bands. In the following, we discuss the models in more detail.  The 400-Myr model fits the UV SED very well. The smooth appearance of the galaxy and the lack of color gradients are consistent with the low dust content required by this model. However, this limiting model has a number of serious problems. First, we note that assuming a standard IMF, this model is expected to fade to an evolved galaxy of about --25.3 mag in $V$ (rest), or about 25$L^*$, a highly improbably luminosity for a present-day galaxy, even for a BCG. The ancestor of such a rare galaxy would have an extremely low probability of being discovered in our survey (see Section 4.3). We note that somewhat younger (e.g. 200 Myr) single-burst models suffer from the same problem. The increase in luminosity from the extinction correction negates most of the compensation from the larger fading expected from the younger age. A flatter IMF, or a low-mass cut off which may have been seen in some low-redshift starburst galaxies (e.g., see Charlot et al.~1993), could lower the expected evolved rest $V$ band luminosity. Preliminary $K$ band photometry (to be reported in a future paper) may be consistent with such a scenario.   A more immediate problem is based on the spectroscopic data. The 400-Myr model, or any single-burst models older than about 10 Myr, cannot explain the entire stellar population. Our spectra show strong evidence of P-Cygni profile in \\CIV~arising from massive winds from early O stars. Hence, star formation must be on-going or has stopped less than 10 Myr ago, and either scenario must have a significant effect on the UV continuum.  The young models are favored by the the existence of C$\\,${\\sb IV}~P-Cgyni profile. The depth of the broad absorption component of \\CIV~and the height of the emission component are qualitatively consistent with line profiles of starburst models in LRH. Since the 3 parameters of age, IMF slope, and upper mass cut-off are correlated, a number of combinations of these 3 parameters produce profiles that are similar to the observed one. In particular, two continuous formation LRH models strongly resemble our line profile: one has a Salpeter IMF with an upper mass cut-off of 40 \\MSUN, and the other has a steeper IMF slope (3.0) with an upper mass cut-off of 80\\MSUN. Both have an age of 9 Myr or greater (at which time the creation and death of massive O stars come into equilibrium). In Figure 9 we plot the \\CIV~profile from the MOS spectrum overlaid by the model profile having a Salpeter IMF with an upper mass cut-off of 40\\MSUN. The agreement is excellent. Models that are significantly younger, with flatter IMFs, or having more massive upper mass cut-offs, have too large  P-Cygni profiles; while models that are steeper in the IMF or lower in upper mass cut-off have too small  profiles.  Similarly, we find that for the single-burst scenario, the line profile is well-described by models with ages of between 5 to 7 Myr for the whole range of IMF slopes and upper mass cut-offs in the LRH models. We note that ``single-burst\" star formation is not physical, but is used as a description for a relatively short star formation episode that has already stopped at the time of observation. As an illustration, in Figure 9 we also plot a LRH model with the standard IMF and an upper mass cut-off of 80 \\MSUN~at an age of 8 Myr after a single star burst. It clearly does not fit the observed profile, indicating that star formation could not have stopped completely for  more than 8 Myr.  The constant star formation and single-burst models in theory can be distinguished by the \\SiIVb,$\\,$1402 lines. According LRH, the 5 Myr single-burst models should have relatively stronger P-Cygni profiles than the constant star formation models, due to enhancement from supergiants in the former. The Si$\\,${\\sb IV} lines in the PG spectrum are heavily dominated by interstellar contributions. There is a possible indication of an emission component of a P-Cygni profile in the \\SiIV~line;  however, there is no evidence of a blue broad wing from the \\SiIVb~line. This may indicate a preference of the continuous star formation models over the very young single-burst models. Higher quality data are required for a more definitive conclusion. However, we note that based on a probability argument, it is not highly plausible that we would observe this galaxies within such a short time since its cessation of star formation.  Although the C$\\,${\\sb IV} P-Cygni profile appears to be in agreement with some LRH models, a strong caveat is that they are based on solar metallicity. Hence, these results should not be taken too literally, but rather as indicative. For lower metalicity massive stars, it is expected that the stellar winds will be less strong due to fewer C$^{3+}$ ions. Hence, qualitatively, the upper cut-off mass could be higher, the IMF could be flatter, or the age could be younger. We also note that if there is a significant additional underlying continuum, the intrinsic P-Cygni profile would be larger relative to the true young star continuum.  For the young models, we can use the 1500\\AA~continuum to estimate the star formation rate based on the models of LRH. The observed $V$ (which is equivalent to almost exactly 1500\\AA~rest) magnitude of 20.6 corresponds to a luminosity of  $4.3 \\times 10^{42}$ erg s$^{-1}$\\AA$^{-1}$, assuming no dust correction. Depending on the upper mass cut-off and the slope of the IMF, LRH predict 1500\\AA~luminosities between $10^{38.95}$ to 10$^{40.6}$ ergs s$^{-1}$ \\AA$^{-1}$ at an age of 9 Myr for a constant star formation rate of 1 ${\\rm M}_\\odot$ yr$^{-1}$. Taking a value of $10^{40.04}$ ergs s$^{-1}$ \\AA$^{-1}$ (for a model with  standard IMF slope and a 40 \\MSUN~cut off), ignoring extinction for the time being, and attributing all of the observed $V$ luminosity to the constant star forming component, we obtain a star formation rate of $\\sim400$ ${\\rm M}_\\odot$ yr$^{-1}$. This rate can be considered as a very generous upper limit for any on-going star formation in a dust-free model. For the constant star formation young model, the flux at 1500\\AA~has to be corrected for extinction by a factor of about 12.4, implying a star formation rate of 4700 ${\\rm M}_\\odot$ yr$^{-1}$. Since a typical $L^*$ galaxy has a stellar mass of $\\sim 5\\times10^{10}{\\rm M}_\\odot$, a galaxy forming stars at this rate will assemble into a relatively massive 2$L^*$ galaxy of 10$^{11}$ ${\\rm M}_\\odot$ in stellar mass in about 20 Myr. However, it should be remembered that the 1500\\AA~luminosity is highly dependent on the parameters used in the star formation model. For example, with a flatter IMF slope of 1.5 and an upper mass cut-off of 80 \\MSUN, the star formation rate can be lowered by a factor of 3.6, requiring about 75 Myr to build up a 10$^{11}$ ${\\rm M}_\\odot$ galaxy. This sets the range of the age of the galaxy in a constant star formation model to be on the order of 10 to 100 Myr.  Similarly, we can apply the LRH models to estimate the size of a young single-burst model. For a single burst of 10$^6$ ${\\rm M}_\\odot$ yr$^{-1}$, the  1500\\AA~luminosity is 10$^{39.0}$ erg s$^{-1}$\\AA$^{-1}$ for the standard LRH model (Salpeter IMF and upper mass cut-off of 80\\MSUN) at an age of 7 Myr. Using the observed 1500\\AA~luminosity, the size of the single burst is estimated to be 4.3$\\times 10^{9}$ ${\\rm M}_\\odot$, assuming no extinction. The continuum again must be corrected by a factor of 12.4 for extinction for the young single-burst model, bringing the burst strength to 5.3$\\times 10^{10}$ ${\\rm M}_\\odot$. Thus, for the young single-burst model, the star formation from the burst alone  amounts to about a $L^*$ galaxy.  In both cases the young models allow the galaxy to be observable relatively easily for about 5 to 10 Myr after the end of star formation, at which time the observed $V$ band photometry will drop by about 1 to 2 mag. Using the extinction required and the K-correction from GISSEL models and assuming passive evolution from the time of observation, we estimate that the single-burst formation model will produce an evolved galaxy about 7 mag fainter in rest $V$, creating a present-day galaxy with an absolute rest $V$ magnitude of --21.4, or about 1.7$L^*$. This luminosity is consistent with the mass from the estimates of the star formation rate.  Although it appears that the young models afford very reasonable fit to both photometric and spectroscopic data, there are, however, a few difficulties that need to be addressed. Several problems are related to the short time scale expected for cB58 to build up a relatively massive galaxy, especially for the lower end of the age range of 10 to 20 Myr. First there is the problem with the creation of metals that we see so prominently in the spectrum. Although this can be explained by postulating a much fainter (at 1500\\AA) older component that is completely hidden by the new star formation, such a model would require the present-day descendant of this galaxy to be even more luminous. Second, the relatively smooth morphology and lack of color gradients also argue against such a young age. It is difficult to imagine such homogeneity in a galaxy after only about 10 Myr into its initial star formation, since  the typical dynamical time for a sizable galaxy is about 10$^8$ yrs. We note that both of these problems can be potentially alleviated by the flat IMF models, which have star formation rates such that it will take about a 100 Myr to assemble a massive galaxy.  The young models are also not consistent with all of the observational data. First, although the observed SED fits the young models reasonably well, it is discrepant at the blue end at the 2$\\sigma$ level. A huge on-going star burst produces a very steep blue continuum with F$_\\lambda ~\\sim~  \\lambda^{-2.5}$. The observed continuum is essentially flat around 1500\\AA~with a roll-over around 1300\\AA. The roll-over is seen in both the fluxed SIS spectrum and the $g-r$ color. The standard dust extinction curve is not able to reproduce the steep continuum shape required by a young model at the $g$ band without having to resort to additional absorption beyond 1300\\AA. However, we note that the amount of extra extinction required is within the uncertainty in the standard extinction curve. Second, there is some evidence that that O stars may not completely dominate the continuum light. In the spectrum, a strong \\AlII~absorption line is seen. This line likely has a significant stellar component as it is broader than the pure interstellar lines, such as \\CII~(see Figure 1). The \\AlII~line could be an indicator of a component of early A type stars in the continuum. Additional data in the IR bands may allow us to put a limit on the size of an older component.  Dust is often thought to be associated with forming galaxies or star bursts. It is likely that our PG candidate does not contain an extraordinary amount of dust. The $E(B-V)=0.3$ required by the young models is not an outlandish amount and is consistent with the average values that are seen towards O star associations in the LMC (LRH). The smoothness of the optical images and the lack of significant color variance argue against significant patchiness due to dust. The lack of Ly$\\alpha$ from the HII regions expected from the young stars requires only a small amount of dust as an explanation.  In summary, our optical photometry is able to limit the age of the galaxy to younger than $\\sim$ 400 Myr, with the age critically dependent on the extinction correction, which is limited to $E(B-V)<0.4$. However, the 400-Myr model is highly unlikely because of the object's very large luminosity. Moreover, spectroscopic signatures of a P-Cygni profile for \\CIV~require that there be a significant number of O stars, and hence a component that is either forming stars or has an age younger than $\\sim$10 Myr must exist. Simple standard models over the whole range of ages all have their respective problems, and none appears to be able to explain all the observations so far. This may be an indication of a non-standard IMF or dust absorption law. Such a conclusion is not surprising, as there seems little reason to believe that star formation and dust absorption in a primeval galaxy should operate in a manner identical to those in nearby galaxies. The most likely scenario is that this galaxy is somewhere between 50 and 100 Myr in age, contains a small amount of relatively uniformly distributed dust, and is forming stars at a high rate, probably with a non-standard IMF.  \\subsection {Estimate of Sky Density}  We have shown that we can interpret the galaxy cB58 as an early-type galaxy observed within 100 Myr of its major star formation episode. Is this object the first of a new, but relatively common, class of high-redshift object? How likely is such a galaxy be observed serendipitously? It is of interest to make an order-of-magnitude estimate of the density of such objects on the sky to see if it is probable to expect the CNOC cluster redshift survey to find one.  There are many factors that enter into the detectability of a PG: the volume density, the luminosity, the recognizability of the object (i.e., observable spectroscopic signatures), and the time interval over which these conditions are favorable. We can estimate the expected number of such objects using our knowledge of galaxy density and the observational parameters of the CNOC survey. However, it should be noted that many of the parameters used are highly uncertain.  For an object similar to cB58, there is a relatively small redshift window in which the object is observable. In order for this object to be identified unambiguously, the spectral region that is rich in absorption lines must be in the observed band, especially when this type of object has not been observed before. This spectral region ranges from \\SiIIa~to \\AlII. Ly$\\alpha$ is not included because of the rapid drop in the continuum there. To the red of \\AlIII, there are no significant absorption lines up to well past 2000\\AA~rest. To be able to detect a similar PG, let us assume that we need to cover the minimum of \\CIV~for the low redshift limit, implying $z\\sim2$. Similarly, we require \\CIV~at the red end of the spectrum for the high redshift limit, implying $z\\sim3$.   The CNOC1 redshift survey covers 0.66 square degrees. The volume between $z=2$ and 3 subtended by 1 sq deg is about  8.4$\\times$10$^6$ Mpc$^3$. Using the present date co-moving density of galaxies derived from integrating the luminosity function of Loveday et al.~(1993), we expect $\\sim$1600  galaxies in the volume sampled by CNOC that are as bright as 2$L^*$, the expected evolved luminosity of the young models. If we assume that galaxy formation is uniform between $z=5$ and 2, then approximately 1/2 of all galaxies would be formed in the interval of $2<z<3$. Of these 800 galaxies, about 20\\% may be expected to be early-types with most of their stars formed in a short-duration large burst, reducing the number to 160. These forming galaxies will be in an ultra-luminous phase for only a short period of time. Assuming this to be $\\sim$ 50 Myr, then about 1/45 of them will be visible at any one time, bringing the number down to 3.5. The CNOC survey has an average completeness of about 40\\% between  $r=20$ and 21 mag (mostly due to the many low-redshift clusters in the sample). Hence, the expected number is about 1.5. The interpretation of this object as a PG {\\it can be} consistent with its serendipitous discovery in the CNOC survey. Using the above estimates, we might expect to find between 10$^{0\\pm1}$ such galaxies per square degree at magnitudes brighter than $r$=21.  Of course, our estimate contains many uncertainties of a factor much larger than 2 each. Within the extreme range of models that we discussed, there may be a factor as large as 100 introduced in the uncertainty of the sky density. As an example, if we use the 200-Myr single burst model, the observable life-time is increased by a factor of about 4, but it is more than offset by the factor of over 1000  fewer galaxies expected to be as bright as 10$L^*$. Other large uncertainties include the volume sampled, which differs by a factor of 3 between $q_0=0.1$ and 0.5, and the fraction of $L^*$ galaxies that were formed at $2<z<3$. For instance, if only 10\\% of $L^*$ galaxies were formed in this redshift period, then we would  expect a factor of 5 fewer such objects. If the bulk of these galaxies were formed at higher redshift, a fainter spectroscopic survey further out in the red is required. Also uncertain is the fraction of galaxies that undergo such a huge burst of star formation. It should be noted that, conversely, additional discoveries such as this should allow us to constrain these important parameters. Our order-of-magnitude estimate here is merely to demonstrate that it is not entirely unreasonable for a survey such as CNOC to discover such a PG candidate serendipitously.   \\subsection {Possible Effects of Lensing}  The main thrust of the interpretation that cB58 is a galaxy at the very early stage of its history is its enormous luminosity, and hence star formation rate. Gravitational lensing can cause the luminosity of an object to increase many fold. An example of a PG candidate being lensed, and hence not being as luminous as was first thought, is IRAS10241+4724 (Rowan-Robinson 1991, Broadhurst \\& Leh\\'ar 1995, Eisenhardt et al.~1996) at $z$=2.29. IRAS10214+4724 was first detected in a ground-based image, with substantially poorer seeing (1.5$''$) than ours. However, higher resolution images, both ground-based IR and HST optical images, show that the object is lensed with an estimated brightening of about a factor of 30 to 100 (Eisenhardt et al.~1996). Another example of a lensed high-redshift galaxy is arc \\#384 of in Abell 2218 with a spectrum similar to cB58 at a redshift of 2.51 (Ebbels et al.~1996), discussed in Section 4.5.  Our PG candidate is 6\\arcsec~from the cD galaxy in MS1512+36, a relatively poor (equivalent to Abell richness 0) cluster at $z=0.373$ with a velocity dispersion of 690 km s$^{-1}$ (Carlberg et al. 1996). The proximity of the BCG to the protogalaxy candidate requires an examination of the possible effect of lensing on the luminosity of the object. However, we note that even a magnification of a factor of several in the luminosity does not severely diminish the interpretation that this object is being observed at a very luminous phase of its life time.  Large magnification usually occurs in the presence of a very strong image shear, as is the case for arcs. Our images of the PG under exceptional ground-based seeing show no sign of distortion, and hence is very unlikely that it is being strongly sheared, or multiply imaged. Any lensing similar to that seen in IRAS10241+4724 would have been easily detected in our $V$ image with 0.65$''$ seeing. The axes of the galaxy are neither tangential nor radial with respect to the center of the cD galaxy, which again argues against the possibility of strong lensing. In addition, an examination of both our images and the deep image in Gioia \\& Luppino (1994) reveals no sign of other gravitationally lensed arcs near the cluster. Based on the axis ratio of the galaxy of $<$2 and the very smooth and regular morphology of the object, an upper limit of a magnification of 2 from weak lensing can be established. Such a low upper limit does not change the conclusion with regard to the high luminosity of the galaxy.  The fact that this object is completely resolved in both directions is the primary argument against gravitational lensing. However, it is possible, although very unlikely, that strong magnification can result without shear. This can occur when there are two subcritical lenses superposed in such a way that the image lies at a location where the shear of one lens is at right angles to the shear of the other. In particular, if the cD is off-center and outside the critical radius of the cluster, then in the directions perpendicular to the cD-cluster center line there will be two positions of considerable magnification, but with little image shear. An alternate possibility is that if the cD and cluster potential have a common center (at least in projection), then a subcritical cluster without a central cusp has a small zone of large magnification, close to the cD and with little image shear and no counter images. (L.~Williams, private communication). In both cases a very precise alignment of two lensing mass profiles possessing a narrow range of parameters is required. At this point, we cannot rule out these possibilities, despite the low probability. However, the poorness of the cluster probably precludes a very large lensing magnification even if the geometry is satisfied.  \\subsection {Comparison to Other Objects at High $z$}  Most objects found at high redshift have been discovered as a result of their non-stellar radiation. These include quasars, high-redshift radio galaxies, and IRAS galaxies. They are typically identified by emission lines, especially Ly$\\alpha$, and stellar spectral signatures have been difficult to observe. Another major class of objects that are inferred to be at high redshift are the objects responsible for the quasar absorption lines; however, little is known about their stellar content.  The object bearing a close spectral resemblance to cB58 is Hawaii-167 at $z=2.33$, discovered by Cowie et al.~(1994). This object was observed as part of their deep $K$-band survey. The optical spectrum shows strong absorption lines in the UV, similar to cB58, and the $I$ band magnitude is nearly identical. However, the similarities end here, and there are major differences. Foremost is that Haw167 is unresolved. Furthermore, it is much redder in the optical: $B-I\\sim3$ mag for Haw167 vs $g-I$=0.7 mag. (Note that $B$ at $z=2.33$ is identical in rest wavelength to $g$ at $z=2.72$.) This color difference is borne out by the drop at about 1800\\AA~ rest in the spectrum of Haw167, which is not found in cB58. Egami et al.~(1995) interpret Haw167 as a dust-enshrouded QSO with a star burst. Morphologically cB58 is clearly not a quasar, and  the difference in color may be due to either dust or a somewhat older stellar population in Haw167.  Cowie et al.~(1995) have also discovered a significant population of strong [O$\\,${\\sb II}]$\\lambda$3727 galaxies at $z$ between 1 and 1.7. They interpret these as massive star forming  galaxies with a star formation rate of $>10{\\rm M}_\\odot$ yr$^{-1}$. It is unlikely that cB58 is a member of this population. More likely, the objects found by Cowie et al.~represent continuous star forming galaxies, (hence, probably later morphological types), whereas cB58 may be representative of galaxies with a massive initial burst of star formation creating most of the stellar mass.  Recent observations of arcs in Abell 2218 have discovered one at $z=2.51$ with an absorption-line spectrum (Ebbels et al.~1996). This arc has a spectrum that  resembles very closely that of cB58, but with weaker absorption features. However, the intrinsic luminosity of this lensed galaxy is several magnitudes fainter than cB58 (Kneib et al.~1996), hence, this galaxy is most likely undergoing a more minor star-burst."
+    },
+    "9602/astro-ph9602135_arXiv.txt": {
+        "abstract": "We investigate in detail the parameter space of active-sterile neutrino oscillations that amplifies neutrino chemical potentials at the epoch of Big Bang Nucleosynthesis. We calculate the magnitude of the amplification and show evidences of chaos in the amplification process. We also discuss the implications of the neutrino chemical potential amplification in the Big Bang Nucleosynthesis. It is shown that with a $\\sim 1$ eV $\\nue$, the amplification of its chemical potential by active-sterile neutrino oscillations can lower the effective number of neutrino species at Big Bang Nucleosynthesis to significantly below 3. ",
+        "introduction": "Neutrino oscillations have been suggested to explain several experimental results and if proven true, they will represent a significant step toward physics beyond the standard model of particle physics \\cite{Langacker}. Mixings between active neutrinos ($\\nue$, $\\numu$ or $\\nutau$) and sterile neutrinos (hypothetical neutrinos that do not interact with known particles via the strong, weak or electromagnetic interactions) are one possible source of neutrino oscillations. The most stringent constraints on the parameters of active-sterile neutrino mixings come from cosmological and astrophysical considerations \\cite{Dolgov,Enqvist,Cline,Shi,Sigl}. In particular, the active-sterile neutrino oscillations at the epoch of Big Bang Nucleosynthesis (BBN) have been investigated extensively and tight constraints have been obtained based on the primordial $^4$He abundance in our universe. Interestingly, it was recently pointed out by Foot, Thomson and Volkas \\cite{FTV} that some parameter space of the active-sterile neutrino mixings can amplify neutrino asymmetries (neutrino chemical potentials) so that the previous constraints on active-sterile neutrino mixings based on BBN can be alleviated \\cite{FTV,Foot}. Also an asymmetry in the electron neutrino sector as large as $\\sim 0.1$ at the time of BBN can change significantly the BBN prediction of the primordial $^4$He abundance.  In section 2 of this paper we expand the original investigation of Foot, Thomson and Volkas \\cite{FTV}, by calculating the parameter space that amplifies neutrino chemical potentials and the magnitude of the amplification. But instead of relying on a simplified equation that only applies outside the resonant regime and when neutrinos are incoherent, we analyze the problem based on the original equations in the density matrix formalism, both analytically and numerically. Our analyses reveal many interesting features of the amplification process that cannot be revealed by the simplified approach. For example, the neutrino asymmetry can be oscillatory long after the initial resonant crossing. There are evidences which suggest that the oscillatory asymmetry is chaotic. As a result, although the order of magnitude of the final neutrino chemical potential is readily predictable, the sign of the chemical potential is very sensitive to the mixing parameters and the input parameters of numerical calculations. This oscillatory behavior and evidences of a chaotic amplification are probed in section 2.  In section 3, we discuss two implications of our results in section 2 and show how active neutrinos as dark matter candidates--having $\\sim 1$eV mass-- can lower the effective number of neutrino species in BBN to significantly below 3 by mixing with a lighter sterile neutrino. ",
+        "conclusions": "In summary, we have calculated the parameter space of active-sterile neutrino mixings that amplifies neutrino chemical potentials, and the size of the amplification. Results are summarized in figure 3. By exploring the sensitivity of the amplification to the initial condition, the mixing parameters and the calculational parameters, and by analyzing the Lyapunov exponent of the system, we showed evidences that the amplification process is chaotic. We have also discussed the implications of our results on BBN. It was shown that a $\\nue$ chemical potential of order 0.1$kT$ could be achieved by either a mixing between $\\sim 1$ eV $\\nue$ and a lighter sterile neutrino, or a mixing between a $\\sim 1$ eV $\\numu$ (or $\\nutau$) and a lighter sterile neutrino coupled with a mixing between almost degenerate $\\nue$ and $\\numu$ (or $\\nutau$). Such a chemical potential in $\\nue$ can lower the effective number of neutrino species to significantly below 3."
+    },
+    "9602/astro-ph9602073_arXiv.txt": {
+        "abstract": "We present a detailed cross--correlation (CCF) and power spectrum re--analysis of the X--ray light curves of the bright Seyfert 1 galaxy NGC5548 obtained with {\\it EXOSAT}.  The 0.05--2 keV and 1--4 keV light curves are cross--correlated with the 1--9 keV and 4--9 keV light curves respectively. We discuss how spurious time lags can be introduced by systematic effects related to detector swapping as well as the switching on and off of the instruments. We also find strong evidence that one of the ME detectors was not working normally during the second part of the March 1986 observation.  When these effects are taken into account, the CCF peaks are in all cases consistent with the absence of delays between X--ray variations at different energies. This is unlike the results found by several authors based on the same data.  The power spectra of the 1--9 keV light curves are calculated and a detailed search for quasi periodic oscillations (QPOs) carried out on these spectra by using a new technique for the detection of periodic (or quasi--periodic) signals even in the presence of source noise variability. No significant peaks are found above the 95\\% confidence detection threshold, except during the second part of the March 1986 observation, most probably as a consequence of the ME detector malfunctioning. We discuss and compare our results with those of Papadakis \\& Lawrence (1993a). ",
+        "introduction": "NGC5548 is a bright, close--by (z=0.017) Seyfert 1 galaxy which was extensively studied in different bands of the electromagnetic spectrum. Large variability of both lines (optical--UV) and continuum have been reported making the source an important laboratory for exploring the mechanisms of spectral formation in AGNs. In particular the study of correlations and time lags between the various spectral components represent an important technique for constraining the geometry of the emitting regions (e.g. Mushotzky et al. 1993, and references therein). >From a systematic study of IUE spectra (1200--3000 \\AA ~ Clavel et al. 1991) a strong correlation between emission lines and continuum variability was established, with lines responding to the continuum variations with delays  of 10--70 days, depending on the degree of ionization of the species (a higher ionization corresponds to a smaller delay). Systematic UV--X--ray observations indicated a strong correlation in the continuum variability in the two bands with un upper limit of $\\leq$ 6 days to any delay (Clavel et al. 1992). This, together with the simultaneous optical UV continuum variations, showed that at least a component of the optical--UV continuum should be generated by reprocessing of the X--rays, rather than by the intrinsic disk variability, which should be characterized by longer time--scales (Molendi, Maraschi \\& Stella  1992). ROSAT observations of a soft X--ray flare with correlated variability in the UV, but without a corresponding change in higher energy X--rays (Done et al. 1995, Walter et al. 1995) made  apparent the complexity of the  processes occurring in the object. New simultaneous observations at various wavelengths are currently being analysed (e.g. Korista et al. 1995 and references therein).  The importance of reprocessing on both cold and warm material is confirmed by the observation of a Fe K fluorescent line  at a centroid energy of $\\sim$~$6.4$~keV ,  and  of a Fe absorption edge  at $\\sim$~$8$~keV superimposed to a rather complex continuum (Nandra et al. 1991). This  rich observational scenario  may be accounted for by models where a hot corona above the accretion disk provides the hard X--ray photons. At the same time these photons are in part reprocessed by an accretion disk which in turn generates the photons  for the Compton cooling of the electrons in the hot corona (e.g. Haardt \\& Maraschi 1993, $\\dot{Z}$ycki et al. 1994).  Because of its $\\sim 4$ day orbital period, allowing long uninterrupted exposures and its wide spectral range (0.05--10 keV), the EXOSAT satellite was particular apt to study variability and reprocessing in the X--ray band on time scales of tens of hours or less. For a systematic study of the variability of Seyferts and QSO contained in the EXOSAT database we refer to Grandi et al. (1992) and Green, Mc Hardy \\& Lehto (1993).  NGC 5548 was observed with EXOSAT 12 times in 1984--86 with a total exposure of $\\sim$200 hrs. Flux variations up to a factor of 4 and 3 between the various observations and within each observation, respectively, were clearly seen. A detailed analysis of the EXOSAT light curves was performed by several authors. Variability on  timescales of hours was studied by Kaastra \\& Barr (1989, hereafter KB), who searched for delays between variations in the soft (0.05--2 keV) and medium (2--6 keV) energy X--ray light curves, suggesting that soft X--rays variations lead by $\\sim 4600\\pm 1200$ {\\it s}. This result, if confirmed, would have profound theoretical implications, favoring models, such as those mentioned above, where the medium energy X--rays are produced by Compton scattering off the UV/soft X--ray photons by electrons in a hot corona. Walter and Courvoisier (1990) re--examined the same data by using a different analysis technique substantially confirming the results of KB.  Papadakis \\& Lawrence (1993a, hereafter PL) performed a power spectrum analysis the EXOSAT ME light curves and reported the likely detection of quasi--periodic oscillations (QPOs) in 5 out of 8 observations. They suggested that the frequency ($\\sim$2~mHz) of the QPOs increases with source intensity, whereas the fractional root mean square amplitude of variability decreases as the source brightens. Since this behaviour is similar to that of e.g. compact galactic X--ray binaries, PL suggest that intensity--correlated QPOs in NGC5548 may also arise from instability or variability in an accretion disk around a massive black hole.  Because of the quality of the EXOSAT light curves and the importance of the physical inferences derived from them, we feel justified in presenting a new and independent analysis of relatively old data. This is done in the light of some systematic uncertainties which may arise in the analysis of the EXOSAT light curves from relatively faint sources (such as most AGNs). These uncertainties are discussed in greater detail in a paper on the BL Lac object PKS 2155--304 (Tagliaferri et al. 1991, hereafter Paper I).  In Section 2.1 we summarise the  characteristics of the EXOSAT light curves of NGC5548.  Our CCF analysis for different X--ray energy bands is described in Section 2.2. Details on our search for QPOs are given in section 2.3. The conclusions are in Section 3. ",
+        "conclusions": "Our re--analysis of the CCFs of EXOSAT ME light curves of NGC5548 {\\co does not confirm} the claim of KB of a $\\sim 5000$~s delay between the medium and soft X--rays variations. This was considered as a strong argument in favour of models where medium energy X--rays are produced by scattering of softer photons. Our results do not exclude this possibility. We note however that the 1990 ROSAT observations (Nandra et al. 1993) detected a variability pattern hardly consistent with very soft X--ray variations (0.1--0.4 keV) preceding the variations of somewhat harder X--rays (1--2.5 keV). This indicates the complexity of physical processes occurring in the source.  Our power spectrum analysis does not confirm the detection of QPOs in the mHz range reported by PL. In particular we have shown that the only power spectrum peak with a significance of $>95\\%$ most probably results from the malfunctioning of one ME detector. The argument of PL according to which the black hole mass of NGC 5548 has an embarrassingly low value of a few hundred thousand solar masses, loses its validity.  While the results of this paper are essentially ``negative\", we hope that our work contribute illustrating subtle effects which may yield spurious results in the analysis of X--ray light curves from AGNs."
+    },
+    "9602/astro-ph9602031_arXiv.txt": {
+        "abstract": "We investigate hydromagnetic turbulence of primordial magnetic fields using  magnetohydrodynamics (MHD) in an expanding universe. We present the basic, covariant MHD equations, find solutions for MHD waves in the early universe, and investigate the equations numerically for random magnetic fields in two spatial dimensions. We find the formation of magnetic structures at larger and larger scales as time goes on. In three dimensions we use a cascade (shell) model, that has been rather successful in the study of certain aspects of hydrodynamic turbulence. Using such a model we find that after ${\\cal O}(10^9)$ times the initial time the scale of the magnetic field fluctuation (in the comoving frame) has increased by 4-5 orders of magnitude as a consequence of an inverse cascade effect (i.e. transfer of energy from smaller to larger scales). Thus {\\it at large scales} primordial magnetic fields are considerably stronger than expected from considerations which do not take into account the effects of MHD turbulence. ",
+        "introduction": "It has been suggested that primordial magnetic fields might arise during the early cosmic phase transitions \\cite{pt}, and recently it has been shown that magnetic fields are indeed a stable feature of  the electroweak phase transition \\cite{davies}. In a first order phase transition magnetic fields could also be generated when bubbles of the new vacuum collide, whence a ring of magnetic field may arise in the intersecting region \\cite{kibble}. It has also been suggested that at very high temperatures the ground state of a non-abelian particle theory is a ``ferromagnetic'' vacuum with a permanent non-zero magnetic field \\cite{us}. The plasma of the early universe has a high conductivity so that a primordial magnetic field would be imprinted on the comoving plasma and would dissipate very slowly \\cite{dissip}. Such a field could then contribute to the seed field needed to understand the presently observed galactic magnetic fields \\cite{fields}, which have been measured both in the Milky Way and in other spiral galaxies, including their halos. Typically the observed present day magnetic field is of the order of ${\\cal O}(10^{-6})$ G.  Locally the primordial field could be very large; it is limited only by the magnetic energy density contribution to primordial nucleosynthesis, which constraints the field at the time of nucleosynthesis by $\\rho_B=B_{\\rm rms}^2/8\\pi\\le 0.3\\rho_\\nu$, where $\\rho_\\nu$ is the energy density of 3 species of massless neutrinos \\cite{dario}. Because flux conservation implies that $B_{\\rm rms}\\sim R^{-2}$, where $R$ is the scale factor, the field could have been much stronger at earlier times. On dimensional grounds, a typical value of the magnetic field fluctuation should be $B_{\\rm rms}\\sim T^2$, so that at the time of the electroweak phase transition one could locally obtain fields as high as $10^{24}\\,$G. Depending on how such a strong, random magnetic field scales at large distances, it could \\cite{poul} be the seed field needed to explain the magnetic fields observed on the scale of galaxies and larger.  However, even assuming that a primordial magnetic field is created at some very early epoch, a number of issues remain to be worked out before one can say anything definite about the role of primordial fields in generating galactic magnetic fields. At earliest times magnetic fields are generated by particle physics processes with length scales typical to particle physics. If the inflation hypothesis proves correct, then after inflation rather long correlation lengths are possible \\cite{ven}. The question is if it is at all possible for the small scale fluctuations to grow to large scales, and what exactly is the scaling behaviour of $B_{\\rm rms}$ or the correlator $\\langle B(r+x) B(x)\\rangle$. Even in an inflationary scenario it would be of interest to see if the relatively large scale can grow even further. To study these problems one needs to consider the detailed evolution of the magnetic field  to account for such issues as what happens when uncorrelated field regions come into contact with each other during the course of the expansion of the universe. In general, turbulence is an essential feature of such phenomena. These questions can only be answered by considering magnetohydrodynamics (MHD) in an expanding universe.\\cite{frozenin} It is the main purpose of this paper to investigate the subsequent development of the primordial field. Expressed in a general way, our conclusion turns out to be that MHD turbulence is operative, and hence the scale of magnetic fields is considerably larger than one would expect if MHD turbulence was ignored. This means that the previous estimates of the strength of the primordial magnetic field ``today\" need to be reconsidered.  We begin by posing the basic equations and consider certain simplified models. A full 1+3-dimensional numerical simulation would be desirable, but is beyond the scope of the present paper. In Sec.~II we derive the relativistic MHD equations for relativistic plasma, which is appropriate for the very early universe. In Sec.~III we discuss the appearance of waves in relativistic MHD. To illucidate the various MHD effects pertaining the early universe, we also present numerical solutions to the MHD equations for a two-dimensional slice. In Sec.~IV we study a cascade model, that reflects important properties of fully three-dimensional turbulence. The cascade model has been rather successful in ordinary hydrodynamics. We find that in the early universe magnetic energy is transferred from small scales to large scales. We also compute the correlation function $\\langle B(r+x) B(x)\\rangle$ in the cascade model. In Sec.~V we offer an interpretation of our results. ",
+        "conclusions": "In the two-dimensional case we found that, starting from a small scale magnetic field, magnetic structures develop at progressively larger scales. This process of self-organization corresponds to an inverse cascade of magnetic energy and helicity. Using then a cascade (or shell) model to study three-dimensional MHD turbulence we were able to follow this inverse cascade over much longer times. Such a cascade model has been rather successful in the study of hydrodynamic turbulence.  The possibility of an inverse cascade means that the scale of fluctuations of the primordial magnetic field increases much beyond its original scale given by particle physics. Taking the parameter $l_0$ as a measure of the coherence length of the magnetic field, we see that there is an increase in the coherence of 4 to 5 orders of magnitude. This means that previous estimates of the field strengths in various mechanisms for generating a primordial field should be revised accordingly. For example, let us consider the estimate by Vachaspati in ref. \\cite{pt}. Taking the area average one has the estimate \\begin{equation} B_r\\sim gT^2/4N, \\label{vas1} \\end{equation} where $N$ is the number of steps needed to reach a given scale in terms of the ``fundamental\" scale at which the field is generated. In this case the fundamental scale is the electroweak scale \\cite{pt}. Proceeding as in ref. \\cite{pt} one has $N\\sim 10^{24}$ today, if the relevant scale is of order 100 kpc. However, due to the MHD corrections $N$ is, from a conservative point of view, reduced because of the effect of turbulence, and one would instead have $N\\leq 10^{19}$, which reduces the stochastic decrease of $B_r$. It should be emphasized that from our calculations one can only say what happens up to a time of order $10^9~ t_{\\rm EW}$, so presumably $N$ is considerably below $10^{19}$ today.  The turbulent nature of the magnetic field may have interesting effects on the various phase transitions in the early universe. Also, the inherent shift of energy from small to large scales may be of interest in connection with the density fluctuations due to the magnetic energy.  Of course, the cascade model is a {\\it model} of the real 1+3 dimensional MHD turbulence. Its successful application in many, widely different nonlinear physical problems suggests, however, that it might also be applicable to the primordial magnetic fields of the early universe. Therefore we believe that its indication of the strong increase in the coherence scale of the primordial  field should be taken seriously, and that further, more detailed studies are warranted."
+    },
+    "9602/astro-ph9602010_arXiv.txt": {
+        "abstract": "We report the first results of a long term program aimed at investigating the photometric properties of the cores of Abell and X-ray selected (EMSS, \\cite{emss}) clusters of galaxies. We observed 77 clusters of galaxies in the redshift range $0.05\\leq z \\leq 0.25$ in the Gunn $g$, $r$ and $i$ filters: in this paper we present the data on 59 clusters with good absolute photometry and on another group of 8 clusters  with acceptable relative photometry. For all these clusters we show color--magnitude diagrams in the two colors: when, as in most cases, the early type galaxy sequence is identifiable, we compare it with the expectation from the Virgo c--m relation (\\cite{vs}) and find that the Virgo relation holds to $z\\sim0.2$. We do not find any sign of active evolution in cluster galaxies since that epoch, nor in the percentage of blue galaxies or in the early type galaxy colors, if we accept that the scatter is of the order of 0.2--0.3 mag with respect to the expectations on the basis of standard k--corrections. We point out the presence of a certain number of anomalously red galaxies in the $r-i$ color, which are too red with respect to their $g-r$ color to be normal field galaxies at a redshift higher than the cluster one. Finally we briefly compare a few properties of the two subsamples of optically selected and X--ray selected clusters. ",
+        "introduction": "\\subsection{The Sample}  The sample we observed comprises: (a) 39 Abell clusters; (b) 21 EMSS clusters; (c) 7 EMSS clusters which are also included in the ACO catalog. ACO and EMSS clusters have been selected in such a way to mantain a similar redshift distribution. Furthermore, Abell clusters cover all morphological types (according to the Bautz--Morgan classification as reported in the ACO catalog) and richness classes. We secured the redshifts for two clusters previously at unknown distance (spectroscopy of galaxies in the cluster sample will be presented in another paper). We are thus able to present data on 67 clusters with measured redshift. \\par Figure 1 a to d gives a pictorial representation of the known properties of our cluster sample: the redshift distributions of the 39 Abell clusters and of the 21+7 EMSS clusters; the X--ray luminosity distributions of the 17 Abell clusters observed by the {\\it Einstein Observatory} and of the 21+7 EMSS clusters; the Bautz--Morgan type and the richness (as defined by the C parameter in ACO) distributions of the Abell clusters. The X--ray luminosity of the Abell clusters has been computed from the IPC count rate, as obtained from the {\\it Einstein Observatory Catalog of IPC X--Ray Sources}, assuming a Raymond--Smith spectrum with a temperature of 6 keV, as it was done in the EMSS. \\par Table 1 reports this same information in unbinned form for all the 67 clusters in the sample. Column 1 gives the cluster name (either the Abell number or the EMSS truncated source coordinates or both), columns 2 and 3 give the catalog cluster center coordinates for the epoch 1950.0, column 4 gives the redshift: the two new redshifts so far unpublished are marked with an asterisk, while redshifts for the other Abell and EMSS clusters are taken from \\cite{aco} and from \\cite{stocke} respectively; columns 5 and 6  the Bautz--Morgan type and the number of galaxies between $m_3$ and $m_3+2$ as reported in ACO (the C parameter) respectively, and finally in column 7 we give the X--ray luminosity, in units of 10$^{+43}$ ergs s$^{-1}$.   \\subsection{The Observations}  The cluster multicolor photometry was carried out in the years 1990--1994 at the 2.1m telescope of the Mexican National Astronomical Observatory at San Pedro Martir, equipped with a CCD of 384x576 23$\\mu$ pixels, resulting in a field of view of $\\sim$ 2x3 arcmin. Since December 1993, a new 1024x1024 22$\\mu$ pixels CCD has been available, thus 12 clusters have been observed with a larger field of view ($\\sim$ 4x4 arcmin). For 10 clusters we made two or more pointings, to explore a larger area. The total number of nights allocated to this program has been 34, 18 of which proved to be photometric. All clusters have been observed through the three {\\it g, r} and {\\it i} Gunn filters with typical exposure times between 20 minutes and 1 hour per filter. When the total exposure time was greater than 20 minutes per filter, the observation was splitted into 2 or 3 exposures, to reduce cosmic ray contamination. The average seeing was $\\sim$ 2 arcsec. \\par Standard stars in the Gunn filters (taken from the lists of \\cite{tg} and \\cite{wade}) were observed all through the nights between cluster exposures to derive flux calibrations. During photometric nights, the accuracy in the zero point is of the order of 0.05 mag. Some clusters were observed in conditions which turned out to be non-photometric. In these cases shorter exposures on the same target were successively made in good weather conditions, and images have been calibrated by matching the magnitudes of stars or compact galaxies in the field. \\par The journal of the observations is given in Table 2, for the 67 clusters in the sample. In column 1, the cluster name is given (an asterisk marks clusters observed in non-photometric conditions), in column 2 the total area explored in Mpc$^2$ at the cluster redshift, in column 3 the number of galaxies detected in all three filters. Columns 4 to 15 contain the information pertaining to each filter separately: filter id, the exposure time in seconds, seeing in arcsec, and the limiting magnitude. Such limiting magnitude has been computed for each image, and represents the fainter apparent magnitude a pointlike source can have to be detected with a signal to noise ratio of three. The images generally allow to detect objects as faint as 22.5 {\\it r} mag in all the three filters.  \\subsection{Data Reduction}  Data reduction was performed in a standard way: after the usual bias subtraction, flat fielding and cosmic ray removal, images in the same filter have been registered and summed. \\par By careful inspection of the images, we produced lists of objects for each filter. These lists were complemented with the ones produced by the MIDAS/INVENTORY search algorithm run with a low detection threshold and crosschecked among the three independent images. Only objects detected in all the three filters were retained in each cluster catalog. As a second step we rejected all objects which were below any of the limiting magnitudes reported in Table 2. This procedure does not produce clear--cut magnitude limited catalogs, because the effective limits are dependent on the colors of the objects, but exposure times were such that normal ellipticals of $M_r=-18$ are generally included in the catalogs of galaxy clusters at $z=0.2$. \\par Finally, bright stars have been removed from the catalogs. Convolution of a Point Spread Function, representative of our images, with exponential and de Vaucouleurs' galaxy profiles have shown that, beyond the complexity of the two parameter dependence, galaxies could be separated from stars if the objects were brighter than $m_r\\sim21$ (\\cite{pocar}). Thus bright stars were removed from the catalogs on the basis of their FWHM. We recall that in the field, stars are a factor 2.5 less numerous than galaxies at $m_r=22$ (see e.g. \\cite{tyson}). \\par Two kinds of magnitudes have been computed for each galaxy: an apparent aperture magnitude, and a metric magnitude. Metric apertures are computed within 10 kpc radius. Magnitudes, and errors, have been computed using the IRAF package {\\it apphot}. In deriving galaxy magnitudes and colors, we have considered both seeing effects and the presence of nearby objects. To include most of the signal in the aperture, angular diameters have been chosen to be twice as large as the largest seeing FWHM in the three filters. In the cases when the aperture magnitude is strongly contaminated by nearby objects, the 10 kpc magnitude has been used. When both magnitudes were affected by nearby objects, the smallest aperture has been used, but a minor weight was given to the point. Magnitudes have been corrected for atmospheric extinction and galactic absorption on the basis of the measured hydrogen column density along the line of sight (\\cite{stark}), converted to A$_g$, A$_r$ and A$_i$ following \\cite{sgh}. Statistical errors on magnitudes (and colors) depend on the signal to noise ratio,and therefore also on apparent magnitude. In Table 3, we illustrate the statistical accuracy of the color determination as a function of magnitude for the whole sample: 90\\% of the galaxies brighter than  m$_r$=21 have errors on $g-r$ smaller than 0.068 mag and on $r-i$ smaller than 0.058 mag; 80\\% of the galaxies brighter than m$_r$=22 have errors on $g-r$ smaller than 0.09 mag and on $r-i$ smaller than 0.076 mag. \\par The number of galaxies in each cluster field for which we have multicolor information ranges from 13 to 175. Colors are available for a total of 3843 galaxies.  \\subsection{Photometric Accuracy}  To check our photometric accuracy, we have compared our magnitudes with those obtained by \\cite{hs} (HS), who observed 185 Brightest Cluster Galaxies in the {\\it g} and {\\it r} Gunn filters. For this comparison, we use $H_0 = 60~km~s^{-1}~Mpc^{-1}$ and magnitudes within 16 kpc radius apertures, as they did. No galactic extinction was applied, as done by  HS. For the 12 galaxies in common with HS, in Figure 2 a and b we plot our magnitudes vs. theirs. The dotted line represents the $y=x$ relation. A linear regression performed on the data of Figures 2 a and b give results that are compatible with the $y=x$ relation, although a slight shift is present in the {\\it g} magnitudes, in the sense that $<g_{this work}> = <g_{HS}> -0.20 \\pm 0.14$ (the same relation for the {\\it r} filter is $<r_{this work}> = <r_{HS}> -0.004 \\pm 0.13$). Such discrepancy could be due to the slightly different filter+detector response we have in the {\\it g} band. We note however that the offset is comparable to the statistical and systematic errors one incurs in when computing magnitudes of large extended objects like the BCGs are. No comparison can be made for the {\\it i} filter, for lack of other observations. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602156_arXiv.txt": {
+        "abstract": "The conventional picture for the origin of the polarization of a supernova is based on a model of Thomson or resonance scattering of photons traveling through an aspherical supernova atmosphere. Positive detection of intrinsic polarization in SN 1987A is then interpretated as evidence of an asymmetrical supernova atmosphere. We show here a different view based on the scattering of the supernova light by a dusty circumstellar material (CSM), or the ``light echo'' effect. At a given epoch after the explosion, the observed photons consist of both those propagating directly from the supernova and those scattered by dust particles in the CSM. Polarized light can be produced if the distribution of the dust particles is aspherical. The model can reproduce both the time evolution of the observed broad band polarization of SN 1987A and major features of the polarization spectra. It is also successful in providing a natural model for the early infrared light curve, in particular the observed 30 day delay of the IR maximum compared to the maximum of the bolometric light curve. ",
+        "introduction": "Polarimetry of SN 1987A was obtained from 2--262 days after explosion (Jeffery 1991a, and references therein). Several groups (Jeffery 1991b; H\\\"oflich 1991) modelled the observations in terms of scattering by the supernova atmosphere (Sutherland \\& Shapiro 1982; McCall 1984). The degree of polarization predicted by these models is generally a decreasing function of time after explosion, as can be anticipated from the fact that the atmosphere gets progressively optically thin as the supernova atmosphere expands. Eventually  electron scattering becomes unimportant and no polarization can be produced. These models are capable of reproducing the observations prior to day 100 after explosion (H\\\"oflich 1991; Jeffery 1991b). The observed degree of polarization, however,  increased sharply from 0.2\\% to about 1.6\\% at around day 100 (Jeffery 1991a; Barret 1988). This behavior is difficult to reconcile in terms of the photospheric scattering model.  Correction for interstellar polarization (ISP) is crucial, but is unfortunately quite uncertain. We base our study on the polarimetric data compiled and corrected for ISP by Jeffery (1991a), where the ISP is estimated from unpublished observations made at La Plata Observatory at around day 600. The data show that the degree of polarization increased from about 0.1\\% at day 2 to about 1\\% at $\\sim$ day 250. The polarization position angle remains fairly constant during all these observations. The time variability of the polarization indicates that the observed polarization is intrinsically related to the supernova event, regardless of the ISP in the direction to the supernova.  We show here that dust scattering plays an important role in producing the polarized light of SN 1987A. In \\S 2.1 the survival of dust particles in the vicinity of the supernova is discussed. The dust scattering process and how that leads to the observed polarization is presented in \\S 2.2. \\S 3 compares the model with various observations. Conclusions and a brief discussion are given in \\S 4. ",
+        "conclusions": "The polarimetry of SN 1987A can be modeled in terms of scattering by a dusty blob. Both the polarimetry curve and the early infrared light curve can be successfully fitted by a CSM dust scattering blob.  This analysis is, however, far from complete. Substantial improvements can be achieved by more detailed modeling. In particular, the polarization position angle also shows clear changes across spectral lines. This suggests that a single scattering blob can not account for all the observed polarization. Two or more blobs may improve the model. Another possibility is that the observed polarization is a combination of the atmospheric scattering model and the CSM dust scattering model presented here.  The hypothesized dust blob may be related with the mystery spot discovered in earlier speckle observations (Cropper et al. 1988). They are located at almost identical position angles of $\\sim\\, 15\\arcdeg$. Further analysis may provide some interesting tests for the authenticity of the speckle sources (Nisenson et al. 1988; Meikle et al. 1987).  Our recent observations show that polarization $\\sim\\ 1\\%$ may be a common phenomena in Type II supernovae (Wang et al. 1996). A correct picture for the origin of the polarization is therefore critical for the understanding of supernova explosion mechanisms in general, and also in establishing Type II supernovae as extra--galactic distance indicators using the expanding photosphere method. If dust scattering is really important in producing polarization, the present study suggests that Type II supernovae are more likely to be associated with CSM dust than Type Ia. It is especially likely that dust scattering plays a role in SN 1993J which was strongly polarized and has a substantial CSM. Polarization measurement can thus set strong constraint on the enviroment of supernovae (Wang \\& Wheeler 1996).  This research is supported by NSF Grant AST 9218635 and NASA grant GO 6020 from the Space Telescope Science Institute which is operated by the Association of Universities for research in Astronomy."
+    },
+    "9602/astro-ph9602142_arXiv.txt": {
+        "abstract": "We have found that the bulge of the large, nearby Sb galaxy NGC~7331 rotates retrograde to its disk. Analysis of spectra in the region of the near-IR Ca II triplet along the major axis shows that, in the radial range between 5$''$ and $\\sim$20$''$, the line of sight velocity distribution of the absorption lines is has two distinct peaks, and can be decomposed into a fast-rotating component with v/$\\sigma$ $>$ 3, and a slower rotating, retrograde component with v/$\\sigma$ $\\sim$ 1 -- 1.5. The radial surface brightness profile of the counter-rotating component follows that of the bulge, obtained from a 2-dimensional bulge-disk decomposition of a near-infrared $K$-band image, while the fast rotating component follows the disk. At the radius where the disk starts to dominate the isophotes change from being considerably boxy to very disky.  Although a number of spiral galaxies have been found that contain cold, couterrotating disk, this is the first galaxy known to have a boxy, probably triaxial, fairly warm, counter-rotating component, which is dominating in the central regions. If it is a bar seen end-on, this bar has to be thicker than the disk. We find that NGC~7331, even though it is a fairly early-type spiral, does not have a conventional, co-rotating bulge. The fact that the inner component is retrograde makes us believe that it was formed from infalling material, in either stellar or gaseous form (e.g. Balcells \\& Quinn 1990). Another possibility however is that the structure has been there since the formation of the galaxy. In this case it will be a challenge to explain the large change in orientation of the angular momentum when going outward radially. ",
+        "introduction": "In the last decade many elliptical galaxies have been found which display complicated multiple-component structure in their kinematics  and, in some cases, also in their photometry. This structure has been interpreted as being due to mergers (e.g. Balcells \\& Quinn 1990), to components which were formed later from processed gas (Bender \\& Surma 1992), or as triaxial systems in projection (Statler 1991). Sometimes these multiple components give rise to disky isophotes, so that they can be found using photometry. However in general they are much better seen in the analysis of the line-of-sight velocity distribution (LOSVDs). Examples of galaxies with multiple components can be found in e.g. Franx \\& Illingworth (1988), Rix \\& White (1992) and Bender \\etal (1994). Typically they consist of a bright, pressure-supported body, and a fainter, fast-rotating central disk, although this is not always the case (e.g. Balcells \\& Carter 1993).  The kinematics of stars in spiral galaxies and lenticulars has been much less studied. Bulges are generally thought to be supported by rotation (Kormendy 1993) and no non-rotating bulge has been found up to now, except maybe for the very small bulge of NGC~4550 (Franx 1993). They rotate in the same direction as the disks, although in general they are also partly pressure-supported. Recently however, a number of cases has been found in which a fraction of the stars in the disk rotates in a retrograde way. These include NGC 4550 (Rubin \\etal 1992, Rix \\etal 1992), NGC 7217 (Merrifield \\& Kuijken 1994) and NGC 3593 (Bertola \\etal 1996). These counter-rotating disks are thought to have originated from infalling gas, accreted in retrograde orbits in a plane close to the equatorial plane (Merrifield \\& Kuijken 1994).  In this paper we present the discovery of a different type of galaxy: one in which the bulge rotates retrograde to the disk. We have found this by analyzing LOSVDs along the major axis of the large, nearby Sb (Sandage \\& Tammann 1981) galaxy NGC~7331. Here two components are seen, of which the slow, counter-rotating one corresponds to the boxy bulge of this galaxy. Explaining the formation of a system like this will be a challenge to galaxy formation theories. Section 2 of this paper gives details about the observations. In  Sections 3 and 4 we describe our analysis of morphology and  the LOSVDs. The discussion and conclusions are given in  Section 5. ",
+        "conclusions": "To summarize the observations: the inner parts of the galaxy consist of a boxy component, dominating the inner 5$''$. It shows position angle twisting, rotates retrograde to the rest of the galaxy, and is rounder. Outside 5$''$ the disk dominates. This component is much colder, is probably flat, and responsible for the disky isophotes. At appr. r=100$''$ the surface brightness profile becomes shallower.  The boxyness and position angle twist make it very likely that we are seeing a triaxial object counter-rotating to the main body of the galaxy. It could be that we have a configuration like in NGC~4736 (M\\\"ollenhoff \\etal 1995), where the bar is oriented end-on towards us, or alternatively, that the galaxy has no bar, and that we are seeing a triaxial bulge. The steep rise in the surface brightness profile and the low ellipticity show that it can't be a classical, flat bar with a Freeman-type profile (Freeman 1966). From this data we cannot distinguish between these two alternatives any further. Whatever the case, the galaxy does not contain a large, co-rotating bulge; if it has a bar a small bulge, as in NGC~4736, can be hidden.  It is likely that the origin of the central component is external. Evidence for this would be the fact that it is retrograde, and maybe the boxyness. Simulations for elliptical galaxies have shown that it is possible to create a central, counter-rotating body as a result of a stellar merger (Balcells \\& Quinn 1990). Another possibility is that the central component is formed as a result of an instability in the disk of counter-rotating stars, like the one of NGC~7217 (Merrifield \\& Kuijken 1994). In that case however the counter-rotating disk must have accreted from outside first. A third possibility might be that the counter-rotating stars were formed at the formation epoch together with the rest of the galaxy. Up to now no models have been proposed where two counter-rotating components were formed from one protogalactic cloud, and where the angular momentum changes so drastically as a function of radius. We might observe here something similar to a counter-rotating bar within a bar, for which Friedli (1996) has shown that it can be very longlived.  Further evidence for a triaxial potential in the center of this galaxy comes from the H$\\alpha$ velocity field by Marcelin \\etal (1994). They find the typical Z-shape in the isovelocity contours, characteristic of non-circular potentials. The galaxy shows some low-level activity (Ho \\etal 1995). On larger scales, the motion of the gas is quite regular (von Linden \\etal 1996, Begeman 1987).  We now ask ourselves how unique this galaxy is. Studies using e.g. the Fourier quotient method (Tonry \\& Davis 1979) generally find that the bulge is dynamically supported by rotation, and rotates in the same direction as the disk (Kormendy 1993). If however the bulge to disk ratio diminishes, it becomes difficult to find any slow-rotating component in this way. To show this, we give the h3 and h4 components, as defined by van der Marel \\& Franx (1993), for NGC~7331 in Fig.~3. They are similar to h3 and h4 profiles of  many elliptical galaxies (e.g. Bender \\etal 1994). Since however a somewhat fainter, co-rotating component could give the same h3 and h4 as we are seeing here, we need to investigate the line profiles in more detail. By now, there are several spiral galaxies and S0's for which the line profiles have been studied, There are many cases with a hot bulge and a cold disk rotating in the same direction (e.g. NGC 3115 (van der Marel \\etal 1994), NGC 4736 (M\\\" ollenhoff \\etal 1995), NGC 4594 (Wagner \\etal 1989)). But as far as we know, no galaxy has been found where the central and outer component rotate opposite to each other. We conclude that NGC~7331 is special in this respect.  We would like to end advancing some speculations. First, we find here a very boxy, slowly rotating inner component, and a large, fast rotating disk. Many giant elliptical galaxies on the other hand have a boxy, even slower rotating, main body, and a small, rotating disk. Apart from the relative scales, the similarity of both systems is striking, and has to imply that the formation scenario of spirals like NGC~7331 and ellipticals cannot be too different. And secondly, we see a discontinuity in the slope of the surface brightness profile at r=100$''$. Since this galaxy is of type Sb, it is rather peculiar that it doesn't have a larger, co-rotating bulge. Could it be that the component that dominates the light between 5$''$ and 100$''$ was the previous bulge? It however is flat, rotates fast, and is disky, but according to Kormendy (1993) many late type spirals have 'bulges' that look like this. This question obviously is still open, and needs much more investigation."
+    },
+    "9602/astro-ph9602004_arXiv.txt": {
+        "abstract": "We discuss the role of turbulence in cloud and star formation, as observed in numerical simulations of the interstellar medium. Turbulent compression at the interfaces of colliding gas streams is responsible for the formation of intermediate ($\\simlt 100$ pc) and small clouds (a few tens of pc), although the smallest clouds can also form from fragmentation of expanding shells around stellar heating centers. The largest cloud complexes (several hundred pc) seem to form by slow, gravitational instability-driven merging of individual clouds, which can actually be described as a large-scale tendency towards homogenization of the flow due to gravity rather than cloud collisions. These mechanisms operate as well in the presence of a magnetic field and rotation, although slight variations on the compressibility and cloud morphology are present which depend on the strength and topology of the field. In summary, the role of turbulence in the life-cycle of clouds appears to be twofold: small-scale modes contribute to cloud support, while large-scale modes can both form or disrupt clouds. ",
+        "introduction": "The interstellar medium (ISM) is a highly compressible turbulent flow (e.g. Dickman 1985, Scalo 1987, Falgarone 1989) in which turbulent density fluctuations are likely to be ubiquitous. With a few exceptions, however, most treatments of cloud dynamics and formation including turbulence have traditionally considered only the contribution of small-scale turbulent motions for cloud support against gravitational collapse (e.g. Chandrasekhar 1951; Shu et al.\\ 1987; Bonazzola et al.\\ 1987; L\\'eorat et al.\\ 1990; Elmegreen 1991; \\VS\\ \\& Gazol 1995). In actuality, turbulent motions at the scale of clouds or larger may have both cloud-forming and disrupting effects, which may be respectively associated with the compressive and shearing modes of the turbulence. Although early contentions were that, being supersonic, it should dissipate rapidly (e.g.\\ Goldreich \\& Kwan 1974), later studies suggested that the energy sources present in the Galaxy (mainly stellar activity and Galactic differential rotation) could be enough to replenish the turbulence (e.g., Fleck 1980; see the review by Dickman 1985). Stellar energy injection originates mostly from OB winds and supernova (SN) explosions, and the average rates for the Galaxy have been determined observationally by a number of authors (e.g., Abbott 1982; Van Buren 1989).  Cloud and star formation (hereinafter CF and SF, respectively) at the interfaces between colliding flow streams have been discussed analytically by Hunter \\& Fleck (1982) and Elmegreen (1993). Low-resolution hydrodynamical simulations of colliding flow streams were performed by Hunter et al.\\ (1986), and the effects of stellar forcing in the large-scale flow have been simulated by Bania \\& Lyon (1980), Chiang \\& Prendergast (1985), Chiang \\& Bregman (1988) and Rosen et al.\\ (1993). Fully turbulent regimes in the vertical direction in the Galactic disk have been explored by Rosen et al.\\ (1993) and by Rosen \\& Bregman (1995). However, all of these calculations have omitted self-gravity and magnetic fields, and have employed somewhat unrealistic SF schemes, thus rendering it impossible to discuss the full energy budget and the life cycles of the clouds that form.  In this paper we discuss the interplay between turbulence and CF and SF in recent numerical simulations of the interstellar medium (ISM) we have performed, which incorporate stellar and diffuse heating with a more realistic SF scheme, parameterized cooling, self-gravity and large-scale shear (\\VS\\ et al.\\ 1995, hereafter Paper I), and magnetic fields and rotation (Passot et al.\\ 1995, hereafter Paper II). In \\S 2 we briefly describe the model system, in \\S\\S 3 and 4 we discuss results without and with magnetic fields, and in \\S 5 we summarize the main conclusions. ",
+        "conclusions": "Turbulence is a major cloud-forming agent in numerical simulations of the ISM at the kpc scale (Papers I and II). The non-magnetic simulations (Paper I) suggest that intermediate ($\\simlt 100$ pc) and small (a few tens of pc) clouds mainly form at the interfaces of colliding flow streams. The smallest clouds can additionally form from fragmentation of expanding shells around stellar heating centers. Large cloud complexes (several hundred pc) appear to form from gravitational instability. However, since the medium is turbulent and contains already sizable clouds, the formation of the large complexes proceeds through gradual merging of the clouds that is almost imperceptible as one watches the simulation evolve, but is noticeable when comparing epochs separated by several $10^7$ yr.  SF injects energy to the turbulence, and the simulations indicate that a self-sustaining cycle may be established, although the caveat exists that the simulations contain an adjustable parameter, namely the threshold density for SF $\\rhoc$, which controls the stability of the cycle.  The magnetic field introduces quantitative variations in these mechanisms, but does not seem to alter them essentially, although the role of the Parker instability cannot be evaluated by the simulations, as they neglect the vertical direction in the Galaxy. The field modulates the density contrast that can be achieved by turbulent compressions, as well as the morphology of the clouds. In turn, the turbulent magnetic energy is also fed by stellar activity, provided a sizable uniform component of the field is present. The stellar energy injection occurs at comparatively small scales, and energy cascades to larger scales. Whether this is an artifact of the two-dimensionality of the simulations remains to be evaluated by resorting to three-dimensional computations.  In summary, the role of turbulence in the life-cycle of clouds appears to be twofold: small-scale modes contribute to cloud support, while large-scale modes can both form or disrupt clouds.  Finally, it should be pointed out that the simulations so far have not included the energy injection from SNe and OB winds, which can inject energy at significantly larger rates than heating from ionizing radiation from main sequence OB stars. The largest turbulent velocity dispersion that can be imparted to the flow by stellar heating is comparable to the thermal velocity dispersion at the temperature of our ``HII regions'', $\\sim 10^4$ K, which roughly balances gravity for these complexes. In the non-magnetic case, this turbulent energy is enough to blow the complexes apart, but not in the magnetic case due to the ``pressure cooker'' effect. The inclusion of SN/wind energy should cause these clouds to be easily disrupted. Work in this direction is in progress (Gazol et al.\\ 1995)."
+    },
+    "9602/astro-ph9602138_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "Fluorine overabundances in giant stars of spectral types MS, S and C have been reported recently (Jorissen, Smith \\& Lambert 1992; Paper~I). These observations reveal a progressive F enrichment of the envelope as the C/O ratio increases, reaching $[\\mathrm{F/O}] \\approx  1.5$ in some C stars. An independent indication of large F overabundances in carbon stars comes from the recent detection of AlF in the inner envelope of the dust-enshrouded carbon star IRC+10216 (Ziurys, Apponi \\& Phillips 1994). These observations shed light on the much debated question of the nucleosynthetic origin of \\chem{F}{19}. They clearly point towards thermal pulses as a likely site for its production.  The first quantitative proof that \\chem{F}{19} can indeed be produced in thermally pulsing AGB stars has been provided by Forestini et al. (1992; Paper~II). However, this early investigation just considered the first 4 thermal pulses of a single (3~\\Msun, $Z=0.03$) star. As a consequence, it left unanswered several important questions, like the dependence of the level of \\chem{F}{19} production on the stellar mass, metallicity, or pulse number.  The present paper aims at exploring some of these questions. It analyzes a dozen pulses in three stars of different masses and metallicities ($M =$ \\mass{3}, $ Z =$ 0.02; $M =$ \\mass{6}, $Z =$ 0.02; $M =$ \\mass{3}, $Z =$ 0.001). Use is made of an improved code which follows more accurately the structural evolution of the AGB stars and the concomitant nucleosynthesis. It also updates some nuclear reaction data of importance in the study of the \\chem{F}{19} production. In particular, a more accurate rate (de Olivera et al. 1995) for the key reaction $\\ag{15}{N}{19}{F}$ is adopted, as well as a recent re-evaluation (Giesen et al. 1994) of the $\\ag{18}{O}{22}{Ne}$ rate. In addition, the surface fluorine enhancement due to the first and second dredge-ups is also investigated.  The nuclear transformations leading to the production of $\\noy{19}{F}$ in thermal pulses are reviewed in Sect.~\\ref{Sect:nucleo}. The evolution code and its key ingredients are briefly described in Sect.~\\ref{Sect:ingredients}, while Sect.~\\ref{Sect:predictions} presents our calculated \\chem{F}{19} abundances. Our expectations are confronted with the observational data in Sect.~\\ref{Sect:observations}, and conclusions are drawn in  Sect.~\\ref{Sect:Conclusions} ",
+        "conclusions": "  \\smallskip \\noindent i) Neither the first nor the second dredge-ups can account for the observed discrepancy between the solar-system fluorine abundance and that of normal K and M giants.  \\smallskip \\noindent ii) Our predictions of surface fluorine abundances resulting from parametrized third dredge-ups in intermediate-mass stars can only account for the lowest F overabundances observed in AGB stars. That conclusion is expected to hold true for low-mass stars when fluorine production results from secondary \\chem{C}{13} only. No concomitant \\chem{O}{18} overabundances are predicted in our intermediate-mass stars, in agreement with the observations.  \\smallskip \\noindent iii) The largest F overabundances observed in C stars call for an additional source of \\chem{C}{13}, of {\\it primary} origin. A larger \\chem{C}{13} supply could account for the observed overabundances of both F {\\it and} s-process elements. Detailed computations are however necessary to confirm that prediction.  \\smallskip \\noindent iv) According to our models, massive AGB stars are not expected to build up large surface F abundances. Moreover, when operating at $T_6 \\ge 70$, HBB leads to a gradual destruction of the envelope fluorine, thereby limiting the surface F overabundance that would result from the third dredge-up. The large fluorine overabundance reported for the super Li-rich star WZ~Cas (where HBB is supposed to be operating) remains therefore unexplained so far.  \\smallskip \\noindent v) Hot thermal pulses in intermediate-mass, {\\it low-metallicity} AGB stars are less efficient in producing F. The lower \\chem{F}{19}/\\chem{C}{12} ratio from these pulses does not lead to large F overabundances at the surface of these AGB stars. Any large overabundances that would be reported by (future) observations of low-metallicity stars, would require a primary source of \\chem{C}{13}."
+    },
+    "9602/hep-ph9602232_arXiv.txt": {
+        "abstract": "Using the O(4) linear $\\sigma$ model, we address the topic of non-equilibrium relaxation of an inhomogeneous initial configuration due to quantum and thermal fluctuations. The space-time evolution of an inhomogeneous fluctuation of the condensate in the isoscalar channel decaying via the emission of pions in the medium  is studied within the context of  disoriented chiral condensates. We use out of equilibrium closed time path methods in field theory combined  with the amplitude expansion. We give explicit expressions for the asymptotic space-time evolution of an initial inhomogeneous configuration including the contribution of thresholds at zero and non-zero temperature. At non-zero temperature we find new relaxational processes due to thermal cuts that have no counterpart in the homogeneous case. Within the one-loop approximation, we find that the space time evolution of such inhomogeneous configuration out of equilibrium is effectively described in terms of a rapidity dependent temperature $T(\\vartheta)=T/\\cosh[\\vartheta]$  as well as a rapidity dependent decay rate $\\Gamma(\\vartheta, T(\\vartheta))$.  This rate is to be interpreted as the production minus absorption rate of pions in the medium and approaches the zero temperature value at large rapidities.  An  initial  configuration localized on a bounded region spreads and decays in spherical waves with slower relaxational dynamics at  large rapidity. ",
+        "introduction": "The dynamics of relaxation of inhomogeneous field configurations out of equilibrium is an important problem and a common theme in cosmology and high energy physics.  In high energy physics the experimental possibility of studying the chiral and quark-gluon phase transition with high luminosity hadron colliders and upcoming heavy-ion colliders makes imperative the understanding of relaxation and transport processes in extreme environments\\cite{mueller,rajagopal}.  A very exciting possibility\\cite{anselm,bjorken,blaizot1,kowalski,bjorken2} that can be studied in high energy-high luminosity hadron collisions or relativistic heavy ion collisions within the energy range of the upcoming RHIC and LHC colliders, is that after the chiral phase transition, regions in which the chiral condensate is misaligned with respect to the vacuum state are formed.  In these ultra-high energy heavy ion collisions ($\\sqrt{s} \\geq 200 \\mbox{Gev/nucleon}$) a large energy density (a few $\\mbox{Gev}/ \\mbox{fm}^3$) is deposited in the collision region corresponding to temperatures above the critical value for chiral symmetry restoration $\\approx 200 \\mbox{Mev}$. In this situation, it is possible that within a volume of a few $\\mbox{fm}^3$ an inhomogeneous coherent field configuration is formed corresponding to a region in space in which the chiral symmetry is restored.   The idea\\cite{anselm,bjorken,blaizot1,kowalski,bjorken2} is that as these regions cool down they relax towards the equilibrium situation emitting a large number of soft pions.  Such a configuration has been dubbed a `disoriented chiral condensate' (DCC).  Bjorken, Kowalski and Taylor proposed a `baked alaska'\\cite{bjorken2} scenario in which this configuration relaxes via the copious emision of pions strongly correlated in isospin. Such a scenario could also explain the Centauro events observed in cosmic rays in which the ratio of charged to neutral pions is different from $1/3$\\cite{rajagopal}.  A microscopic, field theoretical description of the dynamics of relaxation of these field configurations with typically large amplitudes faces at least two major obstacles.  The first is a non-perturbative treatment of inhomogeneous, large amplitude field configurations.  The second is a consistent description of non-equilibrium processes that can allow a {\\em{real time}} calculation of the evolution.  Recently there has been a surge of interest in the description of non-equilibrium processes both in cosmology and high energy physics. We refer to non-equilibrium processes to those corresponding to the the time evolution of quantum states in field theory which {\\bf are not} hamiltonian eigenstates as well as the time evolution density matrices that do not commute with the Hamiltonian. In particular the  relaxation of strong electric fields\\cite{kluger} and the evolution of homogeneous scalar order parameters during phase transitions in Minkowski and cosmological spacetimes \\cite{boyveg}.  The non-equilibrium aspects of disoriented chiral condensates have been studied within different approximation schemes by several authors\\cite{wilczek1,wilczek2,pisarski,cooper,gavin,mueller2,boydcc}, but mainly focusing on homogeneous expectation values or correlation functions.   However, to the best of our knowledge, these studies focused on the evolution of a homogeneous but time (or proper time) dependent mean field and fluctuations around it but did not address the relaxational dynamics of {\\em inhomogeneous} field configurations including {\\em quantum and thermal} effects.  Moreover,  inhomogeneous field configurations appear in cosmology when topological objects, such as textures or cosmic strings involving inhomogeneous field configurations are considered. Their  relaxational dynamics is thought to have a bearing on the spectrum of fluctuations in the cosmic microwave background radiation \\cite{vilenkin,kibble}.  \\bigskip  The focus of this article is to study the {\\em linear} relaxational dynamics of inhomogeneous, coherent field configurations both at zero and finite temperature in the $O(4)$ linear sigma model. As is known such model is the  relevant field theory  for the description of chiral condensates.  We use out of equilibrium closed time path methods \\cite{S,K,maha,CSHY,kapusta} combined with the amplitude expansion \\cite{B1,B3} to study the time evolution of $\\langle\\Phi(\\vec{x},t)\\rangle$. We consider small initial field amplitudes such that we can keep just the linear term in the amplitude expansion. In such approximation the field evolution equations linearize.  Whereas the focus is on the description of the dynamics within the setting of the chiral phase transition, the analysis presented below applies to a wide variety of physical situations in which a inhomogeneous scalar order parameter relaxes to the equilibrium situation via the production of lighter scalar fields in a medium.  The main idea is that after a phase transition (within the framework of DCC's the chiral phase transition) regions are formed within which the scalar order parameter is inhomogeneous. As the zero momentum component rolls towards the equilibrium value, these inhomogeneous configurations will relax in such a way that the spatial gradients of the field configuration become smaller so as to decrease the energy. This relaxation will be accompanied by the production of quanta, such as pions in the case of DCC's.  Here we study the O(4) linear sigma model in the broken symmetry state, at temperatures below the chiral phase transition as an effective low energy theory  of $SU(2)_{\\rm{L}}\\times SU(2)_{\\rm{R}}$ (up and down quarks) which presumably incorporates the effects of strongly interacting QCD on chiral dynamics\\cite{wilczek1}.  Wilczek and Rajagopal\\cite{rajagopal,wilczek1} have argued that the O(4) linear sigma model effectively describes the same {\\em equilibrium} universality class of QCD with two flavors of light quarks and proposed to study it within the context of DCC. Clearly this simple model misses a great number of degrees of freedom, hadrons, vector mesons, etc. and can only be justified as an effective description at temperatures well below the chiral phase transition when the degrees of freedom with much higher masses are not relevant.  The linear sigma model may also be obtained as a Landau-Ginsburg effective theory from a Nambu-Jona-Lasinio model\\cite{klebansky} which is a popular description of the phenomenology of chiral symmetry at the quark level and which has also been used as an alternative description of DCC's\\cite{bedaque}.  The O(4) linear sigma model has also been studied within the context of texture type configurations deemed relevant in certain cosmological models of structure formation\\cite{vilenkin}, hence our study may be relevant in these cosmological scenarios as well.  We will restrict our study to the case in which the condensate occurs in the isoscalar channel, without isospin violation.  In section II we introduce the model, discuss its range of validity and summarize aspects of non-equilibrium field theory relevant to the calculation. Section III presents the bulk of the results, with explicit analysis of the space-time evolution, including the contribution from thresholds and thermal cuts. We summarize our results in the conclusions section. ",
+        "conclusions": "We have studied the relaxational dynamics of an inhomogeneous condensate fluctuation in the O(4) linear sigma model near the broken symmetry state, both at zero and non-zero temperature.  Explicit expressions are obtained for the self-energies at zero and finite temperature and we point out that at finite temperature there are new relaxational processes with origin in thermal cuts and that are {\\bf only} present in the inhomogeneous case with no counterpart in the relaxation of an homogeneous condensate.  For initial Gaussian fluctuations we have given explicit expressions for the asymptotic space-time evolution of the inhomogeneous fluctuation including the effect of thresholds both at zero and finite temperature. At finite temperature we obtained the decay rate of this inhomogeneous  non-equilibrium configuration, in the medium this quantity is the rate of production {\\em minus} the rate of absorption of pions. The space-time evolution is described in terms of an effective temperature and ``decay rate'' that depend on rapidity. We find to one loop that relaxational processes are described in terms of $T(\\vartheta) = T/ \\cosh[\\vartheta]$ and $\\Gamma(\\vartheta, T(\\vartheta))$, and for large rapidities the finite temperature decay rate approaches the zero temperature limit.  We systematically compute the field behaviour for large times and distances compared with the inverse of the typical mass scale ($M$) in the model. This means $t\\gg 1/M, \\; r\\gg 1/M $ {\\bf and} $ \\tau = \\sqrt{t^2 - r^2} \\gg 1/M$. For $ t=r \\gg 1/M $ the behaviour will be quantitatively different.  For large $ \\tau $, the field $\\phi(\\vec x,t)$ propagates as spherical waves for  an initial wave packet of arbitrary shape concentrated around the origin.  Our asymptotic results can be summarized as follows.  \\begin{itemize}  \\item  {\\bf Stable Case: $M_{\\sigma} < 2 M_{\\pi}$}  \\begin{description}  \\item[$T = 0$.]  The pole contribution $ \\phi_{pole}(\\vec x,t) $ dominates, giving an amplitude that decays as $ \\tau^{-3/2} $ and oscillates as a function of $ \\tau $ with  frequency $M_0$ (physical $\\sigma$ mass). The cut gives contributions smaller than  $ \\phi_{pole} $ by a factor  $ \\tau^{-3/2} $ and oscillating with  frequencies  equal to the threshold positions [eqs.(\\ref{asympole}-\\ref{inhD})].  \\item[$T \\neq 0$.]  The pole contribution dominates for large $\\tau$.  $ \\phi_{pole} $ decays with the  same power  $ \\tau^{-3/2} $ as for $T = 0$, but the amplitude and frequency change quantitatively [see eq.(\\ref{asympoleT})] .  A new cut running from $ -ip $ to $ +ip $ appears at non-zero temperature. It gives a  {\\bf non-oscillating} contribution for large $\\tau$ that decays as $ \\tau^{-5} $  [see eq.(\\ref{asinfi}) ]. \\end{description}  \\item {\\bf Unstable Case: $M_{\\sigma} > 2 M_{\\pi}$}  \\begin{description}  \\item[$T = 0$.]  $\\phi$ is now a resonance (a  pole in the second Riemann sheet) yielding an {\\bf exponentially damped} amplitude that oscillates with frequency $ M_0 $. The damping being the width of the resonance [see eq.(\\ref{corteunst})].  Eventually, for very large $ \\tau $, the power-like tail coming from the cut contribution will dominate over the exponentially damped contribution of the resonance.  \\item[$T \\neq 0$.]  It is similar as for $T = 0$, except that the damping rate becomes a non-trivial function of $t,r,T$ and $M_0$ given by  eq.(\\ref{bodrio}).  The thermal cut with its non-oscillating contribution to $\\phi(\\vec x,t)$ is also present here and dominates for $\\tau \\to \\infty$.  \\end{description} \\end{itemize}  All the asymptotic results hold to first order in the field amplitude. However, since the field vanishes for $ \\tau \\to \\infty $, our results are true for {\\bf any initial amplitude} provided $ \\tau $ is large enough.  The relaxation of an initially Gaussian inhomogeneous fluctuation is described in terms of the spreading of the packet and decay in spherical waves, and we found that the time scale for relaxation, production and absorption of pions is a function of rapidity such that for larger rapidities the relaxational and decay processes are slower.  We are currently extending these studies to the case of non-equilibrium fluctuations produced during the stage of parametric amplification as the $\\sigma$ rolls to the ground state\\cite{mueller2}.  We believe that the techniques developed in this work  and the non-equilibrium aspects  found here will be of  use in a variety of physical contexts. One that comes to mind concerns the emission of Goldstone bosons from an axion string\\cite{davisshell,sikivie}. This emission is important in determining the exact upper bound on the axion decay constant, and thus on whether axion models for solving the strong CP problem are still viable. This problem will entail some modifications of the tools described here, the most notable one having to do with the existence of zero modes for the string configuration."
+    },
+    "9602/astro-ph9602127_arXiv.txt": {
+        "abstract": "NGC 4314 is an early-type barred galaxy containing a nuclear ring of  recent star formation. We present CO(1-0) interferometer data of the bar and  circumnuclear region  with $2.3^{\\prime \\prime} \\times 2.2^{\\prime \\prime}$  spatial resolution and 13  km s$^{-1}$ velocity resolution acquired at the Owens Valley  Radio Observatory . These data reveal a  clumpy circumnuclear ring of molecular gas. We also find a peak of CO inside  the ring within 2$^{\\prime \\prime}$ of the optical center that is  not associated with massive star  formation. We construct a rotation curve from these CO kinematic data and the  mass model of Combes et al.\\ (1992). Using this rotation curve, we have identified  the location of orbital resonances in the galaxy. Assuming that the bar ends at  corotation, the circumnuclear ring of star formation lies between two Inner  Lindblad Resonances, while the nuclear stellar bar ends near the IILR. Deviations  from circular motion are detected just beyond the CO and H$\\alpha$ ring, where the dust lanes along the leading  edge of the bar intersect the nuclear ring. These non-circular motions along the minor axis correspond to radially inward streaming  motions at speeds of $20 - 90$ km s$^{-1}$ and clearly show inflowing gas feeding an  ILR ring. There are bright H\\,{\\sc ii} regions near the ends of this inflow region,  perhaps indicating triggering of star formation by the inflow. ",
+        "introduction": "Inner Lindblad resonances (ILRs) are thought to strongly influence  gas motions and star formation in the central regions of galaxies.  Gas  surface densities often peak near the ILR (Combes et al.\\ 1992; Kenney et  al.\\ 1992), lending support to the  idea that the inward flow of gas along  galaxian bars slows down and piles up between the OILR and the IILR  (Combes 1988; Shlosman et al.\\ 1989).  Studies of the CO (1-0) distribution within the central regions of barred spirals (Kenney et al.\\ 1992; Kenney 1995; Regan, Vogel, \\& Teuben 1995; Rand 1995; Sakamoto et al.\\ 1995) show a variety of molecular gas morphologies near  the ILR, including rings, spiral arms, and ``twin peaks'' with two maxima  where the dust lanes along the primary bar cross the ILR.  In  the case of  the starburst galaxy NGC 3504, the CO peaks near the nucleus, probably  inside the ILR (Kenney et al.\\ 1993).  The dynamical and evolutionary differences responsible for this variety are not yet understood, but are  critical for understanding starbursts and the formation and evolution of  compact circumnuclear stellar disks (Kormendy 1993) and stellar bulges  (Pfenniger \\& Norman 1990).  At present, only a handful of barred spirals have been mapped in CO  at sufficient resolution to study the effects of ILRs on the gas distribution and kinematics. The barred spiral galaxy NGC 4314 is a particularly good galaxy to study because it exhibits evidence for many features  associated with resonances and it is nearly face-on ($i \\sim 23\\deg$). It has a large-scale stellar bar of diameter 130$^{\\prime \\prime}$ (6.3 kpc) and a prominent circumnuclear  ring of  star formation of diameter 10$^{\\prime \\prime}$  (500 pc) that is visible in H$\\alpha$  (Pogge  1989), radio continuum (Garcia-Barreto et al.\\ 1991; hereafter G-B), and  optical color maps (Benedict et al.\\ 1992; hereafter P1). CO (1-0) mapping  of NGC 4314 (Combes et al.\\ 1992) at 5$^{\\prime \\prime}$  resolution revealed the presence   of a molecular ring slightly smaller than that of the ring of star formation. Outside of these rings, a blue elliptical feature of diameter 20-25$^{\\prime \\prime}$  (1 kpc)  is seen in optical color maps, which may correspond to a ring of relatively  young but non-ionizing stars (P1).  This feature, which can be seen on the  unsharp masked optical image shown in Fig.\\ 1 (from P1), is elongated perpendicular to the primary bar, suggesting that it extends outward to the  IILR. Combes et al.\\ (1992) suggest that in this galaxy star formation is  propagating inwards. This is in contrast with the evolutionary scenario  proposed by Kenney et al.\\ (1993) for strong starbursts, in which star  formation devours gas most rapidly in the center, ultimately forming a ring of gas near the ILR.  Inside the H\\,{\\sc ii} region ring, a nuclear stellar bar of diameter 8$^{\\prime \\prime}$  (400  pc) is seen in {\\it Hubble Space Telescope} (HST) I band data (Benedict et  al.\\  1993; hereafter P2). This nuclear bar is aligned parallel with the large  scale bar, which suggests that the nuclear bar has the same pattern speed as  the primary bar, and lies within the IILR (Binney \\& Tremaine 1987).  In this paper, we present new CO (1-0) maps of NGC 4314, obtained  at twice the spatial resolution as the Combes et al.\\ (1992) data. The  observations are described in Section 2, while the CO distribution and  kinematics are discussed in Sections 3 and 4, respectively.  In Section 5 we  compare NGC 4314 with other galaxies and discuss star  formation in this galaxy. Some general properties of NGC 4314 are  tabulated in Table 1. ",
+        "conclusions": "The major observational results of this paper include \\begin{enumerate}  \\item We confirm the ring-like CO morphology first reported by Combes et  al.\\ (1992).  \\item We discover significant CO near the center of NGC~4314 without  associated recent star formation.  \\item CO associated with dust along the leading edge of the primary stellar bar  is detected at a 3-$\\sigma$ level.  This association is confirmed for the dust bowl region  by a strong correlation of CO intensity with extinction estimated from optical  color maps.  \\item Much of the CO in and near the ring is coincident with dust.  \\item We identify the nested, orthogonal resonances in NGC~4314. We  identify the IILR with the end of the nuclear bar, OILR with the outer extent of  an elliptical distribution of less recent star formation, corotation near the end of  the large-scale bar, and the OLR with the elliptical outer isophotes oriented perpendicular to the bar. The strongest star formation occurs between the IILR  and OILR and closer to the IILR.  \\item A color gradient, from blue near the IILR to red near the OILR, is  likely due to stellar age differences, not reddening. CO emission is strong near  the IILR and weak near the OILR, hence the expected reddening gradient is the  opposite of the observed color gradient.  \\item Radial inflow at speeds of 20-90 km s$^{-1}$ is detected in two spurs of  molecular gas located just outside the star-forming ring. Significant H$\\alpha$ emission  occurs where the inflow spurs meet the ring. \\end{enumerate}"
+    },
+    "9602/astro-ph9602061_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "\\end{center} \\setcounter{equation}{0}  \\baselineskip=24pt  The origin of the X-ray background (XRB) remains one of the most challenging problems in X-ray astrophysics. Figure 1 summarizes the current results on the XRB energy spectra, $I(\\varepsilon)$ below 10\\keV observed with different satellites.  As is known (McCammon \\& Sanders 1990; Fabian \\& Barcons 1992), $I(\\varepsilon) \\sim 10 \\varepsilon^{-0.4}{\\rm keV}\\cdot{\\rm s}^{-1}\\cdot{\\rm sr}^{-1}\\cdot{\\keV}^{-1}$ is a good empirical fit over the range 3 to 10 \\keV, where $\\varepsilon$ denotes X-ray energy in units of \\keV. Both Einstein IPC (Wu et al. 1991; plotted in diamonds) and Rosat (Hasinger 1992; Shanks et al. 1991; upper-left lines) data suggest a large excess below 2\\keV. More recently the Japanese X-ray satellite ASCA (crosses) reported a modest but significant excess soft component below 1\\keV relative to the extrapolation of the above power-law fit in the higher energy band (Gendreau et al. 1995). Thus the existence of the soft excess is well established although its amplitude is still somewhat controversial.  Since Galactic absorption becomes important in the soft X-ray energy band, the origin of this excess component allows several possibilities including Galactic sources, extragalactic point-like sources (Hasinger 1992; Shanks et al. 1991), a diffuse thermal component (Wang \\& MaCray 1993), the accumulation of the thermal bremsstrahlung emission from distant clusters of galaxies (Kitayama \\& Suto 1996).  Cen et al. (1995) found an excess of soft X-ray background below 1 \\keV from their hydrodynamical simulations which properly incorporate the line emissions as well as the thermal bremsstrahlung emission; the excess is mainly originated from the the low temperature and low density plasma surrounding distant clusters of galaxies since cooler background gas produces much stronger line emissions.  Yet another possibility which we propose here is the emission from an X-ray halo of the Local Group (LG). Since the gaseous halos of clusters of galaxies are known to be strong sources of X-rays, it is reasonable to assume that the LG has its own X-ray halo.  \\newpage  \\begin{center} ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602075_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "The origin of the X-ray background (XRB) remains one of the most challenging problems in X-ray astrophysics. Figure 1 summarizes the current results on the XRB energy spectra, $I(\\varepsilon)$ below 10\\keV observed with different satellites.  As is known (McCammon \\& Sanders 1990; Fabian \\& Barcons 1992), $I(\\varepsilon) \\sim 10 \\varepsilon^{-0.4}{\\rm keV}\\cdot{\\rm s}^{-1}\\cdot{\\rm sr}^{-1}\\cdot{\\keV}^{-1}$ is a good empirical fit over the range 3 to 10 \\keV, where $\\varepsilon$ denotes X-ray energy in units of \\keV. In the soft band, however, the observed spectra do not seem to be in perfect agreement among different satellites; both Einstein IPC (Wu et al. 1991; plotted in diamonds) and ROSAT (Hasinger 1992; Shanks et al. 1991; upper-left lines) data suggest a large excess below 2\\keV. More recently the Japanese X-ray satellite ASCA (crosses) reported a modest but significant excess soft component below 1\\keV relative to the extrapolation of the above power-law fit in the higher energy band (Gendreau et al. 1995). Thus the existence of the soft excess is well established although its amplitude is still somewhat controversial.  \\begin{figure} \\begin{center} \\leavevmode\\psfig{figure=szxrb.eps,height=8cm,angle=-90} \\end{center} \\caption{XRB spectra observed with different X-ray satellites and expected contribution from the halo of the Local Group (Suto et al. 1996). \\label{fig:xrb} } \\end{figure}  Since Galactic absorption becomes important in the soft X-ray energy band, the origin of this excess component allows several possibilities including Galactic sources, extragalactic point-like sources (Hasinger 1992; Shanks et al. 1991), a diffuse thermal component (Wang \\& MaCray 1993), the accumulation of the thermal bremsstrahlung emission from the low temperature low-density plasma surrounding distant clusters of galaxies (Cen et al. 1995; Kitayama \\& Suto 1996).  Figure 2 shows an example of theoretical predictions of the contribution of clusters on the XRB in the standard cold dark matter model ($\\Omega_0=1$, $\\lambda_0=0$, $\\Omega_b=0.06$, $h=0.5$) which is normalized by COBE (Sugiyama 1995).  Here $\\Omega_0$, $\\lambda_0$, $\\Omega_b$, and $h$ are the cosmological density parameter, the dimensionless cosmological constant, the baryon density parameter, and the Hubble constant in units of $100${\\rm km/s/Mpc}.  Three lines indicate the theoretical predictions on the basis of the Press-Schechter theory but with different assumption on the cluster evolution (Kitayama \\& Suto 1996 for details).  Also plotted are the observed data from {\\it ASCA} with statistical errors (Gendreau et al. 1995), and the simulation by Kang et al. (1994; triangles). This suggests that the clusters of galaxies can in fact account for the observed soft-excess in the XRB. \\begin{figure} \\begin{center} \\leavevmode\\psfig{figure=ksxrbeds.eps,height=8cm,angle=0} \\end{center} \\caption{The XRB spectra contributed from clusters of galaxies (Kitayama \\& Suto 1996). \\label{fig:ksxrbeds} } \\end{figure} ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602112_arXiv.txt": {
+        "abstract": "The Press--Schechter description of gravitational clustering from an initially Poisson distribution is shown to be equivalent to the well studied Galton--Watson branching process.  This correspondence is used to provide a detailed description of the evolution of hierarchical clustering, including a complete description of the merger history tree.  The relation to branching process epidemic models means that the Press--Schechter description can also be understood using the formalism developed in the study of queues.  The queueing theory formalism, also, is used to provide a complete description of the merger history of any given Press--Schechter clump.  In particular, an analytic expression for the merger history of any given Poisson Press--Schechter clump is obtained.  This expression allows one to calculate the partition function of merger history trees.  It obeys an interesting scaling relation; the partition function for a given pair of initial and final epochs is the same as that for certain other pairs of initial and final epochs.  The distribution function of counts in randomly placed cells, as a function of time, is also obtained using the branching process and queueing theory descriptions.  Thus, the Press--Schechter description of the gravitational evolution of clustering from an initially Poisson distribution is now complete. All these interrelations show why the Press--Schechter approach works well in a statistical sense, but cannot provide a detailed description of the dynamics of the clustering particles themselves.  One way to extend these results to more general Gaussian initial conditions is discussed. ",
+        "introduction": "This paper is mainly concerned with providing a complete description of gravitational clustering from an initially Poisson distribution.  This is not because one thinks it likely that the initial conditions for gravitational clustering in our universe were Poisson.  Indeed, measurements of the spectrum of temperature fluctuations in the microwave background suggest otherwise.  At present, however, analytic understanding of the growth of clustering from more general Gaussian initial conditions is not as detailed as for the Poisson case studied here.  Thus, the Poisson model serves as a convenient toy model with which to study the evolution of nonlinear clustering.  Many of the results obtained in this paper should provide at least qualitative insight into the evolution of clustering from more general initial conditions.  The Press--Schechter approach allows one to estimate the distribution of the masses of virialized clumps at a given epoch directly from the initial density distribution.  It is based on the hypothesis that, provided the initial velocities are sufficiently small (i.e., that the initial field is sufficiently cold), overdense regions in the initial density field will eventually collapse to form nonlinear structures.  Thus, by studying the statistics of overdense regions in the initial field, one can infer properties of the nonlinear distribution (Press \\& Schechter 1974; Bond \\etal 1991; Lacey \\& Cole 1993). One of the main simplifying assumptions of the Press--Schechter approach is that the overdense regions are assumed to collapse spherically.  While this may be a good approximation in the mean, $N$-body simulations of gravitational clustering from initially scale-free Gaussian random fields show that the collapse of overdense regions is seldom exactly spherical. Nevertheless, the Press--Schechter mass functions for these initially scale-free Gaussian fields have been shown to be in good agreement with the distribution of clump sizes that are measured in relevant $N$-body simulations (Efstathiou \\etal 1988; Lacey \\& Cole 1994).  The Press--Schechter excursion set description of clustering from an initially Poisson distribution of identical particles has been derived (Epstein 1983; Sheth 1995).  The probability that a Poisson Press--Schechter clump has $N$ particles is \\begin{equation} \\eta(N,b) = {(Nb)^{N-1}{\\rm e}^{-Nb}\\over N!}, \\label{borel} \\end{equation} where $N\\ge 1$ and $0\\le b<1$, and $b$ is related to the Press--Schechter overdensity threshold $\\delta_{\\rm c}$: \\begin{equation} b = 1/(1+\\delta_{\\rm c}) . \\label{bdelta} \\end{equation} For an initially cold Poisson distribution, the threshold overdensity $\\delta_{\\rm c}$ decreases as the Universe expands in such a way that $b=0$ initially, and $b\\to 1$ as the clustering develops.  If clumps collapse spherically in a universe with critical density, and both growing and decaying modes are present in the linear perturbation theory, then $\\delta_{\\rm c} = (5/3)\\,1.69/a$, where $a$ is the expansion factor (e.g. Bond \\etal 1991; Lacey \\& Cole 1993). Thus, $b$ changes rapidly at first, but at late times its evolution is slowed by the expansion of the Universe.  (See section~2.3 in Sheth~1995 for another description of the evolution of $b$ that shows this same behaviour.) Since $b$ is known to increase monotonically with time, it will be treated as a psuedo-time variable in the remainder of this paper.  Equation~(\\ref{borel}) is known as a Borel distribution (Borel 1942).  Epstein (1983) derived this distribution by studying the properties of level excursions of a Poisson distribution.  His approach is the discrete analog of that which was used later by Bond \\etal (1991) in their analysis of the initially Gaussian random fields.  The Borel distribution can also be derived using simple `cloud--in--cloud' conditional probabilities for a Poisson distribution (Sheth 1995).  This approach is the discrete analog of Jedamzik's (1995) treatment of the Gaussian case.  If the probability that a randomly chosen clump contains $N$ particles is a Borel distribution, then the probability that a randomly chosen particle is in such a clump is $N\\eta(N,b)/\\langle N\\rangle = (1-b)N\\eta(N,b)$, since the average number of particles in a Borel clump is $\\langle N\\rangle = 1/(1-b)$. In the limit of large $N$ and small $\\delta_{\\rm c}$, Stirling's approximation for the factorial term implies that \\begin{eqnarray} (1-b)\\, N\\, \\eta(N,b) &=& {\\delta_{\\rm c}\\over (1+\\delta_{\\rm c})}\\, \\left({N\\over 1+\\delta_{\\rm c}}\\right)^{N-1} \\, {{\\rm e}^{-N/(1+\\delta_{\\rm c})}\\over (N-1)!} \\nonumber \\\\ &\\to& {\\delta_{\\rm c}\\over\\sqrt{2\\pi N}} \\ \\exp \\left(-{N\\,\\delta_{\\rm c}^2\\over 2}\\right) \\label{limn0} \\end{eqnarray} (Epstein 1983; Sheth 1995).  The final expression is precisely that which obtains for a Gaussian density field with white noise initial fluctuations (e.g. Bond \\etal 1991).  This shows that the Poisson distribution studied in the remainder of this paper can be thought of as the discrete analog of the white noise Gaussian studied by Bond \\etal (1991), and by Lacey \\& Cole (1993, 1994).  Sections 2 and 3 study other derivations of the Borel distribution.  These derivations provide new insight into the accuracy and applicability of the Press--Schechter approach.  An analytic expression that is, essentially, the partition function that describes all possible merger histories of a given Press--Schechter clump is derived in Section 3.  Its properties are consistent with those inferred previously (Sheth 1995).  In particular, it is consistent with the physical requirement that, in the limit of very small time steps, the probability that a clump has two progenitors should be an infinitesimal of smaller order than the probability that it has three, or more, progenitors. Section 3 also shows that the partition function, and so the growth of clustering, satisfies an interesting scaling relation.  In Section 4, this scaling relation is exploited to provide insight into the details of the growth of hierarchical clustering.  Section 4 also shows that this Poisson Galton--Watson, Press--Schechter description is in qualitative agreement with the results of the numerical algorithm developed by Kauffmann \\& White (1993). Section 5 discusses ways in which the branching process extension of the Press--Schechter approach that is developed in this paper can be extended to provide a detailed description of clustering from initially Gaussian random fields.  Appendices A, B and C provide details of some of the calculations.  Appendix D contains a derivation of the distribution of counts in randomly placed cells by extending some of the ideas developed in this paper.  Although all the results of this paper are independent of those derived in Appendix D, it has been included because the final result in it (equation~\\ref{ppsd}) is known to be accurate.  Thus, Appendix D provides one natural way in which the usual Press--Schechter analysis may be extended to provide additional information about the evolution of nonlinear gravitational clustering. ",
+        "conclusions": "The partition function that describes the relative probabilities of all possible merger histories of those Press--Schechter clumps which form as an initially Poisson distribution evolves gravitationally, can be written in closed form (eq.~\\ref{prtfnc}).  The partition function is a function of the initial and final epochs (denoted by $b_1$ and $b_2$, respectively).  Since the time evolution of $b$ can be computed (cf. the Introduction) the temporal evolution of the partition function is known.  The counts in cells distribution associated with this Poisson Press--Schechter distribution can also be computed by extending the branching process analogy (Appendix D).  It, too, is a function of $b$, so its temporal evolution is also known.  Thus, the Press--Schechter description of clustering from an initially Poisson distribution is now complete.  It may be worth pointing out that Appendix B shows that the Poisson Galton--Watson branching process specifies the Borel clump size distribution (equation~\\ref{borel}) and the merger probabilities (equation~\\ref{fjk}) uniquely, and also allows one to describe the merger tree completely (equation~\\ref{prtfnc}).  On the other hand, if equations~(\\ref{borel}) and~(\\ref{fjk}) are treated as the only known constraints on the form of the merger tree, then the branching process considered in this paper is not the only way to construct the merger history trees consistent with these constraints (e.g., Sheth 1995).  For instance, one can always construct an ad hoc Monte-Carlo scheme, like that of Kauffmann \\& White (1993), which satisfies equations~(\\ref{borel}) and~(\\ref{fjk}), and is not necessarily consistent with the branching process description of equation~(\\ref{prtfnc}).  For this reason we have argued that, rather than being completely ad hoc, the branching process is, indeed, a reasonable model for the growth of clustering (Figs.~\\ref{perco} and~\\ref{gwbp} and associated discussion).  A Poisson distribution is similar to a white noise Gaussian random field.  So, this result can be related to the evolution of clustering from an initially Gaussian density field that has a scale free power spectrum with slope $n=0$. Can it be extended to describe the growth of gravitational clustering from Gaussian random fields with arbitrary initial power spectra?  One way in which to do this is as follows.  The Borel distribution reduces to the Press--Schechter multiplicity function in the limit as $N\\gg 1$ and as $b\\to 1$.  Similarly, $f(j|k)$ (eq.~\\ref{fjk}) reduces to the $n=0$ Gaussian result provided that $k\\gg 1$, $j\\gg 1$, $k-j\\gg 1$, and $b_1$ and $b_2\\to 1$ (Sheth 1995).  For a given density threshold $\\delta_{\\rm c}$ [eq.~\\ref{bdelta} shows that $b=1/(1+\\delta_{\\rm c})$], Gaussian fields with different power spectra yield different excursion set Press--Schechter mass multiplicity functions.  However, these differences arise solely because the relation between the variance and the mass depends on the power spectrum.  When written directly in terms of the variance, rather than the mass, the Press--Schechter multiplicity functions and the merger probabilities have a universal form that is independent of the underlying power spectrum (e.g. Lacey \\& Cole 1993). This suggests writing the partition function of equation~(\\ref{prtfnc}) in terms of the variance of the Poisson ($n=0$) distribution.  Then, in the large $l_i$ and $b\\to 1$ limits, equation~(\\ref{prtfnc}) should reduce to the Gaussian result.  The appropriate Jacobian can then transform the partition function (written in terms of the variance) into an expression for the masses of the subclumps.  There may be a more sophisticated way to describe the merger histories of clumps that form from Gaussian random fields.  In this paper, branching processes and queueing theory were both used to formulate and solve a problem that was originally posed in terms of trees associated with random walk excursions on a discrete Poisson distribution.  So the question is, Is it possible to extend any of these relations to the continuous Gaussian case?  For example, is it possible to formulate the excursion sets of random walks associated with Gaussian fields, described by Bond \\etal (1991) in their derivation of the Press--Schechter mass functions, in terms of branching processes?  Trees associated with excursions of random walks on Gaussian fields are the subject of current interest (Neveu \\& Pitman 1989; Le Gall 1993).  The $b\\to 1$ limit is known as the `critical' Galton--Watson process (the $b<1$ case considered in this paper is `subcritical').  Many properties of the tree associated with this critical process have been calculated (Aldous 1993 and references therein).  The application of these ideas to the Press--Schechter description of clustering from arbitrary Gaussian initial conditions is in progress.  One final way in which to extend the ideas of this paper to arbitrary initial conditions is to note that the usual Press--Schechter mass functions provide a good, but by no means perfect, description of $N$-body simulations of clustering.  Thus, one might reconsider the percolation model discussed in Section~2.  First, one would compare the percolation mass functions (obtained numerically) for the Poisson case with those of the branching process (the Borel distribution, equation~\\ref{borel}).  If the percolation mass functions provide an acceptable fit to the simulation results, then one could also compute (numerically) and test the percolation merger probabilities and merger tree.  Of course, it is trivial to apply the percolation model to arbitrary initial conditions.  The discussion at the end of Section~2 shows clearly that this branching process extension of the Press--Schechter approach is correct only in a strictly statistical sense.  This is because, unlike in the percolation model discussed in Section~2, there is no direct correspondence between particles in the initial distribution and particles in clumps.  That is, the Galton--Watson branching process, like the excursion set interpretation of the Press--Schechter mass function, is a purely statistical model for the growth of clustering.  To be useful, it relies on the accuracy and applicability of the ergodic hypothesis.  It cannot, and should not, be expected to work on a particle-by-particle basis."
+    },
+    "9602/astro-ph9602089_arXiv.txt": {
+        "abstract": "Smoothed particle hydrodynamics (SPH) is used to estimate accretion rates of mass, linear and angular momentum in a binary system where one component undergoes mass loss through a wind. Physical parameters are chosen such as to model the alleged binary precursors of barium stars, whose chemical peculiarities are believed to result from the accretion of the wind from a companion formerly on the asymptotic giant branch (AGB). The binary system modelled consists of a $3 \\Msun$ AGB star (losing mass at a rate $10^{-6}~\\Msun~{\\rm y}^{-1}$) and a $1.5 \\Msun$ star on the main sequence, in a 3~AU circular orbit. Three-dimensional simulations are performed for gases with polytropic indices $\\gamma=1$, 1.1 and 1.5, to bracket more realistic situations that would include radiative cooling. Mass accretion rates are found to depend on resolution and we estimate typical values of 1-2\\% for the $\\gamma=1.5$ case and 8\\% for the other models. The highest resolution obtained (with 400k particles) corresponds to an accretor of linear size $\\approx 16\\Rsun$. Despite being (in the $\\gamma = 1.5$ case) about ten times smaller than theoretical estimates based on the Bondi-Hoyle prescription, the SPH accretion rates remain large enough to explain the pollution of barium stars. Uncertainties in the current SPH rates remain however, due to the simplified treatment of the wind acceleration mechanism, as well as to the absence of any cooling prescription and to the limited numerical resolution.  Angular momentum transfer leads to significant spin up of the accretor and can account for the rapid rotation of HD~165141, a barium star with a young white dwarf companion and a rotation rate unusually large among K giants.  In the circular orbit modelled in this paper, hydrodynamic thrust and gravitational drag almost exactly compensate and so the net transfer of linear momentum is nearly zero. For small but finite eccentricities and the chosen set of parameters, the eccentricity tends to decrease. ",
+        "introduction": "\\label{Sect:Intro}  Wind accretion in a binary system influences the surface composition and spin of the accreting star and can change the orbital parameters of the binary.  This paper aims at estimating these effects in a binary system in which one of the components is an asymptotic giant branch (AGB) star which suffers from strong mass loss ($\\dot M \\approx 10^{-6}~\\Msun~{\\rm y}^{-1}$) through a wind.  The companion main sequence (MS) star accretes some of this carbon and s-process element enriched gas and mixes it in its envelope. Stellar evolution will eventually transform the AGB star into a white dwarf (WD) and the initially less massive MS star into a chemically peculiar giant.  This is the evolutionary scenario proposed by Boffin \\& Jorissen (1988) for the formation of barium stars, i.e. late-type giants exhibiting overabundances in carbon and s-process elements (e.g. Lambert 1985).  All barium stars are likely members of a binary system with a WD companion (McClure \\& Woodsworth 1990, Jorissen \\& Boffin 1992).  The theoretical description of the accretion process is complex due to the intricate structure of the flow and the complex non-linear hydrodynamics of the shocked wind material. One can estimate the expected fraction $\\dot M_{acc}/\\dot M$ of material being accreted (Boffin \\& Jorissen 1988) from an interpolation formula due to Bondi \\& Hoyle (1944) and Bondi (1952) for the accretion rate onto a star moving at constant velocity through a gas of uniform density and temperature. This involves the interpolation of the estimated accretion rates from two even more simplified models. In the first model (Hoyle \\& Lyttleton 1939), one neglects gas pressure and estimates the accretion rate $\\dot M_{HL}$ onto a compact gravitating object with mass $M$, moving at constant velocity $v_\\infty$ through a pressureless medium with initially uniform density $\\rho_\\infty$. In this model, gravitationally deflected material which passed on one side of the star collides with material passing on the opposite side, thereby cancelling its transverse velocity and forming an \\lq accretion line\\rq~ behind the star. The material from this accretion line that has a velocity below the local escape velocity from the star will be accreted and one finds (Hoyle \\& Lyttleton 1939): \\begin{equation} \\dot M_{HL} = \\pi R_A^2 \\rho_\\infty v_\\infty, \\label{eq:HL} \\end{equation} where the accretion radius $R_A$ is given by \\begin{equation} R_A = 2GM/v^2_\\infty. \\label{eq:RA} \\end{equation} The second model (Bondi 1952; see also Theuns \\& David 1992 for the closed form solution of the flow pattern) includes gas pressure but neglects instead the motion of the star through the medium. The accretion rate in a steady state is not uniquely determined but the model does provide a maximum accretion rate $\\dot M_{B}$: \\begin{equation} \\dot M_{B} = \\alpha \\pi R_B^2 \\rho_\\infty c_\\infty, \\label{eq:B} \\end{equation} where the Bondi radius is given by Eq.~(\\ref{eq:RA}) but with the sound speed at infinity $c_\\infty$ replacing $v_\\infty$. The efficiency parameter $\\alpha$ is of the order of unity and depends on the value $\\gamma$ of the polytropic index of the gas. The accretion rate $\\dot M_{BH}$ for the Bondi-Hoyle case, where both the velocity of the star as well as gas pressure are included, is based on an interpolation between these two extreme cases: \\begin{equation} \\dot M_{BH} = \\alpha \\pi R_{A}^2 \\rho_\\infty v_\\infty \\left(\\frac{{\\cal M}_\\infty^2} {1 +  {\\cal M}_\\infty^2}\\right)^{3/2}, \\label{Eq:BH} \\end{equation} where ${\\cal M}_\\infty \\equiv v_\\infty/c_\\infty$ is the Mach number.  The theoretically predicted accretion rate $\\dot M_{BH}$ in the Bondi-Hoyle model has been extensively tested against numerical simulations by e.g. Hunt (1971) and more recently by Ruffert (1994) and Ruffert \\& Arnett (1994). The theoretical and numerical rates agree to within 10\\% which must be considered a bit fortuitous since e.g., although transiently present, no accretion line (as envisaged in the Hoyle and Lyttleton picture) actually exists. The latter simulations also show the importance of numerical resolution on the computed accretion rates.  In a previous paper (Theuns \\& Jorissen 1992, Paper~I), we presented three-dimensional (3D) smoothed particle hydrodynamics (SPH) simulations of wind accretion, taking into account the binary motion. Two-dimensional simulations of this problem were previously done by S\\o rensen et al. (1975) and by Matsuda et al.  (1987, 1992). Even a glancing comparison between the structure of the flow in this case (Paper~I) against that in the Bondi-Hoyle model (as in e.g. Ruffert \\& Arnett 1994) shows that the binary motion induces a very different flow pattern. Consequently, using the accretion rate \\MBH\\ in this situation as well might be stretching one's luck too far.  The main aim of this paper is to test whether the mass accretion rates are large enough to produce barium stars through the wind accretion scenario discussed above. The impact of wind accretion on the orbital parameters and on the spin velocity of the accretor will also be evaluated. An especially important diagnostic of wind accretion is provided by the variation of the orbital eccentricity caused by the transfers of linear momentum between the gas and the accreting star. Moreover, if the accreted gas possesses angular momentum, the spin period of the accreting star can be altered. Several aspects considered in the present paper may also be relevant to the study of wind accretion in symbiotic stars or in $\\zeta$ Aurigae systems.  In addition to the complications introduced by the binary motion, uncertainties about the thermodynamic state of the gas call for caution when deducing the accretion rate from oversimplified models. In fact, the simple polytropic equation of state adopted in this paper for the gas may not be appropriate in these systems exhibiting \\lq cooling radiation\\rq~ in the form of UV light.  Interacting systems such as symbiotic or extrinsic S stars, where wind accretion is currently taking place, indeed exhibit a wealth of UV emission lines or continuum UV radiation (e.g. Nussbaumer \\& Stencel 1987; Johnson \\& Ameen 1991). In the interacting S star HD~35155 for example, a total UV continuum emission of about 0.2~L$_\\odot$ is observed (Johnson \\& Ameen 1991). In Paper~I, we identified regions of shocked gas as being responsible for this kind of emission.  Comparison of the accretion rates obtained for different $\\gamma$'s nevertheless allows us to appraise the importance of the assumed polytropic index on the computed rates. ",
+        "conclusions": "The SPH simulations presented in this paper model the accretion of the wind of a mass-losing star by its companion in the case where the wind velocity is comparable to the orbital velocity. The binary motion has then a strong influence on the accretion process, making analytic predictions based on the Bondi-Hoyle model suspect. In fact, we find that the mass accretion rate is at least ten times smaller than the rate predicted by the Bondi-Hoyle prescription, amounting to about 1\\% in a binary system with $M_1 = 3$~$\\Msun$, $M_2 = 1.5$~$\\Msun$, $A = 3$~AU and $v_{\\rm wind} = 15$~\\kms (in the case of an adiabatic gas with $\\gamma=1.5$). Smaller polytropic indices lead to slightly larger accretion rates of the order of 8\\%. These values probably represent upper limits, as the accretion rate decreases with increasing numerical resolution.  Despite the fact that the efficiency of wind accretion appears to have been overestimated in previous studies relying on the Bondi-Hoyle prescription, wind accretion still appears efficient enough to account for the chemical peculiarities exhibited by barium stars with orbital periods up to about 90~y.  The accretion of transverse linear momentum, controlling the eccentricity variation, appears to be negligible, at least in the present simulation involving a circular orbit.  Accretion of spin angular momentum is substantial, with a well-developed accretion disc forming in the isothermal case, contrary to a widespread belief that wind accretion cannot lead to the formation of such a disc. The spin up of the accreting star resulting from wind accretion may even be considered as an important diagnostic of this process. The mild barium star HD~165141 may be one case where such a spin up took place recently. Such fast-rotating post-mass transfer objects must be observed before the various braking processes (among which magnetic braking) slow the star down again. Therefore, we predict a correlation between WD temperature and barium star spin rate: fast spinning barium stars are predicted to have hot WD companions."
+    },
+    "9602/astro-ph9602078_arXiv.txt": {
+        "abstract": "We examine the clustering of galaxies around a sample of 20 luminous low redshift ($z$\\simlt $0.30$) quasars observed with the Wide Field Camera-2 on the Hubble Space Telescope. The HST resolution makes possible galaxy identification brighter than $V=23.5$ and as close as 2$''$ to the quasar. We find a significant enhancement of galaxies within a projected separation of \\simlt 100~\\kpc\\ of the quasars. If we model the qso/galaxy correlation function as a power law with a slope given by the galaxy/galaxy correlation function, we find that the ratio of the qso/galaxy to galaxy/galaxy correlation functions is $3.8\\pm 0.8$. The galaxy counts within $r<15$~\\kpc\\ of the quasars are too high for the density profile to have an appreciable core radius (\\simgt 100~\\kpc).  Our results reinforce the idea that low redshift quasars are located preferentially in groups of 10--20 galaxies rather than in rich clusters. We see no significant difference in the clustering amplitudes derived from radio-loud and radio-quiet subsamples. ",
+        "introduction": "\\label{introduction}  Over the last two decades, it has been well established that quasars are associated with enhancements in the galaxy distribution.  Historically, this provided the first direct observational evidence that quasars were indeed cosmological in origin (Bahcall, Schmidt, \\& Gunn 1969; Bahcall \\& Bahcall 1970; Gunn 1971; Stockton 1978) . Over the years, considerable evidence has accumulated that low redshift, $z$\\simlt $0.4$, quasars reside in small to moderate groups of galaxies rather than in rich clusters (cf. Hartwick \\& Schade 1990,  Bahcall \\& Chokshi 1991, and references therein). At higher redshifts, there is a marked difference in the environments of radio-loud and radio-quiet quasars (Yee \\& Green 1987). At redshifts $z$\\simgt $0.6$, radio-loud quasars are often found in association with rich clusters (Abell richness $R\\ge 1$) while radio-quiet quasars appear to remain in smaller groups, or perhaps in the outer regions of clusters (Boyle \\etal 1988; Yee 1990).  The galaxy environment around quasars provides many important clues as to what triggers and sustains their central engines. As first suggested by Toomre \\& Toomre (1972), mergers and interactions of galaxies can provide an efficient mechanism for transporting gas into the inner regions of a galaxy or quasar. There have been attempts to model the interaction/merger rates of ordinary galaxies in order to explain the luminosity function of quasars (De Robertis 1985; Roos 1985; Carlberg 1990) and the rapid evolution of the merger/interaction rate in clusters with redshift may provide a natural explanation of the strong evolution of clustering observed around radio-loud quasars (Stocke \\& Perrenod 1981; Roos 1981). Knowledge of the galaxy environment around quasars is also important for understanding the large scale distribution of quasars and how it relates to the structure seen in galaxy surveys (Bahcall \\& Chokshi 1991).  In this Letter, we examine the galaxy environment around twenty nearby ($z$\\simlt $0.3$) bright quasars observed with the Wide Field and Planetary Camera-2 (WFPC2) of the Hubble Space Telescope (HST). These fields were imaged as part of an ongoing investigation into the nature of the host environment of quasars (Bahcall, Kirhakos, \\& Schneider 1994, 1995a, 1995b, 1996a). The exceptional resolution of the HST images allows companion galaxies to be detected at very close projected separations, in some cases $r\\sim 3$~\\kpc\\ (\\simlt $2''$), and galaxy/star separation to be performed down to $V\\sim 23.5$. The goal of the work presented here is to quantify the excess of galaxies associated with these quasars.  The outline of the Letter is as follows. A brief description of the quasar sample used in our analysis is given in \\S~2. In \\S~3.1 we argue that galaxy counts are inconsistent with being drawn from a uniform background. In \\S~3.2, we strengthen this conclusion by correcting the counts for the background contamination. We also present the excess galaxy counts above the background in annuli of projected separation. From these counts, we quantify the amplitude of the galaxy clustering around the quasars in \\S~3.3 in terms of a  quasar-galaxy spatial cross correlation amplitude. We discuss our results and their relation to previous work in \\S~4. ",
+        "conclusions": "\\label{sec-disc}  If the quasars were distributed like typical galaxies, then the derived value of $\\langle B_{qg}\\rangle$ would be equal to the amplitude of the galaxy/galaxy correlation function, $\\langle B_{gg}\\rangle\\sim 20$ (Davis \\& Peebles 1983). A higher value of $\\langle B_{qg}\\rangle$ suggests that the quasars lie preferentially in regions of above average galaxy density. Following Bahcall \\& Chokshi (1991), we can convert our values of $B_{qg}$ into an estimate of the  typical richness of quasar galaxy environment. Here, we define richness as the number of $L_\\ast$ galaxies associated with the quasar; this is given by a simple integral over the correlation function, $N= 4\\pi n_\\ast \\int \\xi_{qg}(r)r^2 dr$, where the limits of integration are $r=0$ to 1.5 \\mpc\\ (the traditional Abell radius), and $n_\\ast \\approx 1.5\\times 10^{-2}{\\rm h}^3{\\rm Mpc}^{-3}$ is the number density of $L_\\ast$ galaxies. Using the values of $\\langle B_{qg}\\rangle $ in Table~2, we find the quasars typically reside in groups of 16--25 galaxies. These numbers should be compared to the typical richness of Abell clusters which have 30--49 and 50--79 members for richness classes $R=0$ and $R=1$ respectively.  Moreover, this estimate is most likely an upper limit since there is evidence that the galaxy profile around quasars falls off more steeply than $r^{-1.77}$ for $r$\\simgt 0.25 \\mpc\\ (Ellingson, Yee, \\& Green 1991). In order to test the robustness of our results to a steepening of the galaxy profile, we fitted a double power law model with slope of ${-1.77}$ for $r\\le 0.25$ \\mpc\\ and $-3$ for $r>0.25$ \\mpc. The best fit amplitudes, $B_{qg}$, for this model increased by about 10\\% (still well within the quoted $1-\\sigma$ errors), yet, the number of infered bright galaxies associated with the quasars decreased to 8 due to the steeper profile at large separations.  The derived amplitudes, $\\langle B_{qg}\\rangle $, in the pure power law case are somewhat larger than the estimates by Yee \\& Green (1987). They examined the clustering of 9 radio-loud and 16 radio-quiet quasars in the redshift range $0.15<z<0.30$. On scales of $\\sim$ 20--500~\\kpc, they estimated $\\langle B_{qg}\\rangle \\approx 60\\pm 20$ and $\\langle B_{qg}\\rangle\\approx 42\\pm 14$ for radio-loud and radio-quiet quasars respectively. Hayman (1990) derived galaxy counts around low redshift ($z<0.3$) quasars from the Palomar Sky Survey prints.  He found that the ratio of the quasar/galaxy and galaxy/galaxy angular correlation functions was $3.1\\pm 0.6$; if the quasars and galaxies have similar selection functions, this translates to an estimate similar to Yee \\& Green of  $\\langle B_{qg}\\rangle\\sim 61\\pm 12$. French \\& Gunn (1983) analysed a sample 25 low-redshift quasars ($z\\simlt 0.35$) selected from 1.2-m Palomar Schmidt plates and concluded $\\langle B_{qg}\\rangle= 25 \\pm 12$; they also analysed the data set of Stockton (1978, 27 quasars with redshifts $z\\le0.45$ selected from the red Sky Survey prints) and derived, via the same analysis, $B_{qg}=79\\pm 40$. Our measurements are, with the exception of Yee \\& Green's value for radio-quiet quasars, consistent within the quoted errors.  The slightly higher clustering amplitude we derive for the radio-quiet subsample may be a result of our sample being the subset of the most luminous quasars.  Yee \\& Green (1987) found that the clustering amplitude of galaxies around radio-loud quasars increased by a factor $\\sim 3$ between $z\\sim 0.4$--$0.6$ and at $z\\sim 0.6$ radio-loud quasars are found in environments as rich as Abell class $R=1$. Optical quasars do not evolve as rapidly (Boyle \\etal\\ 1988), perhaps indicating a different formation scenario. It has been suggested the quasars and active galactic nuclei may be triggered by interactions (e.g., Toomre \\& Toomre 1972; Stocke \\& Perrenod 1981; Roos 1981, Yee 1987). This offers a simple explanation for why the quasars are typically not found in rich clusters at low redshifts; the high velocity dispersion of such clusters leads to a low interaction rate.  The HST WFPC2 is an excellent instrument for extending the present analysis to fainter, low-redshift, quasars. This extension would improve the counting statistics while also providing information regarding possible correlations of the quasar environment with luminosity. The imaging of of moderate redshift quasars ($z\\simlt 0.6$) could be accomplished by HST with single orbit exposures. This imaging would provide a more direct comparison with previous ground based work and would increase our knowledge of the evolutionary history of the quasar environment. An increased knowledge of the quasar environment would be useful in using quasar observations to probe large scale structure at higher redshifts."
+    },
+    "9602/astro-ph9602059_arXiv.txt": {
+        "abstract": "{We present {\\em ASCA} results  of the pair clusters A399 and A401. The region between the two clusters exhibits excess X-rays over the value expected with a simple superposition of the two clusters. We see, however, no temperature rise at the merging front; the temperature is near the average of those  in the inner regions of the two clusters. These indicate that the two clusters are really interacting but it is not strong at present. The inner regions of the two clusters show no radial variations of temperature, abundance and absorption values. We set upper-limits of mass deposition rate of cooling flow to be $\\dot{M}<35\\;\\rm M_{\\odot}yr^{-1}$ and $\\dot{M}<59\\;\\rm M_{\\odot}yr^{-1}$ for A399 and A401, respectively. A  hint of azimuthal variation of the temperature is also found.}    \\newpage ",
+        "introduction": "\\label{sec-1} \\baselineskip=20pt \\indent   A  large fraction of clusters of galaxies have been found to contain substructures in the optical galaxy distributions and X-ray surface brightness. This may be understood in the framework of the hierarchical clustering scenario; clusters of galaxies evolve by accreting other clusters or smaller groups of galaxies. By the merging, the volume or density of the X-ray emitting gas may increase, hence X-ray luminosity increases as clusters evolve. Gioia et al.(1990) and Edge et al. (1990), for example, reported  evidence for the luminosity evolution of the  X-ray emitting intracluster gas. Recent observations with {\\em ROSAT} provide more direct evidence for the cluster merging; in some clusters, their X-ray surface brightness distributions are shifted from the galaxy number density distribution, in which local high temperature regions are found. These indicate that the clusters are actually merging with a strong shock at the merging front (Henry \\& Briel 1995 ; Zabludoff \\& Zaritsky 1995 ; Burns et al. 1995). Suppose that the subcluster components are more separated, then the cluster may not be regarded as a single cluster but would be counted as two separate clusters. The distinction may be somewhat ambiguous; it depends on the projected distance of the clusters.  If the system is gravitationally bounded, then the clusters will merge and evolve to a single, richer cluster in the future.  The pair of A399 and A401 is one of the closest and the brightest pair clusters, hence would be the best candidate of pre-merging clusters. Karachentsev and Kopylov (1980) measured the velocity dispersions and mean radial velocities of the two clusters and concluded that these data are consistent with the pair being gravitationally bound. However, although several studies have been made on the pair clusters, whether they are thermally interacting or not is still open problem.  Some authors claimed that they may be interacting. McGlynn \\& Fabian (1984) pointed out that their lack of cooling flows may be due to the interaction. Although both of these two clusters possess luminous cD galaxies, their cooling flows are relatively weak (McGlynn \\& Fabian 1984). By numerical simulations of cluster collisions, they confirmed that, at the first encounter, one cluster can even pass through the other without destroying the optical components, while the intracluster gas is largely affected and hence cooling flows, if any, may be disrupted (McGlynn \\& Fabian 1984). Furthermore, A401 is one of only a few rich clusters that is report to have an extensive cluster radio halo (Harris, Kapahi, \\& Ekers. 1980 ; Roland et al. 1981). Harris et al. (1980) suggested that radio halos form during the coalescence of clusters.  On the other hand, X-ray observation with the {\\em Einstein} satellite indicated that the emission between the two clusters is explained by the simple superposition of the extended emission from the two clusters (Ulmer \\& Cruddace 1981).  We report, for the first time, spatially resolved X-ray spectra of the pair A399 and A401 with the {\\em ASCA} satellite. We found evidence for the interaction between these two clusters; there is an excess X-ray emission in the middle of the two clusters above the expected flux from the simple superposition of the two clusters. We also discuss  large-scale temperature variations in the outer regions of the two clusters, which may be affected by the interaction between them.  We describe the observations and the data reduction in section 2 and the analysis and the results in section 3. Then, in section 4, we discuss the structure of the two clusters and the interaction between them based on the results. Section 5 is devoted to conclusions. We assume $H_{0}=50 {\\rm km s^{-1} Mpc^{-1}}$ and $q_{0}=0.5$ through out this paper. ",
+        "conclusions": "\\label{sec-5} \\indent  We report {\\em ASCA} X-ray observation of the pair clusters A399 and A401. No significant radial variations are found in the central regions of the two clusters. There is an emission excess in the region between them (link-region) over the value expected with a simple superposition of the clusters. The temperature in it is not higher than that in the central regions of the clusters. In the outer regions of A399, signatures of azimuthal variation of temperature are found.  These results suggest that the two clusters are in a pre-merging phase; the interaction is not strong at present.  We cannot exclude a  past strong collision. We encourage more detailed observations and the comparison with numerical simulations.  \\vspace{7mm}  We thank all of the members of the {\\em ASCA} team and the launching staffs of the Institute of Space and Astronautical Science. We thank C. Ishizaka for providing simulation data.  \\newpage"
+    },
+    "9602/astro-ph9602023_arXiv.txt": {
+        "abstract": "\\openup1\\jot We investigate which experiments are better suited to test the robust prediction that cosmic strings do not produce secondary Doppler peaks.  We propose a statistic for detecting  oscillations in the $C^l$ spectrum, and study its statistical relevance given the truth of an inflationary competitor to cosmic strings. The analysis is performed for single-dish experiments and interferometers, subject to a variety of noise levels and scanning features. A high resolution of 0.2 degrees is found to be required for single-dish experiments with realistic levels of noise. Interferometers appear to be more suitable for detecting this signal. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/nucl-th9602029_arXiv.txt": {
+        "abstract": "The reactions $^{10}$B(p,$\\alpha$)$^{7}$Be and $^{11}$B(p,$\\alpha$)$^{8}$Be are studied at thermonuclear energies using DWBA calculations. For both reactions, transitions to the ground states and first excited states are investigated. In the case of $^{10}$B(p,$\\alpha$)$^{7}$Be, a resonance at $E_{Res}=10$ keV can be consistently described in the potential model, thereby allowing the extension of the astrophysical $S$-factor data to very low energies. Strong interference with a resonance at about $E_{Res}=550$ keV require a Breit-Wigner description of that resonance and the introduction of an interference term for the reaction $^{10}$B(p,$\\alpha_1$)$^{7}$Be$^*$. Two isospin $T=1$ resonances (at $E_{Res1}=149$ keV and $E_{Res2}=619$ keV) observed in the $^{11}$B+p reactions necessitate Breit-Wigner resonance  and interference terms to fit the data of the $^{11}$B(p,$\\alpha$)$^{8}$Be reaction. $S$-factors and thermonuclear reaction rates are given for each reaction. The present calculation is the first consistent parametrization for the transition to the ground states and first excited states at low energies. ",
+        "introduction": "The determination of the astrophysical $S$-factor of the reaction $^{10}$B(p,$\\alpha$)$^7$Be at thermonuclear energies is important in several respects. In the search for advanced fusion reaction fuels the reaction $^{11}$B + p $\\rightarrow$ $^8$Be + $\\alpha$ $\\rightarrow$ 3$\\alpha$ is discussed as a promising candidate for a relatively clean fusion fuel~\\cite{rae81,fel88}. However, natural boron contains 19.7\\% $^{10}$B which produces $^7$Be contaminations via the $^{10}$B(p,$\\alpha$)$^7$Be reaction. Therefore, for a full understanding of the feasibility of boron as a fusion fuel one has to consider the rate for the latter reaction as well~\\cite{pet75}.  The importance of the $^{10}$B(p,$\\alpha$)$^7$Be reaction in astrophysical scenarios results from the fact that it is the dominant process for the destruction of $^{10}$B. It was also claimed that it could be usefully incorporated in explaining the abundances of boron isotopes including the present theory of spallative generation of $l$ elements~\\cite{you91,ree71}. Furthermore, in theoretical investigations of primordial abundances of elements, its rate has to be incorporated in the reaction networks employed in nucleosynthesis calculations for inhomogeneous big bang scenarios~\\cite{thi91,app93}.  Experimental results for the $^{10}$B(p,$\\alpha$)$^7$Be reaction at energies below 1 MeV are scarce. There have been measurements of the $^{10}$B(p,$\\alpha_0$)$^7$Be cross sections with sometimes inconsistent results in the energy ranges 220 keV $< E_p <$ 480 keV~\\cite{bur50}, 60 keV $< E_p <$ 180 keV~\\cite{sza72}, 70 keV $< E_p <$ 205 keV~\\cite{bac55}, and, more recently, 120 keV $< E_p <$ 480 keV~\\cite{you91}. However, these measurements do not extend very far into the region dominated by the $J^{\\pi}$=5/2$^+$ resonance ($E_x = 8.701$ MeV in $^{11}$C~\\cite{wie83}) which is of great importance for astrophysics. In our calculations we therefore focused on the description of the resonance and the reproduction of most recent data~\\cite{kna92a,kna92b,ang93} at very low energies. The only available data measuring the $\\alpha _1$ contribution to the $S$-factor are given in Ref.~\\cite{ang93b}.  The reaction $^{11}$B(p,$\\alpha$)$^8$Be has been measured below 1 MeV \\cite{dav79}, \\cite{beck87}, and most recently \\cite{ang93}. It should be noted that the effect of electron screening~\\cite{ang93} increases the very low energy cross sections considerably and thus has a major impact on a reaction's significance as a terrestrial fusion fuel. ",
+        "conclusions": "The reactions $^{10}$B(p,$\\alpha$)$^7$Be and $^{11}$B(p,$\\alpha$)$^8$Be are described well by the DWBA calculations. In the case of $^{10}$B(p,$\\alpha_0$)$^7$Be, we were able to reproduce the resonance structure in a consistent way within our potential model thus suggesting that it can be regarded as a potential resonance. The influence of the resonance on the reaction rate can be seen very clearly in Fig.~\\ref{fig3} (a fairly constant $S$-factor was assumed in Ref.~\\cite{cau88} at low energies). This clearly demonstrates that extrapolations from higher energies have to be done very carefully to include the correct shape of the resonance.  For the transition to the first excited state in $^7$Be the interference effects with a 5/2$^+$ level at about 550 keV have to be taken into account. This was achieved by including a Breit-Wigner term describing the resonance at 550 keV and an interference term between the DWBA and Breit-Wigner contributions. However, the cross sections of the $\\alpha_1$ transition are lower by several orders of magnitude than those of the $\\alpha_0$ transition and therefore it does not contribute significantly to the final reaction rate.  In the case of $^{11}$B(p,$\\alpha$)$^8$Be the inclusion of an interference term between single-level Breit-Wigner and DWBA also reproduces the data acceptably well. Systematic studies at higher energies are being carried out at present \\cite{sch94}."
+    },
+    "9602/astro-ph9602084_arXiv.txt": {
+        "abstract": "We have obtained a {\\it Hubble Space Telescope} ultraviolet image and spectrum of the nearby Wolf-Rayet galaxy NGC~1741. The spatial morphology from the Faint Object Camera image is dominated by two main starburst centers, each being about 100 times as luminous as 30~Doradus. Both starburst centers are composed of several intense knots of recent star formation. A Goddard High Resolution Spectrograph spectrum of a portion of the southern starburst center is consistent with a population of young stars following a Salpeter IMF for masses above $\\sim$15~\\Ms\\ (lower mass stars may also be present), and extending up to $\\sim$100~\\Ms; about $10^4$ O-type stars are inferred from the UV luminosity.  Numerous strong interstellar lines are detected. Although not resolved, their strength suggests that they are formed in individual bubbles and shells with velocities up to a few hundred \\kms. The red wing of the Lyman-$\\alpha$ absorption profile indicates the presence of several neutral hydrogen components, one in our own Galaxy and the others at or close to the distance of NGC~1741. Overall, the stellar and interstellar line spectrum, as well as the continuum shape of NGC~1741, strongly resembles star-forming galaxies recently discovered at high redshift. ",
+        "introduction": "Starburst galaxies are a class of objects experiencing brief but intense episodes of star formation. Massive OB stars are dominant contributors to their ultraviolet and optical spectra in the first 10~Myr after the onset of the burst (e.g., Leitherer et al. 1996). Wolf-Rayet (W-R) galaxies are a subset of starburst galaxies that show {\\em broad} stellar He~II $\\lambda$4686 emission in their spectra (Kunth \\& Sargent 1981; Conti 1991). Wolf-Rayet stars are the descendants of the most massive O stars, which show the products of nuclear burning at their surface due to mass loss and mixing processes (e.g., Abbott \\& Conti 1987). Because only the most massive stars evolve to form W-R stars, their presence in a starburst galaxy might suggest a stellar population with numerous high-mass stars present (Sargent \\& Filippenko 1991), {\\it and/or} a rapid formation timescale, (e.g., Vacca \\& Conti 1992, hereafter VC). The census of W-R stars in the Galaxy and in its OB associations indicates that they have progenitor masses above $\\sim$40~\\Ms\\ (Conti et al. 1983; Humphreys, Nichols, \\& Massey 1985). Evolutionary synthesis models of stellar populations containing W-R stars (e.g., Meynet 1995) predict that about 3 to 6~Myr after the start of the burst the most massive stars evolve into W-R objects and the ratio of W-R to O stars increases dramatically. W-R galaxies therefore offer the unique opportunity to study a burst population whose age is known {\\it a priori}. The determination of the age $t$  and the initial mass function (IMF) of a starburst can in some cases be degenerate, i.e., the absence of 40~\\Ms\\ stars could imply either an upper mass limit M$_{\\rm up}< 40$~\\Ms\\ or $t~>$~5~Myr.  NGC~1741 (=~Mrk1089 =~Arp259) was first described as a W-R galaxy by Kunth \\& Schild (1986). It has a highly disturbed optical morphology with two starburst centers, possibly arising from a galaxy merger, most likely as a result of interaction with other members of Hickson Compact Group 31 (Hickson, Kindl, \\& Auman 1989; Rubin, Hunter \\& Ford 1990). NGC~1741 is one of the most luminous W-R galaxies --- optically and in the far-UV --- in Conti's (1991) catalog: $M_B = -20.3$.  VC obtained optical spectrophotometry of NGC~1741 and inferred the presence of about 700 W-R stars in the southern starburst center (``B'' in their notation). We selected NGC~1741 as one of the targets for our ongoing program of {\\it Hubble Space Telescope} ({\\it HST}) observations of W-R galaxies. Ultraviolet imaging and spectroscopy have been obtained, confirming the presence of numerous OB stars in this galaxy. One purpose of this {\\it Letter} is to draw attention to the spectral similarity between NGC~1741 and starburst galaxies recently discovered at high redshift, e.g., the ``primeval galaxy candidate'' cB58 (Yee et al. 1996), the star--forming galaxies at $z > 3$ (Steidel et al. 1996), and the ``lensed'' galaxy at $z = 2.5$ (Ebbels et al. 1996). ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602101_arXiv.txt": {
+        "abstract": "A new type of gravitational microlensing experiment toward a field where stars are not resolved is being developed observationally and theoretically: pixel lensing. When the experiment is carried out toward the M31 bulge area, events may be produced both by Massive Compact Halo Objects (MACHOs) in our Galactic halo and by lenses in M31. We estimate that $\\sim 10-15\\%$ of the total events are caused by Galactic halo MACHOs assuming an all-MACHO halo. If these Galactic events could be identified, they would provide us with an important constraint on the shape of the halo. We test various observables that can be used for the separation of Galactic halo/M31 events. These observables include the Einstein time scale, the effective duration of an event, and the flux at the maximum amplification, but they cannot be used to separate each population events. However, we find that most high maximum-flux Galactic halo events can be isolated through a satellite-based measurement of the flux difference caused by the parallax effect. For the detection of the flux difference, it is required to monitor events with an exposure time of $\\sim 20\\ {\\rm min}$ by a 0.5 m telescope mounted on a satellite. Such observations could be carried out as a minor component of a mission aimed primarily at events seen toward the Galactic bulge and  Large Magellanic Cloud. In addition, proper motion can be used to isolate Galactic halo/M31 events, but only for $\\sim 5\\%$ of high signal-to-noise ratio M31 events and only 1\\% of Galactic halo events. ",
+        "introduction": "Crotts (1992) and Baillon et\\ al.\\ (1993) have begun a pioneering search for lensing events of unresolved stars in M31 by monitoring individual {\\it pixels} for the time-dependent flux induced by one of the many unresolved stars they contain:\\ `pixel lensing' for short. For a classical lensing event, the light curve is obtained by subtracting the reference flux, $B$, from the amplified flux, $F$: $$ F_{0,i}(A-1)=F-B;\\ \\ B = F_{0,i} + B_{\\rm res}, \\eqno(1.1) $$ where $F_{0,i}$ is the flux of the lensed star before or after the event and $B_{\\rm res} = \\sum_{j\\neq i} F_{0,j}$ is the residual flux from stars that are not participating in the event. Here the amplification is $$ A(x) = {x^2 + 2 \\over x(x^2 +4)^{1/2} },\\ \\ x = [\\beta^2 + \\omega^2 (t-t_0)^2]^{1/2}, \\eqno(1.2) $$ where $\\omega=t_{\\rm e}^{-1}$ is the inverse Einstein ring crossing time, $\\beta$ is the dimensionless impact parameter, and $t_0$ is the time at the maximum amplification. In an uncrowded field where the blending of stars is not important ($B_{\\rm res} \\ll F_{0,i}$), one can directly measure the values of $F$ and $B=F_{0,i}$. On the other hand, it is impossible to measure the absolute values of $F$ and $B$ for the case of pixel lensing since stars cannot be resolved. However, the light curve, $F_{0,i}(A-1)$, can still be obtained by subtracting the reference image from those in which a lensing event is in progress (Ciardullo et al.\\ 1990; Tomaney \\& Crotts 1996).   Pixel lensing has several advantages over classical lensing. First of all, one is not restricted to observing fields where stars are resolvable. Second, one can detect ${\\cal O} (10^2)$ events from a season-long observation of M31 in a single field with moderate resources because a large number of stars are monitored (Han 1996). Finally, one looks through the Galactic halo when an external galaxy is monitored for pixel lensing; some fraction of events may be caused by the Galactic Massive Compact Halo Objects (MACHOs). If the Galactic MACHOs can be distinguished from events caused by the lenses in the external galaxy, pixel lensing will considerably increase the number of Galactic MACHO events. In addition to the increase in the number of events, one can use the data toward these additional lines of sight to constrain the shape of the Galactic halo. However, events caused by one population are just contaminants in the study of the other population unless they are separated. The problem becomes more serious when the fraction of contaminating events is large. It is therefore important to determine the fraction of events caused by each population and to develop methods for distinguishing them from one another.   We estimate that $\\sim 10-15\\%$ of total events seen toward the M31 bulge area are caused by the Galactic halo MACHOs, assuming an all-MACHO halo. The first and second year of data observed by the MACHO group toward the Large Magellanic Cloud (LMC) gives tentative support to the hypothesis that MACHOs constitute a significant fraction of the Galactic halo (D.\\ Bennett 1995, private communication). We test various observables for their ability to distinguish the Galactic halo events from M31 events. We find that most high maximum-flux Galactic halo events can be isolated through a satellite-based measurement of the flux difference caused by the parallax effect. This observational strategy is discussed in \\S\\ 4.1. Proper motion can be measured for $\\sim 5\\%$ of high signal-to-noise ratio ($S/N$) M31 events and can also be used to separate Galactic halo/M31 events. In addition, 1\\% of Galactic events can be isolated by determining the lower limit of the proper motion, $\\mu_{\\rm low}$. Hence, only a small fraction of events can be separated by using the proper motion and proper motion-limit measurements. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602047_arXiv.txt": {
+        "abstract": "We analyze post-refurbishment Hubble Space Telescope images of four globular clusters in M31. The ability to resolve stars to below the horizontal branch permits us to use star counts to extend the surface brightness profiles determined using aperture photometry to almost 5 orders of magnitude below the central surface density. Three of the resulting cluster profiles are reasonably well-fit using single-mass King models, with core and tidal radii typical of those seen in Galactic globular clusters. We confirm an earlier report of the discovery of a cluster which has apparently undergone core collapse.  Three of the four clusters show departures in their outskirts from King model behavior which, based on recent results for Galactic globulars, may indicate the presence of tidal tails. ",
+        "introduction": "Globular clusters have been and continue to be useful tools in the study of the structure, formation, dynamics, and distances of galaxies. The clusters themselves are dynamically simple systems, and are sufficiently luminous to be visible out to $cz \\approx$ 10000 km/sec (Harris 1991). Much recent work has concentrated on the degree to which the properties of globular clusters differ from one galaxy to the next. For example, an important issue currently being addressed is whether the globular cluster luminosity function and its peak are sensitive to metallicity (\\eg Ashman \\etal 1995), and observably different for early- and late-type galaxies. Considerable work is being done to better understand the physical processes which would alter the structure and dynamics of individual clusters, and thereby the global properties of globular cluster {\\it systems}. For example, different orbit shapes, tidal fields, and the presence of disks and bulges may have profound consequences for the ultimate survival of individual clusters, and for the constitution of cluster systems as a whole (Chernoff \\& Weinberg 1990; Kundic \\& Ostriker 1995). While hypotheses are in abundance, the observations necessary to test these ideas remain rather sparse. We need to make detailed studies of globular clusters in differing environments if we are to put useful constraints on the models.  The surface brightness profiles of most globular clusters in our own Galaxy have long been known to be reasonably well represented by King models (King 1966), which require the minimum possible number of parameters to describe the central density, total energy, and external tidal field. Together, the measured core or half-light radii and central densities yield information concerning the formation and long-term dynamical evolution of individual clusters. Conversely, the limiting radii are sensitive, on a relatively short time scale, to the influence of external potentials. The actual form of the observed tidal ``cutoff'' is determined by cluster mass, the tidal field of the galaxy, and the shape and phase of the cluster orbit (Oh and Lin 1992; Grillmair \\etal 1996). While these quantities are not well constrained for any one cluster, obtaining information on a large sample of clusters should enable us to describe the mean orbital properties of the ensemble, as well as to map the potential of the host galaxy. M31 globular clusters are especially interesting in this respect because $(i)$ the M31 cluster system is the most populous cluster system in the Local Group, $(ii)$ we observe the cluster system from a vantage point outside the system, and $(iii)$ M31 is similar in many respects to our own Galaxy, permitting us to gauge the sensitivity of globular cluster properties to detailed structural differences.  The superb resolution afforded by the Hubble Space Telescope (HST) since the installation of the Wide Field/Planetary Camera-2 (WFPC-2) permits study of the stellar content and distribution in extragalactic globular clusters at a level of detail hitherto only possible for clusters belonging to our own Galaxy and its satellites. The observations discussed below were motivated primarily by a desire to calibrate the metallicity dependence of horizontal branch (HB) magnitudes (Ajhar \\etal 1994; Ajhar \\etal 1996, hereafter Paper I). However, the spatial resolution and the depth of the observations also allow a detailed examination of the structure of the clusters themselves. The mass/luminosity ratios derived from velocity dispersion measurements of our sample clusters are discussed elsewhere (Dubath and Grillmair 1996).   Section 2 briefly summarizes the observations. Section 3 describes synthetic aperture photometry, and a star-count analysis is carried out in Section 4. We present the results and discuss their implications in Section 5, and briefly summarize our findings in Section 6. ",
+        "conclusions": "}  In Figures \\ref{fig:k58} through \\ref{fig:k219} we show the surface density profiles derived from both the star counts and the aperture photometry for each of the four clusters. The aperture photometry has been scaled to match the star counts in the region of overlap.  The error bars plotted for the aperture photometry include a 0.3 DN uncertainty in the sky-background level, and the uncertainties for the inner two points have been set to half the difference in DN between them to account for possible miscentering. The error bars shown for the star counts take account of both Poisson statistics and the uncertainties in the completeness corrections. The data are shown only out to the radius at which the surface density first becomes consistent (within the computed error bar) with zero. The apparent agreement between the aperture photometry and the star counts in the region of overlap suggests that our estimates of the sky background are not grossly in error.  Shown as solid lines in Figures \\ref{fig:k58} through \\ref{fig:k219} are the results of 20 deconvolution iterations using the Lucy-Richardson (Richardson 1972; Lucy 1974) algorithm. Simulations have shown that deconvolution of WFPC-2 data to this extent permits reliable recovery of information down to $\\approx0.05\\arcsec$. The deconvolved profiles of both K108 and K219 show a decline (relative to the ``raw'' profile) in the flux measured for the central pixel. This is due to the presence of one or more relatively bright red giants within a few pixels of the center. On the other hand, the inner deconvolved profile for K105 shows a general steepening and a factor of two increase in the flux measured in the central pixel. This is in agreement with the results of Bendinelli \\etal (1993), who used deconvolved, pre-refurbishment Faint Object Camera (FOC) images to determine that K105 (G105 in their paper) has a power-law cusp at the center and has therefore likely undergone core collapse. Indeed, over the range $0.045\\arcsec \\le r \\le 0.23\\arcsec$, we find a power-law slope of $-1.08 \\pm 0.1$, somewhat steeper than their estimated value of -0.75, and strongly suggestive of a past or impending gravothermal catastrophe.    The long-dashed and short-dashed lines in Figures \\ref{fig:k58} through \\ref{fig:k219} are the best-fitting, single-mass King models before and after 2-dimensional convolution with the PC PSF. King-model fitting was carried out in two stages; owing to significant coupling between $r_c$ and $r_t$ and the apparent departures of the observations from King-like behavior at large radii (see below), we chose to fit core and tidal radii separately. Firstly, two-dimensional, single-mass King models were generated and convolved with the PSF appropriate to the position of the cluster on the PC (\\eg Lauer \\etal 1991). The PSF itself was generated using DAOPHOT II to model the morphologies of the brightest $\\sim$30 point sources in each frame. Aperture photometry, in a manner identical to that used for the real data, was carried out to produce a one-dimensional surface density profile, and this profile was then compared with the data to obtain a value of $\\chi^2$.  Using a downhill-simplex method, the core radius and normalization constant were estimated using aperture photometry data within $1\\arcsec$ of the center.  Secondly, the tidal radius was estimated using both aperture photometry and star count data by fixing $r_c$ and searching for the value of $r_t$ which gave the minimum reduced $\\chi^2$. Our best-fit values for $r_c$ and $r_t$ were converted to spatial units assuming a distance to M31 of 770 kpc (Paper I) and are tabulated and compared with those of previous investigators in Table 10.  The uncertainties in the core and tidal radii were estimated using Monte Carlo simulations of the data, based on both the error bars shown in Figures \\ref{fig:k58} through \\ref{fig:k219} and a 5\\% uncertainty in the adopted sky brightness. Irrespective of departures from King-like behavior, the simulations yielded 90\\% confidence intervals of $\\approx1.2\\arcsec$ (4.5pc) in $r_t$ for all clusters. For K58 and K105, the 90\\% confidence intervals on $r_c$ were $\\approx0.01\\arcsec$ (0.04pc), and for K108 and K219 they were $\\approx0.02\\arcsec$ (0.08pc). Of course, these uncertainties apply only in those cases where it is clear that a King model is appropriate. The power-law cusp in the deconvolved profile of K105, for example, is not well fit by the King model which best matches the raw data, and if there is indeed a constant surface brightness core in this cluster, it is likely to be considerably overestimated by the value of $r_c$ given in Table 10. Similarly, the formal fitting uncertainties on $r_t$ are largely irrelevant in those cases where there is a clear discrepancy between the form of the model and that of the data.   The values we obtain for the half-light radii are typical of those found in Milky Way globular clusters ($<r_h> = 4.2$pc, Djorgovski 1993).  It is reassuring that the HST study of K58 by Fusi Pecci \\etal (1994) using deconvolved Faint Object Camera images yields a measurement of the core radius which is in reasonable agreement with ours. On the other hand, the half-light radius found by these authors for K58 (2.1pc) is somewhat smaller than the value of 2.8pc we find here. We attribute this discrepancy to the very small field of view ($11\\arcsec$) of the FOC in its $f/96$ mode. Given the extent of the surface density profile shown in Figure \\ref{fig:k58}, it is quite probable that they overestimated the sky brightness, leading to an underestimate of the total light and extent of K58. Their value of $r_h = 2.4$pc for K105 (compared to our value of 2.9pc) is consistent with this hypothesis.  With the exception of K108 (which, in projection, lies closest to M31 and hence is most severely affected by background uncertainties), all clusters show apparent departures from the best fitting King model at radii approaching $r_t$. These departures take the form of an excess of stars near and beyond the fitted $r_t$, and occur between 3 and 4 orders of magnitude below the central surface density. K105 and K219 illustrate this effect most dramatically. Both the star counts and the aperture photometry show a tidal turn-down beyond $1\\arcsec$. However, rather than following the model profiles to a steep tidal-cutoff, the data continue downwards in a comparatively shallow, power-law fashion.  Due to the distribution of uncertainties, the best-fitting values of $r_t$ are heavily influenced by the inner few data points derived from the star counts. The single-mass King models cannot be made to fit the extended portions of the profiles due to the characteristic, upward inflection of high-concentration, King model profiles.  This is illustrated in Figure \\ref{fig:k219}, where we have plotted a King model with core radius and normalization identical to that of the King model which best fits the data for K219, but with the tidal radius increased to $24\\arcsec$ (90pc). Although this model passes through the outermost points, the reduced $\\chi^2$ is very much greater than for the best-fit case. Clearly, no renormalization of this extended model would improve the overall fit to the data.  Note that the cluster isophotes have ellipticities ranging from $\\sim0.07$ (K58) to $\\sim0.18$ (K219). However, even if we assumed that the tidal cutoff radius varies with position angle like the isophotes, the tidal radius measured along the major axis of the cluster should differ by no more than 10\\% from the tidal radius predicted by a circular model. The differences between the radius of the last measured, non-zero star count datum and the radius at the same surface density of the best-fitting King model exceed what one might have expected purely from cluster ellipticity by factors of $\\sim8$ (K219) to $\\sim15$ (K105). Moreover, K108 shows moderate isophote ellipticities ($\\sim0.12$), but the star counts near $r_t$ are well fit by a King model. Hence we conclude that the departures near $r_t$ between the models and the star counts are not related to the ellipticities seen in the isophotes.  The departures from King models are consistent, both in magnitude and in form, with the findings of Grillmair \\etal (1995), who studied a sample of 12 Galactic globular clusters. They concluded that these extended profiles result from stars which have been stripped from their parent clusters and are in the process of migrating outwards along the tidal tails. The power-law profiles seen in their sample were found to exhibit a variety of slopes, consistent with N-body models of clusters on eccentric orbits (Grillmair \\etal 1996). The magnitude of the effect and the slope of the power-law profile are functions of orbital phase and viewing angle. Owing to the much smaller number of resolved cluster stars in the present work (due to the much brighter limiting absolute magnitude), the uncertainties in the counts near $r_t$ are significantly larger than for the sample of Grillmair \\etal (1995). Similarly, the very small numbers of ``extra-tidal'' stars relative to the background preclude a study of their two-dimensional distribution. However, the systematic behavior of the profiles shown in Figures \\ref{fig:k58} through \\ref{fig:k219} is consistent with the conclusions reached by these authors.  One might be tempted to dismiss the foregoing on the grounds that single-mass King models must necessarily be too simplistic to adequately model the behavior of real clusters. However, we would argue that multi-mass King models are unlikely to be able to account for the observed profile shapes near $r_t$ due to the very small range of stellar masses represented among the counted stars. Likewise, the aperture photometry is dominated by giant branch stars, so much so that stars with $V < 26.5$ (the regime sampled by the star counts) constitute $\\approx80\\%$ of the total light from each cluster. We note that $\\approx50\\%$ of the light is contributed by stars above the horizontal branch, while only $\\approx20\\%$ of the completeness-corrected star counts occur above the same magnitude level. Nonetheless, the difference in mean stellar mass over this magnitude interval is entirely inconsequential from the standpoint of dynamical segregation.  K105 is somewhat analogous to M15 in having both a collapsed core and a pronounced excess of extra-tidal stars. That both phenomena should be observed in each of these clusters may not be surprising; the early onset of core collapse and a rapid rate of repopulation of the tidally stripped region of a cluster both rely on a relatively short 2-body relaxation time."
+    },
+    "9602/hep-th9602066_arXiv.txt": {
+        "abstract": "The classical equation of motion of a charged point particle, including its radiation reaction, implies tunneling. For nonrelativistic electrons and a square barrier, the solution is elementary and explicit. We show the persistance of the solution for smoother potentials. For a large range of initial velocities, initial conditions may leave a (discrete) ambiguity on the resulting motion. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602134_arXiv.txt": {
+        "abstract": "{\\baselineskip 0.4cm Since the appearance of the classical paper of Lifshitz almost half a century ago, linear stability analysis of cosmological models is textbook knowledge. Until recently, however, little was known about the behavior of higher than linear order terms in the perturbative expansion. These terms become important in the weakly nonlinear regime of gravitational clustering, when the rms mass density contrast is only slightly smaller than unity. In the past, theorists showed little interest in studying this regime, and for a good reason: only a decade ago, it would have been an academic excercise -- at scales large enough to probe the weakly nonlinear regime, all measures of clustering were dominated by noise. This is no longer the case with present data. The purpose of this talk is to provide a brief summary of recent advances in weakly nonlinear perturbation theory. We present analytical perturbative results together with results of N-body experiments, conducted to test their accuracy. We compare perturbative predictions with measurements from galaxy surveys. Such comparisons can be used to test the gravitational instability theory and to constrain possible deviations from Gaussian statistics in the initial mass distribution; they can be also used to study the nature of physical processes that govern galaxy formation (``biasing''). We also show how future studies of velocity field statistics can provide a new way to determine the density parameter, $\\Omega$. } ",
+        "introduction": "Tests of theories for the origin of the large scale structure of the universe ultimately depend on a comparison of model predictions with measurable quantities, derived from observations. The statistical measures we will discuss here are the low order N-point correlation functions. They have two clear advantages. First, they can be estimated from galaxy surveys with a reasonable degree of precision and reproducibility (see, e.g. \\cite{{lss1},{lss2},{obs2},{obs3}}, and references therein). Second, correlation functions can be relatively easily related to dynamics (e.g. \\cite{lss1}). Their evolution can be studied by taking moments of the hydrodynamical equations of motion of an expanding self-gravitating pressureless fluid with zero vorticity, which  is a good approximation of the real universe after the hydrogen recombination. Low order correlation functions can also be measured from N-body experiments.  \\subsection{Gravitational instability}  The full description of gravitational instability is nonlinear. The density contrast \\begin{equation} \\delta({\\bf x},t) \\; = \\; {\\rho({\\bf x},t)\\over\\langle\\rho\\rangle} - 1 \\; , \\label{eq:def} \\end{equation} the peculiar velocity ${\\bf v}$ and the gravitational potential $\\phi$ are related by the Euler, Poisson and continuity equations. They can be combined into one expression for the density contrast, \\begin{equation} {\\textstyle \\ddot\\delta + 2H\\dot\\delta - {3\\over2}\\Omega H^2 \\delta \\; = \\; {3\\over2}\\Omega H^2 \\delta^2 + a^{-2}\\nabla\\delta\\cdot\\nabla\\phi + a^{-2}\\nabla_{\\alpha} \\nabla_{\\beta} [(1 + \\delta)v_{\\alpha} v_{\\beta}] \\; . } \\label{eq:motion} \\end{equation} Here ${\\bf x} = \\{ x_{\\alpha} \\}$ and $t$ are, respectively, the comoving spatial coordinates and the cosmological time, the dots represent time derivatives, $\\; \\nabla_{\\alpha} = \\partial /\\partial x_{\\alpha} \\;$, $a(t)$ is the scale factor, $H =\\dot a/a$ is the Hubble parameter, and $\\Omega = 8\\pi G \\langle \\rho \\rangle / 3H^2$ is the density parameter. The brackets $\\langle\\ldots\\rangle$ denote ensemble averaging and $\\rho$ is the mass density. Perturbation theory rests on a conjecture that when the deviations from homogeneity are small, the first few terms of the expansion \\begin{equation} \\delta \\; = \\; \\delta_1 + \\delta_2 + \\delta_3 + \\ldots \\label{eq:series} \\end{equation} provide a reasonable approximation of the exact solution of eq. (\\ref{eq:motion}). The first, linear term is the well known Lifshitz \\cite{el} solution of eq. (\\ref{eq:motion}) with the right-hand side set to zero, \\begin{equation} \\delta_1 \\; = \\; D(t)\\varepsilon({\\bf x}) \\; + \\;\\;{\\rm decaying \\; mode} \\; , \\label{eq:linear} \\end{equation} where $D(t)$ is the standard growing mode (see, e.g., \\S 11 in \\cite{lss1}) and $\\varepsilon$ is a random field with statistical properties defined by initial conditions. The term $\\, \\delta_2 = {\\cal O} (\\varepsilon^2) \\, $ is the solution of the same equation with quadratic nonlinearities included iteratively by using $\\delta_1$ as source terms (as in \\cite{lss1}, \\S 18). For a vanishing cosmological constant ($\\Lambda = 0$) and arbitrary $\\Omega$, the fastest growing mode in this solution is \\cite{j3} \\begin{equation} {\\textstyle \\delta_2 \\; = \\; D^2(t)\\,\\left[{2\\over3}(1+\\kappa)\\varepsilon^2 + \\nabla\\varepsilon\\cdot\\nabla\\Phi + ({1\\over2} - \\kappa) \\tau_{\\alpha\\beta}\\tau_{\\alpha\\beta}\\right] \\; , } \\label{eq:quad} \\end{equation} \\begin{equation} \\;\\;\\; \\textstyle{ \\tau_{\\alpha\\beta} \\; = \\;({1\\over3}\\delta_{\\alpha\\beta}\\nabla^2 -\\nabla_{\\alpha}\\nabla_{\\beta})\\Phi \\; ; \\;\\;\\; \\Phi({\\bf x}) \\; = \\;  - \\, \\int\\,d^3x'\\,\\varepsilon({\\bf x'})/4\\pi |{\\bf x - x'}| \\; . } \\label{eq:tau} \\end{equation} The parameter $\\kappa $ is a slowly varying function of time; its dependence on $\\Omega$ is extremely weak, too. In the range $0.05 \\leq \\Omega \\leq 3$, the $\\Omega$-dependence is well approximated by \\cite{frb2} \\begin{equation} \\textstyle \\kappa\\left[ \\Omega(t) \\right] \\; \\approx \\; {3\\over14}\\,\\Omega^{-2/63} \\; . \\label{eq:kappa} \\end{equation} The higher order terms can be constructed out of the linear solution in a similar way.  \\subsection{Dynamical evolution of correlations}  The main source of difficulties in studies of the dynamics of gravitational instability is the nonlinearity of the equations of motion. Indeed, let us write eq. (\\ref{eq:motion}) for $\\delta({\\bf x}, t)$ and $\\delta' = \\delta({\\bf x'}, t)$, then multiply the first equation by $\\delta'$, and the second by $\\delta$. The result of adding these two equations and averaging is \\begin{equation} \\ddot\\xi + 2H\\dot\\xi - 3\\Omega H^2 \\xi + 2 a^{-2}\\nabla^2 \\xi_v \\; = \\; 3\\Omega H^2 \\langle\\delta^2\\delta'\\rangle \\; + \\; \\ldots \\; , \\label{eq:master} \\end{equation} where $\\, \\xi = \\langle\\delta\\delta'\\rangle \\, $ and $\\, \\xi_v = \\langle {\\bf v \\cdot v'}\\rangle \\, $ are the two-point density and velocity correlation functions, while the right-hand side of the equation involves 3- and 4-point correlations, which we left out except for one term given as an example. The full expression is not the point here; the important message is that we have an infinite hierarchy of equations. Mathematically, this is similar to the BBGKY hierarchy: in order to solve for an $N$-point function, we need to know $(N+1)$- and $(N+2)$-point correlations. Perturbation theory may help, but not without further assumptions, necessary to close the infinite hierarchy of correlation functions. Here we will assume that $\\varepsilon({\\bf x})$ is a random Gaussian field. This approach, also known as the {\\it random phase approximation}, can be justified theoretically for a wide class of models of the early universe \\cite{lss2}. Independently of whether we take ``inflationary'' arguments seriously or not, the random phase assumption is observationally testable. As we will show below, comparing model predictions with observations can help us decide whether the early universe was Gaussian or not. The greatest advantage of our assumption is simplicity. Indeed: all odd-order correlation functions of $\\varepsilon$ vanish, and all even-order moments can be reduced to the two-point function. For example, for the first four moments we have $ \\langle \\, 1 \\,\\rangle \\, = \\, \\langle \\, 123 \\, \\rangle \\, = \\, 0$, $\\, \\langle \\, 12 \\,\\rangle \\, = \\, \\xi_L\\,(|{\\bf x}_1 - {\\bf x}_2|,t) \\, D^{-2} \\,$ , and \\begin{equation} \\langle \\, 1234  \\, \\rangle = D^{-4}\\,\\xi_L \\, (|{\\bf x}_1 - {\\bf x}_2|,t) \\, \\xi_L \\, (|{\\bf x}_3 - {\\bf x}_4|,t) \\;\\; + \\;\\; {\\rm cycl.\\,(two \\; terms)} \\; , \\label{eq:even} \\end{equation} where $\\, \\langle \\, 12 \\ldots N \\, \\rangle \\, \\equiv \\, \\langle \\, \\varepsilon({\\bf x}_1)\\varepsilon({\\bf x}_2) \\ldots \\varepsilon({\\bf x}_N) \\, \\rangle$, while $\\xi_L$ in eq. (\\ref{eq:even}) is $\\, \\langle \\, \\delta_1 \\, ({\\bf x}_1,t) \\, \\delta_1 \\, ({\\bf x}_2,t) \\, \\rangle$, the linear 2-point correlation function. Using these relations and the perturbative expansion, we can express the 3- and 4-point functions on the right-hand side of eq. (\\ref{eq:master}) in terms of $\\xi_L$. This will close the hierarchy, allowing us to calculate weakly nonlinear corrections for the 2-point function $\\xi$. The same procedure can be applied for higher order correlation functions. This simple idea is at the root of all perturbative results, discussed below.  Perturbative calculations are usually easier after applying the Fourier transform. The Fourier transform of $\\xi$ is the power spectrum, $P(k,t)$, where $k$ is the comoving wavenumber. The linear term in the perturbative expansion for $P(k,t)$ is \\begin{equation} P_L\\,(k,t) \\; = \\; \\int \\, \\xi_L\\, ({\\bf r},t) \\, \\exp(i{\\bf k \\cdot r}) \\, {\\rm d}^3k \\; . \\label{eq:power} \\end{equation} The basic properties of linear evolution, as well as the effect of nonlinear interactions can be deduced from the functional form of $\\delta_1$ and $\\delta_2$. Indeed, according to eq. (\\ref{eq:linear}), all fluctuations evolve without changing the shape of their density profiles, so that we can expect that the linear power spectrum evolves without changing its slope (the growth rate is the same for all scales). This is indeed the case. If $P_L\\,(k,t) \\propto k^n$ at some initial time $t_0$, than at later times $P_L\\,(k,t) = P_L\\,(k,t_0)\\,D^2(t)/D^2(t_0) \\, \\propto k^n$. A second characteristic property of a linear field is that its distribution function remains Gaussian; only its width, $\\langle \\delta_1^2\\rangle $ grows in time as $D^2(t)$. Neither of these properties remain valid when nonlinear interactions become important. The nonlinear terms, describing tidal interactions in eq. (\\ref{eq:quad}) redistribute power across the spectrum, and growth rates may be enhanced at some wavenumbers and suppressed at others: the logarithmic slope $n(k,t) = {\\rm d}(\\log \\, P)/{\\rm d}(\\log k)$ is no longer constant, it can now change with time. Moreover, the nonlinear evolution introduces deviations from Gaussian statistics. The lowest-order moment which can detect such deviations is the skewness, $\\langle\\delta^3\\rangle$. The Gaussian distribution is symmetric about $\\delta = 0$, and its skewness is zero. According to eq. (\\ref{eq:quad}), however, the fluctuations with positive initial density contrast $\\delta_1$ grow more rapidly than those with $\\delta_1 < 0$. Hence, nonlinear interactions introduce asymmetry in the distribution, and we can expect that with time, gravity will induce a non-zero skewness in the distribution.  The earliest (to our knowledge) analytic results of weakly nonlinear calculations were obtained by Zeldovich et al. and Peebles \\cite{{zel1},{dor},{lss1}}. They considered power-law initial conditions, $P_L \\, \\propto \\, k^n$, and showed that unless $n < 4$, the nonlinear terms become larger than linear at large scales, and the perturbation theory breaks down. A somewhat more systematic approach, going beyond scaling relations was introduced in the eighties by Juszkiewicz et al. \\cite{{j1},{j2}} and Vishniac \\cite{vis}, who used Eulerian perturbation theory and the random phase approximation to study the first non-trivial, weakly-nonlinear corrections for second moments -- $ P(k,t)$ and $\\xi(r,t)$. Peebles \\cite{lss1} used the same approach to calculate the gravitationally induced skewness, while Fry, in a truly seminal paper, showed how to calculate a reduced correlation function of arbitrary order $N$, using the perturbative expansion, truncated at the $(N - 1)$ -st order \\cite{jf1}.  Unfortunately, N-body simulations, available in the early eighties were too crude to study the range of validity of perturbative methods. Moreover, the galaxy surveys at scales, large enough to probe the weakly nonlinear regime of clustering, were dominated by noise. This made further progress in analytic calculations difficult and the field became dormant for several years. With time, the dynamical range of the simulations improved by more than an order of magnitude \\cite{nbody2}. Simultaneously, the length scales in galaxy surveys, at which correlations can be accurately measured, increased by an order of magnitude \\cite{obs2}, making weakly nonlinear theory useful in their analysis. As a result of these developments, the study of the evolution of the mass distribution in the expanding universe has become a rich topic for analytic perturbative methods. The revival of enthusiasm for analytic calculations produced many new results. It is our purpose to briefly summarize some of these new results here. For a more detailed approach, see e.g. the recent lecture notes of Bouchet \\cite{varenna} and the review by Sahni \\& Coles \\cite{varun}.  \\subsection{Spatial smoothing}  The last issue we need to discuss before moving on to the next section is spatial smoothing. It is useful to employ a field $\\bar\\delta$, such that $\\delta({\\bf x},t)$ is replaced by $\\bar\\delta_1 + \\bar\\delta_2 + \\ldots$, with \\begin{equation} \\bar\\delta_N ({\\bf x},t) \\; = \\; \\int \\, \\delta_N({\\bf x'},t)\\,w(|{\\bf x - x'}|\\,/R)\\,{\\rm d}^3x' \\; , \\label{eq:filter} \\end{equation} where $w(x/R)$ is a spherically symmetric window function. $R$ is the characteristic comoving smoothing length: $w$ decreases rapidly for $x/R \\gg 1$. Its Fourier representation acts as a low pass filter, suppressing the fuctuations with wavenumbers $k \\gg 1/R$. The usual normalization condition requires $\\int \\, {\\rm d}^3r \\,w(r) = 1$. Smoothing over a ball of radius $R$ is also called  ``top hat'' smoothing; another popular filter is a fuzzy ball with a Gaussian profile of width $R$. Note that the effects of smoothing and gravitational evolution are interchangeable only for $\\delta_1$; for terms of second or higher order, these two effects {\\em do not} commute (this is evident from eq. (\\ref{eq:quad}); see also \\cite{gr1}).  For a power-law  initial spectrum, the smoothed linear variance is \\begin{equation} \\langle\\bar\\delta^2_1\\rangle \\, = \\, \\sigma_L^2(R,t) \\, = \\, \\int w(x_1/R)w(x_2/R)\\xi_L(|{\\bf x}_1 - {\\bf x}_2|) \\,{\\rm d}^3x_1{\\rm d}^3x_2 \\; \\propto \\, R^{-(n+3)} \\; , \\label{eq:sigmaL} \\end{equation} so the condition $n > -3$ ensures that the clustering process is hierarchical, and proceeds from small to large scales. (this constraint, together with the Zeldovich-Peebles condition $n < 4$, broadly determines the range of ``reasonable'' values of $n$; a mix of observational arguments and theoretical prejudce shrinks this range to $-2 \\lsim n \\lsim 1$, see e.g. \\cite{lss2}). The smoothing procedure is needed for at least two reasons. First, at any stage of the evolution, in order to be able to use the perturbation theory, we need to smooth-out the strong nonlinearities. In all hierarchical models, this can be done by any a low pass filter of width $R > R_{nl}$, where $R_{nl}(t)$ is the current nonlinear scale, defined by the condition $\\sigma^2_L(R_{nl},t) = 1$. The second reason for introducing the smoothing is practical: all observable quantities, derived from galaxy surveys or numerical experiments, involve spatial smoothing. Hereafter, we will deal with smoothed fields almost without exceptions (which will be explicitly stated), so to simplify notation, we will omit the bars over $\\delta$ and $\\delta_N$. Generally, we will also assume Gaussian initial conditions (again, all exceptions will be properly described). ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602120_arXiv.txt": {
+        "abstract": "We present initial results from a redshift survey carried out with the Low Resolution Imaging Spectrograph on the 10~m W. M. Keck Telescope of a field 14.6 arc-min$^2$ in solid angle.  In the redshift distribution of the 106 extragalactic objects in this sample we find five strong peaks, with velocity dispersions of ${\\sim}500${\\kms}.  There is evidence for a non-uniform areal density of objects in at least two peaks.  These peaks have characteristics (velocity dispersions, density enhancements, spacing, and spatial extent) similar to those of nearby galaxy structures (e.g., walls and clusters), and these are expected in a survey of this kind.  We suggest that the prominence of these structures in our survey relative to that in other surveys can be attributed to our $K$-selection and dense sampling. ",
+        "introduction": "This the first in a series of papers describing the results of a deep survey of faint field galaxies in a single field centered at RA(J2000)~$00^h\\,53^m\\,23^s\\!\\!.20$, Dec.~$+12^{\\circ}\\,33'\\,57''\\!\\!.5$.  The field is from the HST Medium Deep Survey (Griffiths et al 1994), randomly selected on the basis of high Galactic latitude ($b = -50^{\\circ}$) and low reddening ($A_V = 0.13~{\\rm mag}$, Burstein \\& Heiles 1982).  Our redshifts were acquired with the Low Resolution Imaging Spectrograph (Oke et al 1995) on the 10~m W. M. Keck Telescope over a rectangular strip 2 x 7.3 arcmin$^2$ extending north-south, centered on the HST field.  Our primary sample comprises all 155 objects with $K < 20$ mag in the 2 x 7.3 arcmin$^2$ field.  The photometry and the definition of the sample for spectroscopic work is described in Paper II of this series, Pahre et al (1996).  We reach higher galaxy surface densities than the $I$-selected CFRS survey (LeF\\`evre et al 1995); our survey complements and extends to fainter objects the $K$-selected sample of the Hawaii group (Songaila et al 1994).  Paper III in this series (Cohen et al 1996) provides a detailed description of the redshift survey.  Of the 155 objects in the sample, 90 have spectra typical of normal galaxies, three are quasars or broad-lined AGNs and 19 are Galactic stars.  Of the remaining objects, 35 were observed, but no redshift could be determined, and eight were not observed at all. The effects of incompleteness in the sample selection and redshift identification are discussed in Papers II and III; both are irrelevant for the present work.  The median redshift $z$ of the 93 extragalactic objects in the main sample is $z =$ 0.57.  An additional 13 galaxy redshifts were determined for objects in this field which were slightly fainter than the $K < 20$ mag limit or which lie slightly outside the spatial boundaries of the $K$ selected sample. ",
+        "conclusions": "\\subsection{Effects of sample definition decisions}  Although there are many other faint object redshift surveys, none have found as much strong redshift clustering as is presented here.  We believe that this can be attributed to differences in sample definition.  This survey of objects goes to high galaxy number density, $4\\times 10^4$ per square degree at $K=20$, comparable to the ``Hawaii'' survey (Songaila et al, 1995) but with greater completeness at the faint end, and deeper than the CFRS (LeF\\`evre et al 1995) or the $B$-selected LDSS--2 survey (Glazebrook et al 1995).  Infrared selection is not subject to the same biases towards late-type spirals and irregulars as is found in $B$-selected samples. Our redshifts are measured with good precision, allowing good resolution of the peaks and their velocity dispersions.  Finally and most importantly, this sample is not sparse-sampled, i.e., (almost) every object with $K<20$~mag in the field is observed, so we do not miss structures of limited angular extent.  For very sparsely sampled data, one obtains no statistically valid peaks at all in the redshift distribution, c.f.\\ the CFRS (Lilly et al 1995).  As one samples less sparsely, ``walls'' with substantial velocity dispersions, as in the ESO survey (Bellanger \\& de Lapparent 1995), appear.  Suggestions similar to these regarding the effect of different sampling schemes have been offered by de Lapparent et al (1991) and by Ramella et al (1992) to explain why some redshift surveys did not see structures such as the ``Great Wall'' while others, e.g., the CFA survey (de Lapparent et al 1986), did.   \\subsection{Structure morphology}  In general, imaging surveys for galaxy structures find clusters and filaments because these appear as angular patches of higher projected number density, but do not find walls because the projected number density is relatively invariant to collapse from three-dimensional volumes into two-dimensional walls.  On the other hand, redshift surveys like this one, in small angular areas, are unlikely to hit compact clusters and filaments, but should pierce any walls. Because the structures in the present sample are seen in the redshift distribution but not in the imaging, there is some evidence that they are wall-like.  The derived velocity dispersions $\\sigma_v$ are small and reminiscent of those of sparse clusters of galaxies.  Zabludoff et al (1990) found a median $\\sigma_v =740$~{\\kms} for 69 nearby, moderately rich Abell clusters.  Our observed $\\sigma_v$ are much smaller than that found for superclusters; for example, Postman, Geller, \\& Huchra (1988) found $\\sigma_v =1300$~{\\kms} for primary members of the seven Abell clusters in the Corona Borealis supercluster, while Small (1996) found 1800~{\\kms} for the entire supercluster.  If the cluster Cor Bor itself is excluded, then the value drops to $<1000$\\kms (Postman, Huchra, \\& Geller 1992; Zucca et al 1993).  Small (1996) also observed many individual lines of sight through the supercluster and found $400<\\sigma_v <1200$~\\kms; these lines of sight are comparable in extent to our field, while our derived $\\sigma_v$ are among lowest values measured along individual Cor Bor lines of sight.  The velocity dispersion for the Great Wall, however, is 230~{\\kms} measured over 3$^\\circ$ cells (Ramella et al 1992), while that for groups in the SSRS2 (DaCosta et al 1994) is $\\sim 200$~\\kms.  Our $\\sigma_v$ are thus comparable to Great-Wall-size structures, poor clusters of galaxies, or lines of sight through superclusters.  The separations between structures are also comparable: the comoving distances of the five strongest redshift peaks are 915, 981, 1228, 1364, and 1485$h^{-1}$~Mpc, while the distance between the Cor Bor supercluster and the one immediately behind it at $z \\approx$ 0.12 (Goia et al 1982, Sarazin et al 1982, and Small 1996) is 130$h^{-1}$~Mpc.  In this regard our redshift distribution is also reminiscent of that found by Broadhurst et al (1990), although the features in ours are not ``periodic.''  At $z$ = 0.6, our field has a size of $1.7\\times 0.5\\,h^{-2}$~Mpc$^2$. We take the number density of clusters in the universe to be of order of a few times $10^{-5}\\,h^3~{\\rm Mpc^{-3}}$, as reported for Shectman (1985) clusters (Bahcall, 1988) and $R\\geq 1$ Abell galaxy clusters (Postman et al 1996).  The fact that we observe 5 clusters along a randomly selected line of sight implies a typical cluster radius of between $2$ and $5\\,h^{-1}~{\\rm Mpc}$, depending on the exact number density and the world model.  These lengths are on the order of, if a bit higher than, locally measured cluster radii, and are consistent with the marginally non-uniform angular distributions we find above.  That these structures might be sparse clusters is also consistent with their small measured velocity dispersions.  If we ``run the clock backward'' on structure formation, it seems unlikely that the huge walls and rich clusters seen locally would completely disperse into a uniform galaxy distribution by redshift of $1$.  It would be surprising if the high-redshift counterparts of local walls were {\\em not\\/} found in deep redshift surveys. (Rich clusters, on the other hand, only contain a small fraction of all galaxies and are much less likely to be encountered).  Many groups have found similar or related structures.  Bellanger \\& de Lapparent (1995), presenting the first results of the ESO-Sculptor Faint Galaxy Redshift Survey for galaxies with $R<20.5$~mag over 0.28~deg$^2$, have found what they call ``walls'' analogous to the Great Wall seen in the local universe.  The galaxy distribution they see is strongly clustered in the line of sight, consisting of walls and voids, although their redshift histogram does not show structures as strongly peaked in redshift space as those presented here (for the reasons given above).  A structure, interpreted as a normal dense galaxy cluster at high redshift, was found by LeF\\`evre et al (1994), consisting of a group of 12 galaxies around a quasar at $z=0.98$, with velocity dispersion $\\sigma_v=955$~{\\kms} and transverse structure on a characteristic scale $2\\,h^{-1}$~Mpc.  Of course an observation of transverse structure is not an argument against wall morphology because if a sheet becomes self gravitating it must break up into structures that as wide as the wall is thick (Ostriker \\& Cowie 1981). Hutchings et al (1995) have detected compact groups of galaxies with a radius of $\\sim 1$~arcmin (0.25$h^{-1}$ Mpc) around 14 QSOs with $z\\sim 1.1$.  Ellingson et al (1991) (see also Ellingson \\& Yee 1994) find galaxy clusters around QSOs with $\\sigma_v\\sim 400$--$500$~\\kms. Studies of quasar absorption lines in QSO pairs and in individual objects also suggest the existence of superclusters at $z\\approx 2.5$ (Dinshaw \\& Impey 1996).   \\subsection{Critical future observations}  Given these observations and our interpretation, it is possible to predict the outcome of future observations which are crucial in determining the statistical, morphological, and physical properties of these structures.  If they are wall-like, then we expect to see coherence as we extend the survey to adjacent fields.  Even if the clumps are sparse clusters, local observations of Cor Bor and the Great Wall suggest that the clumps will group into large two-dimensional sheets.  Redshift surveys in adjacent fields are essential to answering these questions of morphology.  There are many redshift surveys undertaken by different groups, and it is important to demonstrate that the differences among these surveys in prominence of structures in the $z$ distribution is indeed attributable to differences in sample definition.  It is imperative that the objects in this survey (and adjacent fields) be reselected with photometric and sparse-sampling criteria matched to other surveys; comparison could then be used to confirm that the structures are common and this field is not ``special.''  Also, the morphology-density relation (Dressler 1980) suggests that changing to bluer selection bands should reduce the percentage of galaxies lying inside the structures.  Finally, similar samples in widely separated fields will provide us with statistics for these structures, such as typical separations, velocity dispersions, filling factor, etc.  It is possible that some of the much-discussed field-to-field variations found in galaxy counts, angular correlation functions, and redshift distributions (e.g., Koo \\& Kron, 1992) can be attributed to the field-to-field variations in the walls along the line of sight.  If, as we strongly suspect they will, further observations do substantiate the view that roughly half the old galaxies are localized in walls, then a major challenge will be to to determine whether the galaxies in these walls have virialized.  At present there appear to be no good observational arguments against this view."
+    },
+    "9602/astro-ph9602066_arXiv.txt": {
+        "abstract": "Luminosity functions from theoretical stellar evolution calculations are compared with observed ones of several galactic globular clusters (M30, M92, M68, NGC6397, M4, M80, NGC6352, NGC1851). Contrary to earlier results of Faulkner \\& Stetson (1993) and Bolte (1994) we find no significant discrepancy that could indicate the neglect of important physical effects in the models. However, it is confirmed that the subgiant branch is the most sensitive part and shows the largest deviations in the luminosity function comparison, if parameters are unappropriate. We also find that the main sequence is suited less than the Red Giant Branch for the calibration of theoretical luminosity functions. While for individual clusters different changes in the model assumptions might resolve mismatches, there is no systematic trend visible. It rather appears that the quality of the luminosity function in the subgiant part is insufficient and that improved observations of this particular region are necessary for a better comparison. At the present quality of luminosity functions theory is in agreement with observations and a postulation of WIMPs acting in stellar cores does not seem to be justified. Actually, fits using isothermal core models on the main sequence appear to be worse than those with standard stellar evolution assumptions. ",
+        "introduction": "Globular clusters are compared with the theory of stellar structure and evolution by use of two functions: the colour-magnitude-diagram (CMD) displaying the surface properties of the stars and the luminosity function (LF) measuring the initial mass function (IMF) for the unevolved lower main sequence and the speed of evolution after the turn-off. According to Faulkner \\& Swenson (1993) the agreement with the observed LF of the globular clusters M30, M92, M68 and NGC6397 is sufficiently bad and unresolvable with standard stellar structure physics that an additional physical mechanism acting during the main-sequence evolution has to be postulated. In particular, they found that the observed LFs show an excess of stars on the subgiant branch, which is interpreted as a prolonged subgiant phase resulting from a higher central hydrogen abundance at the turn-off and an initially broader hydrogen burning shell, which is getting very narrow during the subgiant phase. The shell, developing at the end of the main sequence, can be made broader, if some part of the stellar core would be isothermal. Faulkner \\& Swenson (1988, 1993) investigated the evolution of low-mass stars on the main sequence and the subsequent subgiant and giant branch under the assumption of such an isothermal core extending over the innermost 10\\% of the stellar mass and found a better agreement between theoretical and observed LF. Bolte (1994) confirmed these results for M30 using an improved $V$-band LF but the same theoretical models (Bergbusch \\& VandenBerg 1992). As a side-effect of the isothermal cores, globular clusters would be younger by about 20\\% than usually thought (Faulkner \\& Swenson 1993).  The origin of the efficient energy transport needed in the stellar center was suggested to be found in the presence of WIMPs accreted by the star during its main-sequence phase (Faulkner \\& Swenson 1993). Bolte (1994) already mentioned that only a very small region in the WIMPs' mass--cross-section--phasespace is left over from various experiments and theoretical expectations. In addition, solar models without an isothermal core better fit the helioseismological data (Cox, Guzik \\& Raby 1990; Cox \\& Raby 1990; Kaplan et al.~1991). A similar result is reported by Basu \\& Thompson (1996) for a solar seismic model with a reduced central temperature simulating the effect of some additional energy transport by WIMPs. These and additional arguments led Bolte to the conclusion that WIMPs are unlikely to be a major constituent of halo dark matter accreted by stars and affecting their core structure. Since his paper the parameter space for WIMPs -- when assumed that they are neutralinos -- has constantly been shrinking (Fornengo 1994; Mignola \\& Berezinsky, private communication); thus his conclusion is more justified than ever. However, depending on new ideas the nature of WIMPs might be different and they possibly might exist and act in stars.  Independent of this, the result of Faulkner \\& Stetson (1993) and Bolte (1994) remains: that there was an apparent mismatch between theoretical and observed GC luminosity functions, which can be reconciled by an isothermal core. The reason for the isothermality would remain unclear. Since this discrepancy is a severe challenge to stellar structure theory, we re-investigated the quality of both theoretical and observational LFs. In Sect.~2 we will discuss the method of obtaining a theoretical LF and problems and errors associated with the observed LFs. In the following section, we will perform standard comparisons with the Faulkner \\& Stetson and other clusters and discuss the results. In Sect.~4 we will investigate variations in the standard physics and the effect of  isothermal cores. At last, our conclusions follow in Sect.~5. ",
+        "conclusions": "The prime intention of this paper has been to independently investigate whether there indeed exists a systematic and significant discrepancy between observed and theoretical luminosity functions of Globular Clusters as claimed by Faulkner \\& Stetson (1993) and Bolte (1994). We have demonstrated in Sect.~3 that for the clusters used in their papers and for additional data by Piotto \\& Saviane (1996) the agreement between our standard luminosity functions and the observations is very good for reasonable assumptions about cluster composition, distance and age. Actually, we have used values from the literature for these parameters, but exploited the range of uncertainties to find the best fit. All data points can be fitted within $3\\sigma$ of the statistical errors, except for the main sequences. To illustrate the influence of pure number statistics we have performed the following test: we took one of our standard LFs (16 Gyr; $s=2$, $Z=2\\, 10^{-4}$) and constructed a synthesized LF by the rejection method to construct random deviates. Fig.~\\ref{syn-lf} shows the comparison between the synthetic and the theoretical LF for a total number of stars and a number of brightness bins comparable to the cases discussed in this paper. The similarity, for example with Fig.~\\ref{m30-lf1}, demonstrates that our fits are perfect within the statistical errors, and that deviations like those at $M_V = 3.6$ or $3.8$ are to expected. In fact, some of our cases show a better agreement than would be consistent with the statistics. In these cases, the fit parameters might be overdetermined. Of course, the systematic deviations on the main sequences are a clear indication of non-statistical errors. Since all clusters by  Piotto \\& Saviane (1996) are deficient with respect to the main-sequence parts of the theoretical LFs, while those used by Faulkner \\& Stetson (1993) are not, we ascribe this to an underestimate in the completeness-correction applied by the observers.  \\begin{figure} \\centerline{\\epsfxsize=0.65\\hsize \\epsfbox{syn-lf.ps}} \\caption[]{Theoretical and synthetic luminosity functions for standard parameters (16 Gyr; $s=2$, $Z=2\\, 10^{-4}$). 26000 random points in the N-$M_V$ plane were chosen, out of which 6873 fell below the theoretical probability function and were accepted. The number of brightness bins is 30. $1\\sigma$ errors are indicated } \\protect\\label{syn-lf} \\end{figure}  Alternatively, one could speculate that the IMF for these clusters is flatter than assumed. This uncertainty -- and the additional one of the helium content -- prevented us from using the main sequence for the normalization of the LFs, in spite of the much smaller statistical errors. Instead, we used the RGB for normalization, whose LF-slope is an extremely stable quantity independent of all input variables. In the case of M80 we demonstrated that the use of the main sequence would lead to a ``discrepancy'' in the standard LF and its resolution by an isothermal core, just as was the case in the papers triggering the present work.  The main sequence being inadequate for a detailed comparison and the red giant branch being robust against model changes, the subgiant branch is left to reflect model differences, parameter influences and potential problems. However, the present observations do not resolve the LF in this region sufficiently well and with adequate accuracy. We propose that in the future observations aiming at obtaining the LFs of GCs should concentrate exactly on this region.  We found one cluster (NGC6352), which withstands all attempts to fit a theoretical LF. The shape of the observed one is in fact very strange and looks like that of a very low-metallicity system in terms of the missing subgiant break, and like one of high helium content with respect to the slope at the TO (Ratcliff 1987). However, within the parameter range investigated by us, we could not obtain any good fit. We rather suspect that the subgiant break was not resolved properly, and this should be checked by further observations.  Of all quantities we have tested for their influence on the LFs, we found that metallicity and distance (see the example of M80; Figs.~\\ref{m80-lf1a} and \\ref{m80-lf1b}) have the largest influence. Age variations of up to 2 Gyr can be compensated by a luminosity shift smaller than the quoted distance uncertainty (because old clusters change their TO-luminosity hardly with age). Discrepancies in the LF should therefore first raise the question of correct distance or metallicity determinations. In this context information about $\\alpha$-element enrichment is important, too.  In Sect.~4 we repeated the LF-calculations, but with a very efficient energy transport in the innermost 10\\% of our models (following again Faulkner \\& Swenson 1993 and Bolte 1995). Our results show that the fits are not improved at all, but rather get worse. Since we found no evidence for a LF-discrepancy, it is not necessary to discuss properties of hypothetical WIMPs for additional energy transport in the core of main-sequence stars.  To summarize, we have shown that the agreement between observed and theoretical luminosity functions appears to be very good, with problems only arising on the lower main sequence (corrections for completeness might have to be improved) and for the exact shape of the subgiant bump and break (resolution in brightness), which have the potential to yield information about metallicity, age and distance of the cluster."
+    },
+    "9602/hep-ph9602263_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "In presently favoured extensions of the Standard Model, the space of the scalar fields contains many directions in which the potential is almost flat out to large field values.\\footnote{Flat directions occur rather generically in supersymmetric theories as a consequence of a non-renormalization theorem \\cite{GrisaruSiegelRocek}, \\cite{Seiberg}. The renormalizability requirement implies that there are flat directions protected to all orders in perturbation theory. The terms allowed then have couplings suppressed by some power of the Planck mass.} In some of these directions the mass-squared is likely to be negative, leading to a nonzero vev whose magnitude will be large because the potential is flat. Fields of this kind, with  large vevs and almost flat potentials, have been called flatons \\cite{decay}, and it has long been recognized that they may be cosmologically significant \\cite{problem,dinefisch,coughlan,decay,yam,therm,Enqvist,Ross,yam2,interm}. \\footnote {Note the etymology. The term `flaton' refers to the {\\em flat\\/} potential, not to in{\\em flat\\/}ion. Conversely, the familiar word `inflaton' refers to the field which is slowly rolling during ordinary {\\em inflat\\/}ion. The term `flaton' was invented long before it was realized that such a field can give rise to a different type of inflation (thermal inflation).}  Recently a previously almost \\cite{therm} unnoticed aspect of flaton cosmology was studied, and termed thermal inflation \\cite{gutti,therminf}. Thermal inflation is made possible by the flatness of the potential, which near the origin allows the (positive) finite temperature contribution to the mass-squared to dominate the true (negative) mass-squared long after the thermal contribution to the energy density has become negligible. If it occurs, thermal inflation completely alters the standard cosmology, and one purpose of this paper is to examine carefully the features of the finite temperature potential that are supposed to lead to it. Taking into account the validity of the various approximations, we shall verify that thermal inflation indeed takes place, provided that the field is trapped at the origin and that it has sufficiently strong interactions with light fields.  When thermal inflation ends, and the field rolls away from the origin, topological defects may be produced just as in the more familiar case of a non-flat potential. The second purpose of this paper is to see what effect the flatness of the potential has on these defects, and their cosmology. We are particularly interested in the case that the GUT higgs particles are flatons, so that the defects produced are monopoles and (depending on the GUT) cosmic strings, but we shall also study the case of global symmetry breaking.  Recently it has been pointed out \\cite{klslatest} that the interaction of other scalar fields with the flaton might cause inflation even without thermalization. In that case the situation is more model dependent and we shall not study it in this paper, though much of our discussion of topological defects will still be applicable.  Thermal inflation is expected to take place after ordinary (slow-roll) inflation, and to last for fewer $e$-folds. Such a late epoch of inflation can play a welcome role in diluting the abundance of unwanted relics, while leaving unaltered the large-scale density perturbation accounting for the cmb anisotropy and large scale structure. (As we shall discuss this remark applies whether the perturbation is caused by a  vacuum fluctuation during ordinary inflation, or by topological defects such as strings or textures.) There can be more than one epoch of thermal inflation, so that a second epoch can dilute relics left over from a first epoch, and we shall see that this is indeed essential if one is to have a GUT phase transition with a flat potential.  Before getting into more detail we need to be precise about what is meant by a `large' vev, and a potential which is `almost flat'. In this paper we take these terms  to be defined with respect to the energy scale $10^2$ to $10^3\\,{\\rm GeV}$, which is the scale of supersymmetry breaking as defined by the masses of the supersymmetric partners of known particles \\cite{susy}. Thus the vev $M$ of a flaton field is much bigger than the above scale, whereas the the energy scale $|V''|^{1/2}$ defined by the curvature of its potential is taken to be of order that scale. In particular the mass $m$ of a flaton particle is taken to be of that order. Since the potential as well as its first derivative vanishes at the vev, its value $V_0$ at the origin is of order $m^2M^2$.  Each flaton is either a gauge singlet or a GUT Higgs field. A GUT is nowadays considered optional because it is difficult to implement, and can perhaps be avoided by appealing to superstring unification near the Planck scale. On the other hand the observed Standard Model couplings are compatible with the existence of a GUT, with Higgs vevs of order $10^{16}\\GeV$ (plus possible additional intermediate scale vevs \\cite{Mohapatra}). If one accepts the existence of a GUT then it is attractive to suppose that the Higgs fields are flatons because one might then be able to understand the magnitude of the vev without inserting parameters by hand \\cite{flatguts}.  Whether or not there is a GUT, one expects to find some flatons which are gauge singlets. To understand the possibilities here, note first that each flaton will be a complex field.\\footnote {In a supersymmetric theory all scalar fields are complex since the corresponding left- or right-handed spin-half field has two degrees of freedom.}. If there is a single flaton field, and it is charged under a global $U(1)$ symmetry, then its vev will spontaneously break that symmetry. The prime candidate for a global $U(1)$ symmetry is the Peccei-Quinn symmetry associated with the axion, and indeed a flaton model for that symmetry has been proposed \\cite{flataxion}. With more flaton fields one can have higher global symmetries, as we mention at the end of Section V in connection with topological defects.  However there is no need for a flaton field to be charged under a continuous symmetry; on the contrary, it is quite reasonable on the basis of current thinking to expect flaton fields which are charged under at most discrete symmetries. For them, the vacuum manifold consists of discrete points instead of an entire circle. The possibilities that we have in mind are of two kinds. The first are some or all of the moduli fields \\cite{problem,dilaton,Banks,Randall,Steinhardt,Dine}, which are ubiquitous in versions of supergravity motivated by the superstring. For the present purpose, a modulus may be defined as a field whose potential is exactly flat in the limit of unbroken symmetry. It is widely supposed that one or more of the moduli is a flaton, with a vev of order the reduced Planck mass $\\mpl=(8\\pi G)^{-1/2}=2.4\\times 10^{18}\\,{\\rm GeV}$.\\footnote {The origin of a field is defined to be a fixed point of its symmetries. With all relevant fields at the origin the potential has zero slope, where `relevant' means those fields that are coupled to the one in whose direction the slope is being defined. We are assuming that this is the situation for the flatons (or at least that any couplings to other fields displace the minimum by an amount small compared with the vev). In the case of a modulus there is more than one fixed point with the separation between fixed points of order $\\mpl$. In this context the vev may be defined as the distance to the nearest fix point, and the statement that the vev is of order $\\mpl$ simply means that it is not close to any particular fixed point. } (Lower values may be possible so that GUT Higgs fields may be moduli \\cite{flatguts}, but for simplicity, the term `moduli' will be used in what follows to denote only the case where the vev is of order $\\mpl$. It is also possible that moduli have masses much bigger than $1\\TeV$, in which case they are not flatons.) Now consider the opposite possibility, of gauge-singlet flatons whose potential is not flat even in the limit of supersymmetry \\cite{therminf,ewan}. Extensions of the Standard Model can easily contain such fields, and their potentials will resemble Eq.~(\\ref{eq:potential}) below, with couplings perhaps or order 1. As a result, their vev is typically much smaller than the GUT scale, perhaps of order $10^9$ to $10^{11}\\GeV$. ",
+        "conclusions": "The topics discussed in this paper relate to a considerable body of ongoing research at the moment, whose common theme is the form of the effective potential in the early universe, and the consequent cosmological effects of the scalar fields. Some of this research continues to address the traditional, and still very important, question of how to implement an early era of ordinary (slow-roll) inflation. Here we have focussed instead on aspects of `thermal inflation', which occurs if at all after ordinary inflation and lasts for only a few $e$-folds.  The fields which can give rise to thermal inflation are `flaton' fields. By definition, these fields have `flat' potentials and `large' vevs, where these terms refer to the mass scale $10^2$ to $10^3\\GeV$. They arise naturally in currently favoured extensions of the Standard Model, along with the opposite case of fields with `flat' potentials but zero vevs.  A flaton field gives rise to thermal inflation if it is trapped at the origin by virtue of the finite temperature correction to its effective potential, after the thermal contribution to the energy density has become comparatively small. According to a calculation made a decade ago by Yamamoto, this trapping is expected to occur for a flat potential (provided that the flaton field is indeed in the vicinity of the origin), with thermal inflation ending only when the local minimum of the effective potential at the origin almost completely disappears. Part of our objective was to investigate carefully the rather delicate assumptions needed to arrive at this conclusion, using some modern perspectives particularly on the tunneling rate out of the false vacuum. Happily, we confirm the conclusion.  The actual cosmology of flaton fields, which determines whether a given flaton field  will actually find itself in the vicinity of the origin so that thermal inflation can take place, is at present a rather rapidly moving research area. We have not attempted any new advance here, contenting ourselves with a brief snapshot of the situation as it stands at present. Instead we have addressed a question which has received relatively little attention, which is the nature and cosmology of topological defects forming at the end of thermal inflation. One would expect {\\it a priori} that their cosmology might be significantly affected by the flatness of the potential. The answer to this question turns out to depend on the type of defect.  A particular focus for the second part of our investigation has been the possibility that there is a GUT, whose Higgs fields have a flat potential. This possibility is quite natural, and it leads to a cosmology very different from the usually considered case of a non-flat Higgs potential. GUT symmetry breaking, if it occurs, will be preceded by an era of thermal inflation. The Higgs particles produced at the transition are cosmologically dangerous because they are light and long-lived, and to dilute them one needs a second era of thermal inflation. This period of thermal inflation will also dilute the monopole abundance sufficiently, but {\\it will not eliminate cosmic strings} because it lasts only a few $e$-folds. A simple picture of a string network evolution in an inflationary Universe is as follows. First as the Universe expands the strings will quickly reach the density of about one (long) string per horizon volume. After that the network freezes out since there exists no causal process that could incite nontrivial dynamics. This means that from then on the average correlation length grows exponentially as the expansion rate, {\\it i.e.\\/} as $\\xi(t)=\\xi(t_0)\\,{\\rm e}^{H(t-t_0)}$. After about 10 $e$-foldings the correlation length has not grown more than $\\xi\\sim {\\rm e}^{10} /H$; this scale will come back within the horizon after $\\ln (a/a_0)\\sim 10$ which in radiation era means at the temperature $T=T_D\\times {\\rm e}^{-10} \\simeq T_D\\times 10^{-4.3}$; after this strings will quickly reach the radiation era scaling solution. This temperature is far above the relevant scale for onset of cmbr anisotropies and structure formation, keeping the strings cosmologically interesting.  With all this in mind, we expect that cosmological relevance of cosmic strings is specified solely by the mass per unit length of strings. We find that the GUT Higgs vev needed for the strings to give an observable signature in the cmb anisotropy is a few times bigger than in the case of a non-flat potential, but still compatible with the typical estimate $\\simeq 2\\times 10^{16}\\GeV$ of the scale at which the Standard Model gauge couplings become unified. Perhaps the main message is that the flatness of the potential has surprisingly little effect on the string cosmology.  We also discuss global defects, and in particular focus on the case of global textures and monopoles. We arrive to a rather surprising conclusion that global texture and monopole dynamics is not affected at all by the flatness of the potential. This has a simple explanation: the potential is simply not flat enough! The `flatness scale' $m_0\\sim 10^3$GeV is large in comparison to the scale relevant for cosmology: $T\\sim 1$eV, so the flat potential suffices to confine late textures to the vacuum manifold. This leaves both textures and global monopoles as a viable candidate for structure formation with the scales of symmetry breaking given by $M\\sim 1.4\\times 10^{16}$GeV and $M\\sim 1\\times 10^{16}$GeV, respectively. Finally in passing we give an argument that domain walls are still a problem.  In summary, the study of fields with flat potentials and large vevs is becoming another field on which particle physics increasingly meets cosmology. Our hope is that either cosmology will put constraints on particle physics models, or particle physics will offer some interesting cosmological phenomena, at best tell us something about the origin of structures in the Universe. Since this area of research is developing very fast, it is very likely that in near future one will be able to make more definite statements about the formation and nature of defects in theories with flat potentials, and hence make a more definite prediction on their cosmological relevance.  \\subsection*{Acknowledgements}  We thank David Bailin, Beatriz de Carlos, Anne Davis, Andrei Linde, Graham Ross and Ewan Stewart for useful discussions. T.P. acknowledges support from PPARC and NSF. T.B. acknowledges support from JNICT (Portugal)."
+    },
+    "9602/astro-ph9602030_arXiv.txt": {
+        "abstract": "A strong bar rotating within a massive halo should lose angular momentum to the halo through dynamical friction, as predicted by Weinberg.  We have conducted fully self-consistent, numerical simulations of barred galaxy models with a live halo population and find that bars are indeed braked very rapidly.  Specifically, we find that the bar slows sufficiently within a few rotation periods that the distance from the centre to co-rotation is more than twice the semi-major axis of the bar.  Observational evidence (meagre) for bar pattern speeds seems to suggest that this ratio typically lies between 1.2 to 1.5 in real galaxies.  We consider, a number of possible explanations for this discrepancy between theoretical prediction and observation, and conclude that no conventional alternative seems able to account for it. ",
+        "introduction": "Chandrasekhar (1943) showed that a massive object moving though a background ``sea'' of light particles would experience a drag.  The force is the gravitational attraction by the wake produced by the motion of the massive object (see e.g., Binney \\& Tremaine 1987, \\S7.1).  When the mass, $M$, of the perturber is much larger than that of the individual background particles, the acceleration takes the form $$ {d v_M \\over dt} \\propto - \\rho \\; M \\; f\\left( {v_M \\over \\sigma} \\right), \\eqno(1) $$ where $\\rho$ is the background density and $f$ is a function of the ratio of the perturber's velocity, $v_M$, to the (assumed isotropic) velocity dispersion, $\\sigma$, of the background particles.  For a Maxwellian distribution of velocities, this function is a maximum when $v_M \\simeq 1.37\\sigma$.  A similar process must occur as a massive bar rotates inside a halo.  In this case, the bar creates a wake in the halo which lags the bar and the gravitational attraction between the bar and the wake produces a torque which removes angular momentum from the bar and adds it to the halo.  Weinberg (1985), adapting the perturbation theory approach of Tremaine \\& Weinberg (1984a), estimated the magnitude of the frictional drag force.  Assuming a massive, rigid bar rotating in an isothermal halo, he concluded that the spin-down time for the bar would be as little as five bar rotations, for reasonable parameters!  Early low-quality, but fully self-consistent disc-halo simulations (Sellwood 1980) had previously revealed a rapid loss of angular momentum to the halo once a bar formed in the disc, at a rate roughly consistent with Weinberg's prediction.  To our knowledge, no other such simulations have been conducted in the past 15 years; Combes et al.\\ (1990) had only a bulge, not an extensive halo, of live particles, Raha et al.\\ (1991) did not evolve their models for long enough, the ``halo'' particles of Little \\& Carlberg (1991) were confined to a plane, and the bar used by Hernquist \\& Weinberg (1992) was rigid and their model had no disc.  We here report new fully self-consistent, simulations superior in many respects to those of Sellwood (1980), which were designed to reproduce the expected dynamical friction and to determine the secular changes to the bar, disc and halo in the long term. Athanassoula (work in progress) is conducting similar experiments using a direct $N$-body code on a GRAPE device. ",
+        "conclusions": "The pattern speed problem presented by dynamical friction between a bar and bulge/halo is becoming rather insistent.  Most possible solutions seem unattractive, some are excluded and others need to be stretched excessively.  It is becoming increasingly difficult to find a tenable conventional explanation.  \\bigskip \\noindent {\\bf Acknowledgments}  We would like to thank Scott Tremaine for a critical read of the manuscript.  This work was supported by NSF grant AST 93-18617 and NASA Theory grant NAG 5-2803."
+    },
+    "9602/astro-ph9602024_arXiv.txt": {
+        "abstract": "We report the discovery of a substantial population of star--forming galaxies at $3.0 \\simlt z \\simlt 3.5$. These galaxies have been selected using color criteria sensitive to the presence of a Lyman continuum break superposed on an otherwise very blue far-UV continuum, and then confirmed with deep spectroscopy on the W. M. Keck telescope. The surface density of galaxies brighter than ${\\cal R} = 25$ with $3.0 \\simlt z \\simlt 3.5$ is $0.4\\pm0.07$ galaxies arcmin$^{-2}$, approximately 1.3\\% of the deep counts at these magnitudes; this value applies both to ``random'' fields and to fields centered on known QSOs. The corresponding co-moving space density is approximately half that of luminous ($L \\simgt L^{\\ast}$) present--day galaxies. Our sample of $z > 3$ galaxies is large enough that we can begin to detail the spectroscopic characteristics of the population as a whole. The spectra of the $z>3$ galaxies are remarkably similar to those of nearby star--forming galaxies, the dominant features being strong low--ionization interstellar absorption lines and high--ionization stellar lines, often with P-Cygni profiles characteristic of Wolf-Rayet and O--star winds. Lyman $\\alpha$ emission is generally weak ($< 20$ \\AA\\ rest equivalent width) and is absent for $>$50\\% of the galaxies. We assign approximate mass scales to the galaxies using the strengths of the heavily saturated interstellar features and find that, if the line widths are dominated by gravitational motions within the galaxies, the implied velocity dispersions are $180 \\le \\sigma \\le 320$ \\kms, in the range expected for massive galaxies. The star formation rates, which can be measured directly from the far-UV continua, lie in the range $4 - 25$ $h_{50}^{-2}$ $M_{\\sun}$ yr$^{-1}$ (for $q_0=0.5$), with $8.5h_{50}^{-2}$ $M_{\\sun}$ yr$^{-1}$ being typical. Together with the morphological properties of the $z>3$ galaxy population, which we discuss in a companion paper (Giavalisco \\et 1996), all of these findings strongly suggest that we have identified the high-redshift counterparts of the spheroid component of present--day luminous galaxies. In any case, it is clear that massive galaxy formation was already well underway by $z \\sim 3.5$. ",
+        "introduction": "Several years ago we embarked on a program to assess the presence of ``normal'' star--forming galaxies at redshifts $z >3$ using a custom suite of broad--band filters (dubbed $U_n$, $G$, and ${\\cal R}$-- see Steidel \\& Hamilton 1993) designed to reach very faint apparent magnitudes with accurate colors. The technique rests on only two assumptions about the spectra of high-redshift galaxies: {\\it (i)} a roughly flat spectrum ($f_{\\nu} \\propto \\nu^0$) in the far-UV where the flux is dominated by emission from massive stars and, {\\it (ii)} a pronounced spectral ``break'' at 912 \\AA\\ (the Lyman limit in the galaxy rest frame) due to the combination of the intrinsic stellar energy distributions, the inherent opacity of the galaxies to Lyman continuum photons, and the effects of intervening absorbers (see Steidel \\& Hamilton 1992, 1993; Steidel, Pettini, \\& Hamilton 1995 [Papers I,II,III] for an extensive discussion of the technique; see also Madau 1995). The filters are ``tuned'' to selecting objects with redshifts in the range $3.0 \\simlt z \\simlt 3.5$, where the $U_n$ passband falls mostly shortward of the Lyman limit and the $G$ band is not severely blanketed by the Lyman $\\alpha$ forest---at higher redshifts this blanketing can make it very difficult to distinguish a spectral break from an intrinsically red spectral energy distribution.  While earlier work (Guhathakurta \\et 1990) had shown that Lyman break galaxies do not {\\it dominate} the galaxy counts at faint apparent magnitudes, these constraints were too weak to rule out the presence of a substantial population of such objects among the more prominent foreground blue galaxies. Our first searches were directed at the fields of QSOs with known Lyman limit absorption systems at $z > 3$; these  could then be used as ``templates'' to guide one (eventually) to galaxies in the general field, or at least to test the feasibility of detecting normal galaxies at such extreme redshifts. In Paper III we discussed together the properties of both QSO absorbers and candidates in the general field, since QSO absorbers are apparently drawn from the population of relatively luminous (but otherwise normal) field galaxies, at least at smaller redshifts (see Steidel, Dickinson, \\& Persson 1994). We found the surface density to be the same within the errors in all 5 fields studied, and we reported a surface density of ``robust'' candidates for Lyman break galaxies of $\\approx 0.5$~arcmin$^{-2}$ to an apparent magnitude limit of ${\\cal R}=25$.  Since the completion of Paper III, we have been working on extending the $U_nG{\\cal R}$ imaging technique to more random (i.e., non--QSO) fields and, most importantly, on following up the candidates with deep spectroscopy on the Keck telescope. In this paper we report the successful results of our first attempts at spectroscopy of the Lyman break candidates, new $U_nG{\\cal R}$ imaging in (much larger) random fields which confirms the surface density estimates given in Paper III, and the first measurements of $K$--band magnitudes for the $z>3$ population. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602118_arXiv.txt": {
+        "abstract": "A sample of 198 soft X--ray selected active galactic nuclei (AGN) from the {\\it ROSAT} International X--ray Optical Survey (RIXOS), is used to investigate the X--ray luminosity function and its evolution. RIXOS, with a flux limit of \\( 3 \\times 10^{-14}{\\rm erg\\ s^{-1}\\ cm^{-2}}\\) (0.5 to 2.0 keV), samples a broad range in redshift over \\({\\rm  20\\ deg^{2}}\\) of sky, and is almost completely identified; it is used in combination with the {\\it Einstein} Extended Medium Sensitivity Survey (EMSS), to give a total sample of over 600 AGN. We find the evolution of AGN with redshift  to be consistent with pure luminosity evolution (PLE) models in which the rate of evolution slows markedly or stops at high redshifts \\(( z > 1.8 \\)). We find that this result is not affected by the inclusion, or exclusion, of narrow emission line galaxies at low redshift in the RIXOS and EMSS samples, and is insensitive to uncertainties in the conversion between flux values measured with {\\it ROSAT} and {\\it Einstein}. We confirm, using a model independent \\(< V_{e}/V_{a} >\\) test, that our survey is consistent with no evolution at high redshifts. ",
+        "introduction": "X--ray properties are becoming an increasingly important tool for selecting samples of AGN, and sample sizes are approaching that of ultraviolet excess (UVX) selected optical samples.  The largest currently available sample of serendipitous X--ray selected AGN comes from the {\\it Einstein} Extended Medium Sensitivity Survey (EMSS, see Stocke \\etal 1991), with a sky coverage of \\({\\rm  778\\ deg^{2}}\\). However, the EMSS is dominated by low redshift objects (median {\\it z} \\(\\sim\\) 0.2) due to its high limiting flux (typically \\(> 10^{-13}{\\rm erg\\ s^{-1}\\ cm^{-2}}\\) in the energy range 0.5 to 2.0 keV). Deeper {\\it Einstein} surveys (Primini \\etal 1991) identified only 11 AGN. With the arrival of {\\it ROSAT} the possibilities for deeper surveys have been realised: the Cambridge Cambridge {\\it ROSAT} Serendipity Survey covers about \\({\\rm 4\\  deg^{2}}\\) and has a flux limit of \\( 2 \\times 10^{-14}{\\rm erg\\ s^{-1}\\ cm^{-2}}\\) from 0.5 to 2.0 keV (see Boyle \\etal 1995); the survey of Boyle \\etal (1994) contains 107 broad line AGN and probes to flux levels lower than \\( 4 \\times 10^{-15}{\\rm erg\\ s^{-1}\\ cm^{-2}}\\) (0.5 to 2.0 keV), but is only 70\\% complete at this flux limit and covers an area of less than \\({\\rm  1.5\\ deg^{2}}\\) in total. This survey is almost devoid of low redshift objects, with only 3 AGN of \\(z < 0.4\\), and median {\\it z} = 1.5. Even deeper surveys, with even smaller sky coverage, (Hasinger \\etal 1993; Branduardi--Raymont \\etal 1994) are still in the process of optical identification.  The {\\it ROSAT} International X--ray Optical Survey (hereafter RIXOS, see Mason \\etal 1995) occupies  a position between the EMSS and the deeper {\\it ROSAT} surveys, with a sky coverage of over \\({\\rm  20\\ deg^{2}}\\) and a limiting flux of \\( 3 \\times 10^{-14}{\\rm erg\\ s^{-1}\\ cm^{-2}}\\) (0.5 to 2.0 keV). It is constructed from  81 long (\\( > 8000\\ {\\rm seconds})\\) {\\it ROSAT} pointings made with the Position Sensitive Proportional Counter (PSPC) at the focus of the X--ray telescope. The RIXOS AGN sample consists of 198 objects, the largest {\\it ROSAT} selected AGN sample to date, and is particularly useful for investigating evolution models because of its broad sampling of redshifts.  In this paper we present the results of an analysis of the differential X--ray luminosity function (XLF) and its evolution with redshift, using the RIXOS sample. In section 2 we describe briefly the RIXOS survey and present details of the RIXOS AGN sample, while in section 3 we discuss the log {\\it N} -- log {\\it S} of RIXOS AGN compared to that of the EMSS AGN. In section 4 the methods used in this analysis are explained and in section 5 we present our results, which are discussed in section 6. Our conclusions are presented in section 7.  Throughout this paper a Friedmann model universe has been assumed, and a value of  50 \\({\\rm km\\ s^{-1}\\ Mpc^{-1}}\\) has been adopted for the Hubble constant \\( H_{0}\\); two different values for the deceleration parameter, \\(q_{0}\\) = 0 and \\(q_{0}\\) = 0.5 have been used. ",
+        "conclusions": "We have investigated the evolution of the XLF with redshift using a new sample of 198 X--ray selected AGN with a spectroscopic completeness of 93\\% at \\(3\\times 10^{-14} {\\rm erg\\ s^{-1}\\ cm^{-2}}\\) (0.5 to 2.0 keV). We find PLE models, consistent with our data, in which the XLF declines or ceases to evolve beyond {\\it z} \\(\\sim\\) 1.8, and we find no evidence for evolution beyond this redshift from the model independent \\(< V_{e}/V_{a} >\\) test. We have shown that narrow emission line galaxies at low redshift contaminating the AGN sample have little effect on the derived evolution properties. Furthermore, these conclusions are insensitive to the uncertainty in conversion of flux from the {\\it Einstein} to {\\it ROSAT} passbands, although this has a significant effect on the \\(z\\) = 0 XLF and evolution rate for \\(z\\) \\(<\\) 1.8."
+    },
+    "9602/astro-ph9602005_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "On Mar. 7, 1994, Treffers, Filippenko and Van Dyk (1994) discovered the Type Ia supernova (SN~Ia) 1994D   in the SO   galaxy NGC 4526, a member of the Virgo cluster which was was about 12 days before maximum light. Subsequently, both the light curve in Johnson's UBVRI colors and spectra a have been monitored by several groups (e.g. Filippenko et al. private communication, Smith et al., 1995) making this supernova one of the best observed until now. In addition, starting at Mar 12, IR-spectra are reported to be taken at the Royal Greenwich Observatory showing a  P-Cygni line at about 1.0 $\\mu m$ which, likely, is due to He I (Meikle 1995).  It is generally accepted that SN~Ia are thermonuclear explosions of carbon-oxygen white dwarfs (WD) (Hoyle \\& Fowler 1960). However, details of the scenario are still under debate. For discussions of various theoretical aspects see e.g. Wheeler \\& Harkness 1990, Canal 1995, Nomoto et al. 1995, H\\\"oflich \\& Khokhlov 1995, Woosley \\& Weaver 1995). What we observe as a supernova event is not the explosion itself but the light emitted from a rapidly expanding envelope produced by the stellar explosion. As the photosphere recedes, deeper layers of the ejecta become visible. A detailed analysis of the light curves (LC) and spectra gives us the opportunity to determine the density, velocity and composition structure of the ejecta. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602139_arXiv.txt": {
+        "abstract": "Today we have numerous evidences that spirals evolve {\\it dynamically\\/} through various secular or episodic processes, such as bar formation and destruction, bulge growth and mergers, sometimes over much shorter periods than the standard galaxy age of $10-15$ Gyr. This, coupled to the known properties of the Hubble sequence, leads to a unique sense of evolution: from Sm to Sa.  Linking this to the known mass components provides new indications on the nature of dark matter in galaxies.  The existence of large amounts of yet undetected dark gas appears as the most natural option.  Bounds on the amount of dark stars can be given since their formation is mostly irreversible and requires obviously a same amount of gas. ",
+        "introduction": "Until the recent years the concepts prevailing in understanding galaxy evolution have been largely dominated by the ELS scenario (Eggen, Lynden-Bell \\& Sandage 1962).  The views issued from it favour a {\\it synchronous\\/} and rapid ($\\approx 100\\,\\mathrm{Myr}$) formation of the galaxies at a particular early time in the universal expansion. In this context, the properties of present galaxies, such as their typical scale and mass, must depend directly on the initial conditions fixed by the physical state of the Universe at the ``galaxy formation epoch''.  The only subsequent significant evolution in galaxies to be discussed was the slow changes in their stellar populations.  Dynamical processes could be assumed to be settled, and therefore ignored.  Since then major observational and theoretical progresses of direct relevance occurred, which resulted in a gradual shift in the meaning of galaxy evolution.  It suffices to remind that basically all the advances in non-linear dynamics, including the recognition of the fundamental r\\^ole of chaos, and in computer simulations of galaxies have been made in between.  Today, in all its steps the ELS scenario is no longer tenable, as explained in the following.  This is still not perceived as a necessity in fields loosely connected with the recent advances in gravitational dynamics.  An instance of inadequation between the ELS scenario and more recent works appears in cosmological simulations at several $100\\,\\mathrm{Mpc}$ scale (see e.g., White 1994).  In such simulations nothing like homogeneous collapses do occur, instead hierarchical clustering proceeds at all the computable scales with different speeds.  This implies that galaxies and galaxy clusters, exactly as stars, form in different regions of the sky {\\it asynchronously}.  The formation process covers several dex of time-scales, so not only the free-fall time, but merging and later infall participate to it.  For many galaxies the formation/evolution should be considered as not terminated even now. The galaxy age looses its original meaning because the aging, traced by various observables, may occur at widely different speeds.  A central aspect not considered in the ELS picture is the likely coupling of dynamics with stellar activity.  In the recent years the large FIR emission of spirals, which even largely dominates the light in starburst galaxies, was found substantial particularly over the ``late'' part of the spiral sequence.  The FIR emission, coming mostly from UV and visible light absorbed and recycled in the FIR by dust, is consistent with the also recent recognition of the partial opacity of the optical region of spirals (e.g. Davis et al.~1993).  It turns out that the total bolometric luminosity is comparable to the power that dynamics can exchange (Pfenniger 1991b).  This coincidence is best explained by a coupling of star formation and dynamics via a feed-back mechanism (Quirk 1972; Kennicutt 1989).  The interesting aspect of this coupling is that the systematic global properties of galaxies are then no longer necessarily determined by the initial conditions of collapse.  As for stars, the galaxy properties would then be encoded in their internal small scale physics, i.e.~star formation and ISM physics.  This may solve the old problem of the absence of galactic scale at the radiation-matter decoupling epoch, particularly if galaxy formation covers a sizable fraction of the Hubble time.  The galactic scale would result mainly from the balancing of stellar activity effects, particularly during starburst phases, and dynamics.  The Hubble sequence is most probably an incomplete sample.  For example many low surface brightness galaxies may well be missed (e.g.~Bothun et al.~1990).  However, whenever fast morphological changes do occur in normal galaxies (bars, mergers), they must correspond to shifts {\\it within the sequence}, because the already existing stars certainly survive the changes.  The general properties of the Hubble sequence have been progressively determined (e.g.~de Vaucouleurs 1959; Broeils 1992; Roberts \\& Haynes 1994; Zaritsky et al.~1994).  In order to extract useful information of this ``zoo'', we must consider only the most general properties, keeping in mind that galaxies form a variety of objects with different ages and aging speeds.  For example mergers (Schweizer 1993) certainly accelerate morphological changes at speeds which depend on the fortuitous interaction strengths.  The following list summarises the main trends and properties along the spiral sequence {\\bf from Sm to Sa}, that are useful for the discussion.  \\def\\Up{$\\nearrow$} \\def\\Down{$\\searrow$}  \\begin{tabular}{lll} \\hline\\noalign{\\smallskip} Total mass&\\Up&$M_{\\mathrm{tot}} \\approx 1\\to100 \\cdot 10^{10}\\,\\mathrm{M_{\\odot}}$\\\\ Kinetic energy&\\Up&$\\frac{1}{2} V_{\\mathrm{max}}^2 \\approx 25 \\to 500 \\,\\mathrm{eV/nucleon}$\\\\ Bulge-disk ratio&\\Up&$ L_{\\mathrm{sph}}/L_{\\mathrm{tot}}\\approx 0\\to0.6$\\\\ Symmetry&\\Up&\\\\ Metallicity&\\Up&$12+\\log(\\mathrm{ O/H}) \\approx 8.3\\to9 $ \\\\ Detected gas&\\Down&$\\displaystyle\\frac{M_{\\mathrm{HI+H_2}}}{M_{\\mathrm{tot}}}\\approx 0.10\\to0.07$, $\\displaystyle\\frac{M_{\\mathrm{HI+H_2}}}{M_{\\mathrm{stars}}}\\approx1.4\\to 0.1$ \\\\ Dark matter&\\Down&$M_{\\mathrm{dark}}/M_{\\mathrm{lum}} \\approx 10 \\to 1$ \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\vskip-1mm plus 1mm minus 1mm} \\end{tabular} ",
+        "conclusions": "Galactic dynamics forces us to see spirals as structures with possible ``bursts'' of dynamical evolution involving short time-scales.  Taking into account today's observational and theoretical constraints, the only possible sense of evolution is from Sm to S0.  During evolution everything happens as if dark matter in galaxies is transformed into stars, that is dark gas is required.  A possible solution to the dark matter problem in galaxies is that gas in outer HI disks clumps along a fractal structure down to solar system sizes, as explained elsewhere in this volume. The smallest clumps are then very dense and cold which makes them presently hard to detect.  A sizable amount of dark (or dust obscured) stars can also be argued just from the fact that evolved galaxies (Sa-S0's) still contain about 40\\% of dark matter."
+    },
+    "9602/gr-qc9602056_arXiv.txt": {
+        "abstract": "Some recently discovered nonperturbative strong-field effects in tensor--scalar theories of gravitation are interpreted as a scalar analog of ferromagnetism: ``spontaneous scalarization''. This phenomenon leads to very significant deviations from general relativity in conditions involving strong gravitational fields, notably binary-pulsar experiments. Contrary to solar-system experiments, these deviations do not necessarily vanish when the weak-field scalar coupling tends to zero. We compute the scalar ``form factors'' measuring these deviations, and notably a parameter entering the pulsar timing observable $\\gamma$ through scalar-field-induced variations of the inertia moment of the pulsar. An exploratory investigation of the confrontation between tensor--scalar theories and binary-pulsar experiments shows that nonperturbative scalar field effects are already very tightly constrained by published data on three binary-pulsar systems. We contrast the probing power of pulsar experiments with that of solar-system ones by plotting the regions they exclude in a generic two-dimensional plane of tensor--scalar theories. ",
+        "introduction": "Einstein's general relativity theory postulates that gravity is mediated only by a long-range tensor field. It has been repeatedly pointed out over the years (starting with Kaluza \\cite{K21}) that unified theories naturally give rise to long-range scalar fields coupled to matter with gravitational strength. This led many authors, notably Jordan \\cite{J+}, Fierz \\cite{F56}, and Brans and Dicke \\cite{BD61}, to study, as most natural alternatives to general relativity, tensor--scalar theories in which gravity is mediated in part by a long-range scalar field. The motivation for such theories has been recently revived by string theory which contains massless scalars in its gravitational sector (notably the model-independent dilaton).  We shall consider tensor--scalar gravitation theories containing only one scalar field, assumed to couple to the trace of the energy-momentum tensor. The simplest example of such a theory is a scalar field only coupled to the gravitational sector through a nonminimal coupling $\\xi R \\Phi^2$ (see Section \\ref{SEC6} below). For a study of the observable consequences of general tensor--scalar theories (containing one or several scalar fields), see Ref. \\cite{DEF1}.  Actually, one generically expects scalar fields not to couple exactly to the mass but to exhibit some ``composition dependence'' in their couplings to matter. However, a recent study of a large class of viable string-inspired tensor--scalar models \\cite{DP94} has found that the composition-dependent effects represent only a very small fraction ($\\sim 10^{-5}$) of the effective coupling to matter. Such fractionally small composition-dependent effects would be negligible in the gravitational physics of neutron stars that we consider here.  The most general theory describing a mass-coupled long-range scalar contains one arbitrary ``coupling function'' $A(\\varphi)$ \\cite{F56}. The action defining the theory reads \\begin{eqnarray} S & = & {c^4\\over 16 \\pi G_*}\\int {d^4 x\\over c} g_*^{1/2}\\left( R_* - 2 g^{\\mu\\nu}_* \\partial_\\mu\\varphi \\partial_\\nu\\varphi\\right) \\nonumber \\\\ && +S_m[\\psi_m; A^2(\\varphi)g^*_{\\mu\\nu}]\\ . \\label{eq1.1}\\end{eqnarray} Here, $G_*$ denotes a bare gravitational coupling constant, $R_*\\equiv g_*^{\\mu\\nu} R^*_{\\mu\\nu}$ the curvature scalar of the ``Einstein metric'' $g^*_{\\mu\\nu}$ describing the pure spin-2 excitations, and $\\varphi$ our long-range scalar field describing spin-0 excitations. [We use the signature $\\scriptstyle -+++$ and the notation $g_* \\equiv -\\det g^*_{\\mu\\nu}$.] The last term in Eq.~(\\ref{eq1.1}) denotes the action of matter, which is a functional of some matter variables (collectively denoted by $\\psi_m$) and of the ``physical metric'' $\\widetilde g_{\\mu\\nu}\\equiv A^2(\\varphi)g^*_{\\mu\\nu}$. Laboratory clocks and rods measure the metric $\\widetilde g_{\\mu\\nu}$ which, in the model considered here, is universally coupled to matter. The reader will find in Eqs.~(\\ref{eq6.1})--(\\ref{eq6.7}) below an explicit example (nonminimally coupled scalar field) of how an action of the type (\\ref{eq1.1}), involving two conformally related metrics $g^*_{\\mu\\nu}$ and $\\widetilde g_{\\mu\\nu} = A^2(\\varphi) g^*_{\\mu\\nu}$, can naturally arise.  The field equations of the theory are most simply formulated in terms of the pure-spin variables $(g^*_{\\mu\\nu},\\varphi)$. Varying the action (\\ref{eq1.1}) yields \\begin{mathletters} \\label{eq1.2} \\begin{eqnarray} R^*_{\\mu\\nu} & = & 2\\partial_\\mu\\varphi \\partial_\\nu\\varphi + {8\\pi G_*\\over c^4}\\left(T^*_{\\mu\\nu}-{1\\over 2}T^*g^*_{\\mu\\nu}\\right)\\ , \\label{eq1.2a} \\\\ \\Box_{g*}\\varphi & = & -{4\\pi G_*\\over c^4}\\alpha(\\varphi)T_*\\ , \\label{eq1.2b}\\end{eqnarray} \\end{mathletters} with $T_*^{\\mu\\nu}\\equiv 2 c\\, g_*^{-1/2} \\delta S_m/\\delta g_{\\mu\\nu}^*$ denoting the material stress-energy tensor in ``Einstein units'', and $\\alpha(\\varphi)$ the logarithmic derivative of $A(\\varphi)$~: \\begin{equation} \\alpha(\\varphi) \\equiv {\\partial\\ln A(\\varphi)\\over \\partial\\varphi}\\ . \\label{eq1.3}\\end{equation} [All tensorial operations in Eqs.~(\\ref{eq1.2}) are performed by using the Einstein metric $g_{\\mu\\nu}^*$, {\\it e.g.} $\\Box_{g*} \\equiv g_*^{\\mu\\nu} \\nabla_\\mu^* \\nabla_\\nu^*$, $T_*\\equiv g^*_{\\mu\\nu} T_*^{\\mu\\nu}$.] As is clear from Eq.~(\\ref{eq1.2b}), the quantity $\\alpha(\\varphi)$ plays the role of measuring the (field-dependent) {\\it coupling strength\\/} between the scalar field and matter. It has been shown in Refs. \\cite{DEF1,DEF2PN} that all {\\it weak-field\\/} (``post-Newtonian'') deviations from general relativity (of any post-Newtonian order) can be expressed in terms of the asymptotic value of $\\alpha(\\varphi)$ at spatial infinity and of its successive scalar-field derivatives. Let $\\varphi_0$ denote the asymptotic value of $\\varphi$ at spatial infinity, {\\it i.e.}, the ``vacuum expectation value'' of $\\varphi$ far away from the considered gravitating system. Let us also denote: $\\alpha_0\\equiv\\alpha(\\varphi_0)$, $\\beta_0\\equiv\\partial\\alpha(\\varphi_0)/\\partial\\varphi_0$, $\\beta'_0\\equiv\\partial\\beta(\\varphi_0)/\\partial\\varphi_0$. At the first post-Newtonian approximation, deviations from general relativity are proportional to the Eddington parameters \\begin{mathletters} \\label{eq1.4} \\begin{eqnarray} \\overline\\gamma & \\equiv & \\gamma_{\\rm Edd}-1 = -2 \\alpha_0^2/(1+\\alpha_0^2) \\ , \\label{eq1.4a} \\\\ \\overline\\beta & \\equiv & \\beta_{\\rm Edd} - 1 = {1\\over 2}\\beta_0\\alpha_0^2/(1+\\alpha_0^2)^2 \\ , \\label{eq1.4b}\\end{eqnarray} \\end{mathletters} while at the second post-Newtonian approximation there enters, beyond $\\overline\\gamma$ and $\\overline\\beta$, two new parameters \\cite{DEF1,DEF2PN} \\begin{mathletters} \\label{eq1.5} \\begin{eqnarray} \\varepsilon & = & \\beta'_0\\alpha_0^3/(1+\\alpha_0^2)^3 \\ , \\label{eq1.5a} \\\\ \\zeta & = & \\beta_0^2\\alpha_0^2/(1+\\alpha_0^2)^3 \\ . \\label{eq1.5b}\\end{eqnarray} \\end{mathletters} We see explicitly in Eqs.~(\\ref{eq1.4}), (\\ref{eq1.5}) that all deviations from general relativity tend to zero with $\\alpha_0$ at least as fast as $\\alpha_0^2$. This holds true for {\\it weak-field\\/} deviations of arbitrary post-Newtonian order \\cite{DEF2PN}. Therefore, light-deflection or time-delay experiments \\cite{gamma} which set (through Eq.~(\\ref{eq1.4a})) the following limit on the coupling strength of the scalar field, \\begin{equation} \\alpha_0^2 < 10^{-3}\\ , \\label{eq1.6}\\end{equation} tightly constrain the theoretically expectable\\footnote{We assume here the absence of unnaturally large dimensionless numbers appearing in the successive derivatives of $\\alpha(\\varphi)$~: $\\beta_0$, $\\beta'_0$, \\dots} level of deviation from general relativity in all other experiments probing weak gravitational fields. Note that, in many physically motivated models, there are much tighter limits on $\\alpha_0^2$ coming from equivalence principle tests (see {\\it e.g.} \\cite{DVo} which gets $\\alpha_0^2 \\lesssim 10^{-7}$ in string-derived models). These improved limits crucially depend, however, on the detailed structure and magnitude of equivalence-principle-violating effects (and disappear in the sub-class of metrically coupled theories). To stay model-independent, we shall use the post-Newtonian-derived limit (\\ref{eq1.6}) as our standard weak-field limit. As we shall see later, the importance of the nonperturbative effects discussed here is not uniformly decreased when $\\alpha_0$ takes smaller values, but can level off or even be amplified.  In a previous work \\cite{DEF3}, we have shown that experiments involving the {\\it strong gravitational fields\\/} of neutron stars can exhibit a remarkably different behavior from weak-field solar-system experiments. We proved that when a certain mild inequality restricting the curvature of the coupling function $\\ln A(\\varphi)$ was satisfied, namely \\begin{equation} \\beta_0 \\equiv {\\partial^2 \\ln A(\\varphi_0)\\over \\partial\\varphi_0^2} \\lesssim -4\\ , \\label{eq1.7}\\end{equation} nonperturbative strong-gravitational-field effects developed in neutron stars and induced order-of-unity deviations from general relativity, even for arbitrary small values of the linear coupling strength $\\alpha_0^2$. The aim of the present paper is to further study these nonperturbative phenomena and to prepare the ground for a systematic application to binary-pulsar experiments \\cite{DEFT} by computing the observational effects depending upon the inertia moments of neutron stars. One of the main results of the present study will be to show explicitly that binary-pulsar experiments are, in some regions of theory space, much more constraining than solar-system experiments. This will be illustrated in an exclusion plot discussed below.  The organization of this paper is as follows. In Section II, we show how the non-perturbative scalar-field effects discovered in \\cite{DEF3} can be interpreted as a ``spontaneous scalarization'' of neutron stars, analogous to the spontaneous magnetization of ferromagnets. We write in Section III the field equations that must be numerically integrated to study these non-perturbative effects in slowly rotating neutron stars. Section IV discusses the ``gravitational form factors'' governing the physics of neutron stars in tensor--scalar gravity, notably a parameter linked to the variation of a pulsar's inertia moment caused by the presence of an orbiting companion. The constraints imposed by three binary-pulsar experiments on a generic class of tensor--scalar models are then derived in Section V. Finally, the conclusions of our study are given in Section VI. ",
+        "conclusions": "} Before summarizing the main results of the present work, we wish to indicate briefly the frameworks within which our findings might be physically relevant. First, let us note that (barring any unnatural fine tuning) they are not relevant in the case of a strictly massless scalar field $\\varphi$, or at least when its mass $m_\\varphi \\ll \\hbar H_0/c^2 \\sim h_{100} \\times 2.13\\times 10^{-33}\\ {\\rm eV}$, where $H_0 = h_{100}\\times 100\\ {\\rm km}.{\\rm s}^{-1}/{\\rm Mpc}$ denotes Hubble's constant. Indeed, in such a case it was found, both in traditional equivalence-principle-respecting tensor--scalar theories \\cite{DN93} and in string-inspired models with a massless dilaton or modulus \\cite{DP94}, that the cosmological evolution naturally drives the vacuum expectation value of $\\varphi$ toward {\\it minima\\/} of $\\ln A(\\varphi)$. As the latter vacuum expectation value coincides (modulo small {\\it fractional\\/} corrections due to spatial fluctuations of the gravitational potential \\cite{DEF2PN}) with our externally-imposed $\\varphi_0$, a natural prediction of these massless models is that $\\alpha_0= \\partial\\ln A(\\varphi_0)/\\partial\\varphi_0$ is small and that $\\beta_0 = \\partial^2\\ln A(\\varphi_0)/\\partial\\varphi^2_0$ is {\\it positive}. The former result concerning $\\alpha_0$ is in agreement with observational data, but the latter one concerning $\\beta_0$ gives the wrong sign for spontaneous scalarization effects\\footnote{Note, however, that a coupling function of the type $\\ln A(\\varphi)= +\\epsilon^2\\varphi^2 - \\lambda^2\\varphi^4$, with a sufficiently small $\\epsilon$ and a sufficiently large $\\lambda$, would reconcile a (very localized) minimum at $\\varphi=0$ with a mainly concave coupling function leading to nonperturbative effects. We do not wish to consider here such fine-tuned cases.}.  On the other hand, our results are relevant to the wide class of models comprising scalar fields of range\\footnote{The lower bound comes from the requirement that $\\varphi$ be effectively massless on the scale of a typical binary pulsar. If it is violated one must correct our formulas by Yukawa exponential factors.} $10^6\\ {\\rm km}\\lesssim \\lambda = \\hbar/m_\\varphi c \\lesssim c H_0^{-1}$. Assuming that the endemic ``Polonyi problem'' (too much energy stored in the cosmological oscillations of $\\varphi_0(t)$) \\cite{Polonyi1} of such models is somehow solved \\cite{Polonyi2} or fine-tuned to provide $\\Omega = \\rho_{\\rm tot}/\\rho_{\\it cr}=1$ \\cite{Frieman}, the condition for these models to be naturally compatible with solar-system constraints is simply that the location $\\varphi_0$ of the minimum of the $\\varphi$ potential, $V(\\varphi) = (c^4/8\\pi G_*) m_\\varphi^2 (\\varphi-\\varphi_0)^2$, coincide (or nearly coincide) with an {\\it extremum\\/} of the coupling function $A(\\varphi)$. Such a coincidence would, for instance, naturally follow from a discrete symmetry about $\\varphi_0$, say a reflection symmetry $(\\varphi-\\varphi_0) \\rightarrow -(\\varphi-\\varphi_0)$ \\cite{DP94}. Under such a condition, the present value of the weak-field scalar coupling $\\alpha_0 = \\partial\\ln A(\\varphi_0)/\\partial\\varphi_0$ would be naturally extremely small, and the sign of the curvature $\\beta_0= \\partial^2\\ln A(\\varphi_0)/\\partial\\varphi_0^2$ could be expected to be negative with {\\it a priori\\/} 50~\\% probability.  As a simple example of such models, we can consider a finite-range scalar field coupled only to the gravitational sector through a multiplicative coupling to the scalar curvature, say \\begin{eqnarray} S &=& {c^4\\over 16\\pi G_*} \\int {d^4 x\\over c}\\, \\widetilde g^{1/2}\\Bigl( -Z(\\Phi)\\widetilde g^{\\mu\\nu} \\partial_\\mu\\Phi \\partial_\\nu\\Phi - U(\\Phi) \\nonumber\\\\ &&+ F(\\Phi)\\widetilde R\\Bigr) + S_m\\left[\\psi_m;\\widetilde g_{\\mu\\nu}\\right]\\ , \\label{eq6.1}\\end{eqnarray} where we have introduced the possibility of an arbitrary field-dependent normalization of the kinetic term of $\\Phi$. One transforms (\\ref{eq6.1}) into our canonical form (\\ref{eq1.1}) [completed by a $\\varphi$-potential term $V(\\varphi) = (c^4/16\\pi G_*)F^{-2}(\\Phi) U(\\Phi)$] by defining \\begin{mathletters} \\label{eq6.2} \\begin{eqnarray} g^*_{\\mu\\nu} & = & F(\\Phi)\\widetilde g_{\\mu\\nu}\\ , \\label{eq6.2a} \\\\ \\varphi & = & \\int d\\Phi\\left[ {3\\over 4}\\, {F'^2(\\Phi)\\over F^2(\\Phi)} + {1\\over 2}\\, {Z(\\Phi)\\over F(\\Phi)} \\right]^{1/2}\\ . \\label{eq6.2b}\\end{eqnarray} \\end{mathletters} This corresponds to a coupling function \\begin{equation} A(\\varphi) = F^{-1/2}(\\Phi)\\ . \\label{eq6.3} \\end{equation} The simplest example of such models is a massive scalar having a nonminimal dimensionless coupling to curvature: \\begin{eqnarray} S &=& {c^4\\over 16\\pi G_*} \\int {d^4 x\\over c}\\, \\widetilde g^{1/2}\\bigl( -\\widetilde g^{\\mu\\nu} \\partial_\\mu\\Phi \\partial_\\nu\\Phi - m_\\Phi^2 \\Phi^2 \\nonumber\\\\ && + \\xi\\widetilde R\\Phi^2 + \\widetilde R\\, \\bigr) + S_m\\left[\\psi_m;\\widetilde g_{\\mu\\nu}\\right]\\ . \\label{eq6.4}\\end{eqnarray} This corresponds to $Z(\\Phi) = 1$, $U(\\Phi) = m_\\Phi^2 \\Phi^2$, and $F(\\Phi) = 1+\\xi\\Phi^2$. In that case, one can integrate Eq.~(\\ref{eq6.2b}) explicitly. Assuming for instance that $\\xi(1+6\\xi) > 0$ and introducing the notation $\\chi \\equiv \\sqrt{\\xi(1+6\\xi)}$, one gets \\begin{eqnarray} && 2\\sqrt{2}\\,(\\varphi-\\varphi_0) = {\\chi\\over\\xi}\\, \\ln\\left[ 1 + 2\\chi\\Phi \\left( \\sqrt{1+\\chi^2\\Phi^2} +\\chi\\Phi\\right) \\right] \\nonumber \\\\ &&+{\\sqrt{6}}\\, \\ln\\left[ 1 - 2\\sqrt{6}\\,\\xi\\Phi\\, {\\sqrt{1+\\chi^2\\Phi^2} -\\sqrt{6}\\,\\xi\\Phi \\over 1+\\xi\\Phi^2} \\right]\\ . \\label{eq6.5}\\end{eqnarray} Note that $(\\varphi-\\varphi_0) = \\Phi/\\sqrt{2} + O(\\Phi^3)$ when $\\Phi\\rightarrow 0$. From Eqs.~(\\ref{eq6.2})--(\\ref{eq6.3}), one gets easily the following parametric representations \\begin{mathletters} \\label{eq6.6} \\begin{eqnarray} A(\\varphi) &=& \\left( 1+ \\xi\\Phi^2 \\right)^{-1/2}\\ , \\label{eq6.6a}\\\\ \\alpha(\\varphi) &=& {\\partial \\ln A(\\varphi)\\over \\partial \\varphi} \\nonumber\\\\ &=& -\\sqrt{2}\\,\\xi\\Phi\\left[ 1 + \\xi(1+6\\xi)\\Phi^2 \\right]^{-1/2}\\ , \\label{eq6.6b}\\\\ \\beta(\\varphi) &=& {\\partial^2 \\ln A(\\varphi)\\over \\partial \\varphi^2} = -2\\xi\\left( 1+\\xi\\Phi^2 \\right) \\nonumber\\\\ && \\times \\left[ 1+\\xi(1+6\\xi)\\Phi^2 \\right]^{-2}\\ . \\label{eq6.6c} \\end{eqnarray} \\end{mathletters} In particular the value of the curvature parameter around $\\varphi_0$ reads \\begin{equation} \\beta_0 = \\beta(\\varphi_0) = \\left.{\\partial^2 \\ln A(\\varphi)\\over \\partial\\varphi^2}\\right|_{\\Phi=0}= -2\\xi\\ . \\label{eq6.7}\\end{equation} Therefore, positive values of $\\xi$ [which, in the formulation (\\ref{eq6.4}), seem preferred because unable to cause a change of sign of the coefficient $(1+\\xi\\Phi^2)$ of the kinetic term for $\\widetilde g_{\\mu\\nu}$] correspond to the ``interesting'' case where spontaneous scalarization effects can occur\\footnote{In our conventions, the so-called ``conformal coupling'' corresponds to $\\xi = -{1\\over 6}$, and to a coupling function $A(\\varphi) = \\cosh(\\varphi/\\sqrt{3})$. Note, however, that only the action $S_{\\rm conf}\\equiv \\int d^4 x\\, \\widetilde g^{\\,1/2}[-\\widetilde g^{\\mu\\nu} \\partial_\\mu\\Phi \\partial_\\nu\\Phi-{1\\over 6}\\widetilde R\\Phi^2]$ (without bare Einstein term) exhibits a conformal invariance\\cite{Deser}. The action $S_{\\rm conf}$ involves a spin-2 but no spin-0 excitation.}.  Let us now summarize the main results of the present work:  $\\bullet$~We have clarified the physical meaning of the nonperturbative strong-gravitational-field effects discovered in \\cite{DEF3} by interpreting them, by analogy with ferromagnetism, as a phenomenon of ``spontaneous scalarization''. Negative values of the curvature parameter $\\beta_0$ (like negative values of the coupling between magnetic moments $H_{ij} = g \\mbox{\\boldmath$\\mu$}_i \\mbox{\\boldmath$\\mu$}_j$) favor the spontaneous creation of a scalar field when considered in the context of gravitationally compact objects (neutron stars). The critical baryonic mass for the spontaneous scalarization transition is $\\lesssim 1.5\\, m_\\odot$ when $\\beta_0\\lesssim -5$. See Table~\\ref{tab1} and Fig.~\\ref{fig2} for precise values. The presence of some externally-imposed scalar field background $\\varphi_0$ with $\\alpha_0 = \\partial\\ln A(\\varphi_0)/\\partial\\varphi_0\\neq 0$ smoothes the scalarization transition (analogously to the presence of an external magnetic field for the ferromagnetic transition).  $\\bullet$~The development (through the previous mechanism) of strong scalar fields in neutron stars leads to very significant deviations from general relativity. These deviations are measured by some gravitational form factors $\\alpha_A,\\beta_A,K_A^B$ which enter the effects observable in binary-pulsar experiments. In our previous work \\cite{DEF3} we focussed on the effective scalar coupling constant $\\alpha_A$. Here, we gave results for $\\beta_A = \\partial\\alpha_A/\\partial\\varphi_0$ (see also \\cite{gef93}), and for $K_A^B = -\\alpha_B \\partial\\ln I_A/\\partial\\varphi_0$ which enters the parametrized post-Keplerian timing parameter $\\gamma$. To compute $K_A^B$, we generalized to a tensor--scalar context the work of Hartle \\cite{H67} on the general relativistic inertia moments $I_A$ of slowly rotating neutron stars. We found that $K_A^B$ could cause very drastic deviations from general relativity in tensor--scalar theories containing no large dimensionless parameters. This is achieved without fine-tuning and in theories having only positive-energy excitations.  $\\bullet$~We presented a preliminary investigation of the confrontation between scalar models exhibiting nonperturbative effects and actual binary-pulsar experiments. We contrasted the probing power of pulsar experiments to that of solar-system ones by working in the two-dimensional $(\\alpha_0 \\equiv \\partial\\ln A(\\varphi_0)/\\partial\\varphi_0 ,\\ \\beta_0 \\equiv \\partial^2\\ln A(\\varphi_0)/\\partial\\varphi_0^2)$-plane describing a generic class of tensor--scalar models. Using published data on PSRs 1913+16, 1534+12 and 0655+64 (and a specific nuclear equation of state), we found that binary-pulsar experiments exclude a large domain of theories compatible with solar-system experiments (see the exclusion plot, Fig.~\\ref{fig9}). In particular, they constrain $\\beta_0$ (independently of $\\alpha_0$) to \\begin{equation} \\beta_0 > -5\\ . \\label{eq6.8}\\end{equation} Interestingly, this bound can be expressed in terms of the well-known weak-field Eddington parameters \\begin{equation} {\\beta_{\\rm Edd}-1\\over \\gamma_{\\rm Edd} -1} < 1.3\\ . \\label{eq6.9}\\end{equation} The singular ($0/0$) nature of the ratio on the left-hand side vividly expresses the fact that such a conclusion could never be obtained in weak-field experiments (at least until they find a significant deviation from general relativity). It must be kept in mind that the inequality (\\ref{eq6.9}) is one sided only, and that $\\beta_{\\rm Edd} -1$ and $\\gamma_{\\rm Edd}-1$ must be taken with their signs.  Let us note that, on the whole, the fact that pulsar data tend to exclude sufficiently negative values of $\\beta_0$ is nicely compatible with the expectation from cosmological attractor scenarios \\cite{DN93,DP94} that $\\varphi_0$ be dynamically driven toward a {\\it minimum\\/} of $\\ln A(\\varphi)$. Though our results have been derived by assuming a particularly simple coupling function between the scalar field and matter (``quadratic model'': $\\ln A(\\varphi) = {1\\over 2} \\beta_0 \\varphi^2$), we think they hold in general classes of coupling functions containing no small or large dimensionless parameters. In a future publication \\cite{DEFT}, we shall present a more systematic confrontation between tensor--scalar theories and binary-pulsar experiments.  Before ending this paper, we would like to stress some of the limitations of our treatment that we intend to overcome in future work: (i)~We considered here only one specific equation of state, modeled as a simple polytrope; a more complete study should consider a selection of realistic equations of state. (ii)~In the present paper, we did not consider the effect of a cosmological variation of $\\varphi_0$. We are aware that a nonzero $\\dot \\varphi_0$ of order the Hubble parameter could significantly modify the interpretation of some of the pulsar tests discussed above. However, as one disposes of several independent pulsar tests, some of which do not involve $\\dot \\varphi_0$-sensitive observables (such as the $(\\dot\\omega\\mbox{-}\\gamma\\mbox{-}s)_{1534+12}$ test considered above and several ``Stark'' tests \\cite{DS91,A95,W95}), we think that a combined discussion of all pulsar tests will lead to limits on $\\alpha_0$ and $\\beta_0$ similar to the ones we have obtained here, and, in addition, to significant limits on $\\dot\\varphi_0$.  Finally, let us note that the strong scalar field effects discussed in \\cite{DEF3} and the present paper could have a very significant impact on several aspects of the theory of gravitational radiation from compact objects (besides the ones taken into account in Section 6.2 of \\cite{DEF1} and Eq.~(\\ref{eq5.4}) above). Of particular interest would be: (i)~the opening of a new, significant energy-loss channel in (spherically symmetric!) gravitational collapse and neutron-star binary coalescences; (ii)~important modifications in the conditions for the onset of radiative instabilities in fast-rotating neutron stars (Chandrasekhar-Friedman-Schutz instability) \\cite{CFS}. Both issues are particularly worth of further study."
+    },
+    "9602/astro-ph9602011_arXiv.txt": {
+        "abstract": "We report on observations of the Sunyaev-Zeldovich effect towards X-ray ROSAT clusters taken with a double channel (1.2 and 2 {\\it mm}) photometer installed at the focus of the 15m SEST antenna in Chile. This paper describes the first results obtained for the high-z clusters S1077, A2744 and S295. Marginal detections were found for A2744 and at 1 {\\it mm} for S1077. We discuss these data in terms of contamination of sources along the line of sight and give a constraint on the amplitude of the kinematic effect. ",
+        "introduction": "The Sunyaev-Zeldovich (SZ) effect is a characteristic signature on the Cosmic Microwave Background (CMB) radiation spectrum originated from (inverse) Compton scattering of the photons by hot electrons in clusters of galaxies. There has been considerable interest in its detection because of its potential as a diagnostic tool for observational cosmology ({\\it physical process in the earlier universe; determination of H$_0$ and q$_0$ and peculiar velocities of clusters}) and for the study of the intracluster medium (see the original papers by \\cite{SZ69}, \\cite{SZ72}, \\cite{SZ81}). Its detection has been the goal of many observational searches (e.g.  \\cite{B90}), mostly in the Rayleigh - Jeans (R-J) part of the spectrum, where the scattering leads to an intensity decrement. However, interpretation of measurements in the R-J side is plagued both by the possible radio emission from sources within the clusters and by the smallness of the effect. After many attempts to detect this effect (see e.g., the review paper by Birkinshaw 1990) radio observations show the expected decrement at centimeter wavelengths towards A2218, A665, 0016+16, A773 and Coma (\\cite{B91}, \\cite{KL91}, \\cite{J93}, \\cite{GR93}, \\cite{HLR}) and at 2.2{\\it mm} towards A2163 (\\cite{WIL94}).  \\noindent Measurements near the planckian peak and on the Wien side are in principle  more attractive since: (a) the intensity enhancement relative to the planckian value is larger than the magnitude of the R-J decrement; (b) the simultaneous detection of the enhancement (positive) and decrement (negative) on the same cluster provides an unambiguous signature of their presence; (c) sources in the cluster are expected to give a negligible contribution at high frequency; (d) because of the large bandwidth the sensitivity of bolometer systems is excellent.  \\noindent We have built a photometer with two bolometers centered at 1.2 and 2 {\\it mm} to feed the O.A.S.I. (Osservatorio Antartico Submillimetrico Infrarosso) telescope installed at the Italian base in Antarctica (\\cite{DAL92}). This photometer was adapted to the focus of the S.E.S.T. and its performances were tested during an observing run in September 1994. Most of the allocated time was spent in installing the photometer and testing the equipment (\\cite{PIZ95a}) but some useful observations were gathered towards S1077, A2744 and S295. ",
+        "conclusions": "The size of the $y$ value of A2744 derived in table 1 disagrees with the estimated one (see \\S 3) and cannot be reliably ascribed to the sole S-Z thermal effect unless some relevant information is missed in the X-ray data. The resulting data towards A2744 at 1.2 and 2 {\\it mm} and S1077 at 1.2 {\\it mm} could be therefore originated from other sources. Here possible contaminations are briefly discussed.  \\subsection{Contamination of kinematic effect}  Peculiar velocities of clusters could contribute significantly to the signals for values larger than $v_r > $ 1000 km/s. In fact the ratio between the S-Z kinematic effect, $ {\\Delta T \\over T} = {v_r\\over c} \\cdot \\tau $ (where $\\tau = \\int \\sigma _T n dl $ is the optical depth for Thomson scattering along the line of sight and $v_r$ is the peculiar velocity), and the thermal effect depends on the cluster peculiar velocity and the gas temperature. The lack of any other observations of peculiar velocities on these clusters makes the quantification of this contribution impossible and on the basis of the present knowledge we cannot exclude that a fraction of the signals is due to it, even if large peculiar velocities seem quite unlikely in most cosmological scenarios. The present data can be used to constrain the combination $\\tau \\cdot {v_r\\over c}$ which turns out to be $< 3 ~ 10^{-4} $ for A2744, S1077 and S295 (3$\\sigma$ values).  \\subsection{Contamination of cluster sources}  As mentioned in \\S 3 a strong contamination from sources in the cluster seems to be unlikely: (a) clusters do not contain many spirals, (b) late-type galaxies at this redshift would have a millimetric emission below than the expected signals from the S-Z effect and (c) we are looking at the very center of the cluster which, as shown from optical images, is devoid of sources. Only the cluster S295 has three ellipticals very close to the center of the X-ray image and part of their emission can therefore falls in our beam. If there was a significant non-thermal emission we would expect to find positive signals in both channels, contrary to our results (see table 1).  \\noindent Let us suppose that one or more sources fall in one of the reference beams. In this case if a mm-emitting source was located in the right beam, we would find a positive signals in both channels. If it was in the left reference beam the signals would be negative in both channels. This seems to be unlikely since we do not find these systematics (see table 1). However contribution from cluster sources deserves further investigations and will be the goal of future projects.  \\noindent It is also unlikely that both signals are contaminated from the emission of unresolved sources: the estimate of their confusion limit in the beam made by Franceschini et al. (1991) gives at 1.2 {\\it mm} $\\Delta I \\sim$ 0.4 mJy, i.e. $\\Delta T$ $\\sim$ 0.06 mK.  \\subsection{Contamination by other spurious signals}  We have considered other effects affecting the signals: (a) the selected clusters have sizes comparable to our beams and beam dilution cannot significantly lower the detection probability (\\cite{RAP}), but fluxes are reduced up to 30$\\div$40 \\% because of the limited chop throw. (b) Diffraction patterns could enter the beams and alter the signals but constant systematics of this type should be removed by our observing strategy (see \\S 3). However, if the antenna did not track precisely the same paths relative to the local environment, because of a loss of synchronization, some spurious signals survive and we cannot exclude at present that this never happened. (c) We checked if sources are located in the OFF position (15min ahead in R.A.): from the IRAS Faint Source Catalogue and 5GHz NRAO survey it turns out that none fall in these areas. The IRAS 100 $\\mu m$ sky maps do not show any emission from galaxy cirrus at a level greater than their zero level.  \\vskip 1cm"
+    },
+    "9602/astro-ph9602132_arXiv.txt": {
+        "abstract": "We have applied a new method to analyze the horizontal branch (HB) morphology in relation to the distribution of stars near the red giant branch (RGB) tip for the globular clusters M22, M5, M68, M107, M72, M92, M3 and 47 Tuc. This new method permits determination of the cluster ages to greater accuracy than conventional isochrone fitting. Using the method in conjunction with our new high-quality photometric data for RGB and HB stars in the first five of these clusters, we discuss the origins of the spread in color on the HB and its relation to the `second parameter' problem. The oldest clusters in our sample are found to have relatively low ages ($13.5 \\pm 2$ Gyr). A 1$\\sigma$ uncertainty in each of the parameters of mass and helium content combined with the effects of helium diffusion gives a lower limit for the age of the oldest clusters of 9.7 Gyr. ",
+        "introduction": " \\begin{itemize} \\item [i)] The HB morphology is well explained by differential variations in the mass loss efficiency along the RGB among the stars. The distribution of the HB mass is gaussian. In the case of $\\eta$ a value of $<\\eta>=0.4$ is found for the set of GCs studied.  \\item [ii)] Since mass loss is the cause of the HB morphology, the properties of the stars at the RGT can be linked with those at the HB. This gives a powerful method to constrain  the mass of the RGB.   \\item [iii)] The mean mass of the HB and the 2$\\sigma$ value of the HB mass distribution have been used to determine the value of $\\eta$ and the mass at the RGB  \\item [iv)] The HB morphology is explained as variations in $\\eta$, but with a central value of $<\\eta>=0.4$ which is the same for all the clusters. The `second parameter' problem is explained as a mass difference among clusters with identical metallicity, and therefore as an age difference among them.  \\item [v)] The age for the oldest GCs was $(13.5 \\pm 2)$ Gyr. A 1$\\sigma$ uncertainty in each of the parameters of mass and helium content combined with the effects of helium diffusion gives a lower limit for the age of the oldest clusters of 9.7 Gyr. \\end{itemize}",
+        "conclusions": "In this paper we have presented clues for two important questions in stellar evolution and cosmology: the spread in colour of the HB and the age of the oldest known stars in the universe. Using very accurate photometry that we have obtained on five globular clusters we were able to distinguish for three of them the RGB from the AGB with no ambiguity.  With these data we studied the possible scenarios to produce the spread in colour on the horizontal branch. The theory of a delayed helium core flash would produce an extra number of stars above the theoretical helium flash point. We have seen that this is not happening. Also, variations in the core mass would produce a vertical spread in the HB that is not observed, except in cases where there is also a spread in the RGB -- due to metallicity variations -- and these metallicity variations could account for the vertical spread. The only scenario left to explain the spread in colour along the HB is that variations in the mass loss along the RGB produce a different ratio of total mass to core mass. As a consequence of this we have concluded that the explanation for the `second parameter' problem relies on age differences among the clusters that present this effect -- different masses at the RGB. Even though a different value for $<\\eta>$ among these clusters could produce the same result, the question would be, why should these clusters have a different value of $<\\eta>$?  Once the nature of the HB has been explained we have used  the morphology of the RGT to constrain the amount of mass that is lost at this stage of stellar evolution. Using this and the morphology of the HB we have been able to put strong constraints on the mass of the RGB stars. We have calculated ages for the GCs in the sample and found that the oldest clusters have an age of 13.5 Gyr. This estimate of the age is in better agreement with current cosmological models, especially an open universe with $H_{0} \\approx 80 $km s$^{-1}$ Mpc$^{-1}$ (Freedman et al 1994, Pierce et al 1994).   Now we are in the position to answer all the questions that we"
+    },
+    "9602/hep-ph9602233_arXiv.txt": {
+        "abstract": "We consider the influence of non-equilibrium electronic neutrinos (and anti-neutrinos) on the neutron-to-proton ratio. These neutrinos would come from massive $\\nu_\\tau$ annihilations $\\bar \\nu_\\tau \\nu_\\tau \\rightarrow \\bar \\nu_e \\nu_e$. For sufficiently large $\\nu_\\tau$ masses this new effect would strongly enhance the (n/p)-ratio, leading to a very stringent bound on the $\\nu_\\tau$ mass, even adopting a rather weak upper bound on the effective number on neutrino species during nucleosynthesis. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602060_arXiv.txt": {
+        "abstract": "The role that minor mergers have played in the formation and structure of the Milky Way is still an open question, about which there is much debate. We use numerical simulations to explore the evolution of debris from a tidally disrupted satellite, with the aim of developing a method that can be used to identify and quantify signatures of accretion in a survey of halo stars. For a Milky Way with a spherical halo, we find that debris from minor mergers can remain aligned along great circles throughout the lifetime of the Galaxy. We exploit this result to develop the method of Great Circle Cell Counts (GC3), which we test by applying it to artificially constructed halo distributions. Our results suggest that if as few as 1\\% of the stars in a halo survey are accreted from the disruption of a single subsystem smaller than the Large Magellanic Cloud, GC3 can recover the great circle associated with this debris. The dispersion in GC3 can also be used to detect the presence of structure characteristic of accretion in distributions containing a much smaller percentage of material accreted from any single satellite. ",
+        "introduction": "Theories describing the formation of the Milky Way can generally be viewed as variations on or combinations of two scenarios. Based on the kinematics of metal-poor halo field stars, Eggen, Lynden-Bell \\& Sandage (1962; ELS) proposed a model in which the Galaxy began with the free-fall collapse of an approximately uniform, spherical primordial fluctuation. The globular clusters and halo field stars were formed during the free-fall phase, with the bulk of the original Galactic material dissipating energy in the gas phase and falling into a rotationally supported disk. Alternatively, to account for the lack of a metallicity gradient in the halo and the suggestion of a large spread in the ages of outer-halo globular clusters, Searle \\& Zinn (1978; SZ) proposed a scenario in which the Galaxy was assembled through the gradual merging of many sub-galactic sized clouds.  Low-mass stars have lifespans that are greater than the age of the Galaxy and do not dissipate orbital energy. Hence, we have an abundance of ``fossil'' information with which to explore various formation scenarios, through halo stars and clusters, especially their abundances and kinematics. Commonly employed tracers include the kinematics of field stars as a function of [Fe/H] (as in ELS), the age distribution of Galactic globular clusters (e.g. Vandenberg, Bolte \\& Stetson, 1990), trends in cluster age with [Fe/H] and Galactocentric radius (as in SZ), the lack of an abundance gradient with radius or height above the disk, and the persistence of a cold, thin disk (e.g. T\\'oth \\& Ostriker, 1992). (See Larson [1990] and Majewski [1993] for comprehensive reviews.)  Here, we consider the feasibility of probing ``fossil'' signatures of the formation of the Galaxy through structure in the phase-space distribution of halo field stars. Our approach is motivated by numerical simulations, which demonstrate that distinctive features such as tidal tails are a generic consequence of interactions between comparable mass galaxies (e.g. Toomre \\& Toomre 1972; Barnes 1988, 1992; Hernquist 1992, 1993; Hibbard \\& Mihos 1995). Similarly, streamers can be produced along the orbit of a satellite galaxy when stars are torn from it by tidal forces from its host (e.g. McGlynn; 1990, Piatek \\& Pryor, 1995; Johnston, Spergel \\& Hernquist, 1995) . If such tidal debris were to maintain spatial and kinematic coherence for the lifetime of the Galaxy, then a halo formed through the disruption of many SZ fragments would exhibit streakiness in its phase-space distribution, unlike one originating from a smooth, monolithic collapse which ought to be mostly featureless.  The possibility that accretion events may leave observable signatures is supported by observations of satellites  of the Milky Way which display tantalizing evidence for ongoing tidal interactions.  As part of their analysis leading to the discovery of the Sagittarius dwarf galaxy (hereafter Sgr), Ibata, Gilmore \\& Irwin (1994) produced an isopleth map of the overdensity of horizontal branch stars in the region where the dwarf was thought to lie, which shows highly elongated contours with axis ratios $\\sim$ 3:1. Grillmair et al.  (1995) have also reported the detection of distortions in the outer regions of globular clusters.  Both sets of observations are consistent with morphological disturbances  produced by tidal perturbations.  Unfortunately, the low surface density of streamers from tidally disturbed objects makes their detection challenging.  Nevertheless, it is plausible that stars or satellites that appear to be associated spatially and/or kinematically are indeed debris from tidal interactions. From their simulations of Sgr, Johnston, Spergel \\& Hernquist (1995) predict that the tidal streamers associated with this object may be detectable as moving groups in the halo, and that if a similar galaxy had been  destroyed by the Milky Way within the last gigayear (Gyr), its remains should still be detectable as a moving group today.  Observational evidence suggests that there is considerable structure in the phase-space distribution of dwarf galaxy companions in the halo. Lynden-Bell (1976, 1982) was the among the first to note that most of the Milky Way's satellite galaxies lie near two great circles passing close to the Galactic Poles. Three dwarf spheroidal galaxies (Draco, Ursa Minor and Carina) are in the vicinity of the great circle defined by the Small and Large Magellanic Clouds and the Magellanic Stream (the ``Magellanic Plane''), while another five (Fornax, Leo I, Leo II, Sculptor and Sextans) are distinctly aligned along the ``Fornax-Leo-Sculptor Stream''. More recently, Lynden-Bell \\& Lynden-Bell (1995) developed a method to systematically search for coincidences of clusters along great circles, and recovered these and several other possible associations of halo objects. Moreover, Lin \\& Richer (1992), Majewski (1994)  and Fusi-Pecci, Bellazzini, Cacciari \\& Ferraro (1995) have shown that several young halo globular clusters are not far from each of these planes, which further supports an interesting physical explanation for these alignments.  Substructure has also been found in the stellar distribution in the halo. Doinidis \\& Beers (1989) calculated the angular correlation function for the 4400 candidate field horizontal branch stars in a sample covering 2300 deg$^2$ and found a distinct excess of pairs with angular separations less than $10'$. Similar clumpiness has been seen in the phase-space distribution of halo stars in the form of moving groups found in kinematic surveys (e.g Sommer-Larsen \\& Christiansen, 1987; Croswell et al, 1991; Arnold \\& Gilmore, 1992; Majewski, Munn \\& Hawley, 1994).  However, it is not yet clear to what extent (if at all) the current observations of non-uniform distributions of halo matter reflect the epoch of halo formation, and little attention has been given to procedures suitable for evaluating the phase-space structure of the halo as a whole. In this paper we begin to examine the nature of debris from tidal interactions employing numerical simulations of the disruption of satellites by the Milky Way. We use our simulations to determine how long substructure from accretion events can persist in the halo. If this timescale is significant (ie. longer than a few Gyrs) we examine the observable properties of the debris. We test various methods for characterizing the substructure on artificial halo distributions, generated from the simulations.  We describe our computational method and explore the behavior of tidal debris in our simulations in \\S 2. We introduce the method of Great Circle Cell Counts, and apply it to artificial halo distributions in \\S 3. Other methods for measuring structure in the halo are discussed in \\S 4. Finally, we summarize possible implications and limitations of our results in \\S 5. ",
+        "conclusions": "From our simulations of satellite accretion, run assuming a Milky Way model having a spherical halo, we find that debris from interactions can remain aligned in tidal streamers, close to the parent satellite's original orbit, for the lifetime of the Galaxy. If the orbit is near planar, the projected path of the streamers may be closely associated with a single great circle. The method of GC3 exploits this characteristic to recover multiple debris trails from artificial halos containing $\\ge 1\\%$ of material accreted from any single satellite. The dispersion in the counts may be used to quantify the significance of such structure in a distribution containing an even smaller percentage of accreted material, and possibly as an indicator for the role that accretion has played in the formation and evolution of the Milky Way.  Some of these conclusions may depend on our assumption of a spherical halo. There is growing evidence for the oblateness of galactic halos in general (eg Sackett \\& Sparke, 1990; Sackett et al, 1994a, 1994b) and our own in particular (e.g. Larson \\& Humphreys, 1994) and non-planar orbits in such a potential may destroy the alignments with great circles seen in our simulations. However, for orbits in moderately flattened potentials, the total angular momentum is approximately constant and the $z$-component of the angular momentum is exactly conserved (e.g. Binney \\& Tremaine, 1987). Hence, the orbit lies near a plane whose pole maintains a constant angle, $\\theta$, with the symmetry axis, while precessing in $\\phi$ about it (see Figure 7). This suggests that we may still detect structure characteristic of tidal debris on our grid of poles of great circles in ($\\cos(\\theta),\\phi$), either directly, or by summing over cells with constant $\\theta$. Debris from satellites on orbits with pericenters somewhat smaller than the ones in our models, which precess due to the axisymmetric potential of the disk, may also be recovered in this manner.  We base our conclusions on ``all sky surveys'' of our artificial halos. In practice, it may be difficult to use GC3 at low galactic latitudes ($|b| < 15\\deg$) because of the predominance of disk stars and inhomogeneity of absorption by the interstellar medium near the Galactic plane. However, GC3 could equally well be applied to a restricted survey, since the area of intersection of any great circle cell with a given region (or regions) of sky can always be calculated. The average and dispersion in the number of stars to fall in that area from a random distribution of stars in the region can be predicted, and the significance of the counts (from equation 5) still assessed. A survey may have to cover a minimum of a few hundred square degrees to detect local structure since we expect tidal debris to span several degrees, while a survey covering a significant fraction of the sky would be required to look at the global structure of the halo.  Finally, our chances of detecting debris could only be improved with the addition of velocity or distance information. Proper motions could be used to reduce the noise in the cell counts, by applying a stricter membership criterion for a cell, defined not only by a star's position, but also by the alignment of its velocity vector along the great circle. Alternatively, radial positions and velocities could be used to refine the cell counts, using our knowledge that these quantities should oscillate in $\\Psi$ along a great circle, with a period $P < 2\\pi$. The signature of a debris trail in a great circle cell could be detected by systematically searching possible periods for excess power in the quantity $A^2+B^2$ where $$ A=\\sum y \\, \\cos(\\Psi/P), \\,\\,\\, B=\\sum y \\, \\sin(\\Psi/P), $$ the sum is performed over all particles in the cell and $y$ is either radial position or velocity. In all three cases, once an overdense cell has been identified, the extra coordinate (proper motion, radial position or velocity) could be used to recover the parent satellite's content and orbit by identifying coherent streams of stars along this great circle.  To summarize: we propose that the method of GC3 might be used to detect signatures of ancient accretion events in either a spherical or oblate halos, from surveys covering more than a few hundred square degrees. The ultimate test of this assertion is the application of GC3 to a real survey of halo stars, such as the APM survey (eg Maddox et al 1990), the APS survey of POSS-I (eg Pennington et al 1993) or the digitized survey of the POSS-II catalogue (eg Weir, Fayyad \\& Djorgovski 1995)."
+    },
+    "9602/astro-ph9602126_arXiv.txt": {
+        "abstract": "The broad absorption lines (BALs) of QSOs indicate abundances of heavy elements, relative to hydrogen, that are 1 to 2 orders of magnitude higher than the solar values.  In at least one QSO, an especially large enhancement of phosphorus is observed.  These abundances resemble those in Galactic novae, and this suggests that novae may produce the BAL gas. The needed rate of nova outbursts may come from single white dwarfs that accrete gas as they pass through a supermassive accretion disk around a central black hole. ",
+        "introduction": "BAL QSOs  show broad absorption troughs attributed to rapidly outflowing gas, of unknown origin, located outside the continuum source and the broad emission-line region. BALs typically occur in \\lalpha,  C IV \\lam 1549, Si IV \\lam 1400, N V \\lam 1240, O VI \\lam 1034, and sometimes Mg II \\lam 2798 and Al III \\lam 1857 (see reviews by Weymann, Turnshek, and Christiansen 1985, ``WTC''; and Turnshek 1988, 1995).  BALs often set in at the systemic velocity of the QSO, but in some objects ``detached'' BALs begin at a velocities up to $\\sim10^4$ \\kmps.  The absorption often extends to velocities $\\sim 3\\times10^4$ \\kmps, with varying degrees of structure and residual intensity through the line profile.  The derived column densities for the absorbing ions  (e.g., C$^{+3}$) are $\\sim10^{16}$~\\cmmitwo. BALs occur in about 10 percent  of radio quiet QSOs (Weymann \\etal~1991) and rarely in radio loud QSOs . Either most radio quiet QSOs have BAL material covering about 10 percent of the sky as seen from the continuum source, or a subset of QSOs have BAL regions with a larger covering factor. The BAL region may have a  disk-like geometry (Turnshek 1995; Goodrich and Miller 1995).  The absorbing material must have a transverse dimension $\\gtrsim 10^{16}$ cm, sufficient to cover the continuum source.  The N V BAL often strongly absorbs the \\lalpha\\ emission line,  implying a radius $\\gtrsim 10^{18}$ cm (Turnshek \\etal~1996).  The absorbing clouds, which occupy only a tiny fraction of the volume, may be accelerated by radiation pressure involving the resonance lines or by ram pressure of a fast moving wind (WTC). Proposed sources for the BAL gas include winds from red giant stars (Scoville and Norman 1995) or from a supermassive accretion disk (Murray \\etal~1995). ",
+        "conclusions": "In summary, the high chemical abundances derived for the BAL region of QSOs resemble the abundances of Galactic novae. The mass of heavy elements in a nova shell is consistent with constraints on the dimensions and location of the BAL absorbing material.  A very massive nuclear star cluster must be postulated in order for ordinary nova outbursts to explain the observed incidence of BALs in QSO spectra.  Accretion of gas by white dwarfs passing through a supermassive accretion disk can provide sufficient nova outbursts with a more modest star cluster.  Observational consequences include a small BAL region that could show changes over a few years.  Different velocity components in the BALs of a given object could result from different novae and show different abundances.  Odd numbered elements such as Al and P should be especially enhanced in abundance."
+    },
+    "9602/astro-ph9602041_arXiv.txt": {
+        "abstract": "The luminosity functions of E/S0 galaxies are constructed in 3 different redshift bins ($0.2 < z < 0.55$, $0.55 < z < 0.8$, $0.8 < z < 1.2$), using the data from the Hubble Space Telescope Medium Deep Survey (HST MDS) and other HST surveys.  These independent luminosity functions show the brightening in the luminosity of E/S0s by about $0.5 \\sim 1$ magnitude at $z \\sim 1$, and no sign of significant number evolution. This is the first direct measurement of the luminosity evolution of E/S0 galaxies, and our results support the hypothesis of a high redshift of formation ($z > 1$) for elliptical galaxies, together with weak evolution of the major merger rate at $z < 1$. ",
+        "introduction": "The images of faint galaxies taken with the Wide Field and Planetary Camera (WFPC2) on the Hubble Space Telescope (HST) have provided invaluable information on the morphology of galaxies when the universe was about half its present age.  With these data, reliable morphological classification is possible down to a magnitude limit of $I < 22 \\sim 23$, and we have identified the morphological nature of the galaxies which dominate the number counts at faint magnitudes (Im et al. 1995a, 1995b; Griffiths et al. 1994; Casertano et al. 1995; Driver, Windhorst, \\& Griffiths 1995; Glazebrook et al. 1995; Abraham et al. 1996). For E/S0 galaxies, Im et al. (1995c) find evidence for mild luminosity evolution and no strong merging activity at $z < 1$, using the observed size and colour distributions. Likewise, the HST observations of E/S0 galaxies in clusters at moderate and high redshifts show little sign of strong luminosity evolution (Dickinson 1995).  Lilly et al. (1995a) and Ellis et al. (1995) have shown luminosity functions of the total or colour-divided galaxy population at different redshifts and tried to constrain the evolution of faint galaxies up to $z~\\sim~1$.  In particular, Lilly et al. (1995a) divided the galaxies into two subsamples according to their colors (red or blue), and then showed that the luminosity of red galaxies evolves mildly, in contrast with the blue galaxies.  This result on the red galaxies is consistent with the HST observations of ellipticals, but because of the lack of morphological classification for most of the galaxies in the Lilly et al. sample, detailed constraints on the luminosity evolution of E/S0 galaxies have not been obtained in their study.  The HST data can provide such morphological information, and we will thus construct luminosity functions of {\\it E/S0 galaxies} at moderate to high redshifts in order to study their evolution. ",
+        "conclusions": "We have constructed the luminosity functions of elliptical galaxies using data from HST MDS and archived HST surveys.  Redshifts of these E/S0s have been determined using $V-I$ colors and sizes to an accuracy of $\\sim 0.1$ up to $z \\sim 1$.  From the constructed luminosity functions, we find luminosity evolution of about $1 \\pm 0.5$ mags at $z \\sim 1$. On the other hand, we exclude strong number evolution of elliptical galaxies, and our data are consistent with a merger rate of $\\lesssim 0.01~Gyr^{-1}$ with very weak major merger rate unless the mass density of the universe is locally low. This latter result supports our earlier findings (Im et al. 1995c) on the basis of the size and colour distributions of E/S0s."
+    },
+    "9602/astro-ph9602088_arXiv.txt": {
+        "abstract": "We compare optical and near-infrared colors of disks and bulges in a diameter-limited sample of inclined, bright, nearby, early-type spirals. Color profiles along wedge apertures at 15\\deg\\ from the major axis and on the minor axis on the side of the galaxy opposite to the dust lane are used to assign nominal colors for the inner disks (at 2 scale length) and for the bulges ($\\sim$ 0.5 $r_{eff}$), respectively. We estimate  that the effects of dust reddening and the cross-talk between the colors of the two components is negligible. We find that color differences (bulge -- disk) are very small: $\\Delta(U-R)=0.126 \\pm 0.165$, $\\Delta(R-K)=0.078 \\pm 0.165$. Disks tend to be bluer by an amount three times smaller than that reported by Bothun \\& Gregg \\markcite{B84} (1990) for S0's. Color variations from galaxy to galaxy are much larger than color differences between disk and bulge in each galaxy. Probably, the underlying old population of disks and bulges is much more similar than the population paradigm would lead us to believe. Implied age differences, assuming identical metallicities, are less than 30\\%. ",
+        "introduction": "How different are the colors of bulges and disks of spiral galaxies? According to the Population paradigm, disks of spirals, containing Population I stars, are bluer and younger than bulges, the latter made up of Population II stars. Also, the widespread belief that bulges are metal rich (eg. Rich \\markcite{R88} 1988) supports the notion that bulges are redder than disks. This notion needs to be reexamined in the light of the following: (1) We find that bulges of early-type spirals are not metal rich; their mean metallicities are rarely above solar (Balcells \\& Peletier \\markcite{BP94} 1994, hereafter BP94). Also the metallicity of our Bulge nowadays is estimated to be just above solar (McWilliam \\& Rich \\markcite{MR94} 1994). (2) The inner parts of disks are often very dusty (Valentijn \\markcite{V90} 1990, Peletier \\etal \\markcite{PVMFKB95} 1995). Therefore, integrated colors (and colors derived from ellipse fitting, see section \\ref{Data}) trace reddening as much as population characteristics. (3) Colors have become useful age diagnostic tools with the realization by Frogel \\markcite{F85} (1985) that the combination of optical and near-infrared (NIR) color indices allows to partially disentangle age and metallicity. (4) Integrated colors include contributions from bright, localized, blue star forming regions. The colors of the underlying disk contain information on the age of the older components of the disk population.  Excluding regions of ongoing star formation and regions with significant reddening from the color determinations was not possible with aperture photometry, but can now be done using two-dimensional digital photometry and modern data processing techniques.  The recent availability of large-format two-dimensional near-infrared detectors allows the determination of NIR color profiles without the limitations of studies based on NIR aperture photometry data. In this paper we analyze $UBRIJK$ surface photometry of a sample of early-to-intermediate type, inclined disk galaxies. This work is part of a larger study of optical and NIR colors of disks and bulges (Peletier \\& Balcells, in preparation). Here we highlight that the colors of bulges and disks are very similar. Indeed, contrary to what would be expected if bulges are indeed redder than disks, bulges do not appear as distinct morphological components in color index maps derived from optical and NIR images. We find long dust lanes to be the main morphological component of the color maps. A few galaxies show additional features such as  nuclear spots of redder color with a size smaller than the bulge, and progressive  bluing in the outer parts of the galaxy. Color variations from galaxy to galaxy are much larger than color variations within each galaxy.  To put this result in a more quantitative footing, we work with color profiles of disk and bulge determined in a way which minimizes the effects of dust reddening. Our results, from accurate near-infrared array measurements, show that color differences between bulges and disks are very small. Age determinations derived from optical and NIR color-color diagrams with the use of simple population models indicate that the bulk of the (inner) disk and bulge stars are essentially coeval.  At most, disk stars are 2-3 Gyr younger than bulge stars. We disagree with the conclusion of Bothun \\& Gregg \\markcite{BG90} (1990) that disks are bluer than bulges by $\\Delta(B-H) \\sim 0.4$. ",
+        "conclusions": "Can we use the color similarity to put limits on the age differences between bulges and disks? In Fig.~\\ref{TwoGrid} we show the color-color data of Fig.~\\ref{URK} together with single age-metallicity stellar population models by Worthey \\markcite{W90} (1994) and Vazdekis \\etal \\markcite{V95} (1995). Along the solid lines, metallicity varies, while age varies along the dotted lines. It is clear that there are large differences between the color predictions of both models. Color and metallicity are partially disentangled in the models of Vazdekis \\etal and degenerate in those of Worthey. Such differences may be due to difficulties inherent to the modeling of $U$ and $K$ data. In the $U$-band, few stars are available to calibrate the models with. In the $K$-band, model predictions are affected by uncertainties in the treatment of the later stages of stellar evolution, such as the AGB (see Charlot \\etal 1996 for a discussion). Vazdekis \\etal \\markcite{V95} (1995) are not the first to claim that age and metallicity are separable using this diagram: this notion already appears in the models of Bothun \\etal \\markcite{B84} (1984), Peletier \\etal \\markcite{PVJ90} (1990), and Bressan \\etal \\markcite{BCF94} (1994).  \\begin{figure} \\plotone{twogrids.ps} \\label{TwoGrid} \\caption{ Stellar population models for single age, single metallicity stellar population models by Worthey (1994) and Vazdekis \\etal (1995) superimposed on top of the data of Fig.~\\ref{URK}. Drawn lines are lines of constant age, dotted lines of constant metallicity.} \\end{figure}  The agreement among the latter authors prompts us to use the Vazdekis \\etal\\ models to analyze our data. Color differences are so small however, that a separation of age and metallicity will not be attempted. Instead, we use the models to see what age difference is implied by the color differences in the case of constant metallicity, and what metallicity difference is implied by the color differences in the case of constant age. In Table~2 and 3 we give the difference in metallicity while keeping age constant, and the difference in age, while metallicity is left constant. The models used are the two mentioned above at Z=0.02 and age=12 Gyr, plus the models of Rabin \\markcite{R80} (1980) as given by O'Connell \\markcite{O86} (1986). The latter models have been modified slightly using Worthey's values, to convert $U-V$, $B-V$ and $V-K$ to resp. $U-R$, $B-R$ and $R-K$. Table~2 shows again that there are large discrepancies between the 3 models. These show up especially in $R-K$. On the average however, the models agree that one needs a difference of appr. 0.10 in log~Z, or 0.11 in log~t between bulge and disk. At an age of 10 Gyr, this corresponds to a difference of 3 Gyr. It is however much more likely that the difference in color is caused by both metallicity and age. In this case the age difference is smaller. We stress that these results apply to the inner parts of disks only. Photographs of spirals often suggest that the outer parts are bluer, and perhaps the perception many people have about the blueness of disks really applies to the regions at large radii.  \\begin{table} \\begin{center} \\caption{Color - metallicity/age conversions} \\begin{tabular}{ccccc} \\\\ \\\\ \\hline Parameter & & Vazdekis & Worthey & Rabin \\\\ \\hline $\\Delta(B-R)$/$\\Delta(log~Z)$ & = & 0.37 & 0.48 & 0.54  \\\\ $\\Delta(U-R)$/$\\Delta(log~Z)$ & = & 0.79 & 1.28 & 1.46  \\\\ $\\Delta(R-K)$/$\\Delta(log~Z)$ & = & 0.62 & 1.51 & 0.61  \\\\ $\\Delta(J-K)$/$\\Delta(log~Z)$ & = & 0.17 & 0.29 & --    \\\\ $\\Delta(B-R)$/$\\Delta(log~t)$ & = & 0.47 & 0.34 & 0.65 \\\\ $\\Delta(U-R)$/$\\Delta(log~t)$ & = & 1.09 & 0.79 & 1.29 \\\\ $\\Delta(R-K)$/$\\Delta(log~t)$ & = & 0.04 & 0.72 & 0.42 \\\\ $\\Delta(J-K)$/$\\Delta(log~t)$ & = & -0.06 & 0.16 & -- \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table} \\begin{table} \\begin{center} \\caption{Age and metallicity differences between bulges and disks} \\begin{tabular}{cccccccc} \\\\ \\\\ \\hline Parameter & Color & \\multicolumn{2}{c}{Vazdekis} & \\multicolumn{2}{c}{Worthey} & \\multicolumn{2}{c}{Rabin} \\\\ & & Average & Spread & Average & Spread & Average & Spread \\\\ \\hline $\\Delta(log~Z)$ & B -- R & --0.12 & 0.26 & --0.09 & 0.20 & --0.08 & 0.18 \\\\ & U -- R & --0.16 & 0.21 & --0.10 & 0.13 & --0.09 & 0.11 \\\\ & R -- K & --0.13 & 0.26 & --0.05 & 0.11 & --0.13 & 0.27 \\\\ & J -- K & --0.10 & 0.52 & --0.05 & 0.30 & -- & -- \\\\ $\\Delta(log~t)$ & B -- R & --0.10 & 0.28 & --0.13 & 0.28 & --0.07 & 0.15 \\\\ & U -- R & --0.12 & 0.15 & --0.16 & 0.21 & --0.10 & 0.13 \\\\ & R -- K & --2.11 & 4.46 & --0.11 & 0.23 & --0.19 & 0.39 \\\\ & J -- K & 0.28 & 1.50   & --0.10 & 0.56 & -- & -- \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   Our result is different from that of Bothun \\& Gregg \\markcite{BG90} (1990), who find that in the $B-H$ vs. $J-K$ diagram disks of S0's are systematically displaced with respect to bulges. For a given $J-K$ the $B-H$ of a disk is on the average 0.4 mag bluer than a bulge. Here, we don't find any systematic offset between bulges and disks, as they see for example in their Fig.~7. To explain the discrepancy we can only point to the fact that their infrared (JHK) measurements are aperture measurements, where the aperture centers are not centered on the galaxy nucleus. Making these measurements, and comparing them to CCD data is extremely complicated, since results depend for example on the response map of the aperture itself, the wings of the aperture, etc. (see e.g. Peletier \\etal \\markcite{PVJ90} 1990). An error of 0.5 mag seems plausible to us. For the S0's in our sample we find that there is as much, or as little, star formation in the bulge as in the disk.     Our data put constraints on the Eggen \\etal\\ \\markcite{ELS62} (1962) model of galaxy formation by monolithic collapse with progressive enrichment of the galaxy. Our ages make it unlikely that there has been a gap in time between the formation of the bulge and that of the inner disk. The formation of bulge and disk must have been a continuous process. This is the case for example in the continuous infall models of Gunn \\markcite{G82} (1982), where the age of the stars is determined by the free-fall collapse time of the infalling gas. Since this time-scale is almost equal for bulge and inner disk, no stellar population differences will be expected.  The similarities in color between bulges and disks can be believed to stem naturally from models in which the bulge forms from instabilities in the disk (Pfenniger \\& Norman \\markcite{PN90} 1990; Combes et al. \\markcite{C90} 1990; Pfenniger \\markcite{P93} 1993). This is indeed the case in the absence of starburst processes; starbursts imply discontinuous changes which break the age and metallicity link between parent and daughter population. Starbursts are likely to be necessary in order to build up the central space densities of bulges from disk material; they may be unavoidable if the disk contains any gas. Thus, our result does not necessarily support instability-driven bulge formation scenarios.  Our result, that an important part of the disk is as old as the bulge, also means that, at high redshift, a considerable amount of gas must already have been transformed from gas into stars. We predict that the occurrence of \"naked bulges\" surrounded by large amounts of gas is unlikely at any redshift. The statistics of galaxies associated with damped Lyman-alpha systems, still in low numbers, reveal the predominance of disk morphologies (Wolfe \\etal \\markcite{W94} 1994).  How do color gradients in bulge and disk affect this result? Measured gradients in both bulges and disks are generally small (BP94, Peletier \\& Balcells 1995, in preparation), so age or metallicity only differ by a small amount if another position in the bulge or disk is chosen. Fisher \\etal \\markcite{FFI95} (1995) report high values of color and line-strength gradients in S0's. The difference with our data is that their systems are extreme edge-on galaxies. It appears that vertical gradients in bulges are large, contrary to radial gradients. We find that the color gradients in the inner effective radius of the bulges are small, and comparable in size to those in elliptical galaxies (BP94, Peletier \\& Balcells, in preparation), so the discrepancy with Fisher \\etal\\ might be due to the fact that, owing to the almost perfect edge-on aspect, their measurements go further out.  The main conclusion of this work is that, consistently from 4 independent colors, we find that inner disks are only slightly younger than bulges. The difference in age in general lies somewhere between 0 and 3 Gyr."
+    },
+    "9602/astro-ph9602107_arXiv.txt": {
+        "abstract": "The gravitational wave emission by a distorted rotating fluid star is computed. The distortion is supposed to be symmetric around some axis inclined with respect to the rotation axis. In the general case, the gravitational radiation is emitted at two frequencies: $\\Omega$ and $2\\Omega$, where $\\Omega$ is the rotation frequency. The obtained formul\\ae\\  are applied to the specific case of a neutron star distorted by its own magnetic field. Assuming that the period derivative $\\dot P$ of pulsars is a measure of their magnetic dipole moment, the gravitational wave amplitude can be related to the observable parameters $P$  and $\\dot P$ and to a factor $\\beta$ which measures the efficiency of a given magnetic dipole moment in distorting the star. $\\beta$ depends on the nuclear matter equation of state and on the magnetic field distribution. The amplitude at the frequency $2\\Omega$, expressed in terms of $P$, $\\dot P$ and $\\beta$, is independent of the angle $\\alpha$ between the magnetic axis and the rotation axis, whereas at the frequency $\\Omega$, the amplitude increases as $\\alpha$ decreases. The value of $\\beta$ for specific models of magnetic field distributions has been computed by means of a numerical code giving self-consistent models of magnetized neutron stars within general relativity. It is found that the distortion at fixed magnetic dipole moment is very dependent of the magnetic field distribution; a stochastic magnetic field or a superconductor stellar interior greatly increases $\\beta$ with respect to the uniformly magnetized perfect conductor case and might lead to gravitational waves detectable by the VIRGO or LIGO interferometers. The amplitude modulation of the signal induced by the daily rotation of the Earth has been computed and specified to the case of the Crab pulsar and VIRGO detector. ",
+        "introduction": "Rapidly rotating neutron stars (pulsars) might be an important source of continuous gravitational waves in the frequency bandwidth of the forthcoming LIGO and VIRGO interferometric detectors (cf. Bonazzola \\& Marck 1994 for a recent review about the astrophysical sources these detectors may observe). It is well known that a stationary rotating body, perfectly symmetric with respect to its rotation axis does not emit any gravitational wave. Thus in order to radiate gravitationally a pulsar must deviate from axisymmetry. Various kinds of pulsar asymmetries have been suggested in the literature: first the crust of a neutron star is solid, so that its shape may not be necessarily axisymmetric under the effect of rotation, as it would be for a fluid: deviations from axisymmetry are supported by anisotropic stresses in the solid. The shape of the crust depends not only on the geological history of the neutron star, especially on the episode of crystallization of the crust, but also on star quakes. Due to its violent formation (supernova) or due to its environment (accretion disk), the rotation axis may not coincide with a principal axis of the neutron star moment of inertia and the star may precess (Pines \\& Shaham 1974). Even if it keeps a perfectly axisymmetric shape, a freely precessing body radiates gravitational waves (Zimmermann \\& Szedenits 1979). Neutron stars are known to have important magnetic fields; magnetic pressure (Lorentz forces exerted on the conducting matter) can distort the star if the magnetic axis is not aligned with the rotation axis, which is widely supposed to occur in order to explain the pulsar phenomenon. Another mechanism for producing asymmetries is the development of non-axisymmetric instabilities in rapidly rotating neutron stars driven by the gravitational radiation reaction (CFS instablity, cf. e.g. Schutz 1987) or by nuclear matter viscosity (cf. e.g. Bonazzola, Frieben \\& Gourgoulhon 1996 and references therein).  In the seventies it was widely thought that the solid part of neutron stars was large (see Sect.~III of Goldreich 1970). A rotating rigid body whose rotation axis does not coincide with a principal axis of the moment of inertia must have a precessional motion. Therefore, there have been a number of studies about the precession of neutron stars and the resulting gravitational wave emission (Zimmermann 1978\\footnote{the results presented in this reference are erroneous and are corrected in Zimmermann \\& Szedenits (1979)}, Zimmermann \\& Szedenits 1979, Zimmermann 1980, Alpar \\& Pines 1985, Barone et al. 1988, de Araujo et al. 1994). However, modern dense matter calculations reveal that neutron star interiors are completely liquid (see e.g. Haensel 1995). Since the crust represents only 1 to 2\\% of the stellar mass (Lorenz et al. 1993), it appears that most of the neutron star is in a liquid phase. In this case, the star is not expected to precess, or at least its precession frequency must be reduced by a factor $10^5$ as compared with the solid case (Pines \\& Shaham 1974).  The gravitational radiation from a rotating fluid star has not been studied as much as that from a solid star. The only reference we are aware of is the work by Gal'tsov, Tsvetkov \\& Tsirulev (1984) and by Gal'tsov \\& Tsvetkov (1984) about the gravitational emission of a rotating Newtonian homegeneous fluid drop with an oblique magnetic field. These authors have shown that the gravitational waves are emitted at two frequencies: the rotation frequency and twice the rotation frequency. They did not give explicit formul\\ae\\ for the two gravitational wave amplitude $h_+$ and $h_\\times$ but have instead computed the response of a heterodyne detector (consisting in a rotating dumbbell orientated at right angles to the direction of propagation of the wave) directly from the Riemann tensor.  In the present article, we compute the gravitational radiation from a rotating fluid star, distorted along a certain direction by some process (for example an internal magnetic field), the angle between this direction and the rotation axis being arbitrary (Sect.~\\ref{s:generation}). We do not assume that the star is a Newtonian object: it can have a large gravitational field described by the theory of general relativity --- which is relevant for neutron stars. As an illustration we give the specific example of a Newtonian incompressible fluid with a uniform internal magnetic field (Sect.~\\ref{s:incomp}). We consider also the (more general) case where the deformation is due to an internal magnetic field which is also responsible for the $\\dot P$ of the pulsar (electromagnetic braking) (Sect.~\\ref{s:magnetic}). Using the numerical code presented in (Bocquet, Bonazzola, Gourgoulhon \\& Novak 1995) for magnetized rotating neutron stars in general relativity, we compute the gravitational emission for a specific $1.4\\, M_\\odot$ neutron star model and for various configurations of the internal magnetic field. Finally, we derive the response of an interferometric detector to the gravitational signal, taking into account the daily change induced by the Earth rotation in the relative orientation of the detector arms and the source (Sect.~\\ref{s:detect}). ",
+        "conclusions": "We have considered the gravitational radiation emitted by a distorted rotating fluid star. The distortion is supposed to be symmetric with respect to some axis which does not coincide with the rotation axis. The gravitational emission takes place at two frequencies: $\\Omega$ and $2\\Omega$, where $\\Omega$ is the rotation frequency, except in the particular case where the distortion axis is perpendicular to the rotation axis (only the frequency $2\\Omega$ is then present). As an application, the magnetic field induced deformation is treated. If, as usually admitted, the period derivative, $\\dot P$, of pulsars is a measure of their magnetic dipole moment, the gravitational wave amplitude can be related to the observable parameters $P$ and $\\dot P$ of the pulsars and to a factor $\\beta$ which measures the distortion response of the star to a given magnetic dipole moment. $\\beta$ depends on the nuclear matter equation of state and on the magnetic field distribution. The amplitude at the frequency $2\\Omega$, expressed in terms of $P$, $\\dot P$ and $\\beta$, is independent of the angle $\\alpha$ between the magnetic axis and the rotation axis, whereas at the frequency $\\Omega$, the amplitude increases as $\\alpha$ decreases.  Using a numerical code generating self-consistent models of magnetized neutron stars within general relativity, we have computed the deformation for explicit models of the magnetic field distribution and a realistic equation of state. It appeared that the distortion at fixed magnetic dipole moment depends very sensitively on the magnetic configuration. The case of a perfect conductor interior with toroidal electric currents is the less favorable one, even if the currents are concentrated in the crust. Stochastic magnetic fields (that we modeled by considering counter-rotating currents) enhance the deformation by several orders of magnitude and may lead to a detectable amplitude for a pulsar like the Crab. As concerns superconducting interiors --- the most realistic configuration for neutron stars ---  we have studied numerically type I superconductors, with a simple magnetic structure outside the superconducting region. The distortion factor is then $\\sim 10^2$ to $10^3$ higher than in the normal (perfect conductor) case, but still insufficient to lead to a positive detection by the first generation of kilometric interferometric detectors. We have not studied in details the type II superconductor but have put forward some argument which makes it a promising candidate for gravitational wave detection. Due to the complicated microphysics involved in type II superconductors we delay their study to a future paper. We also plan to study the deformation induced by a possible ferromagnetic solid interior of neutron stars, as well as the effects of a strong {\\em toroidal} internal magnetic field.  Regarding the reception of gravitational waves from a pulsar by an interferometric detector, we have computed the amplitude modulation of the signal induced by the diurnal rotation of the Earth. By inspecting the wave form, and assuming the position of the pulsar to be known (the pulsar can be recognized by its period), one can determine the inclination angle of the line of sight with respect to the pulsar rotation axis, as well as the orientation of the pulsar equatorial plane. Moreoever, by comparing the wave forms at the frequencies $\\Omega$ and $2\\Omega$, the angle $\\alpha$ between the rotation axis and the magnetic axis can be determined.  Pulsars may be good candidates for the detection by the forthcoming VIRGO and LIGO interferometric detectors. A frequently invoked mechanism for gravitational emission concerns asymmetries of the neutron star solid crust and the resulting precession. In this article, we have examined instead the bulk deformation of the star induced by its own magnetic field. For some configurations of the magnetic field (stochastic distribution, type II superconductor), the deformation may be large enough to lead to a detectable signal by VIRGO, with the total magnetic dipole moment (or equivalently the surface magnetic field) keeping its (relatively small) observed value. The positive detection of gravitational waves from pulsars would lead to some constraints on the internal magnetic field distribution, which would be of great interest for the theories of pulsar magnetospheres. This would constitute an example of a significant contribution of gravitational astronomy to classical astrophysics."
+    },
+    "9602/astro-ph9602113_arXiv.txt": {
+        "abstract": "The possibility of a correlation between the radio (cm)- and $\\gamma$-ray luminosity of variable AGN seen by EGRET is investigated. We performed Monte-Carlo simulations of typical data sets and applied different correlation techniques (partial correlation analysis, $\\chi^2$-test applied on flux-flux relations) in view of a truncation bias caused by sensitivity limits of the surveys. For K-corrected flux densities, we find that with the least squares method only a linear correlation can be recovered. Partial correlation analysis on the other side provides a robust tool to detect correlations even in flux-limited samples if intrinsic scatter does not exceed $\\sim 40$ \\% of the original $\\gamma$-ray luminosity. The analysis presented in this paper takes into account redshift bias and truncation effects simultaneously which was never considered in earlier papers.  Applying this analysis to simultaneously observed radio- and $\\gamma$-ray data, no correlation is found. However, an artificial correlation appears when using the mean flux. This is probably due to the reduction of the dynamical range in the flux-flux relation. Furthermore, we show that comparing the emission in both spectral bands at a high activity state leads to no convincing correlation.  In conclusion, we can not confirm a correlation between radio and $\\gamma$-ray \\underline{luminosities} of AGN which is claimed in previous works. ",
+        "introduction": "Many attempts have been made in the past to investigate correlations between radio (cm)- and $\\gamma$-ray luminosities of AGN (Stecker et al. 1993, Padovani et al. 1993, Salamon \\& Stecker 1994). By applying regression and correlation analysis to the 2.7 GHz- and 5 GHz-luminosity and the $\\gamma$-ray emission in the EGRET energy band, a (nearly) linear correlation was found. The use of luminosities instead of fluxes, however, always introduces a redshift bias to the data, since luminosities are strongly correlated with redshift. Especially in samples which cover a wide range of distance, a correlation will appear in a luminosity plot even when there is no correlation in the corresponding flux densities (Elvis et al. 1978). Feigelson \\& Berg (1983) show that, if there is no intrinsic luminosity-luminosity correlation, no correlation will appear in the flux-flux relation. On the other hand, though the redshift dependence can be removed, intrinsic correlations between luminosites $L_1$ and $L_2$ may be lost in the flux diagrams $f_1$-$f_2$: if $L_{1} \\sim L_{2}^B$, then $f_{1} \\sim f_{2}^{B}z^{2(1-B)}$. For $B \\neq 1$ each value $f_{2}^{B}$ will be multiplied by a 'random' value $z^{2(1-B)}$ causing any intrinsic $L_{1}-L_{2}$ correlation to be smeared out. The apparent correlation is maintained only when the underlying relationship is linear (Feigelson \\& Berg 1983). Furthermore, any observational uncertainties will hamper the finding of correlations using flux diagrams. It is therefore crucial to estimate the influence of the redshift bias on the correlations between the emission from the two wavebands. Non-parametric partial correlation coefficients (Kendall's $\\tau$ or the Spearman rank correlation coefficient $R_s$) have been used to deal with this problem (Dondi \\& Ghisellini 1995). One limitation of these rank correlation tests is the failure to take into account possible changes in rank due to observational uncertainties.\\\\ Observational flux-limits of the samples are another serious bias which influences any correlation analysis as they restrict the populated region in the luminosity-luminosity diagram to a narrow band. Therefore, Feigelson \\& Berg (1983) proposed to include all upper limits to avoid artificial correlations and incorrect conclusions (Schmitt 1985), and suggest the use of survival analysis. However, if the censored data points are not distributed randomly, but localize a particular area, survival analysis may give misleading results (Isobe 1989). Furthermore, this analysis can not account for a bias caused by misidentified sources or by objects which are completely lost due to the low sensitivity of the instrument. Those truncation effects must be seriously considered when EGRET data are used.\\\\ Another problem arises when considering the nature of the sources. Blazars are known to be strongly variable in the $\\gamma$-ray as well as in the radio band on time scales of days to months (von Montigny et al. 1995). Therefore, simultaneous observations should be the adequate data for a correlation analysis. However, due to the lack of such data, the mean (Padovani et al. 1993) or the brightest flux values (Dondi \\& Ghisellini 1995) of the sources have been used in the past. For highly variable sources this choice reduces the dynamical range in the luminosity-luminosity plot essentially, and can mimic a correlation.  In this paper, we evaluate the reliability of the partial correlation analysis in luminosity-luminosity plots as well as $\\chi^{2}$-fits in flux-flux diagrams for flux-limited truncated samples with strongly variable sources. For this purpose we perform Monte-Carlo simulations of typical data sets with known degrees of correlation and flux sensitivities, which is described in section 2. No censored data are used. The simulations provide a useful test of how correlation analysis methods deal with such selection effects.  The application of a database of EGRET-blazars in the radio (cm)- and $\\gamma$-ray regime is presented in the last section.  Throughout this paper the Hubble constant $H_0=75$ km \\mbox{$s^{-1}$} \\mbox{$Mpc^{-1}$} and the deceleration parameter $q_0 = 0.5$ have been used. ",
+        "conclusions": "In previous work a correlation between the radio (cm)- and $\\gamma$-ray luminosity of blazars seen by EGRET was claimed. For the correlation analysis biases caused by the limited sensitivity of the instruments and the strong redshift dependence of the luminosities must be considered. While the effects of flux-limits can be taken into account by including upper limits, partial correlation analysis deals with the redshift bias. In this paper we considered both effects simultaneously in contrast to earlier works. For this purpose we performed Monte-Carlo simulations of a typical sample of variable blazars with known degree of correlation and completeness. We then analyzed the data sets by partial correlation analysis using the Spearman rank order correlation coefficient and a least square fit in flux-flux relations. The simulations reveal how correlation analysis deals with truncation biases caused by a limited sensitivity of the instrument. While in flux-flux diagrams only linear correlations can be recovered (provided that K-corrected fluxes are used) with the least square method, the partial correlation analysis gives correct results even for flux-limited sample. This is demonstrated for exponents $B=0.8 \\ldots 1.5$ ($L_{\\gamma} \\sim L_R^B$) with random noise factors up to an order of magnitude. Furthermore, in data sets with flux-limits, the partial correlation analysis gives a stronger redshift dependence of the luminosities compared to complete samples.  Applying the $\\chi^2$-method on flux-flux diagrams we found no evidence for a correlation between the simultaneously observed (2.7 GHz and 10 GHz) or flaring radio flux and luminosity (4.8 GHz and 8 GHz) and $\\gamma$-ray luminosities above 100 MeV. The same result is obtained using the partial correlation analysis. The apparently positive signal between the average radio- and $\\gamma$-ray luminosities is caused by the restriction of the dynamical range in the correlation diagrams.  We conclude that a correlation between the radio (cm)- and $\\gamma$-ray luminosities of AGN can not be claimed.  Note however, that there is evidence that both radio- and $\\gamma$-ray emissions from blazars are strongly beamed and not isotropic (von Montigny et al. 1995). The apparent source luminosity is related to the intrinsic luminosity by a factor which depends on the bulk Lorenz factor of the outmoving radiation source region within the jet and the angle between the line of sight and the jet axis. Those quantities are thought to be different for each object and each spectral band. Therefore, the beaming effect may smear out a possible correlation between the intrinsic luminosities, or worse, could cause a misleading correlation slope. Correlating luminosities which are corrected for beaming is unfortunately to date not possible since the Doppler factors in the $\\gamma$-ray band are not known from observations.  The fact that we do not find any radio-$\\gamma$-ray correlation for both simultaneously observed luminosities and high activity state observations has several important implications.  One may ask whether the lack of a correlation would rule out any of the proposed $\\gamma$-ray emission processes in blazars. Especially the (inhomogeneous) synchrotron-self-Compton (SSC) models (Jones et al. 1974, Marscher 1980, K\\\"onigl 1981, Marscher \\& Gear 1985, Ghisellini \\& Maraschi 1989, Marscher \\& Bloom 1992, Maraschi et al. 1992) seem to predict a correlation between the synchrotron emission and the inverse Compton component, as the synchrotron photons are upscattered to $\\gamma$-ray energies by the same beamed relativistic electrons. However, there are several reasons why we do not expect to find any relation between the luminosities in both spectral bands, even if we consider the SSC mechanism as the main $\\gamma$-ray emission process. First, rapid variability, especially in the $\\gamma$-ray regime, causes a large scatter in the luminosity and flux diagrams, and prevents recovering the true correlation slope even if a correlation would exist. Second, different relativistic Doppler factors D for the different objects induce an additional scatter since the observer considers different parts of the spectrum depending on the beaming factor D. Third, different $\\gamma$-ray emission components may contribute differently for each individual object contaminating the relation with the radio band (i.e. the lack of strong emission lines in BL Lac objects indicates that the $\\gamma$-ray emission induced by inverse Compton scattering of photons produced in the broad line region is small (Sikora et al. 1994)). Note also, that the radio emission may be influenced by the contribution of a quiescent component (Valtaoja et al. 1988). Fourth, the different variability behaviour in the radio and $\\gamma$-ray band points to a different spatial origin of the two luminosities which excludes a direct link. Fifth, there is evidence for a time lag between the radio and $\\gamma$-ray flare (M\\\"ucke et al. 1996, Valtaoja \\& Ter\\\"asranta 1995, Reich et al. 1993). Simultaneously observed luminosities are therefore not expected to obey a direct correlation. The same is true for the historical maximum relations between the radio and $\\gamma$-ray regime since those flaring states must not be necessarily physically connected. Instead, the fluence of the two corresponding radio and $\\gamma$-ray flares may be related. However, a single radio outburst can often not be isolated since radio flares often overlap especially at low cm-frequencies. Unfortunately, the lack of long-term $\\gamma$-ray light curves of blazars does not enable a meaningful quantitative study at present. Last, the negative finding of a luminosity correlation could still be compatible with a similar energy content of the radio and $\\gamma$-ray producing particles. Thus, we do not consider the non-correlation between simultaneously observed luminosities as strong evidence against the SSC model.  However, the non-correlation yields other astrophysical consequences. For both, the local luminosity density of blazars and their evolution properties, there is no hint for a direct link between radio (cm)- and $\\gamma$-ray energies. So, the $\\gamma$-ray luminosity function can not be predicted from a known radio luminosity function and its evolution as has been used by several authors (Stecker et al. 1993, Salamon \\& Stecker 1994, Padovani et al. 1993, Stecker \\& Salamon 1996) in order to estimate the contribution of radio-loud AGN to the extragalactic diffuse $\\gamma$-ray background. Instead, we suggest to model the $\\gamma$-ray background induced by unresolved blazars by relating the energy content of particles to the $\\gamma$-ray emission in AGN and take into account cosmological distances and the distribution of beaming angles, duty factors, etc.. This will be subject of future work."
+    },
+    "9602/astro-ph9602055_arXiv.txt": {
+        "abstract": "In view of several conflicting results, we reanalyze the effects of magnetic fields on the primordial nucleosynthesis. In the case the magnetic field is homogeneous over a horizon volume, we show that the main effects of the magnetic field are given by the contribution  of its energy density to the Universe expansion rate and the effect of the field on the electrons quantum statistics. Although, in order to get an upper limit on the field strength, the weight of the former effect is numerically larger, the latter cannot be neglected. Including both effects in the PN code we get the upper limit $B \\le 1\\times 10^{11}$ Gauss at the temperature $T = 10^9~^oK$. We generalize the considerations to cases when instead the magnetic is inhomogeneous on the horizon length. We show that in these cases only the effect of the magnetic field on the electrons statistics is relevant. If the coherence length of the magnetic field at the end of the PN is in the range $10 \\ll L_0 \\ll 10^{11}$ cm our upper limit is $B \\le 1\\times 10^{12}$ Gauss. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602079_arXiv.txt": {
+        "abstract": "We investigate the timescales for stochasticity and chaotic mixing in a family of triaxial potentials that mimic the distribution of light in elliptical galaxies. Some of the models include central point masses designed to represent nuclear black holes. Most of the boxlike orbits are found to be stochastic, with mean Liapunov times that are $3-6$ times the period of the long-axis orbit. In models with large cores or small black holes, the stochastic orbits mimic regular box orbits for hundreds of oscillations at least. However a small core radius or significant black hole mass causes most of the stochastic orbits to diffuse through phase space on the same timescale, visiting a significant fraction of the volume beneath the equipotential surface. Some stochastic orbits, with initial conditions lying close to those of regular orbits, remain trapped in all models.  We estimate timescales for chaotic mixing in the more strongly stochastic models by evolving ensembles of $10^4$ points until their distribution reaches a nearly steady state. Mixing initially takes place rapidly, with characteristic times of $10-30$ dynamical times, as the phase points fill a region similar in shape to that of a box orbit. Subsequent mixing is slower, with characteristic times of hundreds of orbital periods. Mixing rates were found to be enhanced by the addition of modest force perturbations, and we propose that the stochastic parts of phase space might be efficiently mixed during the early phases of galaxy formation when such perturbations are large. The consequences for the structure and evolution of elliptical galaxies are discussed. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602145_arXiv.txt": {
+        "abstract": "We have recently proposed a new candidate for baryonic dark matter: very cold molecular gas, in near-isothermal equilibrium with the cosmic background radiation at $2.73\\,\\rm K$. The cold gas, of quasi-primordial abundances, is condensed in a fractal structure, resembling the hierarchical structure of the detected interstellar medium.  We present some perspectives of detecting this very cold gas, either directly or indirectly. The H$_2$ molecule has an ``ultrafine\" structure, due to the interaction between the rotation-induced magnetic moment and the nuclear spins. But the lines fall in the km domain, and are very weak. The best opportunity might be the UV absorption of H$_2$ in front of quasars. The unexpected cold dust component, revealed by the COBE/FIRAS submillimetric results, could also be due to this very cold H$_2$ gas, through collision-induced radiation, or solid H$_2$ grains or snowflakes.  The $\\gamma$-ray distribution, much more radially extended than the supernovae at the origin of cosmic rays acceleration, also points towards and extended gas distribution. ",
+        "introduction": "The possibility that most of the mass of the Universe could be under the form of gas around or in between galaxies has been widely discussed in the 1960's (e.g. the review by Peebles, 1971). At that time gas was assumed to be distributed in a smooth homogeneous fashion. This was not justified by observations of the interstellar gas already then.  The intergalactic material was supposed to be hydrogen exclusively (with some helium), either atomic, molecular, ionized, or even in a condensed snow. Self-gravity was ignored. Several tests were proposed, such as emission or absorption of HI at the $21\\,\\rm cm$ line, the Gunn-Peterson test in the Ly$\\alpha$ line of HI or the Lyman band in molecular hydrogen; the existence of an intergalactic plasma would have been detected through free-free emission or absorption, recombination lines, chromatic phase lag, or Thomson scattering.  All the tests were negative, constraining the density of any intergalactic gas several orders of magnitude below the closure density.  However, the homogeneity hypothesis is drastic, especially when comparing to the now much better known ISM, and including gravity as a major force.  The Gunn-Peterson test is of course invalidated in case of a clumpy medium.  Even until recently, the hypothesis of cold gas as dark matter was eliminated quickly, without critical reflection, through stability arguments (e.g. Hegyi \\& Olive 1986): for an homogeneous galactic gaseous halo, the virial temperature is of the order of $10^6\\,\\rm K$; the gas cannot remain hot, it would have been seen in X-ray; in any case, at this temperature, cooling processes are violent, the gas collapses and forms stars.  Hydrogen snowballs, massive enough not to collide with each other ($m> 1\\,\\rm g$) quickly evaporate, unless gravitationally bound (but then they join the problem of brown dwarfs).  All the above objections disappear when a realistic hierarchical structure of cold and dense clouds is considered, closely resembling the familiar fractal structure of the detected ISM (Pfenniger et al.~1994). A model was then built to account for the baryonic dark matter around galaxies, composed of basic ``clumpuscules\" of molecular hydrogen, of a Jupiter mass (10$^{-3}$ M$_\\odot$), but with a much smaller density than brown dwarfs, with $10^{9-10}\\,\\rm cm^{-3}$ and radius of $\\approx 20\\,\\rm AU$ (Pfenniger \\& Combes 1994). For these basic units clumpuscule are self-gravitating, statistically in equilibrium with pressure forces, at the limit of the adiabatic regime, since then the radiation transfer time equals the dynamical time.  They compose the smallest scale of a hierarchical structure, that ranges over six orders of magnitude, up to giant molecular clouds of $100\\,\\rm pc$; above this scale, bigger gaseous complexes are torn apart by galactic shear. The ensemble of clumpuscules is in quasi isothermal equilibrium with the bath of photons of the cosmic background radiation at $T=2.73 {\\,\\rm K} (1+z)$.  Due to the fractal structure (of dimension $D<2$), the clumpuscules collide together frequently, with, for such fractal dimension, a rate of the order of the cooling-heating and dynamical times.  But all the fractal structure is far from local thermodynamical equilibrium as explained in another paper in this volume, the equilibrium including gravity is only statistical.  The large fluctuations prevent most of the clumpuscules to cool down and collapse further (were they distributed homogeneously they would not). Through these collision induced fluctuations, the gas maintains exchanges with the background, coalescing in giving back energy, or fragmenting and evaporating in absorbing energy.  The condensed structure resembles closely the well-studied ISM gas, already known as a $D<2$ fractal-like structure over several orders of magnitude in scale (Larson 1981; Scalo 1985; Falgarone et al.~1992).  The main difference is that star-formation in the visible disk has metal-enriched the medium that can then cool down much faster, and the heating sources have partly destroyed the condensed fractal and formed a diffuse intercloud medium.  In galaxy outskirts, the condensed H$_2$ phase is almost only in contact with the intergalactic radiation field, which photodissociates a small fraction of it into HI gas, because at the envisaged column densities ($\\sim10^{25}\\,\\rm cm^{-2}$) H$_2$ can self-shield easily. An even smaller fraction could be ionized.  Since the average surface density of the gas is falling as $R^{-1}$, HI must disappear into ionised gas at a critical radius; this corresponds to the sudden fall off of HI measurements, resembling an ionisation front (e.g.~the case of NGC 3198, van Gorkom et al.~1987, unpublished). From this point of view, the atomic gas in the outer parts serves as a tracer of dark matter.  There is indeed some evidence that the gas and dark matter are intimately related.  From the flat rotation curves, the surface density of dark matter $\\sigma_{\\rm DM}$ varies asymptotically as $R^{-1}$, and as well the HI surface density $\\sigma_{\\rm HI}$ (cf.~Fig.~1). Bosma (1981) was the first to notice a constant ratio of $\\sigma_{\\rm DM}/\\sigma_{\\rm HI}$ as a function of radius in spirals, which has been confirmed by many authors (Sancisi \\& van Albada 1987; Puche et al.~1990), and varies between 10 and 20 according to the morphological type (Broeils 1992; Carignan, this volume).  \\begin{figure}% \\vspace{5cm} \\caption{Ratio of HI to Dark Matter (DM) surface densities in spiral galaxies, adapted from Bosma (1981), Puche et al (1990) and Freeman (1992). } \\end{figure}  The presence of large gaseous extensions around galaxies can explain the widespread detection of absorption lines in front of quasars. The Ly$\\alpha$ forest, and the large frequency of absorptions on a single line of sight (up to 100) remain unexplained. If these absorptions come from gas around galaxies, the derived cross-section of a galaxy corresponds to a radius of $480\\,\\rm kpc$ at $z=2.5$ (Sargent 1988). The gas corresponding to these atomic absorptions remains a small fraction of the critical density.  But already the total contribution of the gas in damped Ly$\\alpha$ systems amounts to about the luminous matter density in present galaxies (e.g.~Lanzetta et al.~1991).  The model of cold gas as dark matter sheds also some light on puzzles such as the presence of huge amounts of hot gas in clusters (the hot gas represents 8 to 10 times the mass in galaxies in some rich clusters, Edge \\& Stewart 1991), the overabundance of gas in interacting galaxies (Braine \\& Combes 1993), the necessity of gas infall for maintaining star-formation and the spiral structure, and the evolution of galaxies along the Hubble sequence (e.g.~Pfenniger et al.~1994).  Recent models have also been proposed, taking up the hypothesis that H$_2$ gas could contribute significantly to the dark matter, based mainly in massive proto globular cluster clouds possibly mixed with brown dwarfs (de Paolis et al.~1995; Gerhard \\& Silk 1995).  In this case, the cold H$_2$ is not {\\it very\\/} cold, but has a temperature between 5 and $20\\,\\rm K$, which makes it easier to detect in emission. To us is unclear, however, how cold gas coexisting with a brown dwarf cluster could be maintained in the mutual gravity at this low temperature, since the brown dwarfs must be subject to a significant dynamical friction, and the cluster must undergo core collapses that inevitably relax and virialize also the gas to a higher temperature. ",
+        "conclusions": "The main argument for the existence of baryonic dark matter comes from the constraints of the Big Bang nucleosynthesis, compared with the observed abundances of primordial elements. The most recent estimates find that the baryon density relative to the critical closure density must lie in the range $\\Omega_B = 0.01-0.09$ (Smith et al.~1993). The visible baryons account for a much lower density, around $\\Omega = 0.002$ (Persic \\& Salucci 1992).  But the galactic dark matter could still be entirely baryonic (e.g.~Carr 1995).  The recent micro-lensing experiments conducted towards the Magellanic Clouds have revealed that Machos can account for about 20\\% of these dark baryons (Aubourg et al.~1993; Alcock et al.~1995), although the constraint is loose.  Molecular hydrogen is one of the least exotic candidate (Pfenniger et al.~1994). Present observations are not incompatible with this hypothesis.  The main difficulty to detect very cold gas in emission is its temperature close to the one of the cosmic background. However, we must search for observational tests to falsify the proposition.  The detection of the ``ultrafine\" structure of the ortho-H$_2$ molecules at km wavelengths raises considerable difficulties for the near future, but could be a means to fix the $N($H$_2$)/I(CO) conversion ratio in our Galaxy.  The LiH rotational lines will be easily detectable by submillimeter satellites.  Absorption lines detection might be the best way if the gas is indeed very cold. Since the surface filling factor of the molecular clumps is low ($f<1\\%$), large statistics are required, but the perspectives are far from hopeless. H$_2$ absorption in the Lyman and Werner bands has already been identified in a damped Ly$\\alpha$ system, at $z=2.8$. For a clumpuscule in our own galaxy falling just on the line of sight of a quasar, we expect a strongly damped and transient absorption over a few months.  Finally, it is not excluded that the cold dust component detected by COBE/FIRAS is tracing the cold H$_2$ component, limited to galactic radii where the cold gas is still mixed with some dust.  Gamma ray data could also be interpreted with the help of radially extended gas distributions."
+    },
+    "9602/astro-ph9602003_arXiv.txt": {
+        "abstract": "In this paper we present new chemo-spectro-photometric models of elliptical galaxies in which infall of primordial gas is allowed to occur. They aim to simulate the collapse of a galaxy made of two components, i.e. luminous material and dark matter. The mass of the dark component is assumed to be constant in time, whereas that of the luminous material is supposed to accrete at a suitable rate. They also include the effect of galactic winds powered by supernova explosions and stellar winds from massive, early-type stars. The models are constrained to match a number of properties of elliptical galaxies, i.e.  the slope and mean colours of the colour-magnitude relation (CMR), V versus (V--K), the UV excess as measured by the colour (1550--V) together with the overall shape of the integrated spectral energy distribution (ISED) in the ultraviolet,  the relation between the ${\\rm Mg_2 }$ index and (1550--V), the mass to blue luminosity ratio ${\\rm M/L_{B}}$ as a function of the ${\\rm B}$ luminosity, and finally the broad-band colours (U--B), (B--V), (V--I), (V--K), etc. The CMR is interpreted as a mass-metallicity sequence of old, nearly coeval objects, whose mean age is 15 Gyr. Assuming the law of star formation to be proportional to ${\\rm M_{g}^{k}(t)}$ with $k = 1$, the rate of star formation as function of time starts small, grows to a maximum, and then declines thus easily avoiding the excess of metal-poor stars found by BCF with the closed-box scheme (the analog of the G-Dwarf Problem in the solar vicinity). Owing to their stellar content, infall models can easily reproduce all the basic data of the galaxies under examination. As far as the UV excess is concerned, the same sources proposed by BCF are found to hold also with the infall scheme. H-HB and AGB manqu\\'e stars of high metallicity play the dominant role, and provide a robust explanation of the correlation between the  (1550--V) colour and the luminosity, mass and metallicity of the galaxies.  Furthermore, these models confirm the potential  of the (1550--V) colour as an  age indicator in cosmology as already suggested by BCF. In the rest frame of a massive and metal-rich elliptical galaxy, this colour suffers from one major variation: at the onset of the so-called H-HB and AGB-manqu\\'e stars (age about 5.6 Gyr). This transition occurs at reasonably small red-shifts and therefore could be detected with the present-day instrumentation. ",
+        "introduction": "Bressan et al. (1994,\\ BCF) have recently presented chemo-spectro-photometric models for elliptical galaxies that were  particularly designed to match the UV excess observed in these systems  (Burstein et al.\\ 1988) together with its dependence on the galaxy luminosity, mass, metallicity, ${\\rm Mg_{2}}$ index,  and age. In brief, BCF models stand on the closed-box approximation and the enrichment law $\\Delta Y/\\Delta Z = 2.5$,  allow for the occurrence of galactic winds powered by the energy input from supernovae and stellar winds from massive stars, and use  the CMR of elliptical galaxies in Virgo and Coma clusters (Bower et al. 1992a,b) as one of the main constraints.  In the BCF models the history of star formation consists of an initial period of activity followed by quiescence  after the onset of the galactic winds, whose duration depends on the galactic mass, being longer in the high mass galaxies and shorter in the low mass ones. With the adoption of the closed-box description of chemical evolution, the maximum efficiency of star formation occurred at the very initial stage leading to a  metallicity distribution in the models (relative number of stars per metallicity interval) skewed towards low metallicities.   The comparison of the theoretical integrated spectral energy distribution (ISED) of the models with those of prototype galaxies with different intensity of the UV excess, for instance NGC~4649 with strong UV emission and (1550--V) $\\sim 2.24$  and NGC~1404 with intermediate UV excess and (1550--V) $\\sim 3.30$, indicated that a good agreement was possible but for the region 2000 $\\AA$ to about 3500 $\\AA$ where the theoretical ISED exceeded the observed one.    \\begin{figure} \\psfig{file=fig1a.ps,height=9.0truecm,width=8.5truecm} \\caption{The observed spectrum of the elliptical galaxy NGC~4649 (filled squares) kindly provided by L. Buson (1995, private communication). The vertical bars show the global error on the (1550--V) colours from Burstein et al. (1988). The two  lines show theoretical spectra for the galactic model of BCF with \\ML = ${\\rm 3\\time10^{12} M_{\\odot} }$ characterized by the parameters ${\\rm k=1}$, $\\nu=20$ and $\\zeta=0.40$ and age of 15 Gyr, obtained with different assumptions concerning the contribution from stars in different metallicity bins. The solid  line is the regular spectrum in which all the components are present. Note the discrepancy in the wavelength interval 2000 - 4000 \\AA.  The dashed line is the same but without the contribution from all stars with metallicity ${\\rm Z \\leq 0.008}$. Note the agreement reached in this case. All the spectra are normalized to coincide in the flux at 5000 \\AA.  } \\label{fig1a} \\end{figure}    The analysis of the reason of disagreement led  BCF to conclude that this is a consequence of the closed-box approximation. Indeed this type of model is known to predict an excess of low-metal stars. The feature is intrinsic to the model and not to the particular numerical algorithm used to follow the chemical history of a galaxy.  In order to check their suggestion, BCF performed several numerical experiments in which they artificially removed from the mix of stellar populations predicted by the closed-box models the contribution to the ISED from stars in different metallicity bins. Looking at the paradigmatic case of the ISED of NGC~4649, BCF found that this can be matched by a mix of stellar populations in which no stars with metallicity lower than ${\\rm Z=0.008 }$ are present. The results of these experiments are shown here in Fig.~\\ref{fig1a} for the sake of clarity.  It goes without saying that one may attribute the discrepancy of predicted and observed spectra in the region 2000 to 3500 \\AA\\ to uncertainties in the theoretical spectra. However, the following considerations can be made. Firstly, those experiments clarified that the main contributors to the excess of flux are the turn-off stars of low metallicity ($\\rm 0.0004 \\leq Z \\leq 0.008 $)  whose effective temperature at the canonical age of 15 Gyr ranges from 6450 K to 5780 K. Secondly, the  Kurucz (1992) spectra (at the base of the library adopted by BCF) fit the Sun and Vega, two stars of different metallicity but whose effective temperatures encompass the above values. Finally,  the disagreement in question implies  that the theoretical spectra overestimate the flux in this region by a factor of about 4, which is hard to accept. On the basis of the above arguments BCF concluded that the excess is real and that the discrepancy in question is the analog of the classical G-Dwarf Problem in the solar vicinity.  In brief, it has long been known that the closed-box model applied to study the chemical history of the solar vicinity fails to explain the metallicity distribution observed among the old field stars, giving rise to the so-called G-Dwarf Problem  (Tinsley 1980a,b). The solution of the G-Dwarf dilemma are models with infall (Lynden-Bell 1975, Tinsley 1980a, Chiosi 1981), i.e. models in which the total mass of the disk is let increase with time at a convenient rate starting from a  much lower value. With the current law of star formation (proportional to the gas mass) the competition between gas accretion by infall and gas depletion by star formation gives rise to a non monotonic time dependence of the star formation rate which instead of steadily decreasing from the initial stage as in the closed-box model starts small, increases to a peak value, and then declines over a time scale which is a sizable fraction of the infall time scale. The advantage of the infall model with respect to the closed-box one, is that the metallicity increases faster, and very few stars are formed at very low metallicity. Since the excess of very low-metal stars is avoided, the G-Dwarf Problem is naturally solved.  Applied to an elliptical galaxy, the infall model of chemical evolution can closely mimic the collapse of the parental proto-galaxy from a very extended size to the one we see today. Plausibly, as the gas falls into the gravitational potential well at a suitable rate and the galaxy shrinks, the gas density increases so that star formation begins. As more gas flows in, the more efficient star formation gets. Eventually, the whole gas content is exhausted and turned into stars, thus quenching further star formation. Like in the disk, the star formation rate starts small, rises to a maximum, and then declines. Because of the more efficient chemical enrichment of the infall model, the initial metallicity for the bulk of star forming activity is significantly different from zero.  In this paper, we present simple models of spectro-photometric evolution of elliptical galaxies whose chemical enrichment is governed by the infall scheme, i.e. models in which the mass of the galaxy is let increase with time. Models of this type are consistent with the chemo-dynamical models by Theis et al.\\ (1992) and are advocated by Phillips (1993) to interpret the galaxy counts at faint magnitudes, which are known to exceed those expected from standard evolutionary models with an initial spike of star formation followed by quiescence.  As in BCF, the models (both chemical and spectro-photometric) are constrained to reproduce at the same time a number of key features of elliptical galaxies, namely the slope and mean colours of the CMR, V versus (V--K), for the galaxies in the Virgo and Coma clusters, the UV excess as measured by the colour (1550--V) together with the overall ISED, the relation between the ${\\rm Mg_2 }$ index and (1550--V), the mass to blue luminosity ratio ${\\rm M/L_B }$ as  a function of the ${\\rm B}$ luminosity, and finally the present broad band colours (U--B), (B--V), (V--I), (V--K), etc.  The CMR is from Bower et al. (1992a,b), the data on the UV excess and ${\\rm Mg_2 }$-(1550--V) relationship are from Burstein et al.\\ (1988), and the data on the ratio ${\\rm M/L_B }$ are from Bender et al.\\ (1992,1993) and Terlevich \\& Boyle (1993) for ${ \\rm H_0=50~ km~ sec^{-1}~ Mpc^{-1}  }$.  There is one aspect of the problem that should be clarified in advance. The bottom line of BCF models is the interpretation of the Virgo and Coma cluster CMR as a mass-metallicity sequence whose slope is driven by the onset of galactic winds (cf. Larson 1974; Saito 1979a,b; Vader 1986; Arimoto \\& Yoshii 1986; Matteucci \\& Tornamb\\'e 1987; Angeletti \\& Giannone 1990; Padovani \\& Matteucci 1993) and whose tightness  is the signature of old ages with small dispersion, $13 \\div 15$ Gyr according to Bower et al. (1992a,b).  In contrast, the CMR for nearby galaxies in small groups and field is more dispersed than the one for the Virgo and Coma clusters suggesting that elliptical galaxies may originate from the mergers of spiral galaxies spread over several billion years (Schweizer \\& Seitzer 1992; Alfensleben \\& Gerhard 1994; Charlot \\& Silk 1994). Also in this case the CMR implies the existence of a suitable mass-metallicity sequence.  Therefore, the view that elliptical galaxies are primordial old objects with a dominant initial burst of activity may co-exist with the view that many if not all of them are the result of mergers  implying  more recent star formation activities that may vary from galaxy to galaxy.  How the  different interpretation of the CMR   affects the use of this as a  constraint on our models ? This question has been addressed by Bressan et al. (1995, \\ BCT), who showed that the slope and the mean colours of the CMR are compatible with both alternatives,  thus validating its ability to  constrain the models. Therefore, owing to its better definition, we adopt the Bower et al. (1992a,b) CMR and assume that it is made of  nearly coeval objects whit  mean age of about 15 Gyr. Strictly speaking the  models below are best suited to cluster elliptical galaxies. The analysis of the merger scenario  is beyond the scope of this study.  Section 2  reports on the new spectral library adopted to calculate the spectro-photometric models presented in this paper. Section 3 describes the properties of Single Stellar Populations (SSP) calculated with the new spectral library above. Section 4 outlines the chemical models in usage. Section 5 contains closed-box models calculated with the new spectral library and compares them with those of BCF. It is found that no closed-box model can match all the imposed constraints. Section 6 is the same but for infall models. Section 7 presents the results of the chemical and spectro-photometric models with infall and describes in detail the properties of the models with different mass. In particular we examine the CMR, the mass to blue luminosity ratio ${\\rm M/L_{B} }$, the UV excess and accompanying (1550--V) versus ${\\rm Mg_{2}}$ relation, the evolution of the broad-band colours as a function of the age, and the potential capability of the (1550--V) colour of being a good age indicator in cosmology. Finally,  Section 8 contains a few concluding remarks. ",
+        "conclusions": "In this paper we have presented one zone chemo-spectro-photometric  models of elliptical galaxies with infall of primordial gas. Galaxies are supposed to be made of two components, i.e. luminous and dark matter. While the mass of the dark component is assumed to be constant with time, that of the luminous material is supposed to fall at a suitable rate into the common potential well.  The main motivation for this type of model was to mimic the collapse of a galaxy and to cope with the difficulty reported by BCF using the closed-box scheme, i.e. the analog of the G-Dwarf problem. The adoption of the infall scheme adds another dimension to the problem, i.e. the time scale of mass accretion $\\tau$. Taking the free-fall time scale as a reference, instead of changing $\\tau$ with the galaxy mass, low mass galaxies would have got closer and closer to the closed-box approximation, we prefer to adopt the unique value $\\tau=0.1$ Gyr, i.e. the free-fall time scale for a typical galaxy with mass 0.5 \\M12.  The models also include the occurrence of galactic winds that have long been considered the key physical process governing the CMR of elliptical galaxies. The galactic winds are supposed to occur when the energy input from supernovae and stellar winds from luminous early type stars equals the gravitational binding energy of the gas. It goes without saying that this is a point of great uncertainty of the models. In this respect our approach does not go beyond other current models of elliptical galaxies with galactic winds taken into account. Both the rate of supernova explosions and cooling of the energy injected by these stars are treated in the standard fashion (see for instance Arimoto \\& Yoshii 1987, Matteucci \\& Tornamb\\'e 1987, Brocato et al.\\ 1990, Angeletti \\& Giannone 1990). However, in agreement with BCF and in contrast with other authors, we confirm that if the CMR of elliptical galaxies has to be understood a mass-metallicity sequence of nearly coeval objects, in addition to supernova explosions another source of energy must be at work in order to trigger galactic winds before the metallicity has grown to such high values that only very red colours are possible.  The adoption of the infall scheme has the advantage over the closed-box description that with the law of star formation proportional to ${\\rm M_{g}^{k}(t)}$, the rate of star formation as a function of time starts small, grows to a maximum and then declines thus easily avoiding the excess of zero metal stars typical of the closed-box scheme. Indeed, infall models have the metallicity distribution of its stellar content skewed towards relatively high metallicities.  As repeatedly said, the bottom line to contrive the various parameters of the models has been the simultaneous fit of many properties of elliptical galaxies. We have found that with the new library of stellar spectra in usage the simultaneous fit of the (V--K) and (1550--V) colours is not possible  with the closed-box description, whereas this is feasible with the infall scheme. The fit of the slope and mean colours of the CMR and UV excess of elliptical galaxies required that the efficiency of star formation per unit mass of gas ($\\nu$) must increase with the galaxy mass. As a consequence of this, the overall duration of the star forming activity (before the onset of galactic winds)  in massive galaxies is shorter  than in the low mass ones. This kind of trend is in agreement with the hint coming from the observational dependence of the [Mg/Fe] ratio with the galaxy luminosity (and hence mass) and the suggestion based on chemical models advanced by Matteucci (1994).  As far as the simultaneous fit of the CMR, (V--K) versus V,  and the (1550--V) versus $Mg_{2}$ relation by our model galaxies as a whole imply that the two observational sets of data refer to the same region of the galaxies. While the data on the UV-excess, colour (1550--V) and $Mg_{2}$ index essentially refer to the central parts of the galaxies (see Burstein  et al. 1988), the integrated magnitudes and colours, such as V, (B--V), (V--K) etc., refer to the whole galaxy. This is particularly true with the data we are using for galaxies in the Virgo and Coma cluster. The presence of spatial gradients of colours and metallicity that are known to exist in elliptical galaxies (Carollo et al. 1993, Carollo \\& Danziger 1994, Davies et al.1993, Shombert et al. 1993) indicates that the one-zone models (either closed or with infall) ought to abandoned in favor of models in which the spatial distribution of mass density and star formation rate is taken into account. It seem reasonable to argue that higher metallicities can be present in the central regions thus facilitating the formation of the right type of stellar source of the UV radiation (H-HB and AGB-manqu\\'e stars) whereas lower metallicities across the remaining parts of the galaxies would make it easier to fit the others integrated colours such as (V--K). Preliminary models of elliptical galaxies with spatial distribution of total mass and gas density, star formation rates, metallicities and colours are in progress (Tantalo 1994).  An alternative scenario is been proposed by Elbaz et al. (1995) aimed at explaining the amount of iron present in the intra-cluster medium. The goal is achieved by means of models for the evolution of elliptical galaxies with bimodal star formation. In brief, most of the heavy elements are supposed to be produced by SNII during a first violent phase during which only high-mass stars are formed. This is followed by a phase of quiescence caused by galactic winds. Finally, star formation is supposed to start again with normal IMF when sufficient gas has been made available by previously borne stars. However, using the Elbaz et al. (1994) models Gibson (1995) has tried to reproduce the present-day photometric and chemical properties of elliptical galaxies with the conclusion that these models while successfully reproducing the observational data for the intra-cluster medium, fail to reproduce the chemo-spectro-photometric properties of the present day elliptical galaxies.  \\vspace{1truecm}  \\noindent {\\bf Acknowledgements} \\\\ This work has been financially supported by the Italian Ministry of University, Scientific Research and Technology (MURST), the National Council of Research (CNR-GNA), and the Italian Space Agency (ASI).  \\vspace{0.5truecm}"
+    },
+    "9602/astro-ph9602017_arXiv.txt": {
+        "abstract": "From the spectra and light curves it is clear that SNIa events are thermonuclear explosions of white dwarfs. However, details of the explosion are  highly under debate. Here, we present detailed models which are consistent with respect to the explosion mechanism, the optical and infrared light curves (LC), and the spectral evolution. This leaves the description of the burning front and the structure of the white dwarf as the only free parameters. The explosions are calculated using  one-dimensional Lagrangian codes including nuclear networks. Subsequently, optical and IR-LCs are constructed. Detailed NLTE-spectra are computed for several instants of time using the density, chemical and luminosity structure resulting from the LCs. The general methods and critical tests are presented (sect. 2).  Different models for the thermonuclear explosion are discussed including detonations, deflagrations, delayed detonations, pulsating delayed detonations (PDD) and helium detonations (sect.3). Comparisons between theoretical and observed LCs and spectra provide an insight into details of the explosion and nature of the progenitor stars (sect. 4 \\& 5). We try to answer several related questions. Are subluminous SNe~Ia a group different from `normal' SN~Ia (sect. 5)? Can we understand observed properties of the LCs and spectra (sect. 4)? What do we learn about the progenitor evolution and its metallicity (sect. 3, Figs. 4,5)? Do successful SN~Ia models depend on the type of the host galaxy (Table 2)?  Using both the spectral and LC  information, theoretical models allow for a determination of the Hubble constant independent from `local' distance indicators such as $\\delta $ Cephei stars. $H_o$ is found to be $67\\pm~9~km~s^{-1}Mpc^{-1}$ and, from SN1988U, $q_o$ equals $0.7 \\pm 1.$ within 95 \\% confidence levels. ",
+        "introduction": "During the last few years it became evident that Type Ia supernovae are a less homogeneous class than previously believed (e.g. Barbon et al., 1990, Pskovskii 1970,       Phillips et al. 1987). In particular, the observation of subluminous SN1991bg (Fillipenko et al. 1992, Leibundgut et al. 1993) raised questions on the SN-rate   and whether we have missed a huge subgroup. The possible impact on our understanding of supernovae statistics and, consequently, the chemical evolution of galaxies must be noted.  It is generally accepted that SNe~Ia are thermonuclear explosions of carbon-oxygen white dwarfs (WD) (Hoyle \\& Fowler 1960). However, details of the scenario are still under debate. For discussions of various theoretical aspects see e.g. Wheeler \\& Harkness (1990), Woosley \\& Weaver (1994), Canal (1995), Nomoto et al. (1995), Wheeler et al. (1995).  What we observe as a supernovae event is not the explosion itself but the light emitted from a rapidly expanding envelope produced by the stellar explosion. As the photosphere recedes, deeper layers of the ejecta become visible. A detailed analysis of the LCs and spectra gives us the opportunity to determine the density, velocity and composition structure of the ejecta. A successful application of these constraints, however, requires both accurate early LC and spectral observations and detailed theoretical  models.  While SN~Ia LC data have enormously increased, until recently (Harkness 1991; H\\\"oflich et al. 1991), models are often hampered by inadequate physical assumptions like a constant opacity, crude gamma-ray deposition schemes and a simplified treatment of the ionization balance, neglect of line scattering and by the use of the diffusion approximation (e.g. Livne \\& Arnett 1995). Based on our detailed models which overcome the approximations just mentioned, we have investigated the validity and influence of the physical assumptions made in LC and spectral calculations and  estimated the accuracy of our models. The importance of a consistent treatment of the explosion mechanism, LCs and spectra became evident. Some of the tests and the basic outline of our approach is given in section 2. In section 3, the basic explosion mechanism are discussed. After investigating the general properties of our LCs and spectra, individual comparisons with observations are presented. Finally, we address the question on $H_o$ and $q_o$. ",
+        "conclusions": ""
+    },
+    "9602/nucl-th9602009_arXiv.txt": {
+        "abstract": "Neutrino-nucleus cross sections for the detection of atmospheric neutrinos are calculated in relativistic impulse and random phase approximations.  Pion production is estimated via inclusive $\\Delta$- hole nuclear excitations.  The number of pion events can confuse the identification of the neutrino flavor. As a further source of pions we calculate coherent pion production (where the nucleus remains in the ground state).  We examine how these nuclear structure effects influence atmospheric  neutrino experiments and study possible improvements in the present detector simulations. ",
+        "introduction": "\\label{sec:cint}  Aside from the puzzling solar neutrino data there exist extensive measurements of atmospheric neutrinos, which are produced in cosmic ray interactions with the earth's atmosphere. Cosmic rays, mostly protons or $\\alpha$ particles, hitting the atmosphere produce pions, which undergo the following decays: \\begin{eqnarray} \\pi^{\\pm}&\\rightarrow &\\mu^{\\pm}+\\nu_{\\mu}\\ (\\bar{\\nu}_{\\mu}) \\nonumber\\\\ & &\\mu^{\\pm}\\rightarrow e^{\\pm}+\\nu_{e}\\ (\\bar{\\nu}_{e})+ \\bar{\\nu}_{\\mu}\\ (\\nu_{\\mu})\\nonumber\\ . \\end{eqnarray} From this, one expects the neutrino flavor ratio, \\begin{eqnarray} r= {\\nu_\\mu +{\\bar \\nu_\\mu} \\over \\nu_e +{\\bar \\nu_e}}\\ , \\label{rat} \\end{eqnarray} to be about 2.  However, two existing experiments, IMB~\\cite{imb} and Kamiokande~\\cite{kamiokande}, consistently measured a ratio of about 1 by studying charged-current neutrino-nucleus reactions [($\\nu_e$, $e$) and ($\\nu_\\mu$, $\\mu$)] in large underground water detectors.  This could mean either there is a depletion of muon-type neutrinos or an enhancement of electron-type neutrinos. While the separate $\\nu_\\mu$ and $\\nu_e$ neutrino fluxes differ in various calculations, their ratio, however, seems to be largely model-independent~\\cite{eng}, and the theoretical results are in clear contradiction with data. This discrepancy is known as the ``atmospheric neutrino anomaly''.  One explanation of this result could be the existence of neutrino oscillations of the muon neutrino into some other neutrino species~\\cite{solar} which would reduce the $\\nu_\\mu$ flux. However, before drawing this conclusion, one should also investigate more conventional sources for uncertainties in the measured flux ratio, which originate from nuclear physics and  detector specifics. One crucial point is the experimental separation of electron and muon events. Most  data  are measured using water detectors. In a charged-current weak interaction with an $^{16}$O nucleus, the incoming neutrino generates an outgoing lepton which is detected through its \\v{C}erenkov radiation. Electrons and muons are distinguished by the characteristics of their tracks. Electrons produce a showering ``fuzzy'' track, whereas muons are identified by their nonshowering track.   Pions constitute a major background in this way of determining the neutrino flavors. In the experimental analysis, two- and more ring events, which signal the production of several charged particles, are rejected. However, a charged pion can be confused with a muon event if the lepton energy is below \\v{C}erenkov radiation threshold, and a neutral pion, with one decay photon missing, can be counted as an electron event.  Therefore, a reasonable determination of the flavor ratio can only be achieved when the pion events are properly taken into account.   One major source of pions is the decay of a $\\Delta$ produced in the neutrino scattering. The delta can subsequently decay into a pion and a nucleon. An estimate of the number of pion events is only possible when the $\\Delta$-h response is properly calculated. Indeed, Monte Carlo simulation of the Kamiokande experiment~\\cite{monte} estimated the number of pion events based on a rather complicated nonrelativistic model~\\cite{fogli}. However, atmospheric neutrinos have a broad energy spectrum ranging from a few hundred MeV to several GeV. It is therefore crucial to have a relativistic formalism that allows the calculation of the cross section for arbitrary kinematics.  As another source of pions, we consider coherent pion production and its role in the detection of atmospheric neutrinos. Coherent pions are produced in both charged- and neutral-current neutrino-nucleus interactions which leave the nucleus in its ground state. Traditionally, the coherent $\\pi^0$'s are distinguished from the resonantly produced pions due to their strongly forward peaked angular distribution.  However, in atmospheric neutrino experiments, due to a lack of directional information of the incoming neutrinos, it is not possible to distinguish the coherent $\\pi^0$ events from incoherent pions.  In the case of coherent charged pions the forward angle dominance of the cross section leads to a small opening angle between two charged particles ($\\mu$ or $e$ and charged pion) in the final state, and these two particles could be detected as one isolated electromagnetic shower.  In contrast to incoherent charged pions whose \\v{C}erenkov tracks are usually identified as muon events, the special characteristics of coherent pions could lead to an increase of identified $\\nu_e$ events and thus to underestimating  the $\\nu_\\mu$ to $\\nu_e$ ratio.  This paper is organized as follows.  In Sec.~\\ref{sec:incl} we present the formalism for neutrino-nucleus cross section including p-h and $\\Delta$-h nuclear excitations within the framework of a relativistic mean-field model of the nucleus. Sec.~\\ref{sec:coh_for} contains the formalism for coherent charged and neutral pion production. The resulting cross sections and production rates are shown in Sec.~\\ref{sec:result}. We close with a discussion  and outlook in Sec.~\\ref{sec:con}. ",
+        "conclusions": "\\label{sec:con}  In this work, we have calculated neutrino-nucleus cross sections using a relativistic formalism to investigate possible improvements for the Monte Carlo simulation of Kamiokande experiment. We have found that RPA corrections reduce the charged-current neutrino events by up to 37\\%  for particle-hole excitations. Because of this reduction, the $\\Delta$-h response could be important and is found to give significant corrections to quasi-elastic nucleon knock-out processes.  We have found that the relativistic mean-fields reduce the $\\Delta$-h responses by  10\\% to 24\\%, which correspondingly reduce the pion events in atmospheric neutrino experiments. Combined with the additional channel of the non-pionic decay of the delta in the medium,  our findings suggest possible uncertainties in the present Monte Carlo simulations of detectors and the interpretation of the measured data.  We have also discussed coherent pion production in neutrino-nucleus scattering. We have shown that coherent pions can produce a sizeable background in the atmospheric neutrino measurements. As neutral pions and collinearly outgoing charged leptons and pions may mimic electron-type events in the detector, this effect can produce uncertainties in determining the flavor ratio of the incoming neutrinos. Coherent pions may provide important systematic errors for the recent Kamiokande experiment in the multi-GeV energy region~\\cite{mult-gev}. In order to get a clearer quantitative estimate of the influence of coherent pions, improved calculations should be done and these should be included in detector simulations of atmospheric neutrino experiments"
+    },
+    "9602/astro-ph9602091_arXiv.txt": {
+        "abstract": "The correlation function of galaxy clusters has frequently been used as a test of cosmological models. A number of assumptions are implicit in the comparison of theoretical expectations to data. Here we use an ensemble of ten large N-body simulations of the standard cold dark matter cosmology to investigate how cluster selection criteria and other uncertain factors influence the cluster correlation function. Our study is restricted to the idealized case where clusters are identified in the three dimensional mass distribution of the simulations. We consider the effects of varying the definition of a cluster, the mean number density (or equivalently the threshold richness or luminosity) in a catalogue, and the assumed normalization of the cosmological model; we also examine the importance of redshift space distortions. We implement five different group-finding algorithms and construct cluster catalogues defined by mass, velocity dispersion or a measure of X-ray luminosity. We find that different cluster catalogues yield correlation functions which can differ from one another by substantially more than the statistical errors in any one determination. For example, at a fixed number density of clusters, the characteristic clustering length can vary by up to a factor of $\\sim 1.5$, depending on the precise procedure employed to identify and select clusters. For a given cluster selection criteria, the correlation length typically varies by $\\sim 20\\%$ in catalogues spanning the range of intercluster separations covered by the APM and Abell (richness class $\\gsim 1$) catalogues. Distortions produced by peculiar velocities in redshift space enhance the correlation function at large separations and lead to a larger clustering length in redshift space than in real space. The sensitivity of the cluster correlation function to various uncertain model assumptions substantially weakens previous conclusions based on the comparison of model predictions with real data. For example, some of our standard cold dark matter cluster catalogues agree better with published cluster clustering data (particularly on small and intermediate scales) than catalogues constructed from similar simulations by Bahcall \\& Cen and Croft \\& Efstathiou.  Detailed modelling of cluster selection procedures including, for example, the effects of selecting from projected {\\it galaxy} catalogues is required before the cluster correlation function can be regarded as a high precision constraint on cosmological models. ",
+        "introduction": "The clustering strength of rich galaxy clusters has long been used as a constraint on models of large scale structure. The two-point correlation function, $\\xcc$, was first estimated for a sample of about 100 rich Abell clusters by Bahcall \\& Soneira (1983) and by Klypin \\& Kopylov (1983) who noted that clusters have a larger clustering amplitude than galaxies. This difference has a natural explanation in theories in which large scale structure grows by gravitational amplification of small fluctuations in an initially Gaussian density field. In such theories, collapsed objects form near peaks of the initial density field and a clustering pattern which depends on the height of the peak is imprinted at the epoch of formation (Kaiser 1984, Barnes \\etal 1985).  Although the statistics of rare peaks provide an appealing explanation for the different clustering strength of galaxies and clusters, it soon became apparent that the early estimates of $\\xcc$ were inconsistent with the predictions of the cold dark matter model, the paradigm of gravitational clustering theories (Davis \\etal 1985). Using N-body simulations of this model, White \\etal (1987) calculated a cluster clustering length, $r_0\\simeq 11$\\hmpc ($r_0$ is defined as the separation at which $\\xcc=1$), whereas Bahcall \\& Soneira (1983) had obtained $r_0 \\sim 25$\\hmpc for Abell clusters of richness class $R\\ge 1$. Bahcall \\& Soneira's estimate helped to motivate an alternative explanation for the formation of structure, based on fluctuations seeded by cosmic strings (Turok 1983, Turok \\& Brandenberger 1986).  Sutherland (1988) pointed out that the apparent clustering amplitude of rich clusters would be artificially enhanced if intrinsically poor clusters which happened to lie near the line-of-sight to a rich cluster -- and thus appear rich in projection -- were included in a rich cluster sample. A signature of this effect is an anisotropy in the correlation function which appears stronger along the line-of-sight than in the perpendicular direction.  Soltan (1988) and Sutherland \\& Efstathiou (1991) showed that such anisotropies were clearly present in Bahcall \\& Soneira's sample and they, as well as Dekel \\etal (1989), argued that correcting for this effect would lower the clustering length of rich Abell clusters to $\\sim 14$\\hmpc. Nevertheless, Postman, Huchra \\& Geller (1992), using a sample of 351 Abell clusters of richness class $R\\ge 0$, supported Bahcall \\& Soneira's original estimate.  Further progress had to await the construction of new cluster catalogues. These finally began to arrive in the early 1990s. Dalton \\etal (1992) and Lumsden \\etal (1992) constructed the first automated cluster catalogues using the positions and magnitudes of galaxies in photographic plates scanned with the APM and Cosmos machines respectively; Lahav \\etal (1993), Romer \\etal (1994) and Nichol, Briel \\& Henry (1994) constructed cluster catalogues using a combination of X-ray and optical data. The ensuing redshift surveys led to new determinations of $\\xcc$. The Oxford group estimated $r_0=12.9 \\pm 1.4$\\hmpc (Dalton \\etal 1992) from a sample of about 200 APM clusters, and $r_0=14.3 \\pm 1.2$\\hmpc from an extended sample of 364 APM clusters (Dalton \\etal 1994); Romer \\etal obtained $r_0=13.7 \\pm 2.3$\\hmpc from an X-ray flux-limited sample of 128 clusters.  Unfortunately the new cluster samples have not fully resolved the debate surrounding $\\xcc$. Bahcall \\& Soneira (1983) argued that the measured values of $r_0$ depend very strongly on cluster richness. Bahcall \\& West (1992) interpreted the discrepancies between different samples as a reflection of their different mean cluster richness, rather than as a result of contamination in Abell's catalogue (see also Peacock \\& West 1992). N-body simulations by Bahcall \\& Cen (1992) seem to support this view, whereas simulations by Croft \\& Efstathiou (1994) suggest that the dependence of $r_0$ on cluster richness is weak. Furthermore, the Oxford group have claimed that the low values of $r_0$ that they obtain for APM cluster catalogues, although much smaller than Bahcall \\& Soneira's (1983) value are still inconsistent with the standard CDM cosmology and favour either CDM models with a low mean density and a non-zero cosmological constant or mixed dark matter models (Dalton \\etal 1992, 1994, Croft \\& Efstathiou 1994). They base this conclusion on Croft \\& Efstathiou's (1994) set of large N-body simulations which, they argue, enable theoretical predictions for $\\xcc$ to be made with better than $10\\%$ accuracy over a wide range of scales.  The work of Bahcall \\& Cen (1992), Croft \\& Efstathiou (1994) and Watanabe \\etal (1994) has highlighted how, as the observational data on $\\xcc$ improves, the need for precise theoretical predictions becomes increasingly important. Making theoretical predictions which are relevant to the interpretation of the data, however, is not straightforward, even for well specified models such as CDM and its variants. In these models the evolution of the {\\it mass} density field on the relevant scales can indeed be predicted quite accurately, particularly through large N-body simulations. The primary difficulty lies in the uncertain identification of clusters in the models with the real clusters found in galaxy surveys. Bahcall \\& Cen, Croft \\& Efstathiou and Watanabe \\etal, like White \\etal (1987), identified ``galaxy clusters'' in their simulations with large mass concentrations in the three-dimensional mass distribution. This is, of course, a very different procedure from that applied to real data where clusters are identified from the {\\it projected } galaxy distribution using relatively complex algorithms. Possible biases in statistics such as $\\xcc$ which might be introduced by this procedure remain largely unexplored.  In this paper, we address the restricted question of how the clustering strength of model clusters identified using the full three-dimensional mass distribution in N-body simulations depends on the details of the cluster finding algorithm and on the way in which the cluster catalogues are constructed. For definitiveness, we consider only one specific cosmological model, the standard CDM model. We find that even in this idealized case, a wide range of values of $r_0$ is obtained from samples selected and analyzed in ways which are in principle equally valid approximations to the real situation. We therefore conclude that much more detailed modelling is required before the present data can be used as a high precision test of currently popular cosmological models.  In the following section we describe our simulations and methods for constructing cluster catalogues. In Section 3 we present estimates of the correlation function for these catalogues. In Section~4 we compare our results with those obtained in previous related studies. We discuss our findings and summarize our conclusions in Section~5. ",
+        "conclusions": "Bahcall \\& Cen (1992) and Croft \\& Efstathiou (1994) compared their simulation results to real data and in both cases concluded that the correlation function of clusters is incompatible with the standard CDM model, but is consistent with a low density CDM model with $\\Omega h \\simeq 0.2$. At first sight this consensus seems rather surprising since, as Fig.~8 shows, the predicted cluster correlation functions in these two studies disagree. The explanation is simply that the comparison in each case was made against different datasets. Croft \\& Efstathiou compared their models to the APM cluster catalogue of Dalton \\etal (1992) whereas Bahcall \\& Cen compared theirs to this catalogue and to Abell's catalogue as well. As may be seen from Fig.~3 of Bahcall \\& Cen (1992), their cluster correlation function for the low density model does not agree particularly well with the APM data and, as can be seen from Fig.~4 of Croft \\& Efstathiou (1994), their low density model strongly disagrees with the Abell cluster data. Thus, the two studies were able to arrive at the same conclusion because they compared different theoretical predictions for the same cosmological models against datasets that exhibit different clustering properties.  We have found here that even in the idealized case where clusters are identified in the three-dimensional mass distribution of a simulation, significantly different outcomes for the cluster correlation function are possible depending on how exactly the clusters are defined and on how the data are analyzed. It is unclear which, if any, of the various possible definitions of clusters in the simulations is appropriate for a comparison with the real data. This difficulty is particularly severe in the case of optical catalogues since the identification of clusters in the projected galaxy distribution is very different from the identification of clusters in the three-dimensional mass distribution of a simulation. Biases in $\\xcc$ arising from projection effects have been shown to be present in Abell's catalogue (Sutherland 1988; Efstathiou \\& Sutherland 1991; Dekel \\etal 1989). The lack of large anisotropies in $\\xcc$ for APM clusters suggests that this catalogue is largely unaffected by biases of this kind, but this important feature by itself does not remove the ambiguity regarding the identification of galaxy clusters in real catalogues with mass clusters in simulations. Identifying cluster populations in both by matching the spatial number density is not a unique procedure since, as we have shown, even at a fixed number density, the cluster correlation function depends, for example, on the statistic used to rank the clusters. In practice, it seems likely that even larger uncertainties will be introduced by the difficulty of determining the richness of clusters in projection and by the associated uncertainties in the estimation of their spatial number density. X-ray selected clusters provide, in principle, cleaner observational samples, but even in this case the comparison with theoretical models is restricted by the lack of reliable predictions.  Figs.~10 and 11 illustrate how some of the uncertainties we have mentioned can affect the confidence with which a specific model is constrained by measurements of $\\xcc$. Here we compare predictions based on the standard cold dark matter model -- the model which Bahcall \\& Cen (1992), Croft \\& Efstathiou (1994) and Dalton \\etal (1994) claim to be strongly excluded by the cluster correlation function data -- with various observational determinations of $\\xcc$. In Fig.~10 we compare our estimates for SOX ``X-ray selected'' clusters with the ROSAT data of Romer \\etal (1994). This catalogue is not volume-limited and thus contains clusters with a range of intrinsic X-ray luminosities.  The simulations, however, indicate that the variation of the correlation length of SOX ``X-ray selected'' clusters with cluster X-ray luminosity (or richness) is small compared with the uncertainties in the data. Within these large statistical errors, the agreement is good except, perhaps, on the largest scales where the observed signal is small and could be affected by systematic uncertainties in the mean number density of clusters. The size of the discrepancy on large scales may be better appreciated in the linear-log plot in the lower panel of this figure.  In Fig.~11, we compare our estimates for SO$_v$ clusters (in redshift space) with data for ``optical'' clusters from the APM (Efstathiou \\etal 1992, Dalton \\etal 1992, 1994) and EDCC (Nichol \\etal 1992) catalogues. The APM sample shown by filled triangles has a mean intercluster separation, $d_{\\rm c}\\simeq 31$\\hmpc, and should be compared with the dashed line which shows our model predictions for the same mean intercluster separation.  The APM sample shown by open triangles has $d_{\\rm c}\\simeq 45$\\hmpc and the EDCC sample has $d_{\\rm c}\\simeq 50$\\hmpc. These should be compared with the solid line which shows our model results for $d_{\\rm c}=50$\\hmpc. On intermediate scales only the denser sample (which has the smallest error bars) is inconsistent with the model, but the discrepancy is quite small and certainly much smaller than the discrepancy found by Dalton \\etal (1994) for the same model (cf their Fig.~4). The reason for this difference is simply that the group-finder applied to the simulations by Dalton \\etal happens to give one of the lowest correlation functions of all the group-finders that we have explored in this paper. On large scales there is an indication that the data are more clustered than the models and, again, the linear-log plot clearly shows that this discrepancy is small.  To summarize, large N-body simulations allow very precise estimates of the cluster correlation function once a specific prescription for identifying clusters is adopted. For a given cosmological model, the statistical uncertainties in these predictions are small compared to the observational errors for existing cluster samples. Unfortunately, they are also small compared to the systematic variations exhibited by cluster samples constructed from the simulations by making different assumptions.  Our analysis shows that the exact form and amplitude of the correlation function of clusters identified in the mass distribution of N-body simulations depends on various factors. In rough order of importance these include: (i) the group-finding algorithm and the statistic used to rank clusters in a catalogue (eg. mass, velocity dispersion, ``X-ray'' luminosity, etc); (ii) the mean density of clusters in the catalogue; (iii) whether clustering is measured in real or in redshift space; and (iv) the assumed value of $\\sigma_8$. These various factors can produce large variations in the resulting clustering length. For example, in the range most relevant to observational data, $30\\le d_{\\rm c} \\le 60 \\Mpc$, our ``X-ray selected\" catalogues in a model with $\\sigma_8=0.5$ have real space clustering lengths varying between 7 and 10\\hmpc, whereas catalogues selected by velocity dispersion in a model with $\\sigma_8=0.63$ have redshift space correlation lengths varying from $10$ to $13 \\Mpc$.  Of the four complicating factors listed above, only the first two refer directly to cluster catalogues. The third one should be straightforward to eliminate but it has sometimes been ignored in comparisons of model predictions with data (eg. Bahcall \\& Cen 1992). Similarly, the value of $\\sigma_8$ appropriate to a given cosmological model is usually fixed from other considerations such as the amplitude of fluctuations in the temperature of the cosmic microwave background or the abundance of galaxy clusters (White, Efstathiou \\& Frenk 1993). There are, however, uncertainties associated with this procedure arising, for example, from possible contamination of the microwave background signal by tensor modes or uncertainties in the masses of clusters.  The resolution of the cluster clustering debate will require further observational and theoretical work. From the observational point of view, progress will come from the analysis of large homogeneous samples of clusters selected entirely from X-ray data or from large redshift surveys such as the forthcoming SDSS and 2df galaxy surveys. From the theoretical point of view, it will be necessary to model in detail the selection procedures employed by observers. Artificial catalogues constructed from cosmological simulations are a valuable aid, but several complications need to be borne in mind. For example, to model the selection of APM clusters, it is necessary to simulate the entire APM galaxy survey and this, in turn, requires modelling the uncertain connection between the distribution of dark matter and the formation sites of galaxies. Modelling cluster catalogues constructed from X-ray data is a simpler problem theoretically since it bypasses the complications associated with galaxy formation. Nevertheless, it requires a better understanding of the mechanisms that determine the total cluster X-ray luminosity than is available at present. It is our view that until the theoretical predictions are brought onto the observational plane, measurements of the cluster correlation function cannot confidently be used to choose amongst competing cosmological models."
+    },
+    "9602/astro-ph9602058_arXiv.txt": {
+        "abstract": "The globular cluster population in M87 has decreased measurably through dynamical evolution caused by relaxation, binary heating and time-dependent tidal perturbation.  For fundamental plane ellipticals in general, cluster populations evolve more rapidly in smaller galaxies because of the higher mass density.  A simple evolutionary model reproduces the observed trend in specific frequency with luminosity for an initially constant relationship.  Fits of theoretically evolved populations to M87 cluster data from McLaughlin et al. (1994) show the following: 1) dynamical effects drive evolution in the initial mass and space distributions and can account for the large core in the spatial profile as well as producing radial-dependence in the mass spectrum; 2) evolution reduces $S_N$ by 50\\% within $16\\kpc$ and 35\\% within $50\\kpc$, implying that $S_N$ was initially 26 in this region.  We estimate that 15\\% of the `missing' clusters lie below the detection threshold with mass less than $10^5 M_{\\odot}$. ",
+        "introduction": "\\label{sec:intro}  Observations of some giant elliptical galaxies reveal globular cluster systems which appear more extended than the host (Harris 1991).  A particularly well-documented example belongs to M87 with a core radius of 7 arc sec and a cluster system with a core radius of 1 arc min (McLaughlin 1995).  However, because a cluster population evolves dynamically due to both internal and external processes, the currently observed population almost certainly differs from the primordial one, complicating the interpretation.  Researchers have attempted to explain the extended core of the M87 cluster distribution as the evolved remnant of an initial profile which more closely resembled the light.  However, neither dynamical friction nor shocking by a compact nucleus can fully account for this feature.  Lauer \\& Kormendy (1986) found that a dynamical friction induced inflow can broaden an initially peaked spatial distribution but not at the observed scales.  Ostriker, Binney \\& Saha (1989, hereafter OBS) subsequently determined that nuclear tidal disruption is viable only if clusters formed exclusively on box orbits.  Another potential mechanism is cluster evaporation through dynamical evolution.  Recent work in this area demonstrates that evaporative mass loss driven by relaxation and heating due to a time-varying tidal field can lead to strong evolution of the Milky Way cluster population in a Hubble time (Weinberg 1994, Murali \\& Weinberg 1996, hereafter MW).  In this paper, we examine these influences on cluster evolution in the dense inner regions of M87 and find that they produce the observed flattened profile from a peaked initial distribution over a wide range of initial conditions.  Direct estimates of initial conditions using dynamically evolved parametric models of the spatial distribution and cluster mass function indicate that roughly 35\\% of the initial population dissolves or evolves below the detection threshold leaving the large core as a result. Furthermore, the decay in the size of the cluster population corresponds to a decrease in the specific frequency of globular clusters, $S_N$, which denotes the number of clusters per unit galaxian luminosity with $L$ measured in units of $M_v=-15$.  The high values of $S_N$ found in giant ellipticals have become a key point in galaxy formation arguments and suggest, for example, that the cluster system formed along with M87 (e.g. Harris 1991; van den Bergh 1995).  Here we show that cluster systems decay more rapidly in less luminous fundamental plane ellipticals; this leaves larger values of $S_N$ in luminous galaxies at the present epoch even if all ellipticals begin with equal $S_N$.  Our results thus provide at least a partial explanation for the observed trend of $S_N$ with $L$.  The plan of the paper is as follows.  We summarize our choices for the cluster population and the mass model for M87 in \\S\\ref{sec:pop}.  The assumptions and method for dynamically evolving the population is presented in \\S\\ref{sec:evol}.  The main results, the statistical comparison of the observed clusters to the theoretical models, is presented in \\S\\ref{sec:results}.  This includes an exploration of the evolutionary trends, best fit spatial profiles and mass functions, and an inference of the primordial population.  In \\S\\ref{sec:disc}, we discuss the discuss the importance of the fundamental plane properties on the observed relation between specific frequency and luminosity A summary is given in \\S\\ref{sec:summary}. ",
+        "conclusions": "\\label{sec:disc}  Our conclusions ignore the possibility that recent merger and accretion events have have strongly contaminated the initial cluster population.  Merging of gas-rich galaxies is expected to produce clusters with strong central concentration (Zepf \\& Ashman 1993; Mihos \\& Hernquist 1994), but the large core of the distribution itself argues against any recent merger which has produced significant numbers of young clusters.  Thus either cluster-producing mergers in M87 have occurred long in the past or not at all.  The recent addition of clusters through satellite accretion is also unlikely to account for a significant fraction of the observed population.  For example, accretion of Milky Way-type spirals can only account for about 25\\% of the observed population, assuming that the $4 L_*$ luminosity of M87 ($R<40 \\kpc$) comes entirely from satellites (Lauer 1988).  Assuming that material is stripped when the mean density of the primary exceeds that in the satellite, accretion will deposit clusters in regions of mean density similar to that in the original environment of the accreted galaxy.  The clusters, then, will remain roughly tidally truncated.  Their new orbits depend on the orbit of the dissolving satellite, but otherwise evolution should be similar and the accreted population may appear coeval regardless of the time of accretion.  The high specific frequency of globular clusters in M87 appears to indicate the exceptional conditions governing the formation and evolution of cD galaxies relative to other ellipticals.  The observation that specific frequency increases with galaxy luminosity suggests that galaxy formation was not an intrinsically hierarchical and homologous process (e.g. Santiago \\& Djorgovski 1993).  However, the current results indicate that density differences between galaxies will lead to differences in the evolution of intrinsic cluster populations.  Since M87 lies approximately on the fundamental plane for elliptical galaxies, we can investigate differences in environment-driven cluster evolution by scaling our results to other ellipticals.  We assume that the initial profile of the cluster population derived above, when scaled homologously, describes the initial population in any elliptical.  The dynamical time scale for a tidally truncated cluster with constant mass and spatial profile is then determined by the mean density of the galaxy: $\\tau\\sim M^{-1/2}R^{3/2}$.  This yields $\\tau\\sim M^{\\beta}$, where $\\beta=0.4$ for the fundamental plane scaling given in Faber et al. (1987), $\\beta=0.26$ for the Djorgovski \\& Davis (1987) results, and $\\beta=0.6$ for a more recent set of parameters from Faber (1995; see also Pahre et al. 1995).  This relation implies that clusters evolve and are depleted more rapidly in smaller, low luminosity ellipticals than in massive, high luminosity ellipticals.  To demonstrate the importance of this effect, we combine the approximate exponential decay rate of the cluster population found in \\S\\ref{sec:evolpop} with this scaling relation.  This gives an expression for the number of clusters remaining in an initially coeval population belonging to a galaxy of luminosity $L$ at time $T$:  \\begin{eqnarray} N_{cl}(L,T)=N_0(L)e^{-T/\\tau_0(M_{M87}/M)^{\\beta}}\\nonumber\\\\ =N_0e^{-T/\\tau_0(L_{M87}/L)^{1.24\\beta}}, \\end{eqnarray} where $\\tau_0\\approx 40 \\Gyr$ for M87, $N_0(L)$ is the initial distribution of cluster population sizes as a function of galaxy luminosity and $M/L\\sim L^{1.24}$ throughout.  Using a power-law distribution $N_0(L)=N_{*}(L/L_{M87})^\\gamma$, we compare the model curve $N_{cl}(L,T)$ to an observed sample of cluster populations in galaxies (Harris 1996). The resulting models show qualitative agreement with the data (Figure \\ref{fig:fpe.1}), falling off at low luminosity more rapidly with increasing $\\beta$ due to the more rapid rate of evolution.  From this we conclude that the observed discrepancies in specific frequency which correlate with galaxy luminosity were smaller in the past.  This is reflected in the smaller exponent $\\gamma$ in $N_{cl}\\propto L^{\\gamma}$ predicted initially compared to the present-day estimate.  \\begin{figure} \\epsfxsize=20pc \\epsfbox[12 138 600 726]{figure8.ps} \\caption{Evolved population size as a function of galaxy luminosity compared to data from Harris (1996) for fundamental plane parameter $\\beta=0.4$ (solid) and $\\beta=0.65$ (dotted) .  The present-day, unevolved fit (dashed) has $\\log N_*=3.89\\pm.05, \\gamma=1.46\\pm.13$.  For $\\beta=0.65$, the initial value of $\\gamma$ is consistent with specific frequency which is independent of luminosity} \\label{fig:fpe.1} \\end{figure}  To conclude, we suggest that cluster evolution may account at least in part for the specific frequency problem.  Because clusters evolve more rapidly in dense environments and because ellipticals are typically denser at low luminosity that at high luminosity, cluster populations diminish more rapidly in low luminosity galaxies. Specific frequencies will therefore correlate more strongly with luminosity at later times.    We have investigated the evolution of the M87 globular cluster system.  The results of this study have enabled us to examine the broader question of population evolution in fundamental plane elliptical galaxies.  Our main conclusions follow:  \\begin{enumerate}  \\item The loss of globular clusters through relaxation and tidally-induced evaporation accounts for the large core and shallow profile in cluster number distribution compared to the light distribution in M87.  \\item Evolution produces a radial dependence in the present-day mass spectrum of clusters such that higher mass clusters predominate in the inner regions.  The models also indicate an initial radial dependence in the mass spectrum.  \\item Likelihood ratio tests reject an initial power law in favor of an initial Gaussian but do not rule out the possibility of an initial core in the profile.  \\item The best-fit model for M87 has an initial population of $7.25\\times10^3$ clusters with projected radius $R<50\\kpc$, about 60\\% more than is currently observed.  Roughly 14\\% of the initial population are now objects of slightly less than $10^5 M_{\\odot}$; dynamical evolution can strongly modify the specific frequency of globular clusters.  \\item Scaling the calculations to fundamental plane elliptical galaxies indicates that cluster evolution in the differing environments qualitatively accounts for the trend in observed population number versus galaxy luminosity.  Smaller galaxies tend to have high densities and thus more rapid evolutionary time scales, so their cluster populations tend to diminish more rapidly.  \\end{enumerate}  Some of these inferences, especially (iii), may be biased by the lack of detailed data in the inner galaxy.  Point source data for $r<1'$ will lead to stronger constraints on the size of the initial core and a deeper survey will pick up low luminosity objects and further constrain the luminosity function."
+    },
+    "9602/hep-ph9602384_arXiv.txt": {
+        "abstract": "We present the  evolution equation describing MSW conversion,  derived in the framework of the Schr\\\"odinger approach, in the presence of  matter density fluctuations. Then we analyse the effect of such fluctuations in the MSW scenario as a solution to the solar neutrino problem. It is shown that the non-adiabatic MSW parameter region is rather stable (especially in $\\delta m^2$) for matter density noise at the few percent level. We also discuss the possibility to probe solar matter density fluctuations at the future Borexino experiment. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602022_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "\\index{jets}\\index{outflows} Narrow, high speed outflows (jets) and less well collimated `bipolar' outflows are observed from very different cosmical objects, ranging from protostars in the solar neighborhood, to galactic X-ray binaries, to the nuclei of active galaxies. The magnetic acceleration mechanism for outflows from accretion disks has gained significant popularity as an explanation for each of these forms of outflow. The model can account for high speeds (for example, Lorentz factors of 10 or higher in AGN and galactic black hole binaries), high degrees of collimation, and large momentum fluxes. Though other processes can also, to varying degrees, account for these properties (e.g.\\ Blandford 1993), the magnetic model combines them in a natural way. While some processes are very different in protostellar disks and the AGN or X-ray binary disks, in particular those that produce the observed radiation, the physics of the magnetic acceleration model is to a large degree independent of these. Progress made by development of the model for explaining observations in one area is therefore likely to have impact for the interpretation of other outflows as well.  In spite of the rather general applicability of the magnetic mechanism, it is unlikely that it is involved in all cases. There are a number of bipolar looking objects in the sky where the explanation may well be purely hydrodynamical, such as the outflows in $\\eta$ Carinae, and in particular the planetary nebulae (Icke et al. 1992).  The protostellar outflows play a key role in supporting the magnet\\-ic/centri\\-fugal model. The momentum flux in these objects can be measured from flow speeds and inferred mass densities, and in many cases turns out to be much larger than can be accounted for by the nearest competing mechanism, radiation pressure from the central star (for a review see K\\\"onigl and Ruden, 1993). The protostellar outflows are also the ones in which observations are most likely, in the near future, to reveal their inner workings by direct imaging, as shown in \\tab{hsta1}. See the contributions elsewhere in this volume for more about this subject.  \\begin{table}[bp] \\caption {\\label{hsta1}Angular size of the accelerating region (assumed to be 100 times the typical size $r_0$ of the inner disk) for jet-producing objects.} \\begin{tabular}{lccc} & inner disk & distance & angular scale ($\"$) \\\\ &  $r_0$ & D & $100\\,r_0/D$ \\\\ nearby protostar & $3 R_\\odot$ & 500 pc & 0.003 \\\\ nearby active galactic nucleus& 100 AU & 10 Mpc & 0.001 \\\\ galactic black hole candidate& 100 km & 2 kpc & $3\\,10^{-8}$ \\end{tabular} \\end{table} ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602065_arXiv.txt": {
+        "abstract": "Assuming  that MSW neutrino oscillations occur and ignoring all solar physics except for  the constraint that nuclear fusion produces the solar luminosity,  we show that new  solar neutrino experiments are required to rule out empirically the hypothesis that the sun shines via the CNO cycle. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602123_arXiv.txt": {
+        "abstract": "Cross-filter K corrections for a sample of ``normal'' Type Ia supernovae (SNe) have been calculated for a range of epochs. With appropriate filter choices, the combined statistical and systematic K correction dispersion of the full sample lies within 0.05 mag for redshifts $z<0.7$.  This narrow dispersion of the calculated K correction allows the Type Ia to be used as a cosmological probe.  We use the K corrections with observations of seven SNe at redshifts $0.3 < z <0.5$ to bound the possible difference between the locally measured Hubble constant ($H_L$) and the true cosmological Hubble constant ($H_0$). ",
+        "introduction": "The homogeneity and brightness of Type Ia SN peak magnitudes have long made it a popular standard candle (Branch and Tammann 1992\\nocite{br:araa}). Recent observations, as discussed in this meeting, now indicate that Type Ia's appear to form a family, rather than a set of identical objects.  However, magnitude corrections based on independent observables can make the Type Ia a calibrated candle. Correlations between peak magnitude and light curve shape (Hamuy et al. 1995\\nocite{ha:hubble}; Riess, Press, \\& Kirshner 1995\\nocite{re:lcs}) or spectral features (Nugent et al. 1995\\nocite{nu:1995}), have made it possible to use the Type Ia as a calibrated candle with $B$ magnitude dispersions of $\\sim 0.2$ mag. Alternatively, the use of Type Ia sub-samples can provide an improved standard candle. Supernovae with high quality data that pass certain color cuts show a low dispersion in $B$ magnitudes of $< 0.28$ mag (Vaughan et al. 1995\\nocite{va:inp}).  In order to use SNe as a tool for cosmology, we use K corrections (Oke \\& Sandage 1968)\\nocite{ok:kcorr} to account for the redshifting of spectra and its effect on nearby and distant flux measurements in different passbands. The sparseness of spectroscopic observations of any individual high-redshift SN requires us to use a statistical approach to K corrections. K corrections must therefore show a small magnitude dispersion to maintain the usefulness of the SN as a standard or calibrated candle.  As of June 1995, the Supernova Cosmology Project had discovered seven SNe lying in the range $0.3<z<0.5$ (Perlmutter et al. 1995; Perlmutter et al. 1996) \\nocite{sn1992bi}.  In a standard Friedmann cosmology, their apparent magnitudes depend on $\\Omega$ and $\\Lambda$ through the luminosity distance. Perturbations on a homogeneous and isotropic universe can cause the locally measured Hubble constant ($H_L$) to differ from the global Hubble constant ($H_0$), as demonstrated by Turner, Cen, and Ostriker (1992)\\nocite{tu:hubble}. In this case if the SN calibrators lie within the local peculiar flow, then the high-z SN apparent magnitudes will additionally depend on $H_L/H_0$. A scenario in which $H_L/H_0 > 1$ has been suggested by Bartlett et al. (1994) \\nocite{ba:h30} and others to reconcile theoretical arguments for a low Hubble constant with the recent observational evidence for a large Hubble constant. We use our seven supernovae, with K corrections, to place bounds on $H_L/H_0$. ",
+        "conclusions": "We conclude that filter-matching K corrections give small enough errors to make SNe useful as cosmological probes. In the future, we will calculate K corrections using more SN spectra and we will search for a relation between light curve shapes and K corrections.  Although we do not see such a correlation for the SNe that we have examined so far, which have only moderate departures from the Leibundgut template lightcurve, such a relation is likely to exist at some level since SN colors appear to be correlated with light curve shape, as discussed at this meeting by Riess et al.(1996) and Suntzeff et al.(1996).  We find that we can rule out significant spatial variation of the Hubble constant for $0 \\le \\Omega_M \\le 2$. We plan to calculate $H_L/H_0$ for cosmologies with a cosmological constant, although positive values of $\\Lambda$ will only strengthen our limit.  This work was supported in part by the National Science Foundation (ADT-88909616) and U.~S. Department of Energy (DE-AC03-76SF000098)."
+    },
+    "9602/astro-ph9602137_arXiv.txt": {
+        "abstract": "The EFAR project is a study of 736 candidate elliptical galaxies in 84 clusters lying in two regions towards Hercules-Corona Borealis and Perseus-Pisces-Cetus at distances $cz$$\\approx$6000--15000\\kms. In this paper (the first of a series) we present an introduction to the EFAR project and describe in detail the selection of the clusters and galaxies in our sample. Fundamental data for the galaxies and clusters are given, including accurate new positions for each galaxy and redshifts for each cluster. The galaxy selection functions are determined using diameters measured from Schmidt sky survey images for 2185 galaxies in the cluster fields. Future papers in this series will present the spectroscopic and photometric observations of this sample, investigate the properties of the fundamental plane for ellipticals, and determine the large-scale peculiar velocity fields in these two regions of the universe. ",
+        "introduction": "This paper is the first in a series reporting the results of the EFAR project studying the properties and peculiar motions of elliptical galaxies and clusters in two volumes of the universe at distances between 6000 and 15000\\kms. The aims of this extensive observational program are: (i)~to study the intrinsic properties of elliptical galaxies in clusters by compiling a large and homogeneous sample with high-quality photometric and spectroscopic data; (ii)~to test possible systematic errors, such as environmental dependence, in existing elliptical galaxy distance estimators; (iii)~to seek improved distance estimators based on a more comprehensive understanding of the properties of ellipticals and how these are affected by the cluster environment; and (iv)~to determine the peculiar velocity field in regions that are dynamically independent of the mass concentrations within 6000\\kms in order to test whether the large-amplitude coherent flows seen locally are typical of bulk motions in the universe.  The EFAR project was conceived in 1986 as a natural progression from the work of the Seven Samurai (7S: Dressler \\etal\\ 1987a; Lynden-Bell \\etal\\ 1988), who studied the peculiar velocity field traced by elliptical galaxies closer than 6000\\kms. The major finding of that work was that the local region of the universe was dominated by large-scale, large-amplitude coherent motions. This result has been substantially confirmed both by further analysis (Faber \\& Burstein 1988; Bertschinger \\etal\\ 1990) and by subsequent observational studies, mostly employing the independent Tully-Fisher distance estimator for spiral galaxies (Aaronson \\etal\\ 1986, 1989; Han \\& Mould 1990; Willick 1990, 1991; Mathewson \\etal\\ 1992; Mould \\etal\\ 1993; Courteau \\etal\\ 1993; Mathewson \\& Ford 1994). Although something is known about the peculiar velocity field within 6000\\kms\\, the nature of the mass concentrations causing the flow remains controversial (see reviews by Bertschinger 1990; Burstein 1990; Dekel 1994; Strauss \\& Willick 1995) and relatively little is known about galaxy motions far away.  Initial comparisons of the velocity field with the predictions of cosmological models (Vittorio \\etal\\ 1986; Bertschinger \\& Juszkiewicz 1988) suggested that the observed motions were difficult to reconcile with the favoured biased CDM models that they considered. The immediate question raised was whether the local volume was indeed typical of other regions of the universe (implying that the standard cosmological models were incorrect) or whether the local motions are merely an unusual statistical fluctuation in the universal velocity field. More recent analyses (Kaiser 1988, 1991; Feldman \\& Watkins 1994; Seljak \\& Bertschinger 1994; Dekel \\etal\\ 1996), suggest that the local motions are consistent with the COBE-normalised standard CDM model. Nonetheless, whether or not the local motions are typical of the universe at large remains an important question. In order to answer this we are measuring the peculiar motions in similarly-large regions at sufficient distances to be dynamically independent of the local volume studied by 7S and most other workers.  During the period it has taken to complete the EFAR observing program, however, a wider variety of questions have arisen. Chief among these has been a more searching inquiry as to the reliability of the $\\Dnsigma$ distance indicator developed by 7S (Dressler \\etal\\ 1987b; Lynden-Bell \\etal\\ 1988) and the physical origin of the fundamental plane (FP) of elliptical galaxies (Djorgovski \\& Davis 1987; Bender \\etal\\ 1992; Saglia \\etal\\ 1993a; J{\\o}rgensen \\etal\\ 1993; Pahre \\etal\\ 1995) of which it is simply a convenient projection.  Various authors have suggested that there may be variations in the FP which correlate with galaxy environment, either directly, through mechanisms such as tidal stripping (Silk 1989), or indirectly via different stellar populations (Gregg 1992, 1995). Such effects could lead to significant systematic errors in any distance estimators based on the FP relations.  Some claims have been made for the detection of such variations (de Carvalho \\& Djorgovski 1992; Guzman \\& Lucey 1993). Other studies, however, find little variation of the FP with environment (Lynden-Bell \\etal\\ 1988; Burstein \\etal\\ 1990; Lucey \\etal\\ 1991), and comparisons of FP and $\\Dnsigma$ distance estimates with those derived from relatively independent and perhaps more accurate estimators, such as the Tully-Fisher and surface brightness fluctuation methods, show good agreement (Jacoby \\etal\\ 1992).  These concerns, and the focus they bring on the formation and evolution of the elliptical galaxy population, have become as important a motivation for this work as the original goal of measuring the peculiar velocity field at large distances. The EFAR project's goal of measuring the peculiar motions of {\\em distant} ellipticals from a {\\em large} and {\\em homogeneous} sample, provides a test of the FP distance estimators that is both {\\em severe} (since systematic errors in peculiar velocities are amplified at large distances) and {\\em fair} (given the difficulties of comparing the FP for differently selected samples observed in different studies). It is worth noting that the 7S dataset is still the largest in the literature for elliptical galaxies, though it was obtained a decade ago and is based on photoelectric aperture photometry and IDS spectroscopy. The EFAR sample is comparable in size to that of 7S and is based largely on CCD imaging and spectroscopy, which confer a number of advantages in the attempt to reduce observational errors.  Even if {\\em systematic} errors prove negligible, the relatively large ($\\sim$20\\%) {\\em random} errors in the $\\Dnsigma$ and Tully-Fisher distance estimators limit exploration of the velocity field beyond about 6000\\kms\\ (the `local' region). The only studies which have attempted to measure the velocity field as far out as 15000\\kms\\ are those of Lauer \\& Postman (1994), using brightest cluster galaxies as distance estimators, and Riess \\etal\\ (1995), using Type Ia supernovae. These sparse, all-sky samples are suitable for measuring the convergence of the Local Group dipole motion to the cosmic microwave background dipole (or lack thereof), but they do not probe the velocity field on scales of tens of Mpc. This requires the use of distance estimators which have greater precision per object and/or apply to clustered objects, allowing denser sampling.  FP-based distance estimators for elliptical galaxies fulfill these criteria. As recent work by J{\\o}rgensen \\etal\\ (1993) has shown, the FP can yield distances with errors as low as 11\\% per galaxy (compared to 17\\% for $\\Dnsigma$ distances based on the same data, and 25\\% for the original $\\Dnsigma$ distances obtained by the 7S). For individual {\\em clusters} it is possible to reduce the distance errors by $\\surd N$, where $N$ is the number of galaxies in each cluster. With 5--15 ellipticals per cluster it is therefore possible in principle to measure individual cluster distances with 3--5\\% precision, corresponding to peculiar velocity errors of 500-750\\kms\\ at a distance of 15000\\kms. Thus with sufficient clusters in a given volume it becomes feasible to measure the peculiar velocity field with enough precision to reliably detect large-scale coherent motions at distances out to 15000\\kms which have amplitudes comparable to those observed in the local volume within 6000\\kms. Provided significant sources of systematic error in the FP distance estimator can be ruled out or corrected for, the potential exists to determine the peculiar velocity fields in distant regions and thus further constrain cosmological models.  Two preparatory papers have discussed the photoelectric photometry and photometric system we use (Colless \\etal\\ 1993) and the methods we apply to correct for seeing in measuring the galaxies' light profiles (Saglia \\etal\\ 1993b), an important effect for galaxies at greater distances. Other aspects of this project and some preliminary results have been reported in Wegner \\etal\\ (1991), Davies \\etal\\ (1993) and Baggley \\etal\\ (1994).  This paper (Paper~I) describes how the galaxy clusters and galaxies belonging to those clusters were selected for this project. Future papers in this series will detail the spectroscopic and photometric data we have obtained, describe the methods used to analyze the luminosity profiles of the galaxies, examine the intrinsic properties of the galaxies and their dependence on environment, derive an optimal distance estimator, and discuss the peculiar motions of the galaxies and clusters and their significance for models of the large-scale structure of the universe.  In \\S2 of this paper we describe the selection of the regions and clusters used in this study, setting them in context with the surrounding large-scale structures. \\S3 gives the procedures used to select candidate elliptical galaxies in each cluster, and gives the master list of basic information on the galaxies in the study. The selection functions for the galaxies in each cluster are quantified in \\S4, and the conclusions we draw from this analysis are given in \\S5. ",
+        "conclusions": "The EFAR project is aimed at measuring the properties and peculiar motions of elliptical galaxies in clusters selected in two regions at distances of 6000--15000\\kms. The primary goals of the project are: first, to study the physical properties of a large sample of elliptical galaxies and derive an optimal distance estimator; and second, to determine the bulk motions in the two selected regions in order to establish whether the large-scale coherent motions seen within 6000\\kms\\ are typical of other regions of the universe, thereby tightening the constraints on cosmological models set by the local velocity field.  There are 84 clusters in our sample, 39 in the region we call Hercules-Corona Borealis (HCB) and 45 on the opposite side of the sky in the region we call Perseus-Pisces-Cetus (PPC). Most of these clusters lie in the redshift range 6000--15000\\kms, with 42 being drawn from the Abell (1958) catalog, 32 from the catalog of Jackson (1982), and a further 10 supplementary clusters found by us on the Sky Survey plates. We give the redshifts for all of these clusters and show where they are located with respect to the superclusters identified by Einasto \\etal\\ (1994). We find that our sample of Abell (1958) clusters in the range 6000-15000\\kms\\ is 92\\% complete in HCB and 87\\% complete in PPC.  In these clusters we have selected 736 candidate elliptical galaxies based on their morphology and apparent size on enlargements of Sky Survey glass copies. Our master list of EFAR galaxies gives the fundamental data, including accurate new positions, for these program galaxies and for 52 calibration galaxies in Coma, Virgo and the nearby field. In order to quantify our selection criteria we have measured visual diameters from the Sky Survey enlargements for all the program galaxies (and for the calibration galaxies in Coma) and for a large number of other galaxies in these clusters---a total of 2185 objects. The visual diameters of the program galaxies turn out to correlate with preliminary estimates of the the photometric sizes established by subsequent CCD imaging. For each cluster we are therefore able to characterise the selection criterion for the sample galaxies in terms of these visual diameters. We find that the samples of galaxies in the program clusters are typically 50\\% complete at 18--20\\arcsec, and drop from 90\\% to 10\\% completeness between 30\\arcsec\\ and 10\\arcsec.  Subsequent papers in this series will report the spectroscopic and photometric observations, estimate distances and peculiar velocities for the galaxies and clusters, and interpret the implied peculiar velocity field."
+    },
+    "9602/astro-ph9602071_arXiv.txt": {
+        "abstract": "We calculate the mean nucleation distance, $d_{nuc}$, in a first order cosmological quark-hadron phase transition. For homogeneous nucleation we find that self-consistent inclusion of curvature energy reduces $d_{nuc}$ to $\\lesssim 2$cm. However, impurities can lead to heterogeneous nucleation with $d_{nuc}$ of several meters, a value high enough to change the outcome of Big Bang nucleosynthesis. ",
+        "introduction": "A number of authors \\cite{witt,hogan,alc1,kunuk} have suggested that inhomogeneities in the baryon number density were produced in the cosmological quark-hadron phase transition and perhaps even persisted to the time of nucleosynthesis, which could alter the abundances of light elements. For this to happen, a few criteria have to be met, though. First of all, only a first order quark-hadron transition seems capable of generating large baryon number fluctuations. Whether the transition is first order is still an unsettled question (see e.g.\\ Ref.\\ \\cite{iwasaki}). Secondly, the mean distance between high and low baryon density regions, denoted {\\it l}, has to be larger than approximately one meter comoving at $ T=100 $MeV \\cite{kunuk}, if inhomogeneous nucleosynthesis results are to differ from the standard homogeneous nucleosynthesis results (otherwise, neutron and proton diffusion will erase the inhomogeneities prior to nucleosynthesis). Other important parameters are the average density contrast between high and low-density regions, the average volume fraction of high-density regions, and the baryon to photon ratio.  We assume that the transition is first order and focus in this paper on determination of the most important of the parameters, the mean distance $ l $.  A first order quark-hadron transition in the expanding Universe can in general terms be described as follows. As the color deconfined quark-gluon plasma cools below the critical temperature, $ T_c \\approx 100$MeV, it becomes energetically favorable to form color confined hadrons (primarily pions and a tiny amount of neutrons and protons due to the conserved net baryon number). However, the new phase does not show up immediately. As is characteristic for a first order phase transition, some supercooling is needed to overcome the energy expense of forming the surface of bubbles of the new hadron phase. When a hadron bubble is nucleated, latent heat is released, and a spherical shock wave expands into the surrounding supercooled quark-gluon plasma. This reheats the plasma to the critical temperature, preventing further nucleation in a region passed by one or more shock fronts \\cite{kuka} (for the range of parameters used in the present investigation, bubble growth is described by deflagrations, where a shock front precedes the actual phase transition front. However, if the shock fronts are very weak the quark-gluon plasma may not be reheated enough to sufficiently prevent new bubble nucleations in such regions \\cite{kuka}. This reduces the mean distance. We will neglect this possibility, because of the very small amounts of supercooling we deal with here). The nucleation stops when the whole Universe has reheated to $ T_c $. This part of the phase transition passes very fast, in about $ 0.05\\mu $s, during which the cosmic expansion is totally negligible. After that, the hadron bubbles grow at the expense of the quark phase and eventually percolate or coalesce; almost in equilibrium at $ T_c $. The transition ends when all quark-gluon matter has been converted to hadrons, neglecting possible quark nugget production \\cite{witt}.  How much the system has to supercool before the first hadron bubbles show up determines the distance between nucleation sites. This distance is called the (comoving) mean nucleation distance, $ d_{nuc} $. It is also roughly the distance between the shrinking, high-baryon-density quark droplets. $ l $ is half that distance, and when we take into account that the Universe expanded about $ 50\\% $ during the phase transition, we have \\begin{equation} \\label{l} l \\simeq 0.8 \\left ( \\frac{T_c}{100\\hbox{MeV}} \\right ) d_{nuc}, \\end{equation} where the factor $T_c/100$MeV scales the value of $d_{nuc}$ from a given critical temperature to $T=100$MeV, where $l$ is defined.  In the bag model we can approximate $ d_{nuc} $ in homogeneous nucleation theory without curvature energy in terms of the bag constant, $B$, by \\begin{equation} \\label{dsurf} d_{nuc} \\approx 13  \\left ( \\frac {145\\hbox{MeV}}{B^{1/4}} \\right )^2 \\left ( \\frac{4B}{L} \\right ) \\left ( \\frac {\\sigma }{0.035B^{3/4}} \\right )^{3/2} \\hbox{cm}, \\end{equation} where $ L $ is the latent heat and $ \\sigma $ the surface tension. Lattice QCD calculations \\cite{surft} suggest that $ \\sigma \\lesssim 0.24T_c^3 $, Ref.\\ \\cite{latent} finds $ \\sigma \\simeq 0.025T_c^3 $ and $L\\simeq 0.4B$. With $ T_c \\simeq 0.68B^{1/4} $, this gives $ d_{nuc} \\simeq 0.1 $m for $ T_c > 100 $MeV.  We calculate the surface tension self-consistently  within the multiple reflection expansion framework of the MIT bag model and get $ \\sigma \\lesssim 0.035B^{3/4} $ and $L\\simeq 4B$, corresponding to $ d_{nuc} \\lesssim 0.1 $m. But we also self-consistently take into account the important curvature term and find a crude approximate formula for $ d_{nuc} $, \\begin{equation} \\label{dcurv} d_{nuc} \\approx 1.8 \\left ( \\frac{145\\hbox{MeV}}{B^{1/4}} \\right )^2 \\left ( \\frac{4B}{L} \\right ) \\hbox{cm}. \\end{equation} Note, that in this approximation the mean nucleation distance depends only on the bag constant (or $ T_c $). Thus the overall effect of including a curvature term is roughly a factor of 7 reduction of $ d_{nuc} $. With these estimates of $ d_{nuc} $ in mind it hardly seems possible to get any significant deviations from standard nucleosynthesis, if homogeneous nucleation theory applies.  However, homogeneous nucleation may not be the prime mechanism responsible for hadron formation. It is well-known from every day life that first order transitions are normally facilitated by impurities, such as dust or ions and container walls. Think e.g.\\ of a charged particle entering a cloud chamber, or boiling water where (half) bubbles are formed at the bottom of the pot. Such nucleations are denoted heterogeneous. Thus it is obvious to consider the effect of such impurities on the quark-hadron phase transition. Candidates are the topological impurities, like primordial black holes, magnetic monopoles, and cosmic strings or relic fluctuations from the electroweak transition.  The idea that impurities could play a role in the cosmological quark-hadron transition is not new (see e.g.\\ \\cite{hoso} for a general discussion of the role of impurities in cosmological phase transitions), but normally it has been argued\\cite{alc1} that the presence of impurities would just act as extra seeds for nucleation sites together with the homogeneous nucleation sites and thus reducing $ d_{nuc} $.  As demonstrated below, this picture is not necessarily correct---there are circumstances where impurities lead to a dramatic {\\it increase\\/} in $d_{nuc}$, maybe even to a value of several meters, resurrecting inhomogeneous Big Bang nucleosynthesis. We find that three limits with, respectively, small, medium and huge amounts of impurities, compared to the number of homogeneous nucleation sites, have to be considered. Certainly, if only very few impurities are present in the early Universe $ d_{nuc} $ will be determined solely by homogeneous nucleation. In the other extreme, for a huge number of impurities, they will determine a $ d_{nuc} $ smaller than the one obtained from homogeneous nucleation. So neither of these cases give large mean nucleation distances.  Any impurity tends to lower the cost in energy when creating the surface of the new phase. Therefore bubbles of the new phase will nucleate easier around an impurity at a smaller amount of supercooling. This means that shock spheres from nucleations around impurities start reheating the Universe earlier. And for some impurity number densities the whole Universe has already been reheated before the amount of supercooling has increased to a value where homogeneous nucleation sets in. In this case $ d_{nuc} $ depends on the impurity number density, $ n $, on the time when heterogeneous nucleation starts, and on the time when homogeneous nucleation becomes important. For a large range of parameters, $d_{nuc}$ can be very big.  We will make this qualitative discussion much more quantitative in Section III, where we also present a self-consistent bag model calculation of homogeneous nucleation, including for the first time the important curvature terms. In Sec.\\ II we describe the thermodynamics of the phase transition and the multiple reflection expansion of the MIT bag model, which we use to calculate the thermodynamical quantities. Section IV contains a general discussion and a summary of our results. ",
+        "conclusions": "We have calculated the mean nucleation distance in a homogeneous nucleation scenario where a curvature term contributes to the surface effects. The framework of the multiple reflection expansion of the MIT bag model has been used to calculate the surface tension and curvature term self-consistently. Inclusion of a curvature term reduces $ d_{nuc} $ by about a factor of 7 to less than 2cm.  Furthermore, we have shown that the presence of impurities, for a broad range of densities, can result in mean nucleation distances above one meter. Such large distances are needed if baryon diffusion should not have smeared out generated baryon inhomogeneities before the epoch of nucleosynthesis. It must be emphasized that this conclusion is qualitatively independent of the model used to describe surface effects. We have only assumed that impurities are present, that they lower the energy barrier and that the reheating mechanism is the same regardless of how the hadron bubbles nucleated. Thus, similar effects could be important in other first order transitions as well (see e.g. Ref.\\ \\cite{hoso}).  Unfortunately there are no reliable theoretical expectations for the nature or density of impurities, but we note again the coincidence, that the horizon distance at the electroweak phase transition has expanded to the interesting range of a few meters at the quark-hadron transition.  Eq.\\ (\\ref{etaapp}) for $ \\eta_f $ in the homogeneous nucleation scenario tells us how uncertainties in the different parameters influence $ d_{nuc} $, which depends linearly on $ \\eta_f $. We believe the curvature coefficient is well-known, because the major contributions come from gluons and the two massless or very light quarks. A significant surface tension contribution from the pions would raise $ \\eta_f $ and thus result in a larger $ d_{nuc,max} $ and an even broader interval of impurity number densities, where the mean nucleation distances is above one meter.  We have throughout used the self-consistently calculated latent heat of $ 4B $, even though some recent lattice calculations \\cite{latent} suggest that the latent heat is about ten times smaller. A temperature dependent bag constant that increases with increasing temperature may reproduce such a result, but no theoretical predictions of such a temperature dependence exist. However, $ \\eta_f $, $d_{nuc}$, and $l$ are all inversely proportional to the latent heat. Thus, a tenfold decrease in $ L $ results in a tenfold increase in $ \\eta_f $, $d_{nuc}$, and $l$.  In Ref.\\ \\cite{alc2} the spectrum of nucleation site separations is calculated. A duality between the nucleation sites and large baryon concentration sites is assumed and the effect on nucleosynthesis calculated. If impurities play a major role, the nucleation site spectrum would look different. For randomly distributed impurities the spectrum would be broad. A sort of regularity in the impurity distribution would result in a very narrow distribution around the mean nucleation distance. Such an approximation is often used in actual inhomogeneous nucleosynthesis calculations.  A serious problem with the bag model in the form used in this paper is the local minimum in the free energy at about 10 fm, even above the critical temperature\\cite{mardor}. One might think about removing it by introducing further corrections in the density of states related to zero point energy, color singlet restrictions, etcetera, but no recipe exists at the moment. We note however, that such effects will be most important for small radii, and since the maximum of the free energy barrier typically lies beyond a radius of 50 fm in the results presented above, the conclusions of the present investigation are not likely to be affected much."
+    },
+    "9602/astro-ph9602050_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "Much of the current controversy concerning the consistency (Copi et al~1995a, 1995b) or inconsistency (Hata et al~1995a) of standard big-bang nucleosynthesis revolves around the chemical evolution of $^3$He. In fact, $^3$He is involved indirectly.  Deuterium plays the crucial role in testing big-bang nucleosynthesis, as its abundance is the most sensitive to the baryon density, decreasing rapidly with increasing baryon density, and its chemical evolution brings in $^3$He.   The chemical evolution of D is straightforward:  it is readily burned to $^3$He but it is not produced in Galactic environments (Epstein, Lattimer \\& Schramm 1976).  This means that the present deuterium abundance can be  used to place an upper limit to the baryon density.  This upper limit, $\\eta \\la 9\\times 10^{-10}$ which implies $\\Omega_B \\la 0.03h^{-2}\\la 0.2$, provides the linchpin in the two-decade-old argument that baryons cannot close the Universe (Reeves et al~1973).  This argument is not questioned in the current controversy.  (As usual, $\\eta$ is the present ratio of baryons to photons, $\\Omega_B$ is the fraction of critical density contributed by baryons, and the Hubble constant $H_0 = 100h\\kms\\Mpc^{-1}$ with $0.4<h<1$.)  Using deuterium to precisely determine the baryon density, or even to set a lower limit to it, is more difficult.  Because deuterium is so easily destroyed in passing through stars, the former is only possible if the deuterium abundance can be measured in a very primitive sample of the Universe.  While there has been much progress toward this goal, with several detections and upper limits based on the D Ly-$\\alpha$ feature in high redshift ($z\\sim 3$), absorption-line systems (York et al~1984; Songaila et al~1994; Carswell et al~1985, 1995; Tytler \\& Fann~1994; Rugers and Hogan 1995; Wampler et al~1996) yielding inferred abundances in the range from $2\\times 10^{-5}$ to $2\\times 10^{-4}$, there is yet no definitive result.  The derivation of a lower limit to the baryon density hinges upon the chemical evolution of $^3$He.  Since D is burned to $^3$He and $^3$He is far more difficult to burn, Yang et al~(1984) proposed using the sum of D + $^3$He for this purpose. Based upon stellar modeling (Iben and Truran 1978), they assumed that at least 25\\% of the primordial D + $^3$He survives stellar processing, which led to the lower limit $\\eta \\ga 2.5\\times 10^{-10}$ and $\\Omega_B \\ga 0.009h^{-2}$. This, together with the upper limit that follows from $^7$Li ($\\eta \\la 6\\times 10^{-10}$ and $\\Omega_B\\la 0.02h^{-2}$), provides the best determination of the baryon density -- between about 1\\% and 15\\% of critical density (allowing $0.4< h < 1$; see Copi et al~1995a) --  and establishes the two dark-matter problems central to cosmology:  most of the baryons are dark (since $\\Omega_{\\rm LUM} \\simeq 0.003h^{-1} < \\Omega_B$) and most of the dark matter must be nonbaryonic, if as several measurements indicate $\\Omega_0 \\ga 0.2$.  Beyond pinning down the baryon density, there is a more fundamental issue: the consistency of the standard model of primordial nucleosynthesis itself and the validity of the hot big-bang model at times as early as 0.01 sec. (By standard model of big-bang nucleosynthesis we mean:  FRW cosmology, uniform distribution of baryons, three light neutrino species, and small neutrino chemical potentials.) The $^7$Li abundance measured in almost 100 old, Pop II halo stars is consistent (``at $2\\sigma$'') with the big-bang prediction provided that $\\eta \\simeq (1-6)\\times 10^{-10}$, which overlaps the D + $^3$He consistency interval (Copi et al~1995a). Of some concern is the primeval $^4$He abundance:  If one accepts at face value the analysis of Olive and Steigman (1995), based upon metal poor, extragalactic HII regions, then their value for $^4$He of $Y_P = 0.232 \\pm 0.003\\ {\\rm (stat)} \\pm 0.005\\ {\\rm (sys)}$ implies $\\eta \\simeq (1-4)\\times 10^{-10}$ (at ``$2\\sigma$'), which is only marginally consistent with the D + $^3$He lower bound. It should be noted, however, that other authors (see e.g., Skillman et al~1994; Sasselov and Goldwirth 1995; and Pagel, private communication) have argued that the systematic uncertainties are at least a factor of two larger, which, owing to the logarithmic dependence of $Y_P$ upon $\\eta$ would enlarge the concordance interval to $\\eta \\simeq (0.6 - 10)\\times 10^{-10}$.  For some time, there has been tension between the measured abundances of $^4$He and D + $^3$He (see e.g., Yang et al~1984; Copi et al~1995a; Walker et al~1991; Olive et al~1995; Scully et al~1996). The resolution could involve a revision of our understanding of the evolution of $^3$He: more stellar destruction than standard stellar models predict would lead to a lower value of $\\eta$ as inferred from D + $^3$He and lessen the tension.  Alternatively, the resolution could involve an underestimation of the primeval $^4$He abundance, by an amount $\\Delta Y_P \\sim 0.01$ (Copi et al~1995b); this would raise the value of $\\eta$ inferred from $^4$He, making it consistent with that inferred from D + $^3$He and conventional stellar evolution of $^3$He.  Hata et al~(1995a) have argued that the discrepancy is real and is evidence for new physics, e.g., an unstable tau neutrino of mass $10\\MeV$ or so or neutrino chemical potentials.  Eventually the deuterium abundance in high redshift Ly-$\\alpha$ clouds will be decisive;  e.g. a value  (D/H)$_P\\sim 10^{-4}$ implies $\\eta \\sim 2\\times 10^{-10}$ and would implicate the chemical evolution of $^3$He, while (D/H)$_P\\sim 3\\times 10^{-5}$ implies $\\eta \\sim 6\\times 10^{-10}$ and would implicate the primeval $^4$He abundance.  Until a definitive determination is forthcoming, continued scrutiny of $^3$He -- both theoretically and observationally -- offers a means of addressing this important issue. Because previous measurements of the abundance of $^3$He have raised as many questions as they have answered -- variations in the abundance measured in HII regions of more than a factor of five (Bania, Rood \\& Wilson 1987; Wilson \\& Rood 1994) with some values {\\it lower than that in the pre-solar nebula} (Black 1972; Geiss \\& Reeves 1972) -- the measurement of the $^3$He abundance in the local ISM by Gloeckler \\& Geiss (1996) is an important development. We will use it to derive a lower bound on the baryon density which is more empirically rooted and less sensitive to the questionable aspects of $^3$He evolution. ",
+        "conclusions": "Gloeckler \\& Geiss' measurement is noteworthy because it allows the evolution of D + $^3$He to be addressed empirically for the first time. The message is simple:  over the past 4.5 Gyr it has not changed dramatically.  This means that those stars that have contributed significantly to the local ISM over this period are not significant producers or destroyers of $^3$He.  This is not a trivial fact, as the increase in $^3$He and the decline in D over this time (almost a factor of two) indicate substantial stellar processing.  According to chemical-evolution models, low-mass stars ($\\sim 1M_\\odot - 1.5 M_\\odot$) have made the dominant contribution to the ISM over the past 4.5 Gyr (Truran \\& Cameron 1971; Rood, Steigman \\& Tinsley 1976; see also the recent discussion by Scully et al~1995). Such a constancy of D + $^3$He implies that low-mass stars cannot be significant producers of $^3$He, which is at variance with the predictions of standard stellar models.  Likewise, there is no evidence to support significant destruction of $^3$He by low-mass stars as predicted with an efficient solar spoon at work (see e.g., Hogan 1995).  However, Dearborn (private conversation) has shown that the slow mixing models that fit the oxygen and carbon isotopic anomalies do not completely deplete $^3$He; they reduce the amount of $^3$He that would have been returned to the ISM by at most 80\\%.  If this is indeed the case, a solar spoon could be consistent with the ISM value of D + $^3$He.  Gloeckler \\& Geiss' result does little to directly constrain the earlier evolution of D + $^3$He.  The stellar mass function at earlier times is expected to favor more massive stars, which deplete $^3$He.  Since the Gloeckler \\& Geiss result constrains low-mass star destruction of $^3$He, massive stars are the only possible way to greatly deplete $^3$He. Massive stars produce heavy elements, and thus there is a limit to the amount of material that could have been processed through massive stars.  In particular, the ejecta of Type II supernovae are about 10\\% oxygen by mass, which implies that only a small fraction of the material in the local ISM -- roughly 10\\% -- could have been processed through massive stars.  Taken together with the apparent constancy of D + $^3$He over the last 4.5 Gyr, this suggests that the primordial value of D + $^3$He cannot differ greatly (about a factor of two for simple closed galaxy models) from the present value. This leads to the lower bound, $[({\\rm D + ^3He})/{\\rm H}]_P \\la 10^{-4}$ which is almost identical to that used by Yang et al~(1984),  but now with firmer empirical roots.  An important assumption underlies the above argument, that all the metals ejected by massive stars make their way back into the ISM.  It is possible that metals produced in the early supernova-active phase of the proto-Galaxy (or the proto-galactessimals that merged to form the Galaxy) were ejected into the surrounding IGM. There is some evidence for this; observations of the x-ray emitting gas in rich clusters show metallicities that are slightly less than solar, distributed in a gas mass that is roughly ten times that in galaxies.  If the Galaxy ejected a comparable amount of metals ten times more material could have been processed through massive stars, depleting  $^3$He dramatically. (Note, material depleted in $^3$He is still returned to the ISM in a pre-supernova stellar wind.)  However, as Copi et al~(1995c) showed, even a relaxed metallicity constraint does not allow the primordial value of $({\\rm D + ^3He})/{\\rm H}$ to exceed about $2 \\times 10^{-4}$.  Finally, let us use the information gleaned from this first measurement of the $^3$He abundance in the ISM to make more quantitative statements about the value of the baryon density and the consistency of standard big-bang nucleosynthesis.  The stochastic history approach of Copi et al (1995c) allows one to use the pre-solar values of $^3$He and D + $^3$He to infer both their primordial values and $\\eta$, while allowing for the heterogeneity of Galactic abundances.  The physical input needed is the mean properties of stellar processing.  Based upon Gloeckler \\& Geiss' measurement, we consider two possibilities for the evolution of $^3$He in low-mass stars, (1)  low-mass stars preserve their $^3$He, but do not produce $^3$He; and (2) low-mass stars destroy 80\\% of the $^3$He they would have returned to the ISM, and two possibilities for metal ejection by massive stars, (a) massive stars return most of the metals they make to the ISM and (b) massive stars only return 10\\% of the metals they make to the ISM (the rest ejected into the IGM). For these four possibilities, 1a, 1b, 2a, and 2b, we have constructed Monte-Carlo likelihood functions for the baryon-to-photon ratio $\\eta$, which are shown in Figure 2. The 95\\% credible intervals are: $\\eta_{1a} = (5 - 7) \\times 10^{-10}$; $\\eta_{2a} = (3 -6) \\times 10^{-10}$; $\\eta_{1b} = (2 - 5) \\times 10^{-10}$; and $\\eta_{2b} = (2 - 5) \\times 10^{-10}$. For reference, the very naive assumption that D + $^3$He has remained unchanged since primordial nucleosynthesis implies $\\eta \\sim 5 \\times 10^{-10}$ and a primeval D abundance $({\\rm D/H})_P \\sim 4\\times 10^{-5}$.  Regarding the consistency of big-bang nucleosynthesis, Model 1a continues to implicate $^4$He as the culprit (or the standard model of big-bang nucleosynthesis itself).  Models 1b, 2a, and 2b lessen the tension between $^4$He and $^3$He and D, with Models 1b and 2b essentially eliminating the tension all together.  The full range for the baryon density based upon these models, $\\eta \\simeq (2-7) \\times 10^{-10}$, is essentially the same as that found previously by Copi et al (1995a). We note that the models that lessen the tension, lead to a stronger upper limit to $\\eta$ and strengthen the case for nonbaryonic dark matter.  For example, for Models 1b and 2b, the joint 95\\% credible region for all the light-elements corresponds to $\\Omega_B = (0.007 - 0.018)h^{-2}$.  In sum, the measurement of the interstellar $^3$He abundance by Gloeckler \\& Geiss (1996) allows the chemical evolution of D + $^3$He to be addressed empirically for the first time, and in turn, tests primordial nucleosynthesis and its prediction for the baryon density. Their measurement indicates little evolution of D + $^3$He over the past 4.5 Gyr, generally confirming the the argument of Yang et al (1984), suggesting that low-mass stars are not significant producers or destroyers  of $^3$He, and calling into question standard stellar models as well as efficient solar spoons. Little can be learned directly from their result about the earlier evolution of D + $^3$He, which is likely to be dominated by high-mass stars.  However, the fact that high-mass stars also produce metals limits the amount of $^3$He depletion, even if 90\\% of the metals they produce are ejected from the Galaxy. We have used this fact together with Gloeckler \\& Geiss' result to establish a more empirically based lower bound to the baryon-to-photon ratio, $\\eta \\ga 2\\times 10^{-10}$.  While only slightly less stringent than the bounds of Yang et al~(1984) and Copi et al (1995a), it suggests the apparent tension between the big-bang abundance of $^4$He and those of D and $^3$He involves the chemical evolution of $^3$He.   \\paragraph{Acknowledgments.} We acknowledge valuable conversations with Robert Rood and John Simpson.  This work was supported in part by the DOE (at Chicago and Fermilab) and the NASA (at Fermilab through grant NAG 5-2788 and at Chicago through NAG 5-2770 and a GSRP fellowship) and by NSF at Chicago through grant AST 92-17969.  \\def\\beginapjbib{\\begingroup"
+    },
+    "9602/astro-ph9602044_arXiv.txt": {
+        "abstract": "The morphological properties of galaxies between $21 {\\rm~mag} < I < 25 {\\rm~mag}$ in the {\\em Hubble Deep Field} are investigated using a quantitative classification system based on measurements of the central concentration and asymmetry of galaxian light. The class distribution of objects in the {\\em Hubble Deep Field} is strongly skewed towards highly asymmetric objects, relative to distributions from both the {\\em HST Medium Deep Survey} at $I < 22 {\\rm~mag}$ and an artificially redshifted sample of local galaxies.  The steeply rising number count-magnitude relation for irregular/peculiar/merging systems at $I < 22 {~\\rm mag}$ reported in Glazebrook \\etal\\ (1995a) continues to at least $I=25~{\\rm mag}$. Although these peculiar systems are predominantly blue at optical wavelengths, a significant fraction also exhibit red $U-B$ colours, which may indicate they are at high redshift. Beyond Glazebrook \\etal 's magnitude limit the spiral counts appear to rise more steeply than high-normalization no-evolution predictions, whereas those of elliptical/S0 galaxies only slightly exceed such predictions and may turn-over beyond $I \\sim 24~{\\rm mag}$. These results are compared with those from previous investigations of faint galaxy morphology with HST and the possible implications are briefly discussed. The large fraction of peculiar/irregular/merging systems in the {\\em Hubble Deep Field} suggests that by $I\\sim 25 {\\rm~mag}$ the conventional Hubble system no longer provides an adequate description of the morphological characteristics of a high fraction of field galaxies. ",
+        "introduction": "The {\\em Hubble Deep Field\\/} (HDF) is a four square arcminute area of sky at RA 12h 36m 49.4s~Dec +62$^\\circ$ 12$^\\prime$ 58.0$^{\\prime\\prime}$ (J2000) observed during December 1995 with the {\\em Hubble Space Telescope} (HST).  The images, obtained during 150 orbits with the Wide Field and Planetary Camera 2 (WFPC-2), represent the deepest optical imaging survey yet undertaken (Williams \\etal~1996).  Four broadband filters were used whose wavelength coverage samples the range between the ultraviolet and near infrared (roughly {\\em UBVI}).  In the $I$-band (F814W), the HDF data reaches $\\sim 3$~mag fainter than the deepest ground-based observations at similar wavelength (Smail {\\em et~al.}  1995), and $\\sim 1$~mag fainter than the deepest $I$-band observations undertaken previously with the HST. The position of the HDF was pre-selected to be a ``typical'' sample of the high galactic latitude sky in terms of galaxy surface density, while also avoiding bright stars or other objects whose presence might compromise the depth of the observations.   \\begin{figure*} \\centerline{\\psfig{figure=figure_1.ps,width=5.3in,angle=-90}} \\caption{Asymmetry vs. central concentration for galaxies in the Hubble Deep Field. RSE's visual morphological classifications are keyed to the plot symbols: E/S0's are shown as ellipses, spirals earlier than Sd are shown as spirals, and irregulars/peculiars/mergers are shown as asterisks.  The sectors subdivide the diagram into regions where each of these morphological types dominates.  Representative error bars are shown. The dotted polygon shows the ``convex hull'' enclosing the artificially redshifted local sample of objects described in the text.} \\end{figure*}   A major objective of the HDF program is to determine the nature of the rapidly evolving faint blue galaxy population seen in ground-based spectroscopic and photometric surveys conducted over the past decade (Broadhurst, Ellis, \\& Shanks 1988; Tyson 1988; Colless \\etal\\ 1990; Ellis 1990; Koo \\& Kron 1992; Lilly 1993; Tresse \\etal\\ 1993; Cowie \\etal\\ 1994; Glazebrook \\etal\\ 1995b). Many of these objects are distant ($z > 0.3$), and currently only the HST provides the resolution needed to characterise their morphological properties. Results to date have proved difficult to reconcile however, in part perhaps because of the inherently qualitative nature of galaxy classification.  Number counts from a sample of 301 galaxies with \\hbox{$I<22$} mag taken from 13 WFPC-2 fields from the {\\em Medium Deep Survey} (MDS) have been reported by Glazebrook~{\\em et al.} (1995a). These authors classified galaxies by eye and found the number-magnitude relationships for elliptical and spiral populations to be consistent with the predictions of high-normalization no-evolution cosmological models, but that the number of irregular/peculiar (possibly merging) galaxies were at least an order of magnitude greater than predicted by the same models.  A similar conclusion was reached by Driver {\\em et~al.} (1995) based on number counts from a single deeper MDS WFPC-2 field to $I$=24.25~mag and by Abraham {\\em et~al.} (1996) from automated morphological classification of 521 galaxies from a larger number (21) of WFPC-2 MDS fields to $I$=22~mag.  In contrast, HST observations reported by Schade {\\em et~al.} (1995) of a small but spectroscopically complete sample of 32 galaxies from the Canada-France Redshift Survey (CFRS) indicate that the morphological mix at $z \\sim0.75$ is similar to that seen locally, except for a population of ``blue nucleated galaxies'' (BNGs) which Schade {\\em et~al.\\/} suggest comprise $\\sim35\\%$ of the faint blue galaxy population. The existence of another major new population of morphologically distinct objects, linear ``chain galaxies'' that may be proto-galactic systems, has recently been reported by Cowie {\\em et~al.} (1995) on the basis of HST observations of faint objects ($I>23$~mag) identified in the Hawaii Deep Survey.  A major advantage of the HDF dataset over the earlier images of Glazebrook \\etal, Schade \\etal, Driver \\etal, and Cowie \\etal~ is the improved signal-to-noise level arising from the considerably longer HDF exposure time. The exposure times available to those authors were 5-8 hours in F814W, which is substantially less than the $\\simeq$35 hours devoted to the HDF in this filter. The improved signal to noise at faint limits enables us to extend the work already begun with the MDS, whose wider spatial coverage nicely complements the HDF data.  Although interesting questions (which the HDF may ultimately answer) concern the very faint ($I>$25 mag) galaxy counts, in this short note we are concerned with understanding the morphology and colour distributions of the $\\sim 300$ objects brighter than $I$=25 mag. As described below, at this limit the bulk of the galaxy population is detectable in all four HDF passbands, and morphological classifications are as robust as was the case in our earlier analyses of the MDS data to $I$=22 mag.  A plan of the paper follows. In \\S2 we review our morphological classification techniques in the context of the HDF data. Section~3 discusses the colour distribution as a function of morphology and, in \\S4, we combine our dataset with that of the MDS, enabling us to extend Glazebrook \\etal~(1995a)'s morphological galaxy count diagram three magnitudes fainter. A number of new trends appear whose implications are briefly discussed. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602102_arXiv.txt": {
+        "abstract": "The unusual properties of the bursting X-ray pulsar GRO J1744-28 are explained in terms of a low-mass X-ray binary system consisting of an evolved stellar companion transferring mass through Roche-lobe overflow onto a neutron star, implying that the inclination of the system is $\\lesssim 18^\\circ$.  Interpretation of the QPO at frequency  $\\nu_{\\rm QPO} = 40$ Hz using the beat-frequency model of Alpar \\& Shaham  and the measured period derivative $\\dot P$ with the Ghosh \\& Lamb accretion-torque model implies that the persistent X-ray luminosity of the source is approximately equal to the Eddington luminosity and that the neutron star has a surface equatorial magnetic field $B_{\\rm eq} \\cong  2\\times 10^{10}\\ (\\nu_{\\rm QPO}/40\\ {\\rm Hz})^{-1} \\ {\\rm G}$ for standard neutron star parameters. This implies a distance to GRO J1744-28 of $\\cong 5\\ (\\nu_{\\rm QPO}/40\\ {\\rm Hz})^{1/6} b^{1/2}$ kpc, where $b<1$ is a correction factor that depends on the orientation of the neutron star. ",
+        "introduction": "GRO J1744-28, a 2.1 Hz X-ray pulsar with unusual bursting properties, was discovered on 2 December 1995 in the direction of the Galactic Center with the Burst and Transient Source Experiment (BATSE) on the {\\it Compton~Gamma~Ray~Observatory}  (Kouveliotou et al. \\markcite{Kouvea96}1996a; Finger et al. \\markcite{Fingera96}1996a).  The source initially bursted about every 3 minutes, then declined in frequency to about 30 per day for the first month, after which it increased in frequency  to about 40 per day in the period 8-18 January 1996 (Fishman et al. \\markcite{Fishm96}1996).   The {\\it Rossi~X-ray~Timing~Explorer} ({\\it RXTE}) (Swank et al. \\markcite{Swank96}1996) localized the source to galactic coordinates $l= +0.02^\\circ, b=+0.3^\\circ$.   Both the persistent and bursting X-ray emissions show pulsations with a 0.467 s period (Kouveliotou et al. \\markcite{Kouveb96}1996b) and have similar hard X-ray and soft gamma-ray spectra (Briggs et al. \\markcite{Brigg96}1996; Strickman et al. \\markcite{Stric96}1996) with a  color temperature $\\approx 10$ keV, characteristic of X-ray pulsars.  In this {\\it Letter}, we show that the observations are best explained by a neutron star accreting matter through Roche lobe overflow from an evolved low-mass stellar companion.  We derive values for the magnetic field stength and accretion luminosity by noting that the beat frequency model (Alpar \\& Shaham \\markcite{Alpar85}1985) for QPOs and the accretion-torque model (Ghosh \\& Lamb \\markcite{Ghosh79}1979) for the pulsar spin-up rate separately imply  relations between these two quantities.  The size and temperature of the polar-cap hot spot are determined.  The observed pulse fraction is used to constrain the orientation and geometry of the neutron star, and to derive the distance to the system from the calculated luminosity. ",
+        "conclusions": ""
+    },
+    "9602/hep-ph9602398_arXiv.txt": {
+        "abstract": "We investigate the effects of random density fluctuations on neutrino oscillations in the Sun environment. We show how the average of certain quantities which can be used to describe the MSW effect can be computed analytically. We examine also the hypothesis commonly accepted that only perturbations inside the resonance layer can have relevance. The average amplitude, which gives the ''coherent probability'', is computed in an analytical and exact way for the case of colored $\\delta$-correlated Gaussian noise: the random perturbation induces a renormalization of the matter density which acquires an imaginary part proportional to the fluctuation magnitude in the resonance region. Integral equations are given for the density matrix of the system in the ''optical'' approximation.  PACS: 96.60.Kx, 02.50.Ey, 14.60.Pq,95.30.Cq, 96.60.Hv,14.60.Gh. ",
+        "introduction": "In this work we investigate matter-enhanced neutrino flavor transformations, the MSW effect \\cite{wol1,mik1}, for the case of a exponentially decaying matter density with a colored Gaussian noise.  There are several   examples of interesting transport and scattering processes induced or modified by the presence of a disordered or random medium: dislocations in crystals and the solid-liquid transition; random impurity potentials produce the localization of quantum wavefunctions which enables one to understand the transition between insulators and conductors; resistance anomalies at low temperature and in presence of magnetic fields of ''weak localized'' electron systems subject to a random potential  (\\cite{ber1,vol1}). The same basic localization phenomena explain also in optics the backscattering enhancement in electromagnetic wave scattering from a randomly rough surface (\\cite{fitz1}). Usually, a strong dependence on the dimensionality of the problem is observed. In some other cases as in nuclear physics the consideration of random matrix hamiltonians allows the simplification of otherwise unmanageable systems.    For neutrino oscillations in presence of random or rapidly twisting magnetic fields considerable  work has been done already \\cite{lor1,ane1,akh1,nic2}. Random density fluctuations have received some attention only recently.   In \\cite{lor1}, a differential equation for the averaged survival probability was derived for the case in which the random noise was taken to be a delta-correlated white Gaussian distribution. The differential equation was solved numerically and the neutrino evolution obtained. Arguments were given which indicate that if the correlation of the matter density fluctuations is small compared to the neutrino oscillation length at resonance, one obtains the same result as for the case of a delta-correlated noise. In \\cite{lor2} the more realistic case of colored noise was considered also in a numerical way and applied to Supernova dynamics. An approximate differential equation for the averaged survival probabilities was obtained using the hypothesis that the fluctuations should not affect the evolution far from the resonance. In \\cite{nun2} , the implications of random perturbations upon the Solar neutrino deficit are considered numerically. It is found that the MSW effect is rather stable under these fluctuations specially in the small mixing case but in anycase the experimental $(\\Delta m^2,\\cos\\theta)$ exclusion curves get modified in an appreciable way.  In this work we try to develop analytical results for neutrino flavor oscillations induced by a random matter density. We are interested mainly in the persistence of the MSW effect. The obtention of concrete values for the survival probabilities and phenomenological consequences for the solar neutrino problem is left for a subsequent work. The obtention of these phenomenological consequences is not an easy task , at least, because the amplitude of possible density fluctuations in the Sun is poorly known experimentally and theoretically. Different arguments can give easily values for it differing by two orders of magnitude (\\cite{nun2}): between $0.1\\%-10\\%$ of the local density. Stronger local inhomogeneities, for example at the near-surface  dark spots should not be discarded.   Our basic starting point will be the exact analytical solution obtained in \\cite{emi1} for the neutrino oscillation amplitudes in presence of an exponentially decaying matter density. The random component is consider as a perturbation to this solution. Inspired by the electromagnetic wave scattering in random media, {\\em coherent } and {\\em incoherent} transition probabilities are defined. The basic result of this work is that the coherent probability, which comes essentially from an averaged amplitude, can be computed exactly in some cases. Different integral equations are given for the incoherent probability using an ''optical'' approximation. Approximate solutions valid in a limited range are obtained for them. By the very nature of the procedure the results are valid or easi    ly generalizable to any number of neutrino species.    The outline of this work is as follows. The first section, after the introduction of some known results, is dedicated to explore diverse quantities whose  stochastic average can be computed exactly or at least by an easy approximation. One of these quantities is the determinant of the evolution operator of the system. We use a ''naive'' argument to show the plausibility that the influence of the random perturbation on the neutrino oscillations can be described by a complex redefinition of the matter density. This suggestion allows to define some {\\em ansatz} probability. Another quantity is the ''total cross section'' of the system, we suggest the interest in further studying this quantity which can contain information on the presence and localization of the MSW resonance. In the same section, the physical supposition that only those random perturbations happening in the resonance layer can have importance  is studied briefly . We argue that this supposition must be taken with care, because, as we show, even small phase shifts {\\em before} the resonance region can have some appreciable importance in the final survival probability.  In the next section we define some perturbative expansion for the density matrix of the system. The {\\em coherent} and {\\em incoherent} parts of this matrix are defined. The coherent part being in some sense the zero-order approximation for the full density. Using an ''optical'' approximation, that is discarding a certain class of terms in the perturbative expansion, a very general integral equation for the averaged density is derived. In this integral equation, basic ingredients are both  the coherent density and the averaged amplitude and the two point correlation of the matter density perturbation.  In Section (\\ref{s4}) different particular cases are considered. A simpler expression for the previous integral equation is given for the case where the random perturbation is $\\delta$-correlated. A further simpler expression is given for the small mixing case of two neutrino species. In this case upper limits for the total averaged probability can be obtained depending on the coherent probability and an autocorrelation integral.  In Section (\\ref{s5}) the coherent part is computed. In a first case the small mixing condition and different approximations are used. Averaging the amplitude amounts to the multiplication by a certain slowly time-varying diagonal matrix. In a second particular case, it is shown how  the average  can be computed exactly. The effect of the averaging is indeed a complex renormalization of the initial matter density as it was  suggested earlier. Even if it is a very particular case, it is shown how it can have relevance in more realistic computations. Finally, making use of these averages, coherent survival probabilities and the ''cross-sections'' previously defined are computed. ",
+        "conclusions": "The most important result of this work is the derivation of analytical exact expressions for the average {\\em coherent} transition probability for a special case of colored $\\delta$-correlated Gaussian noise. We have shown that in this case the consequence of the presence of the noise is a complex renormalization of the matter density. Other approximative expressions for more general cases have been  developed as well. It has been suggested an enhancement of the survival probability when the {\\em incoherent} probability becomes dominant.   It has been proposed for the first time to consider  new scattering ''cross sections''. It has been shown how they are able to describe themselves the MSW effect; clearly   further work has to be done in this respect. The main importance of these quantities is that their statistical average is computable in a simple way.  The general conclusion is that the MSW effect survives the presence of random perturbations at least for small fluctuation values. There are indications that their  effect can become important for large fluctuations.   \\vspace{1cm} {\\Large{\\bf Acknowledgments.}}\\vspace{0.3cm}  I would like to thank to Peter Minkowski for many enlightening discussions. This work has been supported in part by the Wolferman-Nageli Foundation (Switzerland) and by the MEC-CYCIT (Spain).  \\newpage  \\begin{figure}[p] \\centering\\hspace{0.6cm} {\\epsfig{file=p5f003.ps,height=12cm}} \\caption{The ''total cross sections'' \\protect{$\\tWW\\sigma_1^1$} (short dash) and \\protect{$\\sigma_2$} (longer dash)  for a neutrino produced at the Sun (\\protect{$\\rho_0/\\lambda=50,\\cos^2=0.99$}) as a function of \\protect{$\\beta=\\Delta m^2/2E\\lambda$} (see Eqs.(\\protect\\ref{e6024}-\\protect\\ref{e6025})). The continuos line is the survival probability (\\protect{$\\times 10$}).} \\label{f1} \\end{figure}  \\begin{figure}[p] \\centering\\hspace{0.6cm} \\begin{tabular}{c} {\\epsfig{file=p5f001.ps,height=8cm}}\\\\[-2cm] {\\epsfig{file=p5f002.ps,height=8cm}}\\vspace{-1cm} \\end{tabular} \\caption{The survival probability as a function of an arbitrary phase shift introduced in different positions along the trajectory of a neutrino created near the Sun core ( $\\cos^2\\theta=0.99$,$\\beta=1,20$) . For $\\beta=20$ the resonance is situated at $r/r_0\\approx 0.5-0.6$. The continuos line correspond to $r/r_0=1.2$: well ahead the resonance region. The longer dashed to $r/r_0=0.6$: inside the resonance. The largest variation (both shorter dashed lines) correspond to $r/r_0=0.5,0.3$ at the beginning or clearly before the resonance respectively.} \\label{f2} \\end{figure}   \\begin{figure}[p] \\centering\\hspace{0.8cm} \\begin{tabular}{c} {\\epsfig{file=p5f009.ps,height=10cm}}\\\\[-4cm] {\\epsfig{file=p5f010.ps,height=10cm}}\\\\[-1cm] \\end{tabular} \\caption{The $\\nu_e$ coherent survival probability as function of the neutrino creation point. (Top) $P_{ee}^{coh}$ (continuos) and $P^{coh}$ (dashed) computed using $\\Sigma^{eff}$ given by Eq.(\\protect\\ref{e7078}). Bottom, the same probabilities computed  using $U_r(\\rho_0(1-ik\\rho_0/2))$. ($k=10^{-4},\\beta=20,\\cos^2\\theta=0.99$ for both figures.)} \\label{f3} \\end{figure}    \\begin{figure}[p] \\centering\\hspace{0.8cm} {\\epsfig{file=p5f021.ps,height=15cm}} \\caption{Coherent probabilities $P^{coh}$ (longer dash), $P_{ee}^{coh}$ (continuous line) as a function of $k$ and the mixing angle. They are computed using $U_r(\\rho_0(1-ik\\rho_{res}/2))$. $k=10^{-4},10^{-3},10^{-2}$ for upper, middle, lower figures respectively. Right figures: large mixing, $\\cos^2\\theta=0.70$; left, small mixing $\\cos^2\\theta=0.99$. In all the cases $\\beta=20$. It is also depicted $P^M$ (Eq.(\\protect\\ref{e7071})) (short denser dash)  and $C_{min}=1-P^{coh}-k\\rho^2/2$ (short less denser dash).} \\label{f4} \\end{figure}   \\begin{figure}[p] \\centering\\hspace{0.8cm} \\begin{tabular}{cc} {\\epsfig{file=p5f017.ps,height=6cm}} & {\\epsfig{file=p5f018.ps,height=6cm}} \\\\[-1cm] {\\epsfig{file=p5f019.ps,height=6cm}}& {\\epsfig{file=p5f020.ps,height=6cm}}\\vspace{-1cm} \\end{tabular} \\caption{ $\\tWW\\sigma_1^1$ (line) and $\\sigma_2$ (dashed line) for a neutrino with the same characteristics as in Fig.(\\protect{\\ref{f1}}) in presence of a random perturbation. Two top figures: average amplitude computed using $\\Sigma^{eff}$ from Eq.(\\protect\\ref{e7078}). Lower figures: using $U_r(\\rho_{or})$. $k=10^{-4},10^{-3}$ ($\\approx 1\\%, 5\\%$ fluctuations ) respectively left and right.} \\label{f5} \\end{figure}   \\begin{figure}[p] \\centering\\hspace{0.8cm} \\begin{tabular}{c} {\\epsfig{file=p5f011.ps,height=10cm}}\\\\[-1cm] {\\epsfig{file=p5f016.ps,height=10cm}}\\\\[-1cm] \\end{tabular} \\caption{The $\\nu_e$ survival probability as function of the neutrino creation point computed using the {\\em ansatz} of Section \\protect\\ref{s22}, Eq.(\\protect\\ref{e7006}) with $\\rho_{or}=\\rho_0(1-ik\\rho_{res}/2)$. Continuos line: non-random probability ($k=0$). Dashed lines: $k=10^{-4},10^{-3},10^{-2}$ respectively. For both figures $\\beta=20$;  $\\cos^2\\theta=0.99$ (Top), $\\cos^2\\theta=0.70$ (Bottom).} \\label{f6} \\end{figure}   \\newpage"
+    },
+    "9602/hep-ph9602307_arXiv.txt": {
+        "abstract": "We consider the implications of solar matter density random noise upon resonant neutrino conversion. The evolution equation describing MSW-like conversion is derived in the framework of the Schr\\\"odinger approach. We study quantitatively their effect upon both large and small mixing angle MSW solutions to the solar neutrino problem. This is carried out both for the active-active $\\nu_e \\ra \\nu_{\\mu,\\tau}$ as well as active-sterile $\\nu_e \\ra \\nu_s$ conversion channels. We find that the small mixing MSW solution is much more stable (especially in $\\Delta m^2$) than the large mixing solution. The possible existence of solar matter density noise at the few percent level could be tested at future solar neutrino experiments, especially Borexino. ",
+        "introduction": "The long-standing deficit of solar neutrinos (the Solar Neutrino Problem (SNP)) has now been observed by all four operating experiments \\cite{cl,ga,sa,k,val95}.  The main essence of the SNP is the strong deficit of the beryllium neutrinos \\cite{CF}. On the other hand, the high energy boron neutrinos are moderately suppressed, while the low energy ones are almost undepleted.  This strongly suggests that any astrophysical solution fails \\cite{CF,BFL} in reconciling the experimental data with the Standard Solar Model (SSM) predictions \\cite{SSM,turck,CDF}.  It is possible to ascribe the solar neutrino deficit to the existence of two types of neutrino conversion mechanisms, both of which can deplete neutrinos of different energies differently, as required by the experimental data. The long wavelength vacuum oscillations provide a good fit to the most recent results for $\\Delta m^2\\simeq 10^{-10}$eV$^2$ and large neutrino mixing $\\sin^2 2 \\theta \\simeq 1$ \\cite{KP,BR,Cala}. The other scenario is the resonant neutrino conversion due to interactions with constituents of the solar material (the Mikheyev-Smirnov-Wolfenstein  (MSW) effect) \\cite{MSW}. This provides an extremely good data fit in the small mixing region with  $\\Delta m^2 \\simeq 10^{-5}$eV$^2$ and $\\sin^2 2 \\theta \\simeq 10^{-3} \\div 10^{-2}$ \\cite{FIT,smirnov,Cala}. Both of these solutions have been studied against possible changes of the SSM input parameters \\cite{smirnov,BR}. For example, the study of the MSW effect has revealed its stability, especially in the $\\Delta m^2$ parameter.  In this paper we investigate the stability of the MSW solution with respect to the possible presence of random perturbations in the solar matter density, so far not included in the standard MSW picture.  In Ref.\\cite{KS} the effect of periodic matter density perturbations added to an average density $\\rho_0$, i.e. \\be[period] \\rho(r) = \\rho_0  [1 + h \\sin (\\gamma r)] \\ee upon resonant neutrino conversion was investigated. The major effects show up when the fixed frequency ($\\gamma$) of the perturbation is close to the neutrino oscillation eigen-frequency, and for rather large amplitude values ($h \\sim 0.1-0.2$), giving rise to the parametric effects \\cite{KS}. Such effects can either enhance or suppress neutrino conversion in the Sun. There are also a number of papers which address similar effects by different approaches \\cite{AbadaPetcov,BalantekinLoreti}.  Direct observations of solar surface motions, resulting from the superposition of several modes, may indicate a rich spectrum of frequencies. This would suggest the need to consider the effect of random or \"white\" noise matter density perturbations $\\xi(r)$, characterised by an {\\sl arbitrary} wave number $k$, \\beq \\xi (r) = \\int dk \\xi(k)\\sin kr  \\:, \\eeq rather than a periodic or regular perturbation. In such a case the spatial correlation function for a uniform medium \\be[corr] \\langle \\xi (r_1)\\xi (r_2)\\rangle = \\langle \\xi^2 \\rangle_{r_1 -r_2}, \\ee obeys $\\langle \\xi (k)\\xi (k')\\rangle = \\langle \\xi^2 \\rangle_k \\delta (k + k')$ as the averaging rule for the Fourier components, where the wave number $k$ is not fixed. The effect of solar density as well as solar magnetic field fluctuations upon neutrino spin-flavour conversions has also been considered in Ref. \\cite{BalantekinLoreti}, using somewhat different methods.  In this paper, after some discussion (Sec. 2) about the nature of the matter density fluctuations, we derive the most general neutrino evolution equation in random matter, starting from the standard Schr\\\"odinger equation (Sec. 3). This discussion is closer to the particle physics intuition than that of Ref. \\cite{BalantekinLoreti}. Moreover, we consider both the active-active $\\nu_e \\rightarrow\\nu_{\\mu,\\tau}$ as well as the active-sterile $\\nu_e\\rightarrow\\nu_s$ neutrino conversion channels (here $\\nu_s$ is a neutrino state with no standard model interaction). The latter is motivated by the fact that the existence of a sterile neutrino seems to be the only way to simultaneously account for the solar and atmospheric neutrino deficits in the presence of neutrino dark matter \\cite{DARK}.  After an analytical study of the neutrino conversion equations we have investigated the impact of matter density noise upon the MSW scenario in the context of the SNP (Sec. 4). Typically, we find that the presence of matter fluctuations weakens the MSW mechanism, thus reducing the resonant conversion probabilities \\cite{BalantekinLoreti}. We have carried out a fit of the latest solar \\neu data for different values of the noise level, minimising the $\\chi^2$ in the ($\\Delta m^2, \\sin^2 2 \\theta$) plane. As in the noiseless MSW case, we find that the small mixing MSW solution provides a better fit to the data than the large mixing one, both for the case of active, as well as sterile neutrino conversions. We present several plots with the results of our fits in which the effect of the noise is studied in the idealised approximation where all neutrinos are produced at the solar centre. We conclude that the most relevant parameter region corresponding to adiabatic conversion of $^7$Be neutrinos is relatively stable with respect to such density fluctuations, whereas there is a larger effect of the noise for the large mixing MSW solution. We show how the possible existence of solar matter density noise could be tested in the next generation of solar neutrino experiments, especially Borexino. Finally, we comment on how possible solar model uncertainties could affect our results. ",
+        "conclusions": "We have presented a comprehensive study of the effects of the matter density noise upon the MSW solution to the solar neutrino problem. We have adopted the wave function Schr\\\"odinger formalism to re-write the corresponding MSW evolution equations for the neutrino survival probabilities. The fluctuations weaken the efficiency of the MSW suppression in the adiabatic regime, whereas they are much less effective in the non-adiabatic regime. In our data fit we have shown that the MSW solution still exists for realistic levels of matter density noise $\\xi \\lsim 8\\%$. However, our $\\chi^2$ analysis has shown that the quality of the fit gets a little worse if these noisy matter perturbations are present. In any case the mass range determined from our fit for the small mixing MSW solution $4 \\times 10^{-6} \\mbox{eV}^2 <\\Delta m^2< 10^{-5}\\mbox{eV}^2$ is relatively stable at 90\\% C.L., whereas the mixing angle determination appears more sensitive to the assumed level of fluctuations, and shifts $\\sin^2 2 \\theta$ towards larger values up to  $10^{-2}$. These trends also hold for the case of sterile solar neutrino conversion. In the latter case we have found that in the presence of solar density noise the large mixing region gets somewhat improved statistical significance when compared with the noiseless case. However, it is remains highly disfavoured with respect to the small mixing MSW solution.  We have also explored the potential of the Borexino experiment to \"test\" the level of matter density fluctuations in the solar interior through the measurement of the $^7$Be neutrino flux, as depicted in Fig. 4.  Finally, we note that in our analysis we have neglected the details of the neutrino production distribution as a function of the distance to the solar centre. It is well known that this affects mainly the low energy $pp$ neutrinos \\cite{MSW}. As a result, the iso-signal curves we have obtained for the gallium experiments are somewhat less reliable in the position of the kink corresponding to $\\Delta m^2 \\lsim 2 \\times 10^{-6} $eV$^2$ marking the onset of $pp$ neutrino suppression and lying on the gallium non-adiabatic branch. However, this does not substantially affect the determination of the relevant regions where {\\sl all} solar neutrino data are explained through the MSW effect. This includes both small as well as large mixing MSW solutions.   \\vskip 1truecm"
+    },
+    "9602/astro-ph9602068_arXiv.txt": {
+        "abstract": "Very high energy gamma-ray emission from the BL Lac object Markarian 421 has been detected over three observing seasons on 59 nights between April 1992 and June 1994 with the Whipple 10-meter imaging Cherenkov telescope.  During its initial detection in 1992, its flux above 500 GeV was 1.6$\\times$10$^{-11}$photons cm$^{-2}$ s$^{-1}$. Observations in 1993 confirmed this level of emission.  For observations made between December 1993 and April 1994, its intensity was a factor of 2.2$\\pm$0.5 lower.  Observations on 14 and 15 May, 1994 showed an increase over this quiescent level by a factor of $\\sim$10 (Kerrick et al. 1995).  This strong outburst suggests that 4 episodes of increased flux measurements on similar time scales in 1992 and 1994 may be attributed to somewhat weaker outbursts.  The variability of the TeV gamma-ray emission from Markarian 421 stands in contrast to EGRET observations (Lin et al. 1994) which show no evidence for variability. ",
+        "introduction": "The central engines of active galactic nuclei (AGN) are presumed to be massive black holes whose accretion powers radiation over a wide range of the electromagnetic spectrum. More than a score of these have now been reported as sources of multi-hundred MeV/GeV photons (Fichtel et al. 1994), all of which are identified with the sub-class of AGN which are radio-loud and core dominated, and most of which are blazars. (See Dermer \\& Schlickeiser 1992, for a glossary of terms relevant to AGN).  The AGN reported by EGRET vary in brightness by a decade and range in redshift from 0.031 to above 2. The nearest such object is Markarian 421, a BL Lac object extensively observed at radio (Owen et al. 1978; Zhang \\& Baath 1990), UV/optical (Maza, Martin, \\& Angel 1978; Mufson et al. 1990), and X-ray frequencies (Mufson et al. 1990; Mushotzky et al. 1979; George, Warwick, \\& Bromage 1988).  Its parent galaxy has been identified as a giant elliptical (Ulrich et al. 1975; Mufson, Hutter, \\& Kondo 1989) and it was the first BL Lac established as an X-ray source (Ricketts, Cooke, \\& Pounds 1979). Using VLBI techniques, the radio fine structure of the source has been mapped with a resolution of 1 mas to show a core-jet-like structure which exhibits possible superluminal motion (Zhang \\& Bhaath 1990). Variability with a time scale of hours has been observed in X-rays (Giommi et al. 1990) and at lower energies (Xie et al. 1988). Historically the source has exhibited low amplitude variability on the scale of months to years in the radio emission (Aller \\& Aller 1995) but so far there is no evidence for variability in the gamma-ray emission in the high energy regime (100 MeV - 10 GeV, Lin et al. 1994).  A comprehensive summary of recent theoretical work relevant to the blazars is given by von Montigny et al. (1995).  Although the ultimate origin of the blazar power is the central engine, high energy gamma-ray emission is likely to be beamed from a jet of highly relativistic particles.  (If one assumes isotropy for the emission, enormous luminosities, in some cases more than 10$^{49}$ erg/s, would be required.)  Furthermore any gamma rays at GeV-TeV energies that may be created in the inner regions of the source would be readily absorbed via the pair production process.  Because of the likely jet geometry associated with the high energy emission process, time variations observed at a single energy, if present, do not place strong constraints on the region of emission since there are many ways in which that variability may arise.  However, correlated time scale variations at different energies (e.g. hard synchrotron and TeV) are generally useful and are the predicted consequence of models in which the gamma rays are produced by inverse Compton scattering.  It has been pointed out by Gould and Schr\\'eder (1967) and recently reformulated by several authors (Steck\\-er, De Jager, \\& Salamon 1992; Salamon, Steck\\-er \\& De Jager 1994; Dwek \\& Slavin 1994; MacMinn \\& Primack 1995) that the interaction of gamma rays with intergalactic infrared and optical photons produced by galaxy formation may severely attenuate the TeV photon emission from cosmologically distant objects and that evidence for such absorption (or lack thereof) may be used to constrain both the energy density of starlight and the distance of the sources. The dynamic energy range covered by the Whipple Observatory instrument is particularly sensitive to such effects for the assumed distance to Markarian 421 (124 Mpc, H$_0$=75 km s$^{-1}$ Mpc$^{-1}$) and an upper limit on the IR density field from our observation is derived in a separate paper (Biller et al. 1995). ",
+        "conclusions": "\\begin{sloppypar} The average gamma-ray flux observed during 1992 and 1993 corresponds to a photon luminosity of $\\sim 10^{43}$ erg s$^{-1}$ above 500 GeV assuming isotropic emission at a distance of 124 Mpc.  However, the fact that all identified extragalactic sources detected by EGRET can be associated with the blazar class of AGN suggests that the GeV emission is highly beam\\-ed, originating in the relativistic plasma outflow which forms the radio jets. Beamed emission of the gamma rays furthermore overcomes the strong photon-photon absorption implied by the large gamma-ray luminosities for the case of isotropic emission. If we adopt the picture of beamed emission for the TeV component, the true VHE gamma-ray luminosity decreases by a factor of $\\sim$ 10$^{-3}$ for beam opening angles of the order of a few degrees.  \\end{sloppypar}  Much attention has been devoted lately to the discussion of possible absorption of the TeV photons in the intergalactic radiation field (Stecker, De Jager, \\& Salamon, 1992, Salamon, Stecker, \\& De Jager, 1994) i.e. gamma rays with energies of 0.5 TeV - 5 TeV interacting with soft 0.05 - 0.5 eV photons to produce e$^+$e$^-$ pairs.  For a relatively nearby source like Markarian 421 (z=0.031) this is a rather small effect.  (The possible presence of an absorption effect in our original 1992 data (Mohanty et al. 1993) has been used by De Jager et al. (1994) to determine a value for the infrared energy density. A more conservative approach has been adopted by Biller et al. (1995) to establish an upper limit to the energy density.)  The question arises whether or not this process could be occurring in the vicinity of the AGN, that is to say, can gamma rays with TeV energies escape the photon field close to the central engine of the AGN?  For some models this requirement is a severe obstacle as explained for instance by Coppi, Kartje, \\& K\\\"onigl (1993). While their theory nicely models the non-thermal emission component by upscattering radiation (ambient or synchrotron-self-Compton produced) by a relativistic electron/proton beam, they need to place the TeV emission regime further from the core than the GeV emission in order to avoid absorption in the region of high energy density at small distances. In simultaneous observations during the May 1994 TeV outburst (Macomb et al. 1995), the EGRET instrument did not observe an increase in the photon emission from Markarian 421 above the previously detected level (Lin et al. 1994).  Many of the various models that have been proposed to explain the origin of the gamma-ray emission favor production of the gamma-rays by the inverse-Compton process. Beamed relativistic electrons in the jet upscatter low energy photons from synchotron radiation (Ghisellini \\& Maraschi, 1989; Marscher \\& Bloom, 1992) or other ambient sources (Blandford 1993; Dermer, Schlickeiser, Mastichiadis, 1992; Sikora, Begelman, Rees, 1994) to gamma ray energies. These models imply a strong correlation between the production of radiation across the non-thermal waveband. Any change in the electron flow or electron distribution translates into a variation in the emission at other wavelengths. Differences in the relation of variability at given frequencies reflect distinct production regions. The inhomogeneous synchrotron-self-Compton model by Maraschi, Ghisellini, and Celotti (1992) for instance places the production of gamma-rays at the inner part of the jet while hard X-rays (as well as radio and IR emission) are produced in the outer jet part and thus fluctuations in the gamma-ray band are predicted to lead X-ray variability by typically a day.  In a similar model by Sikora, Begelman, and Rees (1994), the entire gamma-ray spectrum is produced within the same jet region through Comptonization of ambient radiation (from dust near the jet) by relativistic electrons. The authors point out that the problem of TeV opaqueness can be overcome if the seed photons come from thermal IR emission by dust rather than UV photons, supported by the fact that while thermal UV excess is common for OVV quasars, it is only observed in very few BL Lac objects. This hints that the non-observation of EGRET detected sources other than Markarian 421 (Kerrick et al. 1995b) above a few hundred GeV may be a combination of an intrinsic source feature (BL Lac's vs. OVV's for example) and the increasing IR absorption towards higher redshifts. Clearly, this needs more observational evidence, some of which can be provided by future measurements of AGN emission spectra in the (10-100) GeV regime where both processes will potentially cut off the power-law spectrum observed at high energies.  In contrast to the leptonic models is the approach by Mannheim (Mannheim 1993; Mannheim \\& Biermann 1992) in which a shock accelerated ultra-re\\-la\\-tivis\\-tic proton population in the jet photoproduces $\\pi^0$'s and the observed gamma rays are then produced by synchrotron cascade reprocessing. In this model, the gamma-ray emission observed at GeV energies extends naturally into the TeV regime. No predictions about the correlation of variability in the TeV gamma-ray emission to the GeV and X-ray emission are made but it was pointed out by von Montigny et al. (1995) that a disturbance in the proton population can translate to large gamma-ray flares downstream from the cascade origin.  Although the Whipple collaboration has recently reported a somewhat weaker TeV AGN source, Markarian 501 (Quinn et al. 1995), Markarian 421, is, so far, the best probe to test models at the highest accessible energies. The detection of a TeV component in the Markarian 421 emission as well as the observation of day scale variability already constrains models of gamma-ray production in relativistic jets. Additional, extended simultaneous observations of Markarian 421 over the optical, UV, GeV, and TeV energy regime are crucial to establish the correlations of non-thermal photon emission with the ultimate goal to understand the gamma-ray production and emission process in blazars."
+    },
+    "9602/astro-ph9602012_arXiv.txt": {
+        "abstract": "By comparing neutrino fluxes and central temperatures calculated  from 1000 detailed numerical solar models, we derive improved scaling laws  which show how each of the neutrino fluxes depends upon the central temperature (flux $\\propto T^m$); we also estimate uncertainties  for the temperature  exponents. With the aid of a one-zone model of the sun, we  derive expressions for the temperature exponents of  the neutrino fluxes. For the most important neutrino fluxes, the exponents calculated with the one-zone model agree to within 20\\% or better with the exponents extracted from the detailed numerical models. The one-zone model provides a  physical understanding of the temperature dependence of the neutrino fluxes. For the $pp$ neutrino flux, the one-zone model explains the (initially-surprising) dependence of the flux upon a negative power of the temperature and  suggests a new functional dependence. This new function makes explicit the strong anti-correlation between the $^7$Be and $pp$ neutrino fluxes.  The one-zone model also predicts successfully the average linear relations between  neutrino fluxes, but cannot predict the appreciable scatter in a $\\Delta \\phi_i/\\phi_i$ versus $\\Delta \\phi_j/\\phi_j$  diagram. ",
+        "introduction": "\\label{sec:intro}  Deciphering the solar neutrino problem offers the combined challenge of understanding the structure of the solar interior and understanding the nature of neutrino interactions. The consensus view at present, in part based upon temperature scalings discussed in this paper, is that the measured solar neutrino fluxes reported in the four operating experiments  cannot be explained by hypothesizing changes in standard solar models (SSMs).  The most plausible explanations, with the currently available data, require some extension of the standard model of electroweak interactions \\cite{b90,f91,p91,b93,bl93,w93,c93,h94,c94,c94pr,ber94,b94,pa95,fol95,h95,b95}.  The long-standing discrepancy between the observed and the predicted neutrino fluxes has motivated the study of many non-standard solar models, which are in  most  cases  {\\it ad hoc} perturbations of  the standard solar model. For many of the proposed  changes of SSM input parameters (e.g. nuclear cross sections, element abundances, and opacities), the predicted neutrino fluxes are approximately  characterized by a single derived model parameter, the central temperature, $T$. For small variations of input parameters, the neutrino fluxes and the central temperature of a detailed solar model can be  related by a power law of the \\hbox{form\\cite{bah88}} \\begin{equation} \\label{phin} \\phi \\propto T^m,~~~~~m = {d  \\log \\phi \\over d \\log T}. \\end{equation} For the fundamental $pp$ neutrinos, $m \\sim  -1$ and for the important ${\\rm ^8B}$ neutrinos, $m \\sim +20$.  The temperature dependences indicated in Eq.~(\\ref{phin}) are obtained from precise calculations with complex stellar evolution codes that solve coupled partial differential equations.  The results of these calculations are well known and have been used in many previous analyses of the implications of solar neutrino experiments.  Our goal is to show how these simple dependences result from the basic physics of the problem and to what extent the parameterization in terms of a central temperature is sufficient to characterize the neutrino fluxes.  The most initially suprising result of the stellar evolution calculations is that the magnitude of the $pp$ neutrino flux dependence depends inversely upon the value of the central temperature.  We shall see that even this result  has a simple, quantitative explanation in terms of the nuclear physics of the energy generation process.   The  scaling  with the central solar temperature can be used to evaluate neutrino fluxes for small deviations from the standard solar model. If a model is known to have a slightly different central temperature than the SSM, the neutrino fluxes can be estimated without detailed numerical calculations. Many classes of non-standard models (involving, e.g., rapid rotation in the solar interior, some variations in  nuclear cross sections, or the existence of a strong magnetic field in the solar core), reduce the central temperature. A quantitative determination of the reduction in the  central temperature implied by a specified change in the input physics requires detailed numerical modeling. With the aid of the temperature scaling laws, neutrino fluxes from non-standard solar models can be investigated over  large parameter ranges.  Using previously determined scaling laws of neutrino fluxes with central temperature, $T$\\cite{c93,bah88,bah89}, several authors\\cite{bl93,c93,h94,c94,c94pr} have studied  non-standard solar  models and have compared the predicted neutrino fluxes with the available experimental data from the four operating solar neutrino experiments, Homestake, Kamiokande, GALLEX, and SAGE\\cite{davis93,suzuki95,abdurashitov94,anselmann94}. These authors\\cite{bl93,c93,h94,c94,c94pr} show  that it is impossible to reconcile the  data from the four operating solar neutrino experiments with the neutrino fluxes predicted  by changes in $T$. They conclude that  non-standard solar models which have different central temperatures than the standard model are unlikely to solve the solar neutrino problem.  In the applications described above, the conclusions depend to some extent upon  the extrapolation of the temperature scalings to a larger temperature range than was covered in the original numerical models from which the scaling laws were derived\\cite{bah88}. We note that Castellani, Degl'Innocenti, and Fiorentini\\cite{c93} (see also references\\cite{c94,c94pr})  find  good agreement between the neutrino fluxes calculated from their non-standard, but detailed,  solar models and the fluxes obtained by scaling with respect to the central solar temperature.  The analysis in the present paper provides a physical  justification for the use of the temperature scalings, even for relatively large changes in the central solar temperature. We present improved determinations for the temperature exponents and estimates for their uncertainties.  We also give exponents for four minor neutrino fluxes for which temperature dependences were not previously available. By construction, the derived scaling law for the fundamental $pp$ neutrino flux is consistent with the observed solar luminosity.  Our results are complementary to the powerful numerical techniques of Castellani et al.\\cite{c93,c94,c94pr}, who have shown that several non-standard solar models have a homologous temperature dependence, and to the insightful physical analysis of Bludman\\cite{bludman96}, who has argued that it is a reasonable approximation to regard the solar interior as a single region that can be largely described by a single parameter, the central solar temperature. Taken together, these previous investigations provide strong motivation for taking seriously  the predictions of a single-zone solar model.   In what follows, we shall refer frequently to different fusion reactions in the $pp$ chain.  For convenience, we summarize in Table~\\ref{ppreaction} the principal reactions in the $pp$ chain.       In Sec.~II of this paper, we report scaling laws (with uncertainty estimates for the exponents) for all the solar neutrino fluxes that were included in the 1000 numerical solar models calculated in the Monte Carlo study  of Bahcall and Ulrich\\cite{bah88}. We present, in Sec.~III, a simple one-zone model for the sun which accounts well for the numerically-derived scaling laws. This model motivates the use of the new functional form for the temperature dependence of the $pp$ neutrino flux.  We discuss in Sec.~IV the Fogli-Lisi sum rule\\cite{fol95} on the temperature exponents. In Sec.~\\ref{sec:correlations}, we display the correlations found between the values of the principal neutrino fluxes in the 1000 solar models and show to what extent the one-zone model can account for these correlations. Finally, in Sec.~\\ref{sec:conclusion} we summarize and discuss our main conclusions. ",
+        "conclusions": "\\label{sec:conclusion}   The most interesting result of our study is the understanding it provides of the negative temperature dependence of the $pp$ neutrino flux.  The empirical fact that the $pp$ flux decreases with increasing central temperature, contrary to the trend found with all other solar neutrino fluxes, has been known since 1988\\cite{bah88,bah89}, but has not previously been explained physically.  At first glance, this negative temperature dependence is counter-intuitive.  In Section~\\ref{subsec:ppTscaling}, we show that the negative temperature dependence is a simple consequence of the fact that at higher temperatures only one $pp$ neutrino is produced per (approximately) $25$ MeV communicated to the star(via fusion of four protons), whereas at lower temperatures two $pp$ neutrinos are produced per $25$ MeV. In other words, at lower temperatures the $^3$He-$^3$He fusion termination reaction(which requires two $pp$ reactions) predominates whereas at higher temperatures the $^3$He-$^4$He reaction is faster (and requires only one $pp$ reaction).  The total energy per unit time communicated to the star must equal the observed solar luminosity, independent of the assumed central temperature. Thus,  as the temperature increases and  more of the nuclear fusion is accomplished by the $^3$He-$^4$He reaction, fewer $pp$ neutrinos are produced (and more $^7$Be and $^8$B neutrinos are created).  In order to obtain a simple physical understanding of the temperature scalings and the correlations between the different neutrino fluxes, we have adopted a one-zone model for the interior of the sun.  This model is characterized by a fixed central temperature, $T_c$, and a total luminosity that is equal to the observed solar luminosity. Given the emphasis in the current literature on calculating ever more precise solar models, with hundreds of different mass shells, it is gratifying and surprising that the one-zone  model accounts semi-quantitatively for some of the most often used results of the detailed model calculations. Moreover, the one-zone model predicts the average correlations found between the different neutrino fluxes, a bonus in insight that was not possible to anticipate without detailed study of the simple model.   Figure~\\ref{functions} and Figure~\\ref{powerlaw} display the dependence upon central temperature of the 1000 detailed solar models used in the Bahcall-Ulrich Monte Carlo study\\cite{bah88,bah89} of theoretical uncertainties in the predicted solar neutrino fluxes.  We have used these data to determine average temperature exponents, $m$, for all of the solar neutrino fluxes, where by assumption $\\phi \\propto T^m$. The exponents determined here are obtained by  a formal best-fitting technique and are to be preferred to the previously-estimated exponents\\cite{bah88,bah89} inferred less formally from these same data; the previously-estimated exponents have been widely used in the literature. We have also estimated, for the first time so far as we know, uncertainties in the inferred temperature exponents.  Figure~\\ref{correlationbe7pp} and Figure~\\ref{correlationb8be7pp} show the correlations, found in the Monte Carlo study, between the different individual neutrino fluxes. These correlations reflect the fact that when one neutrino flux is increased or decreased, there is likely to be a corresponding change in the values of the other neutrino fluxes.  These correlations must be taken into account when comparing the results of theoretical solar model calculations, including their uncertainties, with solar neutrino experiments. The only precise way to include the correlations displayed in Figure~\\ref{correlationbe7pp} and Figure~\\ref{correlationb8be7pp} is to use the complete set of calculated neutrino fluxes in the theoretical analysis(cf. \\cite{haxton89}). Various practical approximations to this rather cumbersome method have been discussed in the literature (see, for example, \\cite{w93,h94,fol95}).    The temperature exponents calculated with the aid of the one-zone model agree with the exponents inferred from the Monte Carlo study of precise solar models to an accuracy of 20\\% or better for the three most important solar neutrino fluxes: $pp$, $^7$Be, and $^8$B.  The results are shown in Table~\\ref{tableexp}, which compares the exponents calculated with the one-zone model with the results obtained from the detailed solar models. Figure~\\ref{scalingvstemp} shows that the scaling exponents calculated in the one-zone solar model are not strongly dependent upon the assumed characteristic central temperature, $T_c$ (taken here to be $T_c = 14\\times 10^6$K).  The quantitative agreement between the results of the one-zone model and the detailed models is impressive given the fact that the temperature exponents vary from $m \\sim -1$ for $pp$ neutrinos to $m \\sim +24$ for $^8$B neutrinos.  The physical insight provided by the one-zone model suggests a new form for the temperature dependence of the $pp$ neutrino flux, which is given in Eq.~(\\ref{preferred}).  In this form, the variation of the $pp$ neutrino flux is, for all temperatures, consistent with the observed solar luminosity, since it was derived by considering the relation between the solar luminosity, Eq. (2), and the $pp$ neutrino flux, Eq. (3).  Moreover, the formula for the $pp$ neutrino flux, Eq.~(\\ref{ppfluxlum}), provided by the one-zone model makes explicit the close correlation between the $^7$Be and $pp$ neutrino fluxes that is manifest in Figure~\\ref{correlationbe7pp}. The expression  used in this paper for the $pp$ neutrino flux, Eq.~(\\ref{preferred}), was derived by considering the relation between the solar luminosity, Eq.~(\\ref{L}), and the $pp$ rate, Eq.~(\\ref{ppnucrate}). Physically, the strong correlation exists because the $^7$Be neutrino flux is proportional to the rate of reaction 5 of  Table~\\ref{ppreaction} and the $pp$ neutrino flux is proportional to a constant minus the rate of reaction 5 (if we neglect the small contribution from CNO neutrinos).  The one-zone model also accounts quantitatively for the average correlation, shown in Figure~\\ref{correlationb8be7pp}, between the $^8$B and $^7$Be neutrino fluxes, and between the $^8$B and $pp$ neutrino fluxes.  No simple model can, however, account in detail for the scatter in the correlation plots shown in Figure~\\ref{correlationbe7pp} and Figure~\\ref{correlationb8be7pp}. The Monte Carlo experiments simulate  uncertainties in many different parameters; the power-law fits in the figures represent only the average response of the neutrino fluxes to the changes in all the individual  parameters. For analyses requiring a precise assessment of the correlations between the different neutrino fluxes, a Monte Carlo study of detailed solar models is required.  What have we learned from this study?  Improved temperature exponents for the neutrino fluxes are now available, with estimates for the uncertainties in the exponents.  A static one-zone model of the sun accounts for the essential features of the temperature scaling of the neutrino fluxes and even describes well the average correlations between the fluxes.  The model does not provide a precise description of the temperature dependences nor of the correlations between the different fluxes. The exponents derived from the one-zone mode model do not satify precisely the sum rule derived from the measured solar luminosity.  The fundamental reason that the one-zone model does not account accuractely for all of the known results is that in precise solar models each  neutrino flux is  produced in a different range of temperatures.  One cannot represent the results of different temperature ranges by a single parameter, $T_c$.  In the future, a new Monte Carlo study must be undertaken to determine the temperature scalings and the correlations between the neutrino fluxes when, as required by helioseismological measurements\\cite{bah92}, diffusion is taken into account in the solar model calculations.  The analysis of Bludman\\cite{bludman96} suggests that the effects of diffusion may alter the inferred temperature exponents by a non-negligible amount when compared to the values given in this paper, which are obtained from detailed solar models that do not include diffusion\\cite{bah88}. A Monte Carlo study is now underway that will create 1000 solar models that include diffusion and other recent refinements of the stellar model\\cite{BPfuture}."
+    },
+    "9602/astro-ph9602140_arXiv.txt": {
+        "abstract": "The cluster A3627 was recently recognized to be a very massive, nearby cluster in a galaxy survey close to the galactic plane. We are reporting on ROSAT PSPC observations of this object which confirm that the cluster is indeed very massive. The X-ray emission detected from the cluster extends over almost 1 degree in radius. The X-ray image is not spherically symmetric and shows indications of an ongoing cluster merger. Due to the strong interstellar absorption the spectral analysis and the gas temperature determination are difficult. The data are consistent with an overall gas temperature in the range 5 to 10 keV. There are signs of temperature variations in the merger region.  A mass estimate based on the X-ray data yields values of $0.4 - 2.2 \\cdot 10^{15}$ \\msu \\ if extrapolated to the virial radius of $3 h_{50}^{-1}$ Mpc. In the ROSAT energy band (0.1 - 2.4 keV) the cluster emission yields a flux of about $2 \\cdot 10^{-10}$ erg s$^{-1}$ cm$^{-2}$ which makes A3627 the 6$^{th}$ brightest cluster in the ROSAT All Sky Survey. The cluster was missed in earlier X-ray surveys because it was confused with a neighbouring X-ray bright, galactic X-ray binary (1H1556-605). The large X-ray flux makes A3627 an important target for future studies. ",
+        "introduction": "In a deep galaxy survey  behind the Milky Way ($|b| \\le 10\\deg $) the cluster A3627 was recently found by Kraan-Korteweg et al. (1996) to be a very rich, nearby cluster of galaxies rivaling the prominent Perseus and Coma clusters in its mass and galaxy content. The observed recession velocity of 4882 km, s$^{-1}$ corresponding to a velocity of 4655  km s$^{-1}$ with respect to the Local Group places A3627 at a distance of about 93 Mpc (we are using a Hubble constant of $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ in this paper). Therefore the cluster is one of the most prominent mass concentrations in the local Universe. Kraan-Korteweg \\et\\ (1996) found a galaxy velocity dispersion of 903 km s$^{-1}$ and attributed a mass of $5 \\cdot 10^{15}$~\\msu\\ to this cluster. Although A3627 has been cataloged by Abell, Corwin \\& Olowin (1989) as a richness class 1 cluster with 59 galaxies it attracted little attention in the past {\\sl because} of its proximity to the galactic plane ($l = 325\\deg $, $b = -7\\deg $). It has also not been detected in previous X-ray surveys (e.g. Piccinotti \\et\\ 1982, Kowalski \\et\\ 1984, Lahav \\et\\ 1989).  We have observed this prominent cluster with the ROSAT PSPC. The results of this observation support the finding of Kraan-Korteweg \\et\\ (1996) that A3627 is a very massive cluster. In this paper we give a detailed report on the X-ray properties of A3627 and describe its structure as revealed by the ROSAT images. Section 2 provides an account of the X-ray morphology. The X-ray spectral analysis is discussed in Section 3. Section 4 describes the mass determination. In Section 5 these results are discussed and Section 6 provides a summary. ",
+        "conclusions": "The low galactic latitude cluster A3627 ($b = 7\\deg $) which received little attention in the past is the 6$^{th}$ brightest cluster in the sky in the ROSAT energy band. The X-ray flux is high enough that it has even been seen in the Uhuru sky survey but it was confused with a galactic low mass X-ray binary which lies close to it and has a similar X-ray flux.  A3627 has a mass around $10^{15}$ \\msu. Therefore it is almost as massive as the prominent Perseus and Coma clusters but, with a redshift of $z =  0.016$, it is even closer. The cluster is an interesting merger system and should be one of the most interesting targets to study cluster mergers with spatially resolved X-ray spectroscopy."
+    },
+    "9602/hep-ph9602369_arXiv.txt": {
+        "abstract": "A brief sketch is made of the present observational status of neutrino properties, with emphasis on the hints from solar and atmospheric neutrinos, as well as cosmological data on the amplitude of primordial density fluctuations. Implications of neutrino mass in particle accelerators, astrophysics and cosmology are discussed. ",
+        "introduction": "\\vskip .2cm  It is beyond any doubt that, although very successful, our present standard \\21 model leaves open too many fundamental issues in particle physics to be an ultimate theory of Nature. One of the most fundamental ones refers to the masses and 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 atmospheric neutrinos, as well as the dark matter in the universe. The implications of detecting nonzero neutrino masses could be very far reaching for the understanding of fundamental issues in particle physics, astrophysics, as well as the large scale structure of our universe.  One interesting aspect of most extensions of the standard model where neutrino have non-vanishing masses is that they may affect the physics of the electroweak sector in a very remarkable way, which can be experimentally tested. Some of the ways to probe the corresponding physics at accelerator as well as underground experiments will be described.  \\subsection{Laboratory Limits } \\vskip .2cm  The most model-independent of the laboratory limits on \\neu mass are those that follow purely from kinematics, given as \\cite{PGG95} \\beq \\label{1} m_{\\nu_e} \t\\lsim 5 \\: \\rm{eV}, \\:\\:\\:\\:\\: m_{\\nu_\\mu}\t\\lsim 250 \\: \\rm{keV}, \\:\\:\\:\\:\\: m_{\\nu_\\tau}\t\\lsim 23 \\: \\rm{MeV} \\eeq The improved limit on the \\ne mass from beta decays was recently given by Lobashev \\cite{Erice}, while that on the \\nt mass comes from the ALEPH experiment \\cite{eps95} and may be substantially improved at a future tau-charm factory \\cite{jj}.  In addition, there are limits on neutrino masses that follow from the non-observation of neutrino oscillations \\cite{granadaosc}. The 90\\% confidence level (C.L.) exclusion contours of neutrino oscillation parameters in the 2-flavour approximation are given in Fig. \\ref{oscil}, taken from ref. \\cite{jjc}. Improvements are expected from the ongoing CHORUS and NOMAD experiments at CERN, with a similar proposal at Fermilab \\cite{chorus}. There are also good prospects for substantial progress at future long baseline experiments using CERN and Fermilab neutrino beams aimed at the Gran Sasso and Soudan underground facilities, respectively. \\bef \\vglue -1.5cm \\psfig{file=oscil.ps,height=11cm,width=7cm} \\vglue -3cm \\caption{Limits on oscillation parameters.} \\vglue -0.5cm \\label{oscil} \\eef  Another important limit follows from the non-observation of neutrino-less double beta decay - ${\\beta \\beta}_{0\\nu}$ - i.e. the process by which an $(A,Z-2)$ nucleus decays to $(A,Z) + 2 \\ e^-$. This process would arise from the virtual exchange of a Majorana neutrino from an ordinary double beta decay process. Unlike the latter, the neutrino-less process violates lepton number and its existence would indicate the Majorana nature of neutrinos. Because of the phase space advantage, this process is a very sensitive tool to probe into the nature of neutrinos. In fact, as shown in ref. \\cite{BOX}, a non-vanishing ${\\beta \\beta}_{0\\nu}$ decay rate requires \\neus to be majorana particles, {\\sl irrespective of which mechanism induces it}. This establishes a very deep connection which, in some special models, may be translated into a lower limit on the \\neu masses. The negative searches for ${\\beta \\beta}_{0\\nu}$ in $^{76}$Ge and other nuclei leads to a limit of about two eV \\cite{Avignone} on the weighted average \\neu mass parameter $\\VEV{m}$ \\beq \\label{bb} \\VEV{m} \\lsim 1 - 2 \\ eV \\eeq depending to some extent on the relevant nuclear matrix elements characterising this process \\cite{haxtongranada}. Improved sensitivity is expected from the upcoming enriched germanium experiments. Although rather stringent, this limit in \\eq{bb} may allow relatively large \\neu masses, as there may be strong cancellations between the contributions of different neutrino types. This happens automatically in the case of a Dirac \\neu as a result of the lepton number symmetry \\cite{QDN}.  \\subsection{The Cosmological Density Limit } \\vskip .2cm  In addition to laboratory limits, there is a cosmological bound that follows from avoiding the overabundance of relic neutrinos \\cite{KT} \\beq \\label{RHO1} m_{\\nu_\\tau} \\lsim 92 \\: \\Omega_{\\nu} h^2 \\: eV\\:, \\eeq where $\\Omega_{\\nu} h^2 \\leq 1$ and the sum runs over all isodoublet neutrino species with mass less than $O(1 \\: MeV)$. Here $\\Omega_{\\nu}=\\rho_{\\nu}/\\rho_c$, where $\\rho_{\\nu}$ is the neutrino contribution to the total density and $\\rho_c$ is the critical density. The factor $h^2$ measures the uncertainty in the determination of the present value of the Hubble parameter, $0.4 \\leq h \\leq 1$. The factor $\\Omega_{\\nu} h^2$ is known to be smaller than 1.  For the $\\nu_{\\mu}$ and $\\nu_{\\tau}$ this bound is much more stringent than the corresponding laboratory limits \\eq{1}.  Recently there has been a lot of work on the possibility of an MeV tau neutrino \\cite{ma1,SF}. Such range seems to be an interesting one from the point of view of structure formation \\cite{ma1,SF}. Moreover, it is theoretically viable as the constraint in \\eq{RHO1} holds only if \\neus are stable on the relevant cosmological time scales. In models with spontaneous violation of total lepton number \\cite{CMP} there are new interactions of neutrinos with the majorons which may cause neutrinos to decay into a lighter \\neu plus a majoron, for example \\cite{fae}, \\beq \\label{NUJ} \\nu_\\tau \\ra \\nu_\\mu + J \\:\\: . \\eeq or have sizeable annihilations to these majorons, \\beq \\label{JJ} \\nu_\\tau + \\nu_\\tau \\ra J + J \\:\\: . \\eeq The possible existence of fast decay and/or annihilation channels could eliminate relic neutrinos and therefore allow them to be heavier than \\eq{RHO1}. The cosmological density constraint on neutrino decay lifetime (for neutrinos lighter than 1 MeV or so) may be written as \\beq \\tau \\ltap 1.5 \\times10^7 (KeV/m_{\\nu_\\tau})^{2} yr \\:, \\label{RHO2} \\eeq and follows from demanding an adequate red-shift of the heavy neutrino decay products. For neutrinos heavier than $\\sim 1 \\: MeV$, such as possible for the case of $\\nu_{\\tau}$, the cosmological limit on the lifetime is less stringent than given in \\eq{RHO2}.  As we already mentioned the possible existence of non-standard interactions of neutrinos due to their couplings to the Majoron brings in the possibility of fast invisible neutrino decays with Majoron emission \\cite{fae}. These 2-body decays can be much faster than the visible decays, such as radiative decays of the type $\\nu' \\ra \\nu + \\gamma$. As a result the Majoron decays are almost unconstrained by astrophysics and cosmology. For a more detailed discussion see ref. \\cite{KT}.  A general method to determine the Majoron emission decay rates of neutrinos was first given in ref. \\cite{774}. The resulting decay rates are rather subtle \\cite{774} and model dependent and will not be discussed here. The reader may consult ref. \\cite{V,fae}. The conclusion is that there are many ways to make neutrinos sufficiently short-lived that all mass values consistent with laboratory experiments are cosmologically acceptable. For neutrino decay lifetime estimates see ref. \\cite{fae,V,RPMSW}.  \\subsection{The Nucleosynthesis Limit} \\vskip .2cm  There are stronger limits on neutrino lifetimes and/or annihilation cross sections arising from cosmological nucleosynthesis considerations. If massive $\\nu_\\tau$'s are stable during nucleosynthesis ($\\nu_\\tau$ lifetime longer than $\\sim 100$ sec), one can constrain their contribution to the total energy density from the observed amount of primordial helium. This bound can be expressed through an effective number of massless neutrino species ($N_\\nu$). Using $N_\\nu < 3.4-3.6$, the following range of $\\nu_\\tau$ mass has been ruled out \\cite{KTCS91,DI93} \\begin{equation} 0.5~MeV < m_{\\nu_\\tau} < 35~MeV \\label{cons1} \\end{equation} If the nucleosynthesis limit is taken less stringent the limit loosens somewhat. However it has recently been argued that non-equilibrium effects from the light neutrinos arising from the annihilations of the heavy \\nt's make the constraint stronger and forbids all $\\nu_\\tau$ masses on the few MeV range.  One can show that if the \\nt is unstable during nucleosynthesis \\cite{unstable} the bound on its mass is substantially weakened translated as a function of the assumed lifetime \\cite{unstable}.  Even more important is the effect of neutrino annihilations \\cite{DPRV}. Fig. 2 gives the effective number of massless neutrinos equivalent to the contribution of massive neutrinos with different values of the coupling $g$ between $\\nu_\\tau$'s and $J$'s, expressed in units of $10^{-5}$. For comparison, the dashed line corresponds to the standard model $g=0$ case. One sees that for a fixed $N_\\nu^{max}$, a wide range of tau neutrino masses is allowed for large enough values of $g$. No \\nt masses below 23 MeV can be ruled out, as long as $g$ exceeds a few times $10^{-4}$. Such values are reasonable in many majoron models \\cite{fae,MASIpot3}. For more details see ref. \\cite{DPRV}. \\begin{figure} \\label{neq} \\psfig{file=bbnjj1.ps,height=6cm,width=7cm} \\vglue -.5cm \\caption{Contribution of a heavy \\nt to nucleosynthesis in terms of the equivalent number of massless neutrinos.} \\end{figure} In short one sees that the constraints on the mass of a Majorana $\\nu_\\tau$ from primordial nucleosynthesis can be substantially relaxed if annihilations $\\nu_\\tau \\bar{\\nu}_\\tau \\leftrightarrow JJ$ are present. More details in the talk by Pastor \\cite{Pastor}.  As a result of the above considerations one concludes that it is worthwhile to continue the efforts to improve present laboratory \\neu mass limits in the laboratory. One method sensitive to large masses is to search for distortions in the energy spectra of leptons coming from $\\pi, K$ weak decays such as $\\pi, K \\ra e \\nu$, $\\pi, K \\ra \\mu \\nu$, as well as kinks in nuclear $\\beta$ decays. ",
+        "conclusions": "\\vskip .2cm  Theory can not yet predict fermion masses, and neutrinos are no exception. Nevertheless neutrino masses are strongly suggested by present theoretical models of elementary particles. On the other hand, they seem to be required to fit together present astrophysical and cosmological observations. Neutrino mass studies in nuclear $\\beta$ decays and peak search experiments should continue. Searches for $\\beta \\beta_{0\\nu}$ decays with enriched germanium could test the quasi-degenerate neutrino scenario of section 2.4.1. Underground experiments at Superkamiokande, Borexino, and Sudbury will shed more light on the solar neutrino issue. Oscillation searches in the \\ne $\\ra$ \\nt and \\nm $\\ra$ \\nt channels at accelerators should soon improve over the present situation illustrated in Fig. 1, while long baseline experiments both at reactors and accelerators are being considered. These will test the regions of oscillation parameters presently suggested by atmospheric \\neu data, shown in Fig. 4. Finally, new satellite experiments capable of measuring with better accuracy the cosmological temperature anisotropies at smaller angular scales than COBE, will test different models of structure formation, and presumably shed light on the possible role of neutrinos as dark matter.  If neutrinos are massive they could be responsible for a wide variety of implications, covering an impressive range of energies. These could be probed in experiments performed at underground installations as well as particle accelerators. For example, we saw in section 3 how neutrinos may produce new signatures at high energy physics collider experiments. Although indirectly, these may test neutrino properties in an important way and will therefore complement the efforts at low energies as well as the non-accelerator studies."
+    },
+    "9602/astro-ph9602094_arXiv.txt": {
+        "abstract": "We continue our study of the astrophysical implications of the linear potential $V(r)=-\\beta c^2/r+\\gamma c^2 r/2$ associated with fundamental gravitational sources in the conformal invariant fourth order theory of gravity which has recently been advanced by Mannheim and Kazanas as a candidate alternative to the standard second order Einstein theory. We provide fitting to the rotation curves of an extensive and diverse set of 11 spiral galaxies whose data are regarded as being particularly reliable. Without the assumption of the existence of any dark matter the model is found to fit the shapes of the rotation curves extremely well, but with a pattern of normalizations which proves to be very instructive. ",
+        "introduction": "During the last few years Mannheim and Kazanas (Mannheim 1990, 1992, 1993a, b, 1994, 1995a, b, c; Mannheim and Kazanas 1989, 1991, 1994; Kazanas and Mannheim 1991) have been exploring conformal gravity (viz. gravity based on invariance of the geometry under any and all local conformal stretchings of the form $g_{\\mu \\nu} (x) \\rightarrow \\Omega (x) g_{\\mu \\nu} (x)$) as a covariant candidate alternative to the standard Newton-Einstein gravitational theory. Their study has entailed both the examining of the formal structure of the theory and the identification of its possible observational astrophysical implications. In particular they found (Mannheim and Kazanas 1989; see also Riegert 1984) the most general, exact all order metric exterior to a static, spherically symmetric source such as a star in the theory, viz. (in standard static coordinates) \\begin{equation} -g_{00}= 1/g_{rr}=1-\\beta(2-3 \\beta \\gamma )/r - 3 \\beta \\gamma + \\gamma r - kr^2 \\label{Eq. (1.1)} \\end{equation} where $\\beta, \\gamma$ and $k$ are three integration constants. Subsequently, they also found (Mannheim and Kazanas 1994) the associated exact interior solution and established its consistency with the exterior one, while also showing that only the coefficients of the $r$ and $1/r$ terms in the general Eq. (\\ref{Eq. (1.1)}) were actually related to properties of the source. As we can thus see, the metric of Eq. (\\ref{Eq. (1.1)}) not only generalizes the potential of Newtonian gravity, it also generalizes the Schwarzschild solution of Einstein gravity as well, so that (for an appropriate choice of the coefficients of the $r$ and $1/r$ terms) the conformal theory can nicely recover the Newtonian potential and its familiar Einstein relativistic corrections on the small distance scale associated with the solar system, and then depart from the standard theory on the much larger distance scale associated with galaxies, this being precisely the distance scale where the standard Newton-Einstein theory can apparently only survive if galaxies contain copious amounts of dark matter. And, moreover, the first fitting (Mannheim 1993a) of the linear potential of Eq. (\\ref{Eq. (1.1)}) to a small set of four characteristic galaxies showed that the conformal theory could actually account for the relevant rotation curve data without the need to invoke dark matter at all.  As regards the actual possible existence of galactic dark matter, we note that neither the vigorous dark or faint matter searches of the OGLE (Udalski et al 1993, 1994), MACHO (Alcock et al 1993) and EROS (Aubourg et al 1993) gravitational microlensing collaborations nor those using the unprecedented optical sensitivity now available to the recently refurbished Hubble Space Telescope (Bahcall et al 1994) have so far been able to confirm the presence of the huge spherical dark matter galactic halo required of the standard theory.  At the very minimum one can say that these searches have certainly not yet achieved their intended goal of validating the standard picture, while at the maximum one can say that they have even actually thrown the entire picture into doubt. While it is of course far too early to contemplate abandoning the standard paradigm, nonetheless the current observational situation does demand a critical reappraisal of its two key components, namely the presumed existence of dark matter and the assumed validity of the Newton-Einstein gravitational theory on distance scales much larger than the solar system one on which it was first established. Moreover, the very assumption of the continuing validity of the standard theory on these much larger galactic distance scales represents a so far unjustified and possibly even dangerous extrapolation; with the very need for dark matter possibly even being an indicator that such an extrapolation is not in fact reliable. Since conformal gravity also reproduces the standard solar system wisdom while leading to a very different galactic extrapolation, it would thus appear to be a legitimately motivated gravitational theory whose eventual ultimate status can only be ascertained through consideration of its observational consequences. The present paper therefore sets out to explore further the observational implications of conformal gravity by applying it to a quite extensive and diverse 11 galaxy rotation curve sample. While this would appear to be a straightforward enough procedure, as we shall actually see, the fitting we present in this paper will lead us to a somewhat unanticipated conclusion. ",
+        "conclusions": ""
+    },
+    "9602/hep-ph9602260_arXiv.txt": {
+        "abstract": "The Hubble expansion of galaxies, the $2.73\\dK$ blackbody radiation background and the cosmic abundances of the light elements argue for a hot, dense origin of the universe --- the standard Big Bang cosmology --- and enable its evolution to be traced back fairly reliably to the nucleosynthesis era when the temperature was of $\\Or(1)$ MeV corresponding to an expansion age of $\\Or(1)$ sec. All particles, known and hypothetical, would have been created at higher temperatures in the early universe and analyses of their possible effects on the abundances of the synthesized elements enable many interesting constraints to be obtained on particle properties. These arguments have usefully complemented laboratory experiments in guiding attempts to extend physics beyond the Standard $SU(3)_{\\c}{\\otimes}SU(2)_{\\L}{\\otimes}U(1)_{Y}$ Model, incorporating ideas such as supersymmetry, compositeness and unification. We first present a pedagogical account of relativistic cosmology and primordial nucleosynthesis, discussing both theoretical and observational aspects, and then proceed to examine such constraints in detail, in particular those pertaining to new massless particles and massive unstable particles. Finally, in a section aimed at particle physicists, we illustrate applications of such constraints to models of new physics. ",
+        "introduction": "} %  There has been interest in problems at the interface of cosmology and particle physics for over thirty years (see Zel'dovich 1965), but it is only in the past decade or so that the subject has received serious attention (see B\\\"orner 1988, Collins \\etal 1989, Kolb and Turner 1990, Ellis 1993). Cosmology, once considered to be outside the mainstream of physics and chiefly of interest to astronomers and applied mathematicians, has become a {\\em physical} subject, largely due to the advances which have been made on the observational front (see Weinberg 1972, Peebles 1993). It has become increasingly clear that particle physicists can no longer afford to ignore the cosmological ``laboratory'', which offers a powerful probe of new physical phenomena far beyond the reach of terrestrial laboratories (see Steigman 1979, Dolgov and Zel'dovich 1981). Cosmological phenomena have thus come under detailed scrutiny by particle physicists, prompting deeper theoretical analyses (see Weinberg 1980, Wilczek 1991) as well as ambitious observational programmes (see Sadoulet 1992, Kolb and Peccei 1995).  The increasing interaction between particle physics and cosmology has largely resulted from the establishment of `standard models' in both fields which satisfactorily describe all known phenomena but whose very success, paradoxically, establishes them as intrinsically incomplete pictures of physical reality. Our reconstruction of the history of the universe in \\fref{cosmohist} emphasizes the interdependence of these models. The familiar physics of electromagnetism, weak interactions and nuclear reactions provide a sound basis for the standard Big Bang cosmology up to the beginning of the nucleosynthesis era, when the universe was about $10^{-2}\\sec$ old. The Standard $SU(3)_{\\c}{\\otimes}SU(2)_{\\L}{\\otimes}U(1)_{Y}$ Model (SM) of particle physics (see Cheng and Li 1984, Kane 1987), brilliantly confirmed by all experiments to date (see Burkhardt and Steinberger 1991, Bethke and Pilcher 1992), allows us to extrapolate back further, to $t\\sim10^{-12}$ sec. Two phase transitions are believed to have occurred in this interval, although a detailed understanding of their dynamics is still lacking (see Kapusta 1988, Yaffe 1995). The first is associated with the confinement of quarks into hadrons and chiral symmetry breaking by the strong interactions at $T_{\\c}^{\\qh}\\sim\\Lambda_{\\QCD}\\approx200\\MeV$, and the second with the spontaneous breaking of the unified electroweak symmetry to electromagnetism, $SU(2)_{\\L}{\\otimes}U(1)_{Y}\\,\\rightarrow\\,U(1)_{\\rm em}$ at $T_{\\c}^{\\EW}\\sim250\\GeV$, when all known particles received their masses through the Higgs mechanism. To go beyond this point requires an extension of the SM; indeed, the very success of the model demands such new physics.   Similarly, the standard cosmological model of an adiabatically expanding, homogeneous and isotropic universe requires extreme fine tuning of the initial conditions of the Big Bang, as emphasized by Dicke and Peebles (1979). The problem essentially consists of explaining why the universe is as old ($\\geqsim3\\times10^{17}\\sec$) or as large ($\\geqsim10^{28}\\cm$) as it is today, relative to the Planck time ($5.39\\times10^{-44}\\sec$) or the Planck length ($1.62\\times10^{-33}\\cm$), which are the appropriate physical scales governing gravitational dynamics. Although a resolution of this may have to await progress in our understanding of quantum gravity (see Penrose 1979, 1989), there has been enthusiastic response to the simpler solution proposed by Guth (1981), viz. that there was a period of non-adiabatic accelerated expansion or `inflation', possibly associated with a phase transition in the early universe (see Linde 1990). This has the additional advantage that it naturally generates a nearly scale-invariant `Harrison-Zel'dovich' spectrum of scalar density fluctuations (see Mukhanov \\etal 1992) which can seed the growth of the observed large-scale structure in the expanding universe (see Efstathiou 1990, Padmanabhan 1993). Another fundamental problem of the standard cosmology is that the observed abundance of baryonic matter is $\\sim10^{9}$ times greater than the relic abundance expected from a state of thermal equilibrium, while no antimatter is observed (see Steigman 1976), thus requiring a primordial asymmetry between matter and anti-matter. To generate this dynamically requires new physics to violate baryon number ($B$) and charge-parity ($CP$) at high temperatures, in an out-of-equilibrium situation to ensure time asymmetry (Sakharov 1967, see Cohen \\etal 1994, Rubakov and Shaposhnikov 1996). More recently, it has been recognized that baryons are probably a minor constituent of the universe, since all observed structures appear to be dominated by dark matter (see Binney and Tremaine 1987, Peebles 1993) which is probably non-baryonic. The growing interest in the early universe stems from the realization that extensions of physics beyond the SM naturally provide the mechanisms for processes such as inflation and baryogenesis, as well as new particle candidates for the dark matter. These exciting developments have been discussed in a number of schools and conferences (see e.g. Gibbons \\etal 1983, Setti and Van Hove 1984, Kolb \\etal 1986b, Piran and Weinberg 1986, Alvarez \\etal 1987, Hinchliffe 1987, De R\\'ujula \\etal 1987, Unruh and Semenoff 1988, Yoshimura \\etal 1988, Peacock \\etal 1990, Nilsson \\etal 1991, Sato and Audoze 1991, Nanopoulos 1991, Sanchez and Zichichi 1992, Akerlof and Srednicki 1993, Astbury \\etal 1995).  Such new physics is in fact necessary to address the theoretical shortcomings of the Standard Model itself (see Ross 1984, Mohapatra 1992). Its phenomenological success requires that the Higgs boson, which gives masses to all known particles, cannot itself be much more massive than its vacuum expectation value (VEV) which sets the electroweak (`Fermi') scale, $v\\equiv(\\sqrt{2}G_{\\F})^{-1/2}=246\\GeV$. This creates the `naturalness' or `hierarchy' problem, viz. why is the Higgs mass not pushed up to the Planck mass ($M_{\\Pl}\\,\\equiv\\,G_{\\rm N}^{-1/2}=1.221\\times10^{19}\\GeV$) due to the {\\em quadratically} divergent radiative corrections it receives due to its couplings to all massive particles?\\footnote{By contrast, it is `natural' for fermions to be light relative to the Planck scale since letting their masses go to zero reveals a chiral symmetry which tames the radiative corrections to be only logarithmically divergent; there is no such symmetry to `protect' the mass of a scalar Higgs boson.} Supersymmetry (SUSY) addresses this problem by imposing a symmetry between bosons and fermions which makes such radiative corrections cancel to zero. This requires all known particles (boson/fermion) to have supersymmetric (fermion/boson) partners distinguished by a new quantum number called $R$-parity; if it is conserved, the lightest supersymmetric particle would be stable. Supersymmetry must be broken in nature since known particles do not have supersymmetric partners of the same mass.  However the Higgs mass would still be acceptable if the scale of SUSY breaking (hence the masses of the supersymmetric partners) is not much beyond the Fermi scale. When such breaking is realized {\\em locally}, as in gauge theories, a link with general coordinate transformations, i.e. gravity, emerges; this is supergravity (SUGRA) (see Van Nieuwenhuizen 1981, Wess and Bagger 1993). Technicolour is an alternative approach in which the offending {\\em elementary} Higgs particle is absent (see Farhi and Susskind 1981, Kaul 1983); electroweak symmetry breaking is now seen as a dynamic phenomenon (see King 1995, Chivukula \\etal 1995), akin to the breaking of chiral symmetry by the strong interactions. However no realistic technicolour model has been constructed satisfying all experimental constraints, in particular the small radiative corrections to SM parameters measured at \\LEP (see Lane 1993).  Another conundrum is that $CP$ is known to be well conserved by the strong interactions, given the stringent experimental upper limit on the neutron electric dipole moment, whereas QCD, the successful theory of this interaction, contains an arbitrary $CP$ violating parameter. An attractive solution is to replace this parameter by a field which dynamically relaxes to zero --- the axion (see Kim 1987, Cheng 1988). This is a pseudo-Goldstone boson generated by the breaking of a new global $U(1)$ `Peccei-Quinn' symmetry at a scale $f_{a}$ (see Peccei 1989). This symmetry is also explicitly broken by QCD instanton effects, hence the axion acquires a small mass $m_{a}\\,\\sim\\,f_{\\pi}^2/f_{a}$ when the temperature drops to $T\\sim\\Lambda_{\\QCD}$. The mixing with the pion makes the axion unstable against decay into photons; negative experimental searches for decaying axions then constrain $f_{a}$ to be beyond the Fermi scale, implying that axions are light enough to be produced in stellar interiors. Considerations of stellar cooling through axion emission imply $f_{a}\\geqsim10^{10}$ GeV, which {\\em requires} the axion (if it exists!) to have an interesting cosmological relic density (see Raffelt 1990, Turner 1990).  Yet another motivation for going beyond the Standard Model is the unification of forces. Grand unified theories (GUTs) of the strong and electroweak interactions at high energies also provide a physical need for inflation in order to dilute the embarrassingly large abundance of magnetic monopoles expected to be created during the breaking of the unified symmetry (see Preskill 1984). Unification naturally provides for baryon and lepton number violation (see Langacker 1981, Costa and Zwirner 1986) which allows for generation of the cosmological baryon asymmetry (see Kolb and Turner 1983) as well as masses for neutrinos (see Mohapatra and Pal 1991). Recent data from \\LEP on the evolution of the gauge interaction couplings with energy indicate that such unification can occur at $M_{\\GUT}\\approx2\\times10^{16}$ GeV, but only in a (broken) supersymmetric theory with superparticle masses at around the Fermi scale (see Dimopoulos 1994, Ellis 1995). Moreover in such unified models, electroweak symmetry breaking via the Higgs mechanism is driven quite naturally by supersymmetry breaking (see Ib\\'a\\~nez and Ross 1993). A dynamical understanding of how supersymmetry itself is broken is expected to come from the theory of superstrings, the most ambitious attempt yet towards a finite quantum theory of gravity and its unification with all other forces (see Green \\etal 1987). Following the initial euphoria over the discovery of the anomaly-free heterotic superstring, progress has been difficult due to the problems of relating low energy physics to the higher dimensional world which is the natural domain of the string. However explicit examples of compactified four-dimensional strings have been constructed which reduce to a supersymmetric version of the Standard Model at low energies and also contain additional gauge bosons and gauge singlets which have only gravitational couplings to matter (see Dine 1988, 1990, Ib\\'a\\~nez 1994). Furthermore, the recent exciting discoveries of the special `duality' properties of string theories (see Giveon \\etal 1994, Polchinski 1996) have begun to provide important insights into the issue of supersymmetry breaking (see Zwirner 1996).  It is thus a common feature of new physics beyond the Fermi scale to predict the existence of new particles which are unstable in general but some of which may be stable by virtue of carrying new conserved quantum numbers. Moreover their generic weak interactions ensure a cosmologically significant relic density (see Primack \\etal 1988, Turner 1991). In addition, known particles such as neutrinos, although strictly massless in the Standard Model, may acquire masses from such new physics, enabling them also to be candidates for dark matter. Conventionally, particle physicists look for new physics either by directly attempting to produce the new particles in high energy collisions at accelerators or by looking for exotic phenomena such as nucleon instability or neutrino masses. In this context, the standard cosmology, in particular Big Bang nucleosynthesis (BBN), provides an important new testing ground for new physics and, indeed, in many cases, provides the {\\em only} ``experimental'' means by which the properties of new particles may be restricted (see Sarkar 1985, 1986). Whether or not one finds this satisfactory from a philosophical point of view, it is essential for this enterprise that we have the best possible understanding of the cosmological laboratory. This review presents the current status and is primarily aimed at particle physicists, although it is hoped that astrophysicists and cosmologists will also find it useful.  A decade or more ago, it was possible for reviewers (e.g. Steigman 1979, Dolgov and Zel'dovich 1981) to give a comprehensive discussion of all constraints on fundamental physics from cosmological considerations and many of the key papers could be found in one collection (Zee 1982). Subsequently several hundred papers on this subject have been published. For reasons of space we will restrict ourselves to a discussion of the constraints which follow from BBN alone but, for completeness, present all the neccessary cosmological background --- both theory and observations. Rather than engage in a detailed critique of every published work, we present a pedagogical discussion of the basic physics, together with a summary of the key observational inputs, so that readers can assess the reliability of these constraints. Raffelt (1990) has presented a model review of this form which deals with astrophysical methods for constraining novel particle phenomena. A similar discussion of all types of cosmological constraints, including those deduced from the observed Hubble expansion and the $2.73\\dK$ blackbody radiation as well as other radiation backgrounds, will appear in Sarkar (1997).  We begin by outlining in \\sref{stdcosm} the basic features of the standard Big Bang cosmological model and then discuss the thermodynamics of the early radiation-dominated era. In \\sref{primnucl} we present the essential physics of the BBN era and then discuss the observational data in some detail, highlighting the sources of uncertainty. We argue for the consistency of the standard model and briefly mention possible variations. This sets the stage for deriving general constraints in \\sref{cons} on both relativistic and non-relativistic hypothetical particles which may be present during nucleosynthesis. Finally, for the benefit of particle physicists we illustrate in \\sref{appl} how such cosmological arguments have complemented experimental searches for physics beyond the Standard Model, particularly in the neutrino sector, and also provided entirely new probes of such physics, e.g. technicolour and supersymmetry. We also discuss the implications for the nature of the dark matter.  It appears to be a widely held belief that cosmological data are not particularly accurate and the associated errors uncertain, so that the derived constraints cannot compare in reliability with those obtained in the laboratory. Although not entirely incorrect, this view is being increasingly challenged by modern observations; for example measurements of the background radiation temperature and anisotropy, the cosmic abundance of helium \\etc are now routinely quoted to several significant figures. Correspondingly there has been a growing appreciation of the systematic effects involved in the analysis of cosmological observations and careful attempts at their estimation. More importantly, cosmological data, even if more imprecise than accelerator data, are often much more {\\em sensitive} to novel particle phenomena; for example, even a crude upper limit on the present energy density of the universe suffices to bound the masses of relic neutrinos to a level which improves by several orders of magnitude over precise laboratory experiments. Nevertheless, one should be cautious about rejecting an interesting theoretical possibility on the basis of a restrictive cosmological constraint (e.g. the bound on the number of neutrino-like particles present during BBN) without a critical appreciation of the underlying assumptions. We have tried wherever possible to clarify what these assumptions are and to refer to expert debate on the issues involved. (In writing down numerical values where errors are not quoted, the symbols $\\sim$, $\\approx$ and $\\simeq$ indicate equality to within a factor of 10, factor of 2 and $10\\%$, respectively.)  Due to space limitations, the references are not comprehensive but do include the seminal papers and recent reviews from which the intervening literature can be traced. We have used `natural' units ($\\hbar=c=k_{\\B}=1$) although astronomical units such as year, megaparsec or Solar mass are given where convenient. (For reference, $1\\,{\\GeV}^{-1}=1.973\\times10^{-14}\\cm=6.582\\times10^{-25}\\sec$, $1\\,\\GeV=1.160\\times10^{13}\\dK=1.783\\times10^{-24}\\gm$, $1\\,\\Mpc=3.086\\times10^{24}\\cm$, $1\\,\\yr=3.156\\times10^{7}\\sec$, $1M_{\\odot}=1.989\\times10^{33}\\gm$.) Astronomical quantities are listed in Allen (1973) and Lang (1992), while clarification of unfamiliar astrophysical terms may be sought in the excellent textbooks by Shu (1981), Mihalas and Binney (1981) and Longair (1981). ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602080_arXiv.txt": {
+        "abstract": "We present an essentially complete, all-sky, X-ray flux limited sample of 242 Abell clusters of galaxies (six of which are double) compiled from ROSAT All-Sky Survey data. Our sample is uncontaminated in the sense that systems featuring prominent X-ray point sources such as AGN or foreground stars have been removed. The sample is limited to high Galactic latitudes ($|b| \\geq 20^{\\circ}$), the nominal redshift range of the ACO catalogue of $z \\leq 0.2$, and X-ray fluxes above $5.0 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the 0.1 -- 2.4 keV band. Due to the X-ray flux limit, our sample consists, at intermediate and high redshifts, exclusively of very X-ray luminous clusters.  Since the latter tend to be also optically rich, the sample is not affected by the optical selection effects and in particular not by the volume incompleteness known to be present in the Abell and ACO catalogues for richness class 0 and 1 clusters.  Our sample is the largest X-ray flux limited sample of galaxy clusters compiled to date and will allow investigations of unprecedented statistical quality into the properties and distribution of rich clusters in the local Universe. ",
+        "introduction": "Being the largest gravitationally bound systems to have decoupled from the Hubble expansion of the Universe, clusters of galaxies are ideal tracers of the formation and evolution of structure on the largest mass scales.  In the past, galaxy clusters were detected and classified mainly on the grounds of their optical appearance, the largest compilations of this kind being the catalogues of Abell (1958), Abell, Corwin \\& Olowin (1989, from here on ACO), and Zwicky and co-workers (1961--1968). In particular the Abell and ACO catalogues have been used extensively for statistical studies of cluster properties, but served also as finding lists for more detailed investigations of individual clusters.  The major drawback of optical compilations of this kind lies in the high risk of projection effects corrupting the sample. Since the cluster detection and selection is performed on the two-dimensional distribution of galaxies as it appears in projection on optical plates, fluctuations in the surface density of field galaxies as well as superpositions of poor clusters along the line of sight can lead to an overestimation of a system's richness.  On the other hand, poor clusters can be missed completely as they often do not contrast strongly with the background field. Several studies have investigated the statistical completeness and contamination of the Abell catalogue emphasizing the importance of projection effects for optically selected cluster samples (Lucey 1983, Sutherland 1988, Struble \\& Rood 1991).  The serious problem of projection effects can, however, be overcome completely by selecting clusters in the X-ray rather than in the optical. X-ray emission from a diffuse, gaseous intra-cluster medium (ICM) trapped and heated to temperatures of typically a few $10^7$ K is not only a sure indication that the system is indeed three-dimensionally bound. Being caused by ion-ion collisions in the ICM, the X-ray flux is proportional to the square of the ion density and thus much more peaked at the gravitational centre of the cluster than the projected galaxy distribution. This property requires clusters to be almost perfectly aligned along the line of sight in order to be mistaken for a single, more luminous entity, so that projection effects caused by such superpositions can be effectively neglected in the X-ray.  X-rays thus provide a very efficient and unbiased way of compiling cluster samples which, as far as statistical studies are concerned, will eventually supersede the present, optical catalogues. Early statistical cluster samples have been compiled from the X-ray data taken by the {\\sc Uhuru} (Schwartz 1978), {\\sc Ariel V} (McHardy 1978), and {\\sc Heao 1 A-2} satellites (Piccinotti et al.\\ 1982), all of which comprised about 30 clusters. Including clusters detected during the {\\sc Exosat} mission, Lahav et al.\\ (1989) and Edge et al.\\ (1990) presented an X-ray flux limited sample of 55 clusters, very similar in size to that compiled by Gioia et al.\\ (1990) from X-ray sources detected in the {\\sc Einstein} Medium Sensitivity Survey. Projects aimed at the compilation of % purely X-ray selected cluster samples based on the X-ray data collected during the ROSAT All-Sky Survey (RASS, Voges 1992) are well under way in both the northern (Allen et al.\\ 1992, Giacconi \\& Burg 1993, Crawford et al.\\ 1995) and the southern hemisphere (Guzzo et al.\\ 1995, De Grandi 1996, see also Pierre et al.\\ 1994); the two-point correlation function for RASS selected cluster samples compiled in smaller sky areas has been presented recently (Nichol et al.\\ 1994, Romer et al.\\ 1994) showing again the advantage of X-ray over optical selection procedures.  An overview of the various RASS cluster projects is given in B\\\"ohringer (1994).  In the following, we shall pursue an approach that investigates the X-ray properties of optically selected clusters and leads to an all-sky sample of the X-ray brightest Abell-type clusters of galaxies (XBACs)\\footnote{An early, if incomplete, version of this sample has been presented by Ebeling (1993).}  which is more than four times larger than the largest X-ray flux limited cluster sample published to date.  In this first article of a series we present the sample and discuss its statistical properties; in subsequent papers we shall establish the X-ray luminosity function of Abell-type clusters of galaxies (Ebeling et al., in preparation) as well as the cluster-cluster correlation function (Edge et al., in preparation), and investigate correlations between X-ray and optical properties of the XBACs.  Throughout this paper, we assume an Einstein-de Sitter Universe with $q_0 = 0.5$ and $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$. ",
+        "conclusions": "\\begin{figure*} \\epsfxsize=\\textwidth \\vspace*{2cm} \\hspace*{2cm} \\epsffile{flow.ps} \\mbox{}\\\\*[-3.5cm] \\caption[]{Flow diagram summarizing the compilation of the XBACs sample} \\label{xbacs_flow} \\end{figure*}  In this first paper of a series investigating the properties of the X-ray brightest Abell-type clusters detected in the ROSAT All-Sky Survey (RASS) we present the sample and describe its compilation.  The flow diagram in Fig.~\\ref{xbacs_flow} gives on overview of the compilation procedure, from the master list of RASS X-ray sources detected by the SASS to the final XBACs sample selected from VTP detections.  The VTP re-analysis of the RASS photon data is a crucial step in the compilation of the XBACs sample. VTP is a source detection and characterization algorithm developed by one of us (HE) to detect overdensities of essentially arbitrary shape in the exposure normalized surface density of RASS photons. Since VTP makes no assumptions about the source geometry (such as sphericity or radial flux profile) it is particularly well suited for the detection and characterization of extended and potentially irregular emission from clusters of galaxies.  The VTP analysis includes a correction of the initial count rates for low-surface brightness emission that has escaped direct detection. This correction takes into account the PSPC point spread function in the RASS and assumes a King law for the generalized radial surface brightness profile. The contribution from point sources embedded in the diffuse cluster emission is corrected for by adopting the geometrical mean of the raw and the corrected count rate values as the final count rate from diffuse ICM emission.  A comparison of these final VTP values with the best count rate determinations from archival, pointed PSPC observations for 100 selected clusters shows no systematic differences for count rates ranging from 0.1 to 10 s$^{-1}$; the total $1\\sigma$ uncertainty in the VTP count rates is about twice as high as that from photon statistics alone.  When comparing the VTP cluster count rates with the respective SASS values we find the SASS count rates to underestimate the true cluster brightness by typically a factor of 0.60 (median) with the 25$^{\\rm th}$ and 75$^{\\rm th}$ percentiles being 0.43 and 0.75.  Although the source extent is not used in the compilation of the XBACs sample, we investigate its significance for future compilations of purely X-ray selected samples. A comparison of the extent parameters returned by the SASS and VTP shows that, at redshifts greater than 0.12, both algorithms successfully detect some 70 per cent of our ACO clusters as extended X-ray sources. At lower redshifts, VTP has a higher efficiency (approaching 100 per cent at $z\\sim 0.05$). If the extent information from both algorithms are combined, essentially all XBACs are classified as extended by at least one method irrespective of redshift.  The sample compiled from VTP detections after non-cluster sources have been discarded consists of 277 clusters (or subclusters) above a flux limit of completeness (95 per cent) of $5.0 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$. Six more sources with X-ray fluxes below this value are included as subclusters of systems whose brightness exceeds the flux limit when both components are combined.  In order to remain consistent with previous studies, we limit our statistical sample to clusters at $|b| \\geq 20^{\\circ}$ and within the nominal volume of completeness of the ACO catalogue, i.e.\\ to redshifts lower than 0.2. This leads to our final sample of the 242 X-ray brightest Abell-type clusters (XBACs) from the ROSAT All-Sky Survey, 92 per cent of which have measured redshifts. The sample's overall completeness is 80 per cent; in view of the complicated selection procedure, we believe this figure to come close to the maximum value achievable for such a sample. As, at its high redshift end, our X-ray flux limited sample is intrinsically dominated by very X-ray luminous and optically rich clusters, the volume incompleteness in poor clusters of the underlying optical catalogue does not affect the XBACs.  The XBACs constitute the largest X-ray flux limited sample of clusters of galaxies compiled to date. By design, the sample is not only essentially complete but also free from contamination by non-cluster X-ray sources. In the following papers of this series the XBACs' properties are investigated in detail."
+    },
+    "9602/astro-ph9602049_arXiv.txt": {
+        "abstract": "Ratnatunga et al. (1995) have recently proposed that an object, HST 14176+5226, in the ``Groth-Westphal\" HST survey strip is an ``Einstein cross'' gravitational lens.  By chance, this object has been previously observed in the Canada-France Redshift Survey.  The candidate lens, catalogued as CFRS14.1311, is an elliptical galaxy at z = 0.81.  In addition, the spectrum shows a strong, spatially-extended, emission feature at 5342\\AA\\ that almost certainly originates from two of the four ``lensed'' images. We tentatively identify this emission line as Ly $\\alpha$ at z =3.4. A less prominent emission feature at 6822\\AA\\ may be C IV 1549.  Our data thus support the identification of this system as a new quadruple-image lens. ",
+        "introduction": "Ratnatunga et al. (1995; hereafter ROGI) have recently proposed that two objects found in archival Hubble Space Telescope images are new examples of ``Einstein cross'' gravitational lenses produced by distant elliptical galaxies.  In each case four very faint blue compact images are symmetrically placed around a brighter red galaxy.  One of these objects, HST 14176+5226, is in the ``Groth-Westphal\" GTO survey strip (Groth et al. 1994). By chance, this object has been previously observed during the Canada-France Redshift Survey (CFRS) as CFRS14.1311.  This fact is not as improbable as it may seem, since the 1415+52 CFRS field and the ``Groth-Westphal'' survey strip were both (and independently) co-located with the ultradeep radio field of Fomalont et al.  (1991) to allow identifications of these radio sources (see Hammer et al. 1995; CFRS VII).  As well as pointing out the previously-published redshift for the candidate lens (z = 0.81, Lilly et al. 1995b; CFRS III), we show in this {\\it Letter} that our spectrum of CFRS14.1311 also displays a spatially extended emission line at 5342 \\AA\\ that almost certainly originates from two of the four faint blue images. This line is likely to be Ly $\\alpha$ at z = 3.4. A less prominent emission feature at 6822\\AA\\ may be C IV at this redshift.  Our data thus strongly support the identification of this system as a new quadruply-imaged gravitational lens.  The secure redshift of the lens and the proposed redshift of the lensed object will allow more detailed modelling of this system. ",
+        "conclusions": "Our spectra thus indicate that the lensing elliptical galaxy CFRS 14.1311 has a redshift z = 0.809, and that the lensed object is a quasar or a strong emission-line galaxy at z = 3.4.  At these redshifts a simple singular isothermal lens model using the parameters derived by ROGI implies a lens dispersion of $\\sigma = 230$ km s$^{-1}$  and a mass-to-light ratio M/L $=$ 14h$^{-1}$ M$_{\\sun}$/L$_{\\sun}$. This indicates that CFRS14.1311 is probably a straightforward case of a simple, isolated lens.  Although Refsdal and Surdej(1994) list eight four-image lenses in their recent review, several of these have complex geometry of either the lensed images and/or the lensing object(s). CFRS14.1311 appears to be one of simpler cases involving an isolated normal elliptical galaxy acting as a lens for a background AGN, and thus could, in principle, provide cosmological constraints through further monitoring (e.g., Refsdal and Surdej 1994, Schneider 1995). We emphasize that confirmation of our source redshift should be undertaken.  ROGI point out that their discovery rate of candidate lenses, two per 50 WFPC2 fields, is consistent with the surface density of bright foreground elliptical galaxies and the surface density of $I < 26$ objects in the survey fields. Thus we might expect that the lensed object should be broadly ``typical'' of the compact objects seen at $I \\sim 26$ in deep HST images.  If we assume that the lensed object discussed here is indeed typical, then a large fraction of all of the faint background objects (8000 per WFPC2 field) should have strong Ly $\\alpha$ emission.  However, neither the extensive searches that have been undertaken for Ly $\\alpha$ from primeval galaxies (e.g., Pritchet 1994), nor a simple extrapolation of the observed quasar number counts (e.g., Schade et al.  1996) to $I = 26$, only $\\sim$3000  deg$^{-2}$, would support such a conclusion.  Finally,  this example of gravitationally lensed objects also demonstrates that lensing effects must be taken into account when deep HST images of field galaxies are being used to derive conclusions about numbers of companions, merger rates, and evolution in the high redshift universe."
+    },
+    "9602/astro-ph9602056_arXiv.txt": {
+        "abstract": "An equilibrium model is presented for a hydrostatic isothermal hot gas distribution in a gravitational well of a giant elliptical galaxy immersed in a massive dark halo. The self-consistently determined gravitational potential of the system provides the optical surface brightness distribution of the galaxy, matching the de Vaucouleurs $R^{1/4}$-law, if the sound speed in the gas $c$, stellar velocity dispersion $\\sigma_*$, and velocity dispersion of dark matter $\\sigma_{\\rmn{DM}}$ are related by $\\sigma_{\\rmn{DM}}\\simeq c\\simeq\\sqrt{2}\\sigma_*$. While the first equality follows from virial considerations, the second one relies on the observed similarity of optical and X-ray brightness profiles. The thermal equilibrium of the gas is described in the framework of a one-zone model, which incorporates radiative cooling, stellar mass loss, supernova heating, and mass sink due to local condensational instability. The sequence of {\\em saddle-node} and {\\em saddle-connection} bifurcations provides a first-order phase transition, which separates the stable hot phase in the cooling medium for a reasonably high efficiency of condensation and precludes catastrophic cooling of the gas. The bifurcations occur naturally, due to the shape of the radiative cooling function, and this holds for a wide range of gas metallicities. The one-zone model implies a scaling relation for the equilibrium stellar and gas densities $\\rho_*\\propto\\rho^2$, which allows a hydrostatic, thermally stable distribution for the hot isothermal coronal gas. The resulting homologous recycling model reproduces the basic optical and X-ray properties of relatively isolated giant elliptical galaxies with hot haloes. Implications for cD galaxies in X-ray dominant clusters and groups are briefly discussed in the general context of galaxy formation and evolution. ",
+        "introduction": "The analysis of the extended X-ray emission of early-type galaxies has led to an extension of the cooling flow model [see Fabian, Nulsen \\& Canizares \\shortcite{fabian..84} and Fabian \\shortcite{fabian94}] applied to the class of bright elliptical galaxies \\cite{thomas....86}. The typical X-ray luminosities of such galaxies are in the range $10^{39}-10^{42}$ erg s$^{-1}$, the gas temperature is $5\\times10^6-2\\times10^7$~K, and the electron number density in the centre is $0.01-0.1$ cm$^{-3}$. The galaxies usually contain $10^9-10^{10}$M$_{\\odot}$ of the gas in their extended hot haloes \\cite{forman..85}. Observations reveal a tight connection of X-ray and optical characteristics of the galaxies, including a strong correlation between the X-ray and optical luminosities, although with a considerable scatter \\cite{canizares..87,donnelly..90}, and a similarity of X-ray and optical light distributions (Trinchieri, Fabbiano \\& Canizares 1986; Killeen \\& Bicknell 1988). \\nocite{trinchieri..86} \\nocite{killeen.88}  The source of the gas in hot haloes of early-type galaxies is usually attributed to the normal stellar mass loss with a specific rate $\\dot\\rho_*/\\rho_*\\sim1-2\\times10^{-12}$ yr$^{-1}$. Gas injection from stars and supernova events in the galaxy make the gas inhomogeneous \\cite{mathews90} and thermal instability causes matter to cool out at radii $r<r_{\\rmn{cool}}$, where $t_{\\rmn{cool}}<H_0^{-1}$ \\cite{nulsen86}. The dynamics of the gas, therefore, is determined by the relative rates of the processes, controlling its mass and energy balance. Theoretical work done so far has demonstrated a variety of possible behaviours for the gas, ranging from supersonic galactic winds to subsonic inflows [see Thomas \\shortcite{thomas86}, Hattori, Habe \\& Ikeuchi \\shortcite{hattori..87}, David, Forman \\& Jones \\shortcite{david..90}, Ciotti et al. \\shortcite{ciotti...91} and references therein]. However, due to the non-linear nature of radiative cooling, numerical modelling of inflows usually ends up with a cooling catastrophe in the centre of the galaxy, which overproduces the observed X-rays and precludes the study of further evolution. As there are no certain observational indications of such a catastrophe, it is important to understand the nature of possible physical damping mechanisms, which are able to prevent it.  It is an objective of this paper to present a purely hydrostatic, thermodynamically stable model for gas recycling in an elliptical galaxy, which satisfies the observational constraints. The model incorporates all the same physical processes as the `standard' multiphase cooling flow model, and may also be useful in interpreting the X-ray emission from the central `cooling flow' regions in dynamically quiet X-ray dominant clusters or groups, where the outskirts of the central dominant optical galaxy extend to the cooling radius. ",
+        "conclusions": "A purely hydrostatic model for the hot gas in normal elliptical galaxies has been presented as an alternative to the `standard' cooling flow picture. The model independently develops an option suggested by Thomas et al. \\shortcite{thomas....86} that perhaps there are no radial gas flows, so that all gas present at some radius originated there (e.g. from stellar mass loss) and cools and is deposited {\\em in situ}. Energy injection from SNe and feedback heating due to star formation from deposited material have been also incorporated in the description of equilibrium gas recycling. Formally, the model is quite similar to the steady-state cooling flow formulation of Sarazin \\& Ashe \\shortcite{sarazin.89} but with zero bulk motion of the gas. (Note that the notion of the cooling flow itself implies highly subsonic velocities and there are no observational data that would directly confirm the existence of the {\\em flow}.) Therefore, the recycling model, representing a `minimal cooling flow', gives the exact lower estimates of the mass deposition rate in elliptical galaxies and fills the gap in a family of models for hot gaseous coronae, being a fully self-consistent hydrostatic model which for the first time takes into account the whole set of sources and sinks of mass and energy. It has the advantage of an analytical description of gas equilibria, which allows the study of their linear stability. However, the global thermal-convective stability of such a hydrostatic recycling atmosphere of a galaxy must be the subject of a separate numerical investigation.  Summarizing the results of this study, it is worthwhile to stress that the galaxy, located in the centre of the cooling region, may play an important role, providing a complex mechanism for hydrostatic support of thermally unstable gas. The moderate amount of SN Ia heating in combination with the distributed mass sink is able to compensate the radiative cooling and the stellar mass loss in such a way that quasi-isothermal hot gas is kept in both {\\em hydrostatic} and {\\em thermodynamic} equilibrium. The latter becomes possible due to a first-order phase transition, which separates the stable hot gas phase for a sufficiently high efficiency of condensation. Having no energy input and/or mass sink it is impossible to avoid a cooling catastrophe in the centre of the `cooling flow' region, where radiative cooling and stellar mass loss are efficient."
+    },
+    "9602/astro-ph9602110_arXiv.txt": {
+        "abstract": "We present bright galaxy number counts in the blue ($16<\\Bj<21$) and red ($15<\\R<19.5$) passbands, performed over 145 sq. degrees, both in the northern and southern galactic hemispheres.  The work was conducted on Schmidt plates digitized with the MAMA machine, and individually calibrated using an adequate number of CCD sequences. We find a relatively large density of bright galaxies implying a ``high'' normalization of the local luminosity function. Our counts and colour distributions exhibit no large departure from what standard no-evolution models predict to magnitudes $\\Bj<21$, removing the need for evolution of the non-dwarf galaxy population in the optical, out to $z \\approx 0.2$. This result disagrees with that of Maddox et al. (1990) on the APM catalog. We show that the APM and similar catalogs may be affected by a systematic magnitude scale error which would explain this discrepancy. ",
+        "introduction": "Since Hubble (\\cite{hubble}), galaxy number counts have been widely used as a statistical tool for probing the distant universe, with the hope of constraining both its geometry and the evolution of its content. Modern, high efficiency instruments and detectors reach impressive magnitude limits in the optical: $\\Bj \\approx 28$ (Tyson \\cite{tyson}, Metcalfe et~al. \\cite{metcalfe:alb}, Smail et~al. \\cite{smail:al}), although on a limited area. Galaxy counts done at brighter magnitudes ($\\Bj<21$) are equally useful; their interpretation is much less model-dependant because of the smaller lookback-times and weaker cosmological effects. They thus constitute the link between models and deeper counts by providing a normalisation of both space densities and colours at low redshift.  The problem with bright galaxy counts is that they obviously require substantial solid angles to be surveyed in order to provide statistically significant samples. Until very recently, such areas could only be surveyed in a reasonable time using photographic Schmidt plates. Since the mid-eighties, fast microdensitometers coupled with image analysis computer programs have been employed to produce automatically highly complete catalogs from large areas of the high galactic latitude sky like the COSMOS (Heydon-Dumbleton et~al. \\cite{heydon:al}), the MRSP (Seitter \\cite{seitter}), or the APM (Maddox et~al. \\cite{maddox:alb}) galaxy surveys.  The main difficulty when dealing with photographic material concerns flux measurements. Many scientific issues typically require the systematic errors to be kept $\\la 0.1$~mag.. This is quite a difficult task to achieve on large scales, and can be reached only if a large number of {\\em galaxy} standards per Schmidt plate, over the whole magnitude range of the counts, are observed to calibrate the data (Metcalfe et~al. \\cite{metcalfe:alc}). Having such a high density of photometric standards is practically not possible with very large Schmidt plate surveys, because of the huge observing time it would require. The catalogs extracted from these surveys are therefore likely to be hampered by systematic errors in their photometry, rather than by limited statistics.  In this paper we re-examine galaxy number counts in the blue and red photographic passbands over the magnitude range $16<\\Bj<21$ and $15<\\Rf<19.5$. Our surveyed area is modest (140 sq. deg.), but has been carefully calibrated using a fair density of CCD standards ($\\approx 700$, half of them beeing galaxies), as an attempt to keep photometric systematic errors $\\la 0.1$~mag. {\\em over the whole magnitude domain} of the survey. The procedure is described in detail in \\S \\ref{par:dataproc}, as well as the general data processing. The galaxy number counts are derived and compared to previous studies in \\S \\ref{par:numbercounts}. Model predictions about number counts and galaxy colour distributions are tested in \\S \\ref{par:modcomp}. We finally discuss the implications of our results in \\S \\ref{par:discuss}. ",
+        "conclusions": "We have presented a new photometric survey of bright galaxies with $15<\\Bj<21$ and $14<\\Rf<19.5$. Comparison with our own CCD galaxy standards, as well as others from the literature, allowed us to estimate systematic errors in the photometry to be $\\la 0.1$~mag over the whole magnitude range of the survey. Our data reveal no evidence of strong galaxy evolution as had been reported by Maddox et~al. (\\cite{maddox:ald}) with the APM survey, although the slope of our number counts is still somewhat higher than what is predicted by no-evolution models brightward $\\Bj\\approx 17$. A comparison with the APM data and other arguments suggest that the APM catalog is affected by a large magnitude scale error, underestimating by about 0.4 mag the flux of galaxies with $\\Bj\\la 17$. This would then justify the choice of a ``high'' normalization of the field luminosity function (concerning $M^*$ as well as $\\phi^*$) already used by several authors to improve the fit of models at $\\Bj>19$ (e.g. Metcalfe et~al. \\cite{metcalfe:ala},\\cite{metcalfe:alb}, Glazebrook et~al. \\cite{glazebrook:al}).  But, given the difficulties inherent to photographic photometry, it would be good to confirm our analysis with linear detectors in order to conclude definitively. Forthcoming digital large scale surveys like the SDSS (e.g. Kent \\cite{kent}) should provide this possibility."
+    },
+    "9602/astro-ph9602070.txt": {
+        "abstract": "We calculate the lowest-order non-linear contributions to the power spectrum, two-point correlation function, and smoothed variance of the density field, for Gaussian initial conditions and scale-free initial power spectra, $P(k) \\sim k^n$. These results extend and in some cases correct previous work in the literature on cosmological perturbation theory. Comparing with the scaling behavior observed in N-body simulations, we find that the validity of non-linear perturbation theory depends strongly on the spectral index $n$. For $n<-1$, we find excellent agreement over scales where the variance $\\sigma^2(R) \\la 10$; however, for $n \\geq -1$, perturbation theory predicts deviations from self-similar scaling (which increase with $n$) not seen in numerical simulations. This anomalous scaling suggests that the principal assumption underlying cosmological perturbation theory, that large-scale fields can be described perturbatively even when fluctuations are highly non-linear on small scales, breaks down beyond leading order for spectral indices $n \\geq -1$. For $n < -1$, the power spectrum, variance, and correlation function in the  scaling regime can be calculated using dimensional regularization. ",
+        "introduction": "\\label{sec:intro}   Several independent arguments suggest that the growth of cosmological density perturbations on large scales can be described by perturbation theory, even when the density and velocity fields are highly non-linear on small scales. For example, analytic and numerical work showed that linear perturbation theory describes the evolution of the large-scale density power spectrum $P(k)$, provided the initial spectrum falls off less steeply than $k^4$ for small $k$~(\\cite{Zel'dovich65,Peebles74,PeGr76,Peebles80}).  In addition, for Gaussian initial conditions, leading-order non-linear perturbative calculations of higher order moments of the density field, e.g., the skewness and kurtosis, agree well with N-body simulations in the weakly non-linear regime, where the variance of the smoothed density field $\\sigma^2(R) \\equiv \\langle \\delta^2(R) \\rangle \\la 0.5 - 1$~(Juszkiewicz, Bouchet \\& Colombi 1993, \\cite{Bernardeau94c,GaBa95}, Baugh, Gazta\\~naga \\& Efstathiou 1995). Thus, leading-order perturbation theory has been shown to work surprisingly well in the cases where comparisons with numerical simulations have been made.  The success of leading-order cosmological perturbation theory raises questions: since the perturbation series is most likely asymptotic, what happens beyond leading order---does the agreement with simulations improve or deteriorate? More generally, what sets the limits of perturbation theory, beyond which it breaks down? These questions have become more urgent since it has been shown that leading-order perturbation theory appears to provide an adequate description even on scales where {\\it next-to-leading order} and higher order perturbative contributions would be expected to become important.  To address these questions, one must calculate loop corrections, {\\it i.e.}, corrections {\\it beyond} leading order, in non-linear cosmological perturbation theory (NLCPT). This is the second paper of a series devoted to this topic. In the first paper~(\\cite{ScFr96}), we applied diagrammatic techniques to calculate loop corrections to 1-point cumulants of unsmoothed fields, such as the variance and skewness, for Gaussian initial conditions.  Here, we calculate the 1-loop (first non-linear) corrections to the power spectrum, the volume-averaged two-point correlation function, and the variance of the smoothed density field for scale-free initial spectra, $P(k) \\sim k^n$.  While the linear power spectrum for the Universe is not scale-free (on both observational and theoretical grounds), scale-free spectra are useful approximations over limited ranges of wavenumber $k$; they also have the advantage of yielding analytic closed form results. In a forthcoming paper, we will present 1-loop corrections to the bispectrum (the three-point cumulant in Fourier space) and the skewness of the smoothed density field for scale-free and cold dark matter spectra.  One-loop corrections to the two-point correlation function and power spectrum have been previously studied in the literature ~(\\cite{Juszkiewicz81,Vishniac83}, Juszkiewicz, Sonoda \\& Barrow 1984, \\cite{Coles90,SuSa91}, Makino, Sasaki \\& Suto 1992, \\cite{JaBe94,BaEf94}). Multi-loop corrections to the power spectrum were considered by~\\cite{Fry94}, including the full contributions up to 2 loops and the most important terms at large $k$ in 3- and 4-loop order. Some of our results overlap in particular with the analytic results for the 1-loop power spectrum reported by Suto \\& Sasaki (1991) and  in Makino et al. (1992). However, since we found that some of their expressions contain errors, we present complete corrected expressions for the 1-loop power spectrum.  One-loop corrections to the variance were studied numerically for Gaussian smoothing by \\L okas et al. (1995).  We correct some numerical errors in their results, extend them to include top-hat smoothing and to the average two-point correlation function, and provide analytic derivation of some of the numerical results.  The limiting behavior of the 1-loop corrections on large scales leads us to reconsider the issue of self-similarity in perturbation theory ~(\\cite{DP77,Peebles80}). It is commonly accepted that the evolution of density perturbations from scale-free initial conditions in an Einstein-de Sitter universe is statistically self-similar. N-body simulations have generally found self-similar scaling for $-3 < n < 1$, although the results for $n<-1$ have been somewhat ambiguous due to problems of dynamic range in the simulations~(\\cite{EFAL88,BG91,RG91,G92,CBH95,Jain95,PCOS95}). The issue of self-similarity in NLCPT has recently been investigated by Jain \\& Bertschinger (1995), who present arguments that the perturbative evolution is also self-similar.  While we do not disagree with their calculations, we start from a different premise regarding cutoffs in the initial spectrum, which leads us to conclude that loop corrections break self-similar scaling if the spectral index $n \\geq -1$.  In the regime where we do find scaling, $n < -1$, we use dimensional regularization to calculate analytically the asymptotic behavior of the 1-loop power spectrum, variance, and average correlation function; these results agree quite well with the `universal scaling' extracted from numerical simulations~(\\cite{HKLM91,PeDo94}, Jain, Mo \\& White 1995).  The paper is organized as follows.  In Sec.~\\ref{sec:dynstat} we discuss non-linear perturbation theory and review the diagrammatic approach developed in~(\\cite{Fry84,GGRW86,ScFr96}) to the calculation of statistical quantities.  Section~\\ref{sec:results} presents the results of 1-loop calculations for the power spectrum, the smoothed variance, and average two-point correlation function.  Self-similarity in perturbation theory is the subject of Sec.~\\ref{sec:selfsimpt}.  We discuss the conditions under which 1-loop corrections exhibit self-similarity and compare our results with the universal scaling hypothesis, based on numerical simulations.  Section~\\ref{conclusions} contains our conclusions. ",
+        "conclusions": "\\label{conclusions}  We have calculated the one-loop (first non-linear) corrections to the power spectrum, average two-point correlation function, and variance of the density field, including smoothing effects for top-hat and Gaussian filters, for scale-free Gaussian initial conditions. For the power spectrum, these results should replace the expressions given in Makino et al. (1992) for the $p_{22}$ contribution on lengthscales smaller than the inverse ultraviolet cutoff of the linear spectrum ($k_c<k<2k_c$). Our results for the one-loop corrections to the variance extend those of \\L okas et al. (1995) to top-hat smoothing and correctly go over to the unsmoothed values found in Scoccimarro \\& Frieman (1996) when the smoothing radius approaches zero.   We found that, when formulated with ``fixed'' comoving cutoffs,  non-linear perturbation theory beyond tree-level does not obey self-similar scaling when the spectral index $n \\geq -1$. As a consequence, in this spectral range, NLCPT disagrees with the results of N-body simulations embodied in the self-similar universal scaling ansatz (USH). For $-3 < n < -1$, however, we found that one-loop corrections do obey self-similar scaling and are in excellent agreement with the USH down to lengthscales where the dimensionless power spectrum $\\Delta_e(k) \\simeq \\sigma^2(R\\sim 1/k) \\approx 10$; below this scale, the N-body results make the transition to the highly non-perturbative regime described approximately by stable clustering.  We interpret the breaking of self-similarity for $n \\geq -1$ as a signature of the breakdown of the fundamental assumption that the large-scale evolution of the density field  can be calculated perturbatively when there are highly non-linear fluctuations on small scales. In this instance, self-similarity breaking can be thought of as arising from an ultraviolet divergence in the 1-loop corrections as the small-scale cutoff $k_c$ of the linear power spectrum goes to infinity. This problem is more severe as the spectral index $n$ increases, which is the expected behavior due to the increase in small-scale power.    Since for $n \\geq -1$, the relative power on small scales is large, the question arises of whether the disagreement between the perturbative results and numerical simulations is due to the breakdown of perturbation theory (as we argued above) or to the use of the single-stream approximation, which neglects the effects of pressure gradients due to velocity dispersion. When these effects are included, one expects that the anisotropic velocity dispersion (non-radial motions) associated with small-scale substructure can inhibit the collapse of large-scale perturbations---the ``previrialization'' effect (\\cite{DP77,EvCr92,LJBH95,Peeb90}). In fact, as $n \\rightarrow -1$ from below, $s^{(1)}$ changes sign and becomes negative, showing that such an effect is already present at some level in the single-stream approximation. In addition, the terms neglected in the single-stream approach have been included by Bharadwaj (1995), who nevertheless finds exactly the same results as in the single-stream approximation. Therefore, if velocity dispersion terms indeed become important, it appears that their effects cannot be treated perturbatively. This suggests that the interpretation given above is the correct one, i.e., for $n \\geq -1$ the non-linear evolution of the density power spectrum is inherently non-perturbative, and perturbative methods are of little use.   Given the excellent agreement between 1-loop NLCPT and N-body results for $n \\la -1$, one is naturally led to ask whether this agreement can be improved upon, i.e., extended to smaller lengthscales, by going to next order (2-loop corrections). We regard this to be unlikely, since the USH suggests that  the physical picture involved in this small-scale regime is well described by stable clustering, which cannot be obtained perturbatively starting from linear solutions. Another interesting question is whether the validity of loop corrections in NLCPT for higher order cumulants (such as the bispectrum, skewness, etc.) will follow the same pattern with spectral index. We will address this issue in future work."
+    },
+    "9602/astro-ph9602104_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "\\label{sec:intro}  Neutrinos with masses and mixings are one of the simplest extensions of the standard model of elementary particles. All direct searches for neutrino oscillations using neutrino beams from accelerators and reactors have produced only upper limits, with the possible exception of the Liquid Scintillator Neutrino Detector experiment (LSND) at Los Alamos \\cite{lsnd}. The claimed LSND signal is still somewhat controversial; in Ref. \\cite{hill95} the LSND data are interpreted as yielding only an upper limit. However, for the purposes of this paper we will accept the interpretation of the LSND data as a neutrino oscillation signal, except where explicitly stated otherwise.  In spite of a dearth of direct evidence, neutrino mixing is a popular explanation for a number of measured neutrino phenomena that appear to be at variance with predictions based on massless, non-mixing neutrinos. These phenomena include the so-called solar and atmospheric neutrino ``deficits'' (\\cite{solrev} and \\cite{atrev}, respectively). Massive neutrinos can affect cosmological evolution \\cite{omeganu} and large-scale structure formation \\cite{lssrev,primack95}. Neutrino flavor mixing could affect supernova dynamics \\cite{fuller87,fuller92} and nucleosynthesis \\cite{qian93,qian95,sigl95}, and big bang nucleosynthesis \\cite{enqvist92,shi93}. These cosmological/astrophysical settings sometimes suggest stricter limits on neutrino masses and mixings than those obtainable with earth-based experiments.  For simplicity, neutrino oscillations are often analyzed in a two-flavor framework, with one neutrino mass eigenvalue difference and one mixing angle. In the wake of the LSND result, interest has grown in constructing models of neutrino mixing that accommodate ``everything.'' In several of these models, results from two-flavor interpretations of various physical effects are combined to make a consistent composite model \\cite{neumodels,raffelt95,chun95}. In addition, some of these models use three independent mass differences: one each for solar neutrinos, atmospheric neutrinos, and LSND.   Two aspects of these models might be considered ``unnatural.'' First, the use of three independent neutrino mass eigenvalue differences requires the introduction of a fourth neutrino. In light of measurements of the width of the $Z_0$ at LEP \\cite{zwidth}, this fourth neutrino must be taken to be ``sterile'' (an $SU(2)$ singlet). Second, these models sometimes employ an ``inverted'' neutrino mass hierarchy. In these inverted schemes, the neutrino mass eigenvalue most closely associated with $\\nu_e$ is heavier than those associated with $\\nu_{\\mu}$ or $\\nu_{\\tau}$, or the mass eigenvalue most closely associated with $\\nu_{\\mu}$ is heavier than that associated with $\\nu_{\\tau}$.  One might hope that these unnatural features could be removed by employing a genuine three-neutrino mixing scheme. It is apparent that non-trivial mixing among three or more generations of neutrinos has the possibility of a richer phenomenology than models in which two-generation mixings are ``stitched together.'' In particular, the excess of electron- over muon-induced events observed in atmospheric neutrinos could be due to the $\\nu_{\\mu}$ oscillating into {\\it both} $\\nu_e$ and $\\nu_{\\tau}$, not just one or the other.  As noted earlier, several previous models have used three independent neutrino mass eigenvalue differences. However, a three generation scheme has only two independent mass differences, which we label $\\delta m^2_1$ and $\\delta m^2_2$. We define $\\delta m^2$ to be difference of the squares of two vacuum neutrino mass eigenvalues, and take $\\delta m^2 > 0$. In the Mikheyev-Smirnov-Wolfenstein (MSW) mechanism \\cite{msw,bethe},  matter effects can enhance or suppress neutrino mixing and may lead to flavor conversion. This mechanism is a popular explanation of the solar neutrino problem. Since a mass-level crossing is the basis of the MSW effect, we are forced to take one of the independent neutrino mass differences from a relatively narrow range determined by solar parameters. An MSW solution to the solar neutrino problem thus determines, for example, $\\delta m^2_1$.  Note, however, that a fair range of neutrino mass differences can be employed in vacuum oscillation explanations of the LSND and atmospheric neutrino data. In a ``last resort'' effort to find a natural three-neutrino mixing scheme, we here consider the possibility of explaining both the atmospheric neutrino data and the LSND data with the other independent neutrino mass difference, $\\delta m_2^2$. Unfortunately, this common neutrino mass eigenvalue difference would have to lie in the range $\\delta m_2^2 \\approx 0.2-0.4$~eV$^2$. This particular range of values for $\\delta m_2^2$ would be an order of magnitude {\\em lower} than the neutrino mass difference most commonly associated with LSND, while it would be an order of magnitude {\\em larger} than the most popular value associated with atmospheric neutrinos. Further, this single common neutrino mass eigenvalue difference turns out to be narrowly specified: as we shall see, it is bounded from above by supernova $r$-process considerations (for non-inverted neutrino mass hierarchies), and more strictly by the compatibility of atmospheric neutrinos with laboratory limits; and from below by the compatibility of LSND with other laboratory limits.  Is the use of the same neutrino mass difference ($\\delta m_2^2$) to give a vacuum neutrino oscillation solution to {\\em both} the LSND data and the published atmospheric neutrino results warranted? A wide range of neutrino mass differences appears to be capable of providing a neutrino oscillation explanation of the LSND data (see Fig. 3 of Ref. \\cite{lsnd}). With one of the neutrino mass eigenvalues set to zero, the ``favored'' LSND value of $\\delta m^2 \\approx 6$ eV$^2$ \\cite{primack95} yields neutrino masses that are convenient from the perspective of cold+hot dark matter models. Note, however, that neutrino oscillation interpretations of the LSND data only probe mass {\\em differences}. We can compensate for a smaller mass difference, $\\delta m_2^2 \\approx 0.2-0.4$ eV$^2$, by offsetting {\\em all} of the neutrino mass eigenvalues from zero \\cite{neumodels,raffelt95}. Such an offset would allow the {\\em sum} of the neutrino mass eigenvalues to provide the requisite contribution of hot dark matter in the models of Ref. \\cite{primack95}. Furthermore, the ($\\sin^2 2\\theta$, $\\delta m^2$) plot of the allowed LSND oscillation parameters \\cite{lsnd} shows that compatibility with KARMEN \\cite{armbruster95}, BNL E776 \\cite{borodovsky92}, and CCFR \\cite{mcfarland95} is readily achieved for $\\delta m_2^2 \\approx 0.2-0.4$~eV$^2$ \\cite{lsnd}.  The range of $\\delta m_2^2$ allowed by a neutrino oscillation solution to the atmospheric neutrino anomaly is a more subtle issue. The isotropy of the sub-GeV data makes them amenable to an oscillation solution for any $\\delta m_2^2 \\gtrsim 10^{-3}$ eV$^2$ (see e.g. \\cite{acker94} and references therein). At one point it was argued that data from upward-going muons restricted $\\delta m_2^2$ to values around $10^{-2}$ eV$^2$ \\cite{becker92}. However, this conclusion relied on calculations of {\\em absolute} neutrino fluxes and cross-sections---not just ratios of these quantities, as can be used in analyzing contained events. It has now been pointed out that calculations of absolute neutrino fluxes and cross-sections different from those used in Ref. \\cite{becker92} permit all of the parameter space allowed by the sub-GeV contained events, in particular ``high'' ($\\sim 10^{-1}$eV$^2$) values of $\\delta m_2^2$ \\cite{frati93,atrev}.  Potentially more damaging to ``high'' $\\delta m_2^2$ values in this context are the Kamiokande multi-GeV data, especially the claim that the data show zenith angle dependance \\cite{fukuda94}. The Kamiokande group's best fit to this data would imply $\\delta m_2^2 \\approx 10^{-2}$ eV$^2$, with 90\\% C.L. upper limits at $\\delta m_2^2 \\sim 0.1$ eV$^2$.  We note that 95\\% C.L. contours could extend the allowed range of mass differences to $\\delta m_2^2 \\sim 0.3$ eV$^2$. In addition, the statistical significance of the Kamiokande group's best fit has been questioned \\cite{saltzberg95,fogli295}.  We will now explore what would be possible if new analyses of the atmospheric neutrino data, or future data with better statistics, were to allow $\\delta m_2^2 \\sim 0.3$ eV$^2$. By assigning the other independent mass difference to be $\\delta m_1^2 \\sim 10^{-5}$ eV$^2$ to use for solar neutrinos, our scheme falls into the category that the authors of Ref. \\cite{fogli95} call ``one mass scale dominance'' (OMSD), in which $\\Delta_{32}, \\, \\Delta_{31} \\gg \\Delta_{21}$ (where $\\Delta_{ij} \\equiv |m^2_i - m^2_j|$, and $m_i$ and $m_j$ are neutrino mass eigenvalues). Great simplification occurs in this limit \\cite{omsd}. Here we take the vacuum mass eigenstates 1, 2, and 3 to be those most closely corresponding to the electron, muon, and tau neutrino flavor eigenstates, respectively.  In Ref. \\cite{fogli95} the OMSD limit is used for neutrinos propagating in vacuum, and it is shown that interpretations of experiments based on this scheme can easily be related to two-flavor interpretations. Additionally, Ref. \\cite{fogli95} offers new interpretations of the available accelerator and reactor data in terms of this simple three-generation framework. In this OMSD scheme, the CP-violating phase, which is inherent in a three-neutrino mixing framework, and the mixing angle $\\theta_{12}$ drop out of the problem. With this simplification,  neutrino oscillation effects in vacuum can be described in terms of the two mixing angles $\\theta_{13},\\,\\theta_{23}$ and one mass-squared difference, $\\Delta \\equiv \\Delta_{32} \\approx \\Delta_{31}$.  Previous authors have studied the three-flavor MSW effect in the limit of well-separated mass scales \\cite{kuo87}. For the matter density scales relevant to the solar neutrino problem, a decoupling to an effective two-flavor mixing problem occurs, allowing an MSW solution employing only $\\theta_{12}$ and $\\Delta_{12}$. For density scales relevant to supernovae, a similar decoupling occurs, leading to an effective two-flavor mixing described in terms of $\\theta_{13}$ and $\\Delta$. In both of these cases, the CP-violating phase does not appear in the final results.  In Sec.\\ \\ref{sec:super} we argue that, for the supernova hot-bubble/$r$-process environment, the survival probability $P(\\nu_e \\rightarrow \\nu_e)$ is roughly the same in the OMSD three-neutrino mixing case as in the two-neutrino mixing case. This allows the general results of calculations of the effects of two-neutrino mixing on the supernova $r$-process to be directly applied to OMSD three-neutrino mixing. In Sec.\\ \\ref{sec:solution} we discuss the chances for obtaining a natural solution to ``everything.'' Concluding remarks, along with a discussion of the prospects of future experiments to clarify the issues discussed in this paper, are contained in Sec.\\ \\ref{sec:concl}. ",
+        "conclusions": "\\label{sec:concl}  Some models which have sought to account for currently  available clues about neutrino properties have employed neutrino oscillations with sterile neutrinos, and/or an inverted mass hierarchy \\cite{neumodels,raffelt95,chun95}. These devices are required when two-flavor interpretations of various physical effects are ``stitched together'' to make a consistent composite model. Here we have sought a ``natural'' three-neutrino mixing scheme---without sterile neutrinos or an inverted mass hierarchy---that satisfies: 1. Accelerator and reactor data, including LSND; 2. Atmospheric neutrinos; 3. Solar neutrinos; 4. Supernova $r$-process nucleosynthesis; 5. Cold+hot dark matter models. Along the way we have argued that, under some circumstances, putative supernova $r$-process nucleosynthesis bounds on two-neutrino flavor mixing can be applied directly to three-neutrino mixing in the case of one mass scale dominance.  We have found a possible ``natural'' solution, and it is quite restricted. The mass differences in this putative solution are $\\Delta_{12} \\approx 7 \\times 10^{-6}$ eV$^2$ (for an MSW solution to the solar neutrino problem), and $\\Delta_{13} \\approx \\Delta_{23} \\approx 0.3$ eV$^2$ (for LSND and atmospheric neutrinos). Here $\\Delta_{ji} \\equiv |m^2_j - m^2_i|$. To be of use to cold+hot dark matter models, we take the masses themselves to be $m_1 \\approx m_2 \\approx 1.5$ eV, and $m_3 \\approx 2$ eV. The mixing angles we require are $\\sin^2 \\theta_{12} \\approx 2 \\times 10^{-3}$, $\\sin^2 \\theta_{13} \\approx 10^{-2}$, and $\\sin^2 \\theta_{23} \\sim 0.5$ (see Figs. 1a-1d for $\\theta_{13}$ and $\\theta_{23}$, and Figs. 4 and 6 of Ref. \\cite{fogli94} for $\\theta_{12}$). While this ``natural'' solution exists, it is rather fragile and has some arguably unnatural features, as discussed in Sec. \\ref{sec:solution}.  Perhaps the main difficulty with this scheme is finding a common mass difference suitable for a vacuum neutrino oscillation solution for both LSND and the zenith-angle dependance of the Kamiokande multi-GeV atmospheric neutrino data. As the statistics for both of these experiments are not compelling at this stage, future results may shed light on the matter. Super-Kamiokande \\cite{superkam} will have much better statistics for the atmospheric neutrino deficit. Furthermore, the LSND experiment continues and may provide better statistics in the future. The KARMEN experiment \\cite{armbruster95}, which probes regions of parameter space similar to LSND, hopefully will also report further results.  Proposed terrestrial experiments could definitively eliminate the solution presented in the last section. New reactor experiments at San Onofre and Chooz will have increased sensitivity to $\\delta m^2$, but apparently will not have increased sensitivity to the oscillation probability at the mass difference we are interested in \\cite{sanonofre}. Conversely, the CHORUS and NOMAD experiments have high sensitivity to the oscillation probability, but will probably not reach small enough $\\delta m^2$ to convincingly eliminate the putative ``natural'' solution \\cite{cernexp}. However, proposed long-baseline accelerator experiments will probe the entire region of parameter space favored by the atmospheric neutrino sub-GeV data \\cite{fogli95,michael95} and could thus provide the crucial test.  Assuming the validity of the data and astrophysical arguments we have attempted to satisfy, the convincing experimental elimination of the last remaining ``natural'' solution presented here would be an important development. It would be significant evidence for the existence of sterile neutrinos and/or a neutrino mass hierarchy of a different nature than that of the charged leptons (i.e. an ``inverted'' hierarchy, in which the mass eigenvalue most closely associated with $\\nu_e$ is heavier than those associated with $\\nu_{\\mu}$ or $\\nu_{\\tau}$; or a hierarchy in which no one neutrino mass eigenvalue difference dominates the other neutrino mass eigenvalue differences)."
+    },
+    "9602/astro-ph9602042_arXiv.txt": {
+        "abstract": "%  We complete our study of pulsars' non-uniform surface temperature and of its effects on their soft X-ray thermal emission. Our previous work had shown that, due to the effect of gravitational lensing, dipolar fields cannot reproduce the strong pulsations observed in the four nearby pulsars for which surface thermal radiation has been detected: PSR 0833-45 (Vela), PSR 0656+14, PSR 0630+178 (Geminga), and PSR 1055-52. Assuming a standard neutron star mass of 1.4 \\Msol, we show here that the inclusion of a quadrupolar component, if it is suitably oriented, is sufficient to increase substantially the pulsed fraction, $Pf$, up to, or above, the observed values if the stellar radius is 13~km or even 10~km. For models with a radius of 7~km the maximum pulsed fraction obtainable, with (isotropic) blackbody emission, is of the order of 15\\% for orthogonal rotators (Vela, Geminga and PSR 1055-52) and only 5\\% for an inclined rotator as PSR 0656+14. Given the observed values, this may indicate that the neutron stars in Geminga and PSR 0656+14 have radii significantly larger than 7~km and, given that very specific quadrupole components are required to increase $Pf$, even radii of the order of 10~km may be unlikely in all four cases. However, effects not included in our study may possibly seriously invalidate this temptative conclusion.   We confirm our previous finding that the pulsed fraction always increases with photon energy, below about 1 keV, when blackbody emission is used and we show that it is due to the hardenning of the blackbody spectrum with increasing temperature. The observed decrease of pulsed fraction may thus suggest that the emitted spectrum softens with increasing temperature and that this observed effect must be of atmospheric origin.   Finally, we apply our model to reassess the magnetic field effect on the outer boundary condition used in neutron star cooling models and show that, in contradistinction to several previous claims, it is very small and most probably results in a slight reduction of the heat flow through the envelope.    \\vspace{1.cm} \\noindent Submitted to {\\em The Astrophysical Journal}. ",
+        "introduction": "}  In a previous paper (Page 1995a; hereafter Paper I) we had presented a simple but realistic model of temperature distribution at the surface of a magnetized neutron star. This model was used to study the effects of temperature inhomogeneity on the neutron star thermal emission and compare them with the observed properties of the four neutron stars for which this emission has been detected by the {\\em ROSAT} satelite (\\\"Ogelman 1995a; Page 1995b): PSR 0833-45 (Vela; \\\"Ogelman, Finley, \\& Zimmermann 1993), PSR 0656+14 (Finley, \\\"Ogelman, \\& Kizilo\\mbox{\\u{g}}lu 1992), PSR 0630+178 (Geminga; Halpern \\& Holt 1992) and PSR 1055-52 (\\\"Ogelman \\& Finley 1993). As a first step we did not take into account magnetic effects in the atmosphere which are also very large. We also restricted ourselves to dipolar magnetic fields and the three main results we obtained were: \\newline $-$ Magnetic effects on heat transport in the neutron star envelope do induce very large temperature differences at the surface but when gravitational lensing is taken into account and only dipolar fields are considered the predicted pulsed fractions are smaller than those observed in the soft band where the surface thermal emission is seen. In particular, in the case of 1.4 \\Msol neutron stars with small radii $\\sim 7.0 - 9.0$ km, the pulsed fractions $Pf$s are below 1\\%. With very small radii $ \\leq 6.5$ km, and the same mass, the gravitational beaming can however increase $Pf$ up to 6 - 7\\%. Observed values of $Pf$ range between 10 - 30\\%. \\newline $ - $ With dipolar fields the observable light curves are very symmetrical and their shapes do not correspond to the observations. Together with the previous result, this lead us to conclude that the inclusion of dipolar fields only is not sufficiently adequate to model the surface magnetic field of the four observed neutron stars. \\newline $ - $ The amplitude of the pulse profile increases with increasing photon energy. This result does not correspond to what is observed in the soft X-ray band in the cases of PSR 0656+14, Geminga and PSR 1055-52.    We complete here this study and consider also the effects of these surface temperature distributions on the cooling of neutron stars. Since the completion of the work presented in Paper I, complementary results considering the magnetic effects in the atmosphere, where the emerging spectrum is generated, have been presented (Pavlov {\\em et al.} 1994; Zavlin {\\em et al.} 1995) and clearly showed that they can be as important as the magnetic effects in the envelope that we consider here and in Paper I. A complete study must obviously include both aspects of the problem, envelope and atmosphere. However, as long as the exact chemical composition of a pulsar's surface is not known, analysis with blackbody (BB) emission will still remain a mandatory first step and a reference to which atmosphere models will be compared. For this reason it is important to characterize BB emission and determine what can and what cannot be obtained with it. As important as the successes of our model will thus be its failures which can guide us toward the correct atmosphere, or more generally surface, model (Page 1995b; Page, Shibanov, \\& Zavlin 1995). Nevertheless, several of our results are, we hope, sufficiently general and robust to be valid despite of the limitations of BB emission. We restrict ourselves to study general properties of the model and refrain from detailed analyses of the data which we leave for future work. Our method to generate surface temperature distributions and the observable X-ray fluxes is described in detail in Paper I.   In \\S~\\ref{sec:data} we present some complements to our summary of the {\\em ROSAT} data given in Paper I. In \\S~\\ref{sec:quadr} we consider the effects of the inclusion of the quadrupolar components of the magnetic field and show that when they are superposed to a dipolar field they provide a sufficiently general configuration for our study. Some basic and straightforward results are presented in \\S~\\ref{sec:strength} and \\ref{sec:spectra}. We use these configurations to discuss the reliability of our surface temperature model in \\S~\\ref{sec:reliability} and also to study the possibility of putting constraints on the neutron star (NS) size through gravitational lensing in \\S~\\ref{sec:lens}. We then discuss the energy dependence of the light curves amplitude in \\S~\\ref{sec:pf}, this section will give a clear indication of the inadequacy of BB emission for understanding detailed features of the observations. The next section, \\ref{sec:caps}, adds a second component of emission in order to model the higher energy tail of Geminga. Section~\\ref{sec:cooling} studies the effect of the surface temperature distribution on the boundary condition used for modeling the thermal evolution of neutron stars. Finally, \\S~\\ref{sec:concl} presents our conclusions. ",
+        "conclusions": "}%   \\subsection{Modeling pulsars' surface temperature distribution}%  We have completed our general study of thermal emission from magnetized neutron stars, within the framework of blackbody (BB) emission. Our original intent was to determine if the inhomogeneous surface temperature distribution induced by the anisotropy of heat flow in the neutron star envelope is strong enough to produce the pulsed fraction observed by {\\em ROSAT} in Vela, Geminga, PSR 0656+14, and PSR 1055-52. We have shown explicitly in \\S~\\ref{sec:reliability} that most of the effect is in the geometrical factor $\\cos^{1/2} \\Theta_B$ (Eq.~\\ref{equ:Ts2}), as presented originally by Greenstein \\& Hartke (1983), and fortunately the actual value of $\\chi_0$ (Eq.~\\ref{equ:chi}) is of little importance. Moreover, models of magnetized neutron star envelopes allow us to relate the interior temperature $T_b$ to the maximum surface temperature $T_s(\\Theta_B~=~0^\\circ)$ quite reliably as long as the latter is higher than about $3~\\cdot~10^5$~K. The minimum surface temperature $T_s(\\Theta_B~=~90^\\circ)$ is however poorly determined but its actual value, equivalent to $\\chi_0$, is not really important. In short, {\\em the surface temperature distribution of a magnetized neutron star can be modeled well enough for present practical purposes}.    Since we do not take into account the effect of the magnetic field on the emitted spectrum but only on the local surface temperature, our results are not sensitive to the actual strength of the field. As long as the magnetic field is stronger than a few times 10$^{11}$~G its effects saturate and we are mostly dealing with warm plates whose size is determined by the field orientation $\\Theta_B$. When atmospheric effects are included one can expect a real dependence on the field strength due to the presence of absorption edge(s) and a field dependent anisotropy of the emission. As in the case of dipolar fields (Paper I), the composite BB spectrum produced by the temperature nonuniformity is close to a BB at the star's effective temperature with some excess in the Wien tail (figure~\\ref{fig:comp-spectrum}). The quadrupole+dipole configurations which induce strong pulsations show an excess similar to the pure dipolar case since they have warm and cold regions of similar areas but many of the configurations which produce little pulsations give a composite BB spectrum closer to a single temperature BB spectrum since they have many warm regions and smaller cold regions.   \\subsection{The quadrupolar component}%  We showed in Paper I that dipolar surface fields do not allow to reproduce the amplitude of the observed pulse profiles; however, we have shown here that the inclusion of a quadrupole component is sufficient to raise the pulsed fraction up to or above the observed values. Gravitational lensing is crucial in this result, and the effect of the quadrupole, in the cases where it does increase the pulsed fraction, is basically to push the two magnetic poles, and the warm regions surrounding them, closer to each other: the areas of the warm and cold regions do not change substantially, compared to the purely dipolar case, but gravitational lensing is then less effective in keeping constantly a warm region in sight. This shift of the magnetic poles location is in fact immediately seen in the cases of Geminga and PSR 1055-52 which are considered to be orthogonal rotators but show only a single wide peak in the soft X-ray band while a dipolar field would produce two peaks (see also discussion in Halpern \\& Ruderman 1993). If we want to stick to a multipolar description, strong octupole components would make the matter worse, inducing still more warm region and flattening even more the light curves. Thus, the simple observation of significant pulsations imply that the surface field is dominated by the dipole with a significant quadrupole component, but nevertheless weaker than the dipole, and that higher order components must be weaker. It has been proposed that the presence of the quadrupole is also felt in the anomalous dispersion measure of the high frequency vs. low frequency radio pulses which has been observed in several pulsars (Davies {\\em et al.} 1984; Kuz'min 1992). {\\em Similar observations of PSR 0833-45, 0656+14, and 1055-52, would allow a direct comparison with the strengths of the quadrupoles needed for the interpretation of the radio and X-ray emission of these pulsars.} No confirmation of the presence of a strong quadrupole component could help constraining the size of these neutron stars through gravitational lensing. Another possibility is that the multipole expansion is inappropriate to describe the surface field (i.~e., would require the inclusion of many strong high order components) as would be the case if the field has a sunspot like structure: plate tectonics (Ruderman 1991a, b, c) does predict such a structure which may be appropriate for cases like Geminga (Page {\\em et al.} 1995).    Most dipole+quadrupole configurations produce little observable pulsations and it is surprising that all four observed pulsars have chosen an {\\em a priori} highly improbable configuration which does increase $Pf$. If we visualize the quadrupole as a pair of coplanar dipoles of equal strengths but opposite directions, then the dipole+quadrupoles configurations which are successful in increasing $Pf$ are almost coplanar: there must be a good physical reason for this to happen in all four cases. Possible identification of a new young neutron star in a supernova remnant close to CTB 1 has been recently reported (Hailey \\& Craig 1996): this object appears to be similar in age and temperature to the Vela pulsar and also shows pulsations at a 10\\% $\\pm$ 10\\% level. If the detected X-rays are from the surface thermal emission and $Pf$ is above 10\\% this would raise to five the number of neutron stars in which a quadrupole with adequate orientation is present. The theory of plate tectonics (Ruderman 1991a, b, c) predicts that the crust of young fast spinning pulsars should break under the stress produced by the internal superfluids and superconductors. Plate motion, towards the equator, will ensue and the mutual attraction of the two magnetic poles may slowly pull them toward each other (as is apparently seen in Geminga and PSR 1055-52) inducing a strong quadrupole component as we need in the present study.   \\subsection{Constraints on the neutron star size.}%   We found that, within our present study, the observed pulsed fractions of Geminga and PSR 0656+14 cannot be reproduced if these stars have a radius of 7~km for a mass of 1.4 \\Msol, or, more generally, if their radius is of the order of $1.7~\\cdot~R_S$, $R_S$ being the star's Schwarzschild radius. Moreover, a 10~km radius for a 1.4 \\Msol star, i.e., $R~=~2.4~R_S$, can be `rejected at the 99\\% confidence level' in the sense that less than 10 models out of 1000 generated for figure~\\ref{fig:stat} can reproduce the observed $Pf$ for these same two neutron stars. Similarly, for Vela and PSR 1055-52, $R~\\sim~1.7~R_S$ can be also be `rejected at the 99\\% confidence level' for the same reason but $R~=2.4~R_S$ is acceptable. Our results thus favor rather large radii, above 10~km for a 1.4 \\Msol mass, for all four neutron stars. These conclusions must be taken with extreme caution for several reasons: 1) the dipole + quadrupole description of the surface field may not be adequate, 2) anisotropic emission due to the magnetic field may increase significantly the observable pulsed fraction, 3) the emitted flux may be partially absorbed in the magnetosphere, 4) the surface temperature may be controlled by factors other than heat flow from the hot interior. Each of these factors may, however, increase or reduce the observable pulsed fraction. Thus, we consider our results as mostly illustrative but notice that it is encouraging that the observed pulsed fractions fall within the range where we could potentially rule out large classes of neutron star models through gravitational lensing of the surface thermal emission. We finally mention that gravitational lensing of the hot polar cap emission may also provide complementary estimates of the neutron star radius: Yancopoulos, Hamilton, \\& Helfand (1994) find $R~=~(2.4~\\pm~0.7)~R_S$, or $10~\\pm~3$~km at 1.4~\\Msol, for PSR 1929+10 but the anisotropy of emission and the lack of information on the location of the second polar cap make this estimate also uncertain.   \\subsection{Limitations of blackbody emission}%    We confirmed our previous result that with BB emission the observable pulsed fraction $Pf$ is unavoidably increasing with photon energy, at low energy where the surface thermal emission is detected, and explained that this is due to the hardening of the BB spectrum with increasing temperature. The three pulsars Geminga, 0656+14, and 1055-52 show a decrease of $Pf$ with energy in the soft band, irreproducible with BB emission. The model of Page {\\em et al.} (1995) produced such a decrease of $Pf$ by superposing two different spectra, a non magnetic hydrogen atmosphere at low temperature on most of the stellar surface and a magnetized hydrogen atmosphere on two warm magnetized plates: the magnetic spectrum is much softer than the non magnetic one and $Pf$ decreases strongly with energy. Extrapolating from this result we may speculate that the decrease of $Pf$ with energy is due to a softening of the emitted spectrum with increasing temperature. We also showed in \\S~\\ref{sec:caps} that the separation of the soft thermal component and the hard tail is large enough to make the contribution of the latter in the soft band quite small: the observed decrease of $Pf$ is an intrinsic feature of the surface thermal emission and not due to an interference of these two components as proposed by Halpern \\& Ruderman (1993) in the case of Geminga.   A second clear indication for the inadequacy of BB, or BB-like, emission is directly found in the spectral fits: for a reasonable neutron star size the required distance $D$ must agree with other estimates as, e.~g., the distance obtained from the pulsar's dispersion measure or from parallax measurements (Geminga). However, the distance $D$ and stellar radius $R^\\infty$ obtained from spectral fits have to be taken with caution due to the uncertainty in the calibration of the PSPC in the lower channels band; the fit in this band determines the column density $N_H$ which strongly affects the estimate of the emitted flux and the PSPC calibration uncertainty translates into uncertainty on $D$ and $R^\\infty$ (Meyer, Pavlov, \\& M\\'{e}sz\\'{a}ros 1994). In the case of Vela, BB emission is disastrous since it implies a 3~-~4~km radius neutron star at 500~pc or, alternatively, a distance of $\\sim$~1500~pc for a standard size neutron star (\\\"Ogelman {\\em et al.} 1993). The same is apparently happening in the case of the newly discovered neutron star RXJ 0002+6246 (Hailey \\& Craig 1996) for which BB spectral fits also require an effective blackbody radius $R^{\\infty}~\\sim~2.5~-~4.5$~km for the estimated distance of 3~kpc. However, magnetized hydrogen atmosphere models do pass the distance test and can fit the Vela spectrum with the standard distance of 500~$\\pm$~125~pc (Page, Shibanov, \\& Zavlin 1996) and a reasonable hydrogen column density. In the case of Geminga, with a parallax measured distance of $157^{+59}_{-34}$~pc (Caraveo {\\em et al.} 1995), available magnetized hydrogen atmosphere models do pretty bad since they need a distance of about 20~pc (Meyer {\\em et al.} 1994; Page {\\em et al.} 1995) while BB is acceptable with a required distance between 150 and 400~pc (Halpern \\& Ruderman 1993). This low distance needed for magnetized hydrogen atmosphere spectral fits may, however, be partly due to the uncertainties in the PSPC calibration at low energies (Meyer {\\em et al.} 1994). The crucial difference between Vela and Geminga is in the temperature: at 10$^6$ K (Vela) the hydrogen atmosphere is fully ionized and the models are reliable while in Geminga, at $T~\\sim~3~-~5~\\cdot~10^5$~K, the atmosphere is only partially ionized and the available atmosphere models are not yet very accurate. Atomic thermal motion (Pavlov \\& M\\'{e}sz\\'{a}ros 1993) and its effects on the bound-bound (Pavlov \\& Potekhin 1995) and bound-free (Pavlov \\& Potekhin 1996) opacity will affect significantly the absorption edge present at $T_e$ below 10$^6$~K, but inclusion of these effects into atmosphere models has not yet been done and the results are hard to predict. The field and temperature dependence of the absorption edge may naturally induce a decrease of $Pf$ with increasing photon energy as observed (Pavlov 1995).   These limitations of BB emission may cast doubts on the validity of our results. As we stated in the introduction we consider that modeling NS thermal emission with BB emission is a mandatory first step and our study does allow us to delimitate more clearly the limitations of this approach. BB is certainly not adequate for reliable estimates of pulsars' surface temperature, as well as size and distance, from spectral fits. However, modeling of the pulse profiles is probably less dependent on the exact atmosphere characteristics since the anisotropy of the emission from the atmosphere is not very strong at low energy and it is unlikely to affect dramatically our conclusions on the effect of gravitational lensing and neutron star sizes. Our results obviously have to be taken only as a first step which hopefully can indicate the direction to be taken for future work(s).    \\subsection{Neutron star cooling}   We have finally applied our model to assess the effect of the crustal magnetic field on the outer boundary condition used in neutron star cooling models. We have shown that the global effect of the magnetic field, compared to the non magnetic case, is very small and most probably consists in a slight reduction of the heat flux in the envelope which results in a lower effective temperature $T_e$ for a given inner temperature $T_b$. This is in sharp distinction to the case where magnetic effects are naively applied under the assumption of a uniformly magnetized surface, implying an increase of $T_e$ as compared to the non magnetic case.  \\subsection{Warning about the relevance of the model}    Finally, we must repeat the warning expressed in Paper I about the relevance of our study: strong heating of the surface by magnetospheric hard X-rays and/or substantial magnetospheric absorption of the emitted flux, as proposed by Halpern \\& Ruderman (1993), could invalidate seriously our results (Page 1995b). The possible detection of pulsed $\\gamma$-rays from PSR 0656+14 (Ramanamurthy 1995) would mean that all four pulsars we study are copious $\\gamma$-ray emissors and heating of the neutron star surface by such energetic magnetospheres is not an unreasonable hypothesis. In the original Halpern \\& Ruderman (1993) model for Geminga the flow of electrons/positrons onto the polar caps produced a luminosity $L_p~\\sim~2.6~\\cdot~10^{32}$~ergs s$^{-1}$, radiated away as hard X-rays which would then be back scattered onto the neutron star surface and reemitted as soft surface thermal X-rays. The X-ray luminosity of Geminga is about $2~\\cdot~10^{31}$~ergs~s$^{-1}$ for the soft thermal component and $8~\\cdot~10^{29}$~ergs~s$^{-1}$ for the hard tail (for a distance of 160 pc as measured by Caraveo {\\em et al.} 1995) so that even a small proportion of the hard polar X-rays hitting back the surface could explain the observed temperature. The problem of this model was that the {\\em total} predicted X-ray luminosity is more than one order of magnitude higher than observed. The argument has been recently revised and the polar cap luminosity scaled down by Zhu \\& Ruderman (1996) who now obtain $L_p~\\sim 8~\\cdot~10^{30}$~ergs~s$^{-1}$. This is much closer to what is observed and may be the factor controlling Geminga's surface temperature. However, the soft X-ray flux is about 30 times larger than the hard one so that in this model almost all the polar cap hard X-rays must be scattered back onto the stellar surface. On the other side, neutron star cooling models are also perfectly able to explain the observed temperature (Page 1994). We can only hope that future work(s) will elucidate this dilemma."
+    },
+    "9602/hep-ex9602004_arXiv.txt": {
+        "abstract": "We suggest the use of moderately superheated liquids in the form of superheated droplet detectors for a new type of neutralino search experiment. The advantage of this method for Dark Matter detection is, that the detector material is cheap, readily available and that it is easily possible to fabricate a large mass detector. Moreover the detector can be made \"background blind\", i.e. exclusively sensitive to nuclear recoils.\\\\ ",
+        "introduction": "Current models explaining the evolution of the universe and the measured slight anisotropy of the cosmic background radiation have in common, that they predict an appreciable contribution of non-luminous, non-baryonic matter in the form of a mixture of relativistic, light particles and non-relativistic, massive particles (so-called Hot and Cold Dark Matter). Accelerator experiments and results from the first round of Dark Matter experiments explored up to now only a small range of masses and types of possible candidate particles. In particular they leave room for an interpretation of Cold Dark Matter in terms of weakly interacting massive particles with masses between 20 GeV and a few hundred GeV. Here a fitting candidate is the neutralino of the \"minimal supersymmetric standard model\"(ref. [1]). The cross section of these particles are expected to be of electroweak strength, with coherent or axial coupling. These particles are supposed to be concentrated in a self gravitating, spherical halo around our galaxy with a Maxwellian velocity distribution in the galactic frame with a mean velocity of 300 km$/$sec and a local density at the solar system of 0.3\\, GeV/cm$^{3}$. The detection reaction of neutralinos would be elastic scattering with a detector nucleus and the nuclear recoil energy would be the measurable quantity. Under the given assumptions the mean recoil energy would be\\\\ \\begin{equation} <E_{R}> \\approx 2\\,keV M_{N}\\,(GeV)\\; [M_{X}/(M_{N} + M_{X})]^2 \\end{equation} where $M_{N}$ and $M_{X}$ are the masses of the detector nucleus and of the neutralino respectively. The recoil spectrum falls off exponentially with \\begin{equation} dN/dE \\approx \\exp ( -E / <E_{R}> ) \\end{equation}  For all detector materials the recoil energies are expected to be smaller than 100 keV and depending on more detailed assumptions on the cross section the expected event rates in the range between 4 to 20 keV are between 0.01 to 100 events/kgd. Therefore in order to ensure a reasonable countrate a high target mass of more than 50 to 100 kg is needed, especially if one wants to detect the 5\\% annual variation in count rate due to the motion of the earth around the sun. The latter would be the decisive signature for the detection of Dark Matter candidates. Due to background limitations, however, current experimental sensitivities are still far away ($>$factor 100) to reach the small interaction rates predicted. We propose a new approach for Cold Dark Matter detection, which has the potential of substantially improved sensitivity. ",
+        "conclusions": "We suggest the use of moderately superheated liquids in the form of a superheated droplet detector for a new type of dark matter search experiment. The advantage of this method for Dark Matter detection is, that the detector material is cheap, readily available and that it is easily possible to fabricate a large mass detector. From its easy operating conditions and its suitable isotopic composition ($^{19}F$ is a spin-$1/2$$^{+}$ isotope) CCl$_{2}$F$_{2}$, i.e. Freon 12, is an interesting active material. Even more attractive because of its higher mass is CF$_{3}$Br. It has a similar vapor pressure curve as Freon 12 and is non-inflammable. Compared to alternative techniques our method is insensitive to $\\alpha$, $\\beta$ and $\\gamma$ radiation, and avoids the need of complex cryogenics."
+    },
+    "9602/astro-ph9602019_arXiv.txt": {
+        "abstract": "\\noindent \\rightskip=0pt We study the uniqueness and robustness of acoustic signatures in the cosmic microwave background by allowing for the possibility that they are generated by some as yet unknown source of gravitational perturbations. The acoustic {\\it pattern} of peak locations and relative heights predicted by the standard inflationary cold dark matter model is essentially unique and its confirmation would have deep implications for the causal structure of the early universe.  A generic pattern for isocurvature initial conditions arises due to backreaction effects but is not robust to exotic source behavior inside the horizon.  If present, the acoustic pattern contains unambiguous information on the curvature of the universe even in the general case. By classifying the behavior of the unknown source, we determine the minimal observations necessary for robust constraints on the curvature.  The diffusion damping scale provides an entirely model independent cornerstone upon which to build such a measurement.  The peak spacing, if regular, supplies a precision test. ",
+        "introduction": "\\label{sec-intro}  Cosmic microwave background (CMB) anisotropies provide a unique window into the early universe through which we receive information on both the model for structure formation in the universe (e.g.~\\cite{EfsBonWhi,Goretal,OstSte,CriTur,Albetal,Duretal,DodGatSte}) and the background cosmology (e.g.~\\cite{Bonetal,HuSugSil,Jungman,Sel,ScoWhi}). In the simplest models for structure formation, based on the gravitational instability of initial density perturbations, the acoustic signature in the anisotropy power spectrum provides a clean and unambiguous means of measuring all of the parameters of a Friedman-Robertson-Walker cosmology. In particular, it offers a standard ruler with which to make a classical test for curvature in the universe (\\cite{DZS,SugGou,KamSS,HuWhi}). However, structure formation may have proceeded by a more complicated route. Recent investigations have begun to probe the acoustic signature in texture (\\cite{CriTur,Duretal}) and string (\\cite{Albetal,Magetal}) models for structure formation.  Their predictions differ strongly from the standard inflationary case and suggest that perhaps without prior knowledge of the correct model for structure formation, the information contained in the acoustic signature cannot be extracted. In this paper, we focus on two general questions:  does the acoustic signature uniquely specify the model for structure formation? how robust is a measurement of the curvature to changes in the underlying model?  As discussed in the appendix, the question of uniqueness is especially interesting in the case of the standard inflationary paradigm.  Inflation is the unique {\\it causal} mechanism for generating correlated curvature perturbations above the horizon (\\cite{Lid}). All other causal mechanisms generate significant curvature perturbations only near horizon crossing.  We will refer to these alternate possibilities as {\\it isocurvature} models. A unique signature of superhorizon curvature fluctuations can be used as a test of inflation.  By generalizing the external source formalism of Hu \\& Sugiyama (1995a,b; hereafter \\cite{HSa,HSb}) to include backreaction effects in \\S \\ref{sec-physical}, we find that a large class of isocurvature models carry a distinct acoustic signature that can be easily distinguished from the inflationary case, independent of the curvature and other cosmological parameters. The distinction lies in the gross properties of the spectrum, not small or subtle shifts in the peaks heights which require very high resolution measurements to see. As shown in \\S \\ref{sec-signatures}, tell-tale features include the harmonic series of peak locations, the alternating relative peak heights due to baryon drag, and the diffusion damping tail. From this study, we conclude in \\S \\ref{sec-inflation} that {\\it the standard inflationary model with big bang nucleosynthesis (BBN) baryon-to-photon ratio bears an essentially unique signature}.  If acoustic oscillations are present in the CMB, as is the case in all but models with significant reionization (see e.g.~\\cite{Pee,EfsBon,Couetal}) or late formation of perturbations (\\cite{Jafetal}), their signature will provide information on the curvature. The damping tail contains the most robust information, but its location alone can constrain but not precisely fix the curvature. Additional restrictions such as a standard recombination history and a near BBN baryon content are required to make this a sensitive probe of the curvature. Furthermore, in models where the acoustic signature is sufficiently regular, the spacing of the peaks provides a precision test of the curvature even if the baryon content is anomalously low or somewhat high. It can also be combined with the damping tail to discriminate against truly exotic models. In \\S \\ref{sec-curvature}, we discuss specifically what regularities must be observed before the curvature can be unambiguously measured if the correct model for structure formation is not assumed to be known {\\it a priori}.  The outline of this paper is as follows. The next section contains the derivation of our principle results. Their impact on the features in the angular power spectrum of the CMB is described in \\S \\ref{sec-signatures}. We discuss our main points in \\S\\S \\ref{sec-inflation} -- \\ref{sec-curvature}, and present our conclusions in \\S \\ref{sec-conclusions}. ",
+        "conclusions": "\\label{sec-conclusions}  We have generalized the formalism of \\cite{HSa} to include backreaction effects and examined the uniqueness and robustness of acoustic signatures in the CMB. By clarifying the role of compensation and feedback in the evolution of fluctuations, we have shown that the phase of the oscillation, and hence the ratios of peak locations, distinguishes inflation from ``typical'' isocurvature models. Specifically, two robust test are the ratio of third to first peaks and first to peak-spacing. Our analysis also provides a better understanding of the structure of the peaks in these models. However since it is possible to imagine isocurvature models which are tuned to mimic the inflationary pattern of peaks, we have stressed the importance of baryon drag, which allows us to distinguish compressions from rarefactions in potential wells. This can help lift the confusion between adiabatic models and these contrived isocurvature models. Although the level of drag does not make this distinction clear for all possible baryon densities, it is observable for the value predicted by big bang nucleosynthesis. Further understanding of the diffusion damping of anisotropies will allow us to extend the lower limit of baryon densities for which we can distinguish the rarefaction and compression peaks.  We focus on the importance of the damping tail as a measure of spatial curvature which is independent of the model for structure formation, and discuss the robustness of curvature measurements from the CMB.  Even in the case where no acoustic peaks are seen (e.g.~possibly string models) the damping scale can be estimated under mild assumptions about the thermal history and baryon content. A more precise test is possible if the acoustic peaks are regularly spaced as indeed expected of models without extreme  behavior at small scales. Either approach allows one to infer the physical scale of the acoustic feature(s). Its projection onto the sky allows us to perform a classical angular diameter distance test to determine the curvature of the universe.    In the process of carrying out these tests of the curvature, the general nature of the model for the fluctuations can be reconstructed as well as the baryon content of the universe.  All of these studies focus on the tale told by the CMB spectrum taken as a whole. In particular,  the acoustic pattern, which arise from forced oscillations in the photon-baryon fluid before recombination, leaves a distinct signature from which we may begin to reconstruct the cosmological mdoel.   \\bigskip"
+    },
+    "9602/astro-ph9602125_arXiv.txt": {
+        "abstract": "We present the first measurement of the rate of Type Ia supernovae at high redshift. The result is derived using a large subset of data from the Supernova Cosmology Project as described in more detail at this meeting by Perlmutter et al. (1996). We present our methods for estimating the numbers of galaxies and the number of solar luminosities to which the survey is sensitive, the supernova detection efficiency and hence the control time. We derive a rest-frame Type Ia supernova rate at $z\\sim0.4$ of $0.82\\ {^{+0.54}_{-0.37}}\\ {^{+0.42}_{-0.32}}$ $h^2$ SNu where the first uncertainty is statistical and the second includes systematic effects. ",
+        "introduction": "Beginning with the discovery of SN 1992bi (Perlmutter et al 1995), we have developed search techniques and rapid analysis methods that allow systematic discovery and follow up of ``batches'' of high-redshift supernovae. The observing strategy developed compares large numbers of galaxies in each of $\\sim$50 fields observed twice with a separation of $\\sim$3 weeks. This search schedule makes it possible to precisely calculate the ``control time,'' the effective time during which the survey is sensitive to a Type Ia event.  For this analysis, we have studied a set of 52 similar search fields observed in December 1993 and January 1994. The data were obtained using the ``thick'' $1242\\times 1152$ EEV5 camera at the 2.5m Isaac Newton Telescope, La Palma.  The projected pixel size is $0.56''$, giving an image area of approximately $11'\\times 11'$. Exposure times were 600s in the $R_{Mould}$ filter, and the images typically reach a $3\\sigma$ limit of $R=23$mag. Seeing was typically around $1.4''$. Many of the fields were selected due to the presence of a high-redshift cluster ($z\\sim 0.4$). Suitable clusters and their redshifts were taken from Gunn, Hoessel \\& Oke (1986) and from the ROSAT catalog. The total useful area covered in this study is 1.73 sq deg.  The analysis procedure and method for finding SNe can be summarized as follows.  For most fields, two first-look ``reference'' images were obtained and for all fields two second look ``search'' images were obtained $2-3$ weeks after the reference images.  The images were flat-fielded and zero-points for the images were estimated by comparison with E (red) magnitudes of stars in the APM (Automated plate measuring facility, Cambridge, UK) POSSI catalog (McMahon \\& Irwin 1992).  The search images were combined (after convolution to match the seeing of the worst of the four images) and the combined reference images were subtracted from this after scaling in intensity. The resulting difference image for each field was searched for SN candidates.  In this subset of the search data three SNe were found with redshifts 0.354, 0.375 and 0.420 determined from spectra of the host galaxies. For the purposes of this analysis, we assume these are all Type Ia SNe (see Perlmutter et al. 1996). The method used to calculate the rate can be divided into two main parts: (i) estimation of the number of galaxies and the total stellar luminosity (measured in $10^{10}$\\Lbsun) to which the survey is sensitive and (ii) estimation of the SN detection efficiency and hence the control time. ",
+        "conclusions": "This measurement is the first direct measurement of the Type Ia rate at high redshift.  In their pioneering work searching for high redshift supernovae, Hansen et al. (1989) discovered a probable Type II event at $z = 0.28 $ and a Type Ia event at $z = 0.31 $ (N\\o rgaard-Nielsen et al., 1989), however no estimates of SN Ia rates were published based on this discovery.  Nearby Supernovae rates have been carefully reanalyzed recently (Cappellaro et al. 1993a \\& 1993b, Turatto et al., 1994, Van den Bergh and McClure, 1994, Muller et al. 1992) using more precise methods for calculating the control times and correcting for inclination and over-exposure of the nuclear regions of galaxies in photographic searches.  The rate obtained for Type Ia SNe are now consistent among these groups and vary between 0.2 $h^2$~SNu and 0.7 $h^2$~SNu depending on the galaxy types (E, Sa etc., higher rates are found in later type galaxies).  Taking into account the facts that at $z\\sim$0.4 the ratios of galaxy type are different and using the Type Ia rates reported in Cappellaro et al (1993b), we calculate a local rate of $0.53 \\pm 0.25$ $h^2$ SNu for the mix of galaxies expected at $z\\sim0.4$.  Our measured value of $0.82 {^{+0.68}_{-0.49}}h^2$ SNu (where statistical and systematic uncertainties have been combined), although slightly higher, agrees with this value within the uncertainty and indicates that Type Ia rates do not change dramatically out to $z\\sim$0.4.  Note, however that correcting for host galaxy extinction and inclination may change this conclusion.  Theoretical estimates of Type Ia SNe rates have been derived from stellar and galaxy evolution models. Calculations were done mostly for elliptical galaxy type.  Recent calculations, based on evolutionary models of elliptical galaxies, predict rates of $\\sim 0.1 h^2$ SNu (Ferrini \\& Poggianti 1993). Assuming a factor of $\\sim 2$ higher rate in non-elliptical galaxies compared to ellipticals (Capellaro et al. 1993b) and a mix of galaxy types as above, we convert this to an overall rate of Type Ia SNe at $z \\sim 0.4$ in all galaxy types, and derive a value of $\\sim 0.37 h^2$ SNu.  Our total uncertainty of $^{+0.68}_{-0.49}$ in the measurement presented in this paper does not allow any firm conclusion but our observed rate seems to lie above this theoretical prediction. There may be an increase of Type Ia rate with redshift.  Ruiz-Lapuente, Burkert \\& Canal (1995) predict significant increase in rate for redshifts between 0.4 and 0.8 depending on the specific model they consider.  In the near future, our ongoing high-$z$ SN search and others should provide enough data to constrain the theoretical calculations.  This work was supported in part by the National Science Foundation (ADT-88909616) and the U.S. Dept. of Energy (DE-AC03-76SF000098).  We thank the La Palma staff \\& observers for carrying out service observations. We also thank Simon Lilly and Caryl Gronwall \\& David Koo for providing their galaxy counts prior to publication, and Richard Kron for useful discussions. I.M. Hook acknowledges a NATO fellowship. R. Pain thanks Gerard Fontaine of CNRS-IN2P3 and Bernard Sadoulet of CfPA, Berkeley for supporting this work."
+    },
+    "9602/astro-ph9602063_arXiv.txt": {
+        "abstract": "Building a self-gravitating hydrodynamic code as a combination of a hydrodynamic solver and a gravity solver is discussed. We show that straightforward combining those two solvers generally leads to a code that does not conserve energy locally, and instead a special Consistency Condition ought to be satisfied. A particular example of combining Softened Lagrangian Hydrodynamics (SLH) with a P$^3$M gravity solver is used to demonstrate the effect of the Consistency Condition for a self-gravitating hydrodynamic code. The need to supplement the SLH method with the P$^3$M gravity solver arose because the Moving Mesh Gravity solver, used in conjunction with the SLH method previously, was found to produce inaccurated results. We also show that most existing cosmological hydrodynamic codes implicitly satisfy the Consistency Condition. ",
+        "introduction": "Numerical cosmological simulations are now widely used to assess properties of various cosmological models and to study theoretical aspects of large-scale structure evolution. The variety of numerical methods employed ranges from high resolution collisionless $N$-body simulations (e.g.\\ Davis et al.\\markcite{Dea85} 1985; Park\\markcite{P90} 1990; Bertschinger \\& Gelb\\markcite{BG91} 1991; Gelb \\& Bertschinger\\markcite{GB94a} 1994a,\\markcite{GB94b} 1994b; Xu\\markcite {X95} 1995) to sophisticated gas dynamics methods (e.g.\\ Cen \\& Ostriker\\markcite{CO92a} 1992a; Katz, Hernquist, \\& Weinberg\\markcite{KHW92} 1992; Evrard, Summers, \\& Davis\\markcite{ESD94} 1994; Bryan et al.\\markcite{Bea94} 1994; Summers, Davis, \\& Evrard\\markcite{SDE95} 1995) to numerical algorithms that at some level include galaxies together with collisionless dark matter and intergalactic gas as a distinct third component of a simulation (Cen \\& Ostriker\\markcite{CO92b} 1992b,\\markcite{CO93a} 1993a,\\markcite{CO93b} 1993b; Frenk et al.\\markcite{Fel95} 1995; Gnedin\\markcite{G96a} 1996a,\\markcite{G96b} 1996b).  However, not all of these simulation methods are reliably tested and well understood. Historically, collisionless $N$-body methods have been subjected to the most detailed study and testing. The theory of $N$-body methods is well developed and their errors and limitations are well understood (Hockney \\& Eastwood\\markcite{HE81} 1981; Efstathiou et al.\\markcite{Eea85} 1985). Cosmological hydrodynamics methods range from Eulerian hydrodynamic techniques borrowed from engineering applications (Cen et al.\\markcite{Cea90} 1990; Ryu et al.\\markcite{Rea93} 1993; Bryan, Norman, \\& Ostriker\\markcite{Bea95} 1995) to quasi-Lagrangian methods that include the ``Smooth Particle Hydrodynamics'' (SPH) method (Evrard\\markcite{E88} 1988; Hernquist \\& Katz\\markcite{HK89} 1989) and ``Moving Mesh approach'' (MMA) (Gnedin\\markcite{G95} 1995; Pen\\markcite{P95b} 1995b). Unfortunately, currently there is no full understanding of errors and specifics of various hydrodynamic methods. The first comparisons between various cosmological methods (Kang et al.\\markcite{Kea94} 1994; see also Gnedin\\markcite{G95} 1995) demonstrated that differences between various methods can be so substantial that only a few statistical quantitative results are consistent between different methods. Further work on comparison of various numerical techniques is therefore required to support quantitative results of cosmological hydrodynamical simulations.  Since the Moving Mesh approach is the most recent and therefore least understood, it requires special attention in comparing it with other existing numerical techniques. In this paper we concentrate on comparing an implementation of the Moving Mesh approach called Softened Lagrangian Hydrodynamics (SLH; Gnedin\\markcite{G95} 1995) with existing $N$-body techniques. One possible advantage of the Moving Mesh approach is that it offers a way to calculate gravitational forces with high resolution without employing a ``particle'' approach like P$^3$M or TREE. Since the moving mesh represents a single-valued coordinate transformation from quasi-Lagrangian space into real (physical) space, the Poisson equation in real space can be reduced to an elliptic partial differential equation in quasi-Lagrangian space, which can then be solved by the standard numerical techniques. In the SLH method thus calculated the gravity force has also the virtue of having {\\it exactly\\,} the same resolution as the hydrodynamic solver has (we elaborate below what this actually means). There are also efforts under way to develop a purely $N$-body version of the Moving Mesh approach (Pen\\markcite{P95a} 1995a).  The Moving Mesh approach looks very promising since it has higher spatial resolution in dense regions than Eulerian codes for the same amount of computational resources and is substantially faster than SPH methods for the same spatial resolution. We, therefore, undertook to investigate its accuracy by comparing SLH with the P$^3$M gravity solver. Our detailed tests however have uncovered two serious problems with the Moving Mesh gravity solver which have not been discussed in the literature; the problems we found may appear in a wide range of algorithms, so we report them in this paper. The first problem is the discreteness noise (shot noise) in high density regions. It is well known that the shot noise does not represent a serious problem in traditional gravity solvers (Hernquist, Hut, \\& Makino\\markcite{HHM93} 1993), where force errors effectively average out in high density regions; however, in the Moving mesh approach the relationship between the density and the gravitational force is nonlinear, with the nonlinearity coming from the nonlinear dependence of the volume of a mesh element on the force acting on it. In this case the shot noise does not necessarily average out as in the linear case, leading to spurious numerical heating of high density regions. Secondly, a solution to the finite-difference realization of the Poisson equation on a strongly deformed mesh does not necessarily converge to the true solution, giving the incorrect force law. We propose a solution to the first problem in \\S2, but can not resolve the second one. We have therefore been forced to abandon the Moving Mesh gravity solver and have turned instead to a P$^3$M gravity solver that is known to work correctly and accurately. However, while developing the new SLH-P$^3$M code, we encountered another problem which is generic for (and potentially present in) any self-gravitating hydrodynamic code, not only a Moving Mesh code: since the gravity force solver and the hydrodynamic solver may treat spatial gradients differently, they are not a priori consistent with each other. This inconsistency may lead to a serious local nonconservation of energy even if, globally, energy is well conserved. We discuss this effect in detail in \\S3, where we also derive the ``Consistency Condition'' that every cosmological hydrodynamic code must obey in order to be energy conserving and apply it in \\S4 to develop a new SLH-P$^3$M code. Finally, we repeat  some of the Kang et al.\\markcite{Kea94} (1994) tests for the new code in \\S5. We summarize our conclusions in \\S6. ",
+        "conclusions": "We have demonstrated that Moving Mesh gravity solvers suffer from two serious limitations: the shot-noise in high density regions and incorrect Green functions. We have proposed a way to solve the former problem and failed to completely identify the source of the latter. Since our Moving Mesh gravity solver can not be trusted now, we have proceeded further by building a combined SLH-P$^3$M code, where the SLH hydrodynamic solver is combined with the P$^3$M method for computing gravity.  In order for a self-gravitating hydrodynamic code to conserve energy, a special ``Consistency Condition'' ought to be satisfied. We have derived the Consistency Condition for the SLH-P$^3$M code and showed what errors would appear should this condition be violated. Ignoring this condition can lead to serious errors in dense regions.  Finally, we have compared the SLH-P$^3$M code with existing cosmological hydrodynamic codes including the original SLH code. The SLH-P$^3$M approach offers a substantial increase in resolution, making it comparable to the SPH codes in the highest density regions, while maintaining reasonable accuracy in resolving shocks and hydrodynamic features. In particular, we have repaired the biggest problem of the original SLH method: failure to simulate clusters accurately, which stem out of a faulty Moving Mesh gravity solver.  We, therefore, conclude that the new SLH-P$^3$M code is a good numerical tool to investigate the formation of nonlinear structure in the universe including both gas and dark matter."
+    },
+    "9602/astro-ph9602077_arXiv.txt": {
+        "abstract": " ",
+        "introduction": "This lecture is concerned with the three questions raised by the Mayor of Toledo at the reception held during this conference:  {\\bf ?`De d\\'onde venimos?} Namely, what is the origin of the structure we see in the Universe?  {\\bf ?`Qu\\'e somos?} Namely, what is the nature of the Dark Matter around us? and  {\\bf ?`Ad\\'onde vamos?} Namely, what lies beyond the Standard Model? ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602131_arXiv.txt": {
+        "abstract": "We examine the statistical properties of defects formed by the breaking of a U(1) symmetry when the Higgs field has a power spectrum $P(k) \\propto k^n$. We find a marked dependence of the amount of infinite string on the spectral index $n$ and empirically identify an analytic form for this quantity. We also confirm that this result is robust to changes in the definition of infinite string. It is possible that this result could account for the apparent absence of infinite string in recent lattice-free simulations. ",
+        "introduction": "Cosmic strings formed at a primordial phase transition are strong candidates as seeds for large-scale structure in the universe and anisotropies in the microwave background \\cite{Rev,Vil}. Firm predictions are difficult to extract from string models, due to both the uncertainty in the statistical properties of the string network at formation and the non-linearity of network evolution. Information from numerical simulations has been the main tool for calculations to date. The observational consequences of the string scenario may be strongly dependent on whether, in addition to a scale-invariant distribution of loops, any infinite string exists. The presence or absence of infinite string may lead the string network into significantly different scaling solutions \\cite{Ped}, and in particular to solutions that studies of large-scale structure will be able to distinguish between.  Lattice-based simulations have traditionally been employed to predict the initial fraction as infinite string, although recent work has used a lattice-free method to simulate string formation in a first order phase transition \\cite{Julian}. The simplest approach, however, is to assume that the string-forming field takes up random values at each site on a regular lattice, each representing a causally disconnected volume, and look for strings at each face of the lattice. Vachaspati and Vilenkin \\cite{VV} first performed this on a cubic lattice and found that approximately 80\\% of string was in `infinite' form. Further work on a tetrahedral lattice by Hindmarsh and Strobl demonstrated a slightly lower figure \\cite{MarkKarl}. Furthermore, Vachaspati \\cite{V} has shown that for theories without an exact U(1) symmetry (i.e. one field orientation is statistically favoured) it is possible to reduce the fraction as infinite string to zero. Another approach \\cite{KibbleYates}, retaining the symmetry of the theory but introducing a variation in the ranges over which the field is correlated by means of a simple domain-laying algorithm, showed a dependence of the fraction as infinite string on the variance of domain volume. Here we present a new single-parameter method for varying the scale over which the string field is correlated to investigate the behaviour of the amount of infinite string present. It is a controlled means of varying the string-forming field configuration and allows us to employ analytic results concerning defect densities. In addition, the method makes it possible to explore a range of physically-motivated theories in which string-like defects arise \\cite{RayGlyk}, and in particular the role of causality in determining the properties of a cosmic string network on very large (super-horizon) scales.  In section 2 we discuss our method for generating string configurations and present our basic results regarding the fraction as infinite string. In section 3 we compare the total density of string to a theoretical prediction of the density of zeroes of the string-forming field. In section 4, we suggest a functional form for the relation between the  power spectrum of the string field and the density of infinite string. In section 5 we confirm that our results are insensitive to changes in the definition of infinite string, and section 6 presents our conclusions. ",
+        "conclusions": "We have demonstrated a strong dependence of the fraction as infinite string formed during a $U(1)$ symmetry breaking phase transition on the power spectrum of the Higgs-field configuration. For a power spectrum $P(k)\\propto k^n$, we find agreement with an analytic formula for variation of the total density of string with $n$. We find a strikingly simple closed-form expression which fits the density of infinite string and is closely related to the analytic expression for the total string density. We find that the fraction as infinite string falls with $n$, and vanishes at $n=-2$ if the formula for the density of infinite string is to be believed. Between $n=-2$ and $n=-3$, string only exists in loop form, and below $n=-3$ the total string density is zero. We have also shown that the dependence of the fraction as infinite string with $n$ is robust to reasonable changes in our definition of infinite string.  For $n=0$, the value of the Higgs field at each lattice site is independent, and as expected we reproduce the Vachaspati-Vilenkin result of approximately 80\\% infinite string \\cite{VV}. In a cosmological phase transition, we expect values of the Higgs field to be completely independent on scales larger than the causal horizon \\cite{kibble}. Such a field configuration has a spectral index $n=0$ on the largest scales. Provided string we classify as infinite in our lattice simulation does indeed correspond to infinite string in the continuum limit, we can confirm that a non-zero fraction as infinite string will be formed in a cosmological phase transition. The existence of this population of infinite string may have important consequences for the scaling solution of the evolving network, and hence for the observational consequences of cosmic string.  Our findings show that if we want to compute the fraction as infinite string in the universe correctly using numerical simulations, the field configuration must have the correct $n=0$ form on the largest scales. Failure to fulfil this condition could severely under- or over-estimate the fraction as infinite string. In particular, the results of recent lattice-free simulations of a first-order phase transition \\cite{Julian} appear to be consistent with an infinite string fraction equal to zero. It is important to ensure that on large scales the power spectrum of the field configuration does have a spectral index $n=0$, before applying this result in a cosmological context."
+    },
+    "9602/astro-ph9602092_arXiv.txt": {
+        "abstract": "The UV continuum spectrum of quasars and AGN is assumed to originate from an accreting disk surrounding a massive rotating black hole. We discuss the structure and emission spectra of a disk which drives a powerful jet. Due to the large efficiency of extracting energy from the accreting matter in the inner part of the disk close to the massive object, all the energetic conditions for the formation of jets are fulfilled.  The total energy going up into the jet depends strongly on the Kerr black hole parameters, on the disk features and on the mass flow and thickness of the jet. The shape of UV spectra of the AGN can be explained by a sub-Eddington accretion disk which drives a jet in the innermost parts. ",
+        "introduction": "The general concept of an AGN (active galactic nucleus) contains three important ingredients. The first constituent which binds the space-time geometry and dictates how the matter has to move in its vicinity is a black hole. The great luminosities observed depend on the masses of the black holes $M\\in[10^8 M_{\\odot},10^{10} M_{\\odot}]$. The second one is the disk of gas. Matter flows towards the massive object following nearly circular orbits. The disk-like accretion onto a black hole is the most plausible explanation for the strong emission in the ultraviolet (UV), the so-called \"big blue bump\" (Shields 1978) and for the X-rays in AGN. The third component of an AGN is the jet.  The notion that the jets and the accretion disks are symbiotically related originated since the first radio and optical observations revealed the parsec-scale of the AGNs. A geometrically thin disk with a mass accretion rate lower than the limit imposed by the Eddington accretion rate can drive outflows in the dense regions in the vicinity of a black hole. The flow can be ejected from the regions of the disk where the forces of the radiation pressure, gas pressure and gravitation are comparable. Relativistic jets can extract high amounts of energy from the central object (Falcke and Biermann 1995). Pudritz (1986) proposed that an MHD jet could extract the angular momentum of the accretion disk, so that the jet can control the accretion process in the vicinity of the last marginally stable orbit. The jet-disk system is governed by the conservation laws of mass, angular momentum and energy.   \\section {The inner boundary conditions}  We assume that the spacetime geometry where the disk develops is described by the Kerr metric. The mass of the black hole is M and its specific angular momentum is $a$.  The geometry of the accretion disk is described by the radius R and its thickness H.  The disk is geometrically thin ($ H/R \\ll 1$) and we work with the basic assumption that the accretion is quasisteady (Novikov and Thorne 1973, hereafter NT73) with a constant rate of mass accretion $\\dot{M}$. The disk consists of fully ionized hydrogen and the effects of magnetic fields on the disk structure are neglected in our calculations. We consider the case that the disk is not self-gravitating and its rotation law is keplerian.  \\begin{figure} \\iffigures \\centerline{\\psfig{figure=fig1.ps,height=6.0cm,width=9.0cm}} \\else \\picplace{6.0cm} \\fi \\caption[]{\\label{fig_one}Angular velocity via the radius of the disk. The shadowed zone indicates the innermost region of the disk where the jet is created. The thin-dashed line describes the behavior of the infalling matter if the disk can not drive the jet at its boundaries. } \\end{figure}  We assume that the jet starts between $R_{ms}$, the last, marginally stable orbit in the absence of a jet and the radius $R_{jet}$ with $R_{jet}>R_{ms}$, extracting mass, energy and angular momentum from the disk.  The presence of the jet will modify both the behaviour of infalling matter across the radius $R_{jet}$ and the structure of the keplerian disk.  The jets are placed in the potential well of a Kerr black hole. The rotating central object supplies, through its gravitational potential well, the energy necessary for driving and maintaining a stable jet in the AGNs. We consider the case that the gravitational potential energy released between the $R_{ms}$ and the outer radius of the jet $R_{jet}$ goes into the jet.   On the other hand, the geometry and the total energy of the jet are connected with the physical properties of the interior regions of the disk. Using the above mentioned premises we have to deal with a jet-disk symbiosis (Falcke and Biermann 1995): a ``standard disk'' (Shakura and Sunyaev 1973, hereafter SS73) and a powerful jet (Blandford and K\\\"{o}nigl 1979). Both the disk and the jet structure are governed by the mass and the spin of the black hole.   The angular velocity of the accreting matter at the innermost boundaries will sharply decrease from $\\Omega(R_{jet})$ to the angular velocity of the black hole, if we consider that a jet exists. The angular velocity of the black hole is unknown because the surface of infinite redshift (the horizon) hides the image of the massive black hole. Therefore, the massive rotating object has no solid surface on which the accreting matter can follow. The radius of the horizon is considered to be the surface of the black hole. We believe that $\\Omega(R)$ will drop rapidly compared with the simple case of a disk and no jet (see Fig.1). The problem of the inner boundary conditions for the disk structure calculations is solved by using an approximate condition: $\\frac{d\\Omega(R)}{dR} = 0$ at $R = R_{jet}$. We assume that viscous stresses do not exist across the new boundary of disk.  In a first exploration of the physical consequences of the jet-disk symbiosis we use $R_{jet}$ as our key parameter; it is obvious that a full treatment would have to model the detailed magnetohydrodynamics. However, whatever the details of any model, energy and mass conservation have to hold, and so we parametrize a range of possible models with $R_{jet}$, the radius where the transition from disk to jet occurs.  \\section { The connection between jet and disk}  The unification scheme for all quasars and AGN suggests that one can not treat the disk and jet as two separate objects. The concept of a jet-disk symbiosis has been introduced by Falcke and Biermann (1995). In order to describe the reciprocal influences between the three principal objects : the black hole, the disk and the jet, we shall use a standard set of parameters (Falcke et al. 1995a), defining:  \\begin{equation} Q_{accr} = \\dot{M} c^2, \\end{equation}  \\begin{equation} q_m = \\frac{\\dot{M}_{jet}}{\\dot{M}}, \\end{equation}  \\begin{equation} q_{j} = \\frac{Q_{jet}}{\\dot{M} c^2}, \\end{equation}  \\begin{equation} q_{l} = \\frac{L_{disk}^{jet}}{\\dot{M} c^2}. \\end{equation}  Here $\\dot{M}$ is the rate of mass accretion rate in the disk. The mass flow rate into the jet is $\\dot M_{jet}$. We assume that a jet exist on both sides of the disk, and use this symmetry in our quantitative modelling. $ Q_{jet}$ is the total power of the jet -- including the rest energy of the expelled matter -- and is expressed as:  \\begin{equation} Q_{jet} = L_{disk} - L_{disk}^{jet}. \\end{equation}   $L_{disk}^{jet} $ is the total luminosity of a disk modified by the presence of the jet and $L_{disk} $ is the total luminosity of the disk if there are no physical conditions to drive the jet. A large fraction of the total power of the jet is in magnetic fields and relativistic particles.   The dimensionless parameters $q_{j}, q_{l}$ and $q_{m}$ have to be less than 1, because the jet can not consume more matter and energy than is provided by the accretion disk. pushed out into the jet will be swallowedby the black hole. The structure of the disk is strongly connected with a new characteristic length scale $r_{j}$ defined as follows:  \\begin{equation}R_{jet}=R_{ms} (1+r_{j}).\\end{equation}  The parameter $r_{j}$ is dimensionless and the null values for the $q_{m}$ and $r_{j}$ require the simple system with two components: the black hole and the disk. ",
+        "conclusions": "We have studied a symbiotic jet-disk system stationary in time. The jet is fed by a geometrically thin disk and we assumed that the inner part of the \"standard\" disk contains all energetic conditions necessary to form and to maintain a stable jet's configuration. The jet energy has the upper bound corresponding to the gravitational potential energy lost by the infalling gas between $R_{ms}$ and $R_{jet}$.   The structure of the accreting disk is modified by the presence of the jet.  The mass and angular momentum extraction from the disk makes it thinner and denser. The disk becomes cooler, especially in the regions close to the jet.  Between the radius of the marginally stable circular orbit and the radius where it is assumed that a jet has the exterior edge $R_{jet}$ we do not have a disk-structure any longer. A mass rate $q_{m} \\dot M$ is pushed out into a powerful jet. We neglect the influence of the magnetic field in the structure of the disks and the jets. A very important step remains to be made. We would like to study in more details the vertical structure of a relativistic disk with a jet starting at its inner edge. A detailed analysis of a vertical structure of a relativistic disk surounding a Schwarzschild black hole has been done by D\\\"orrer (1995).   We demonstrate that our disk model driving a jet in the innermost dense regions very close to a maximal Kerr black hole can replace all the simple models of the accretion thin disk surrounding a Schwarzschild black hole.  The thickness of the jet given by the dimensionless parameter $r_{j}$ is limited by the necessity to have sufficient UV photon flux in order to explain the high luminosity of the disk and the unsaturated componized photons in the base of jet.  The presence of the jet cuts off the high energy part of the spectrum of the disk without any jet (in the soft X-ray range) and its hot base comptonizes the photons coming from the disk. The total energy carried out by the jet is strongly dependent on the mass and angular momentum of black hole.  We have shown that there is a strong correlation between the disk and the jet guided by the conservation laws of the mass, angular momentum and energy. The jet-disk connection leads to interesting models which explain the spectra of the quasars over the whole range of frequency. The assumption that the disk accretion is the machine which provides the energy in AGNs is the basic ingredient for all models mentioned."
+    },
+    "9602/astro-ph9602021_arXiv.txt": {
+        "abstract": "The term ``violent relaxation'' was coined by Donald Lynden-Bell as a memorable oxymoron describing how a stellar dynamical system relaxes from a chaotic initial state to a quasi-equilibrium. His analysis showed that this process is rapid, even for systems with many stars, and that it leads to equilibria which may plausibly be related to bounded isothermal spheres. I review how numerical simulations have improved our understanding of violent relaxation over the last thirty years. It is clear that the process leads to equilibria which depend strongly on the initial state, but which nevertheless have certain common features. A particularly interesting case concerns objects formed in an expanding universe through dissipationless hierarchical clustering from gaussian initial conditions; these may correspond to galaxy clusters or to the dark halos of galaxies. While such objects display a wide range of shapes and spins, the distributions of these properties depend only weakly on the cosmological context and on the initial spectrum of density fluctuations. Halo density profiles appear to have a universal form with a singular central structure and a characteristic density which depends only on formation epoch. Low mass halos typically have earlier formation times and thus higher characteristic densities than high mass halos. ",
+        "introduction": "By the mid-1960's it was known that the luminosity profiles of elliptical galaxies are smooth, symmetric, and well fitted by theoretical models based on modifications of the isothermal sphere (see, for example, King 1966). This was perceived as paradoxical since Chandrasekhar's (1942) calculation of the effects of ``collisional'' relaxation showed the relevant timescales for elliptical galaxies to be much longer than the age of the universe. Furthermore such two-body relaxation should lead to equipartition of energy and so to radial segregation of stars by mass; this would produce unacceptably large colour gradients in an elliptical galaxy. Thus a relaxation process with a short timescale and no dependence on stellar mass is required to explain the observations.  Although Donald Lynden-Bell was not the first to realise that rapid changes in gravitational potential during protogalactic collapse could provide such a process, his 1967 paper gave a clear description of the mechanisms involved, calculated the relevant timescale, showed how evolution might lead to a state resembling an isothermal sphere but with no mass segregation, gave a first discussion of the reasons why real galaxies would be unable to reach the ``most probable'' state, and above all gave the process a name that everyone can remember. This paper made strong claims in the inimitable Lynden-Bell style, all of them stimulating, but not all, perhaps, fully justified. After deriving a fourth kind of equilibrium statistics to set beside those of Maxwell-Boltzmann, Einstein-Bose, and Fermi-Dirac, and arguing that incomplete relaxation produces rotationally flattened objects resembling real galaxies, the text concludes by throwing doubt on its own major assumptions and careering off into a discussion of the gravothermal catastrophe, the heat death of the universe, and the origin of the Seyfert phenomenon!  The forthright and provocative style of Lynden-Bell (1967) is undoubtedly responsible both for its success and for its rather mixed reputation among professional dynamicists. The latter can be inferred from the slightly hostile (and in my opinion misguided) reworking in Shu (1978), and from the absence of any discussion of the statistics of violent relaxation either in the standard textbook of Binney \\& Tremaine (1987) or in Tremaine, H\\'enon \\& Lynden-Bell's (1986) treatment of mixing processes during violent relaxation. There can be no doubt, however, that the 1967 paper did much to shape our current understanding of how stellar dynamical systems come to equilibrium, and even today it remains both frequently cited and controversial. (See Earn's contribution to this volume for some illuminating statistics).  Numerical experiments in the 1970's showed that violent relaxation from realistic initial conditions does {\\it not} lead to a universal final structure, even after accounting for the restriction already discussed explicitly by Lynden-Bell (1967), namely that the finite mass of any real system implies that the most probable statistical distribution can only be realised for orbits with periods shorter than the collapse time of the system, and so for energies more negative than some truncation threshold. The discrepancy was first clearly demonstrated by the collapse simulations of White (1976) and Aarseth \\& Binney (1978). These showed that the violent collapse of stellar dynamical systems can produce strongly aspherical equilibria, whereas Lynden-Bell's arguments imply a final state which is spherical, at least in the absence of significant rotation. In numerical experiments the final state clearly remembers nonspherical aspects of the collapse and relaxation phases. Unfortunately the corresponding constraints are difficult to formulate and have so far been ignored in attempts to go beyond the Lynden-Bell's original analysis and to obtain a deeper understanding of the observed uniformity of elliptical galaxies ({\\it e.g.} Stiavelli \\& Bertin 1987; White \\& Narayan 1987; Tremaine 1987). The only clear area of progress since 1967 is in understanding the light profiles at large radii. A number of authors have shown that asymptotically a relation $S\\propto r_p^{-3}$ should hold between surface brightness $S$ and projected radius $r_p$ (Aguilar \\& White 1986; Jaffe 1987; Tremaine 1987); this results from the expected continuity in the distribution of stellar binding energies across the value corresponding to escape from the system.  In the current contribution I will discuss recent work on violent relaxation in a different context, that of the formation of dark halos by hierarchical clustering of dissipationless dark matter in an expanding universe. Most past studies of this process have concentrated on differences in the predicted halo structure as a function of the density of the Universe, and of the nature of the initial fluctuations from which structure forms. Here I will instead highlight certain regularities that emerge from the available numerical data, and argue that in this context violent relaxation actually produces a ``universal'' density structure which is independent of halo mass, of cosmological parameters, and of the initial fluctuation spectrum. Unfortunately, we do not yet have a physical understanding of this universal structure of the kind which Lynden-Bell (1967) attempted to provide. ",
+        "conclusions": "The material discussed in the preceding sections suggests that there is indeed a certain universality in the properties of the objects which form by dissipationless hierarchical clustering. Although it is generally agreed that the shapes and spins of dark halos are predicted to have distributions which depend at most weakly on halo mass, on the power spectrum of initial density fluctuations, and on the density of the universe, our claim that the density profiles of objects are also independent of these quantities seems, at first sight, to conflict with earlier work. This discrepancy is, however, only superficial. Previous numerical studies fitted power laws to halo density profiles over a limited range of overdensities, typically $\\sim 10^2$ to $\\sim 10^4$. As a result, the systematic dependence of formation redshift on power spectrum and $\\Omega$, clearly visible from the point distribution in Figure 2, produced systematic variations in the measured slopes. Furthermore, most studies surveyed too small a range of halo masses to notice the systematic mass dependence. Direct comparison of our predictions with the numerical results of earlier studies shows good agreement in almost all cases.  It seems, therefore, that violent relaxation in hierarchical clustering leads to objects with universal density profiles and universal distributions of shape and spin. Gravitational potential fluctuations during the collision and merging processes which characterize hierarchical clustering are evidently strong enough to cause convergent evolution of the kind envisaged by Lynden-Bell (1967). Since Donald's detailed statistical arguments provide at best a partial explanation of this process, further work is needed to achieve a deeper understanding of the simplicity which emerges from numerical experiment."
+    },
+    "9602/astro-ph9602035_arXiv.txt": {
+        "abstract": "We propose a new approach to study the dynamical implications of mass models of clusters for the velocity structure of galaxies in the core. Strong and weak lensing data are used to construct the total mass profile of the cluster, which is used in conjunction with the optical galaxy data to solve in detail for the nature of galaxy orbits and the velocity anisotropy in the central regions. We also examine other observationally and physically motivated mass models, specifically those obtained from X-ray observations and N-body simulations. The aim of this analysis is to understand qualitatively the structure of the core and test some of the key assumptions of the standard picture of cluster formation regarding relaxation, virialization and equilibrium. This technique is applied to the cluster Abell 2218, where we find evidence for an anisotropic core, which we interpret to indicate the existence of a dynamically disturbed central region. We find that the requirement of physically meaningful solutions for the velocity anisotropy places stringent bounds on the slope of cluster density profiles in the inner regions. ",
+        "introduction": "Studying the velocity structure of the cores of clusters of galaxies promises to provide new insights into the physics of the formation of clusters. The crucial physical consequence of the cumulative dynamical history of a cluster is its underlying mass distribution. We propose a new approach to study the dynamics of cluster galaxies using the mass profile measured from gravitational lensing. Lensing provides the most accurate determination of the mass profile, and is independent of assumptions as to the kinematics of the cluster.  The dramatic arcs and multiple images produced by rich clusters tightly constrains the mass of the cluster in the inner-most regions, on the scale of the Einstein radius, with typical range $30\\,<\\,r_E\\,<\\,200\\,h^{-1}_{50}$ kpc \\cite{kneibsoucail95}. Current progress in observational techniques have made it possible to map the cluster mass out to large radii from the weak shearing of faint background galaxies. While uncertainties arise due to the correction for the point-spread function and the unknown redshift-distribution of background galaxies, reliable mass profiles for an increasing number of clusters should be available in the near future \\cite{kaiser95b}. Combining this with knowledge of the spatial distribution and line-of-sight component of the velocities of cluster galaxies, and assuming them to be good tracers of the cluster potential well, we can solve in detail for the variation of the velocity anisotropy parameter $\\beta$ with distance from the cluster centre.  It is an observational challenge to perform a comprehensive redshift survey of distant clusters, $z\\,>\\,0.1$ in sufficient detail to interpret the velocity histogram and efficiently disentangle the effects of substructure, existence of velocity anisotropy and axisymmetric infall. However, securing 200 to 800 velocities for distant cluster galaxies is now becoming possible with the newly developed MOS multi-object imaging spectrographs (\\citeN{yee96} and \\citeN{lefevre94}). This would enable constructing secure line-of-sight velocity dispersion profiles for high redshift clusters.  Recent optical surveys by \\citeN{colless95} and X-ray studies of clusters by \\citeN{briel95} seem to indicate that while the X-ray isophotes of cluster cores have an overall smooth appearance most of them are not only dynamically young but also quite disturbed. The nature of orbits is therefore an important indicator of the dynamical state of both the inner and outer regions of the cluster. Principally radial orbits, expected at the outskirts, are the signature of a region dominated by infall, whereas isotropic orbits imply the existence of a well-mixed region.  Previous studies of the velocity dispersion profile and estimates of the degree of anisotropy in clusters have provided ambiguous results, primarily due to the lack of knowledge of the underlying mass distribution. The velocity structure of galaxies in clusters have been studied in detail (Coma and A2670) by \\citeN{kent82}; \\citeN{the86}; \\citeN{merritt87}; \\citeN{sharples88}; \\citeN{colless95} and several other groups. These analyses used the observed galaxy positions and velocities to constrain the distribution of total mass and simultaneously find consistent and physically meaningful solutions for the velocity anisotropy.  \\citeN{the86} examined the uncertainties in the virial mass profiles derived for Coma from observational data. They showed that a wide range of mass models are consistent, consequently permitting a large range of orbital structures. Mass models that were more compact (but had low overall masses) implied circular orbits for the galaxies whereas the higher mass models implied predominantly radial orbits. \\citeN{merritt87} examined the relative distribution of the dark and luminous matter in Coma and showed that it was impossible to distinguish between models where the galaxies trace the total mass and were on isotropic orbits versus those in which the dark matter was very concentrated and the galaxies were on primarily transverse orbits. Therefore, sufficient constraints on the total mass distribution and the velocity anisotropy cannot be obtained simultaneously using only the luminous tracers. In the recent survey of the Coma cluster by \\citeN{colless95}, they find that the velocity distribution is highly non-Gaussian. The dynamics can be better interpreted in terms of an on-going merger between two sub-clusters, thus indicating that the system is not in virial equilibrium.  Similar studies of the cluster A2670 have also been inconclusive; in the faint photometric and spectroscopic survey by \\citeN{sharples88}, both extremely anisotropic models and nearly isotropic ones were indistinguishable in terms of goodness of fit with respect to the available data.  In these clusters and others that have been studied, the principal uncertainty in the determination of the anisotropy arises from ignorance of the distribution of the total mass. In our analysis, independent data from lensing that constrains the overall mass distribution allows the elimination of this largest source of uncertainty.  The plan of this paper is as follows: in Section 2 we review observational probes of the internal dynamics of clusters and construct mass models from X-ray data, lensing data, and from N-body simulations. The mathematical formalism that forms the basis of our approach is presented in Sections 3 \\& 4. The robustness of the method is demonstrated for several fiducial forms of the total mass profile (Section 5), and the technique is then applied to the cluster A2218 (Section 6). Finally, we discuss our results and their implications for the physical state of the core of A2218. Throughout this paper, we assume $H_{0} = 50\\,$km s$^{-1}$Mpc$^{-1}$, $\\Omega = 1$ and $\\Lambda = 0$. ",
+        "conclusions": "Gravitational lensing provides the `cleanest' way to construct the total mass profile for a cluster independent of the kinematic details; additionally, combining strong and weak lensing removes the scaling ambiguity allowing the calibration of other independent mass models. With `good' data for an individual cluster, the requirements for consistency on the smallest to the largest scales are stringent enough to constrain the slope of generic density profiles for rich clusters. Accurate mass profiles are crucial to settling many important issues such as the baryon fraction problem and in understanding the discrepancies and biases arising in the X-ray, lensing and virial mass estimates for clusters.  In this paper, we have demonstrated that the dynamics and velocity structure of the core of galaxy clusters can be probed given an independently inferred total mass profile. The future applications of our method to study cluster cores are promising, given the prospect of collecting more spectro-photometric data of galaxies in cluster lenses (e.g. \\citeNP{yee96}).  With current data, we find strong evidence for the existence of an anisotropic central region. This is consistent with the picture of on-going relaxation, wherein anisotropies in the velocity tensor can arise naturally as a consequence of the initial conditions coupled with evolution. Given the range of complex physical processes that operate in cluster cores that could alter galaxy orbits, for instance; dynamical friction (dynamical friction in an aspherical cluster can induce and amplify the velocity anisotropy as demonstrated by \\citeN{binney77}, a possible origin for the inferred velocity anisotropy, specially in the case of A2218), potential fluctuations arising due to the presence of substructure and the frequent presence of a cD galaxy at the centre of the cluster potential it is not surprising that the core is not isotropic.  Distinguishing between the dynamical effects of the various physical mechanisms in order to model them satisfactorily, in addition to requiring from the observations more accurately determined line-of-sight velocity dispersion profiles for clusters would enable the application of this technique more effectively. Further extension of this analysis to incorporate the dynamics of the intra-cluster gas with the lensing model self-consistently, is required in order to understand the possible role of baryons in the dynamics of cluster cores."
+    },
+    "9602/astro-ph9602109_arXiv.txt": {
+        "abstract": "VLA\\footnote{The Very Large Array is a part of the National Radio Astronomy Observatory, which is operated by Associated Universities Inc, under contract with the NSF} B-configuration snapshots have been obtained at 6cm and 20cm of 43 radio quasars with a range of known galaxy cluster environments. The radio sources are characterised by measures of morphology, flux, and size. These are studied as a function of cluster density, characterised by published B$_{gq}$ numbers. We find no strong dependence of radio source properties on this measure of cluster density. We find $\\sim2\\sigma$ evidence that sources in higher cluster densities have more lobe-dominated morphology with higher lobe luminosities and lower core luminosities. The 7 unresolved sources are found in low density clusters, as are the most bent sources. We find that QSO clusters have fewer other radio sources within a radius of 3 arcminutes than do the central galaxies of a sample of rich X-ray selected clusters. ",
+        "introduction": "We have obtained radio images of 43 QSOs of redshift  0.7 or less with a range of known galaxy environments from optical studies. The program was planned to study the connection between radio source properties and the galaxy cluster density around the QSO.The cluster density of the QSOs is measured by the B$_{gq}$ number derived from the galaxy companion data largely published separately by Yee and Ellingson (1993 and references therein). These  numbers have units of Mpc$^{1.77}$. We have optical images of some of the clusters, but do not yet have a uniform dataset that allows us to measure other cluster properties, such as size, or the location of the QSO within the cluster.  The radio observations were made in the B-configuration of the VLA, in December 1991, taking $\\sim$5 minute snapshots of ~50 sources at 6cm and 20cm, interspersed with calibration sources, and determining the flux density scale using 3C48 and 3C138. The images were self-calibrated where possible. The synthesised beam in the final images is $\\sim$4 arcsec at 20cm and $\\sim$1.3 arcsec at 6cm. The largest-scale structure fully determined in the B-configuration observations is $\\sim$100 arcsec at 20cm, and $\\sim$35 arcsec at 6cm, and care was taken when discussing the radio structure of QSOs with angular sizes comparable with or larger than these values. We have also used the A-configuration maps of Price et al (1993) and Gower and Hutchings (1984) of 10 of the sample QSOs, and also of 11 more QSOs with B$_{gq}$ measures, for which we have A-configuration measures by Hutchings, Price and Gower (1988), or Gower and Hutchings (1984). This gives us a total sample of 54 QSOs with radio data, of which we have B$_{gq}$ values for 45.  A similar investigation has been published by Rector, Stocke, and Ellingson (1995), based on published radio maps of QSOs with known B$_{gq}$ values. Since the B$_{gq}$ value is the common parameter, there is an overlap of 16 objects in our specifically observed sample, and another 9 for which we had earlier maps, within their sample of 30 objects. We discuss the comparison of conclusions at the end of this paper.  The maps and related details will be published separately. The list and measures relevant to this paper are listed in Table 1. The sources were characterised by  measures of flux density, morphology, size, and shape, also shown in Table 1. The bend angles are defined by the core and hot-spots in the lobes and are usually well-defined. The source types are defined in the Table 1 footnotes, and are based on whether one or two lobes are present, and the relative flux densities of core and lobes at 20cm. The classification depends on the resolution of small-scale structure (hence redshift), and dynamic range where the lobe flux per beam area is very low. The system is the same as that used in papers on QSO and radio galaxy structure from A-configuration maps, as referred to above. Many of the quantities in the A-configuration maps are taken from the measures in the references given. We have also listed a source `complexity' measure assigned independently by two of us from visual inspection of the maps. This is meant mainly to quantify non-symmetry which is not measured by source type and bend.   We also noted where other faint radio sources were detected in the field, to investigate the presence of other cluster sources, and also to study the possible effect of the cluster on background sources. The map dynamic range was usually in the range 500 to 1200 (peak flux density to RMS noise). Values outside this range were found in exceptionally weak sources where thermal noise dominated at 0.1 mJy at 20cm and 0.04 mJy at 6cm, or strong sources where dynamic range dominated ($\\sim$2000 for peak fluxes above 2Jy: 4 sources have noise above 1mJy at 6cm and 3 have noise above 2mJy at 20cm). ",
+        "conclusions": "Rector et al (1995) have published tables and analysis similar to ours on the QSO source structure. Their conclusions are in agreement with ours, but there are some detailed differences we should note. The individual source sizes and bend angles differ in detail, presumably as a result of map noise, dynamic range, resolution, and occasionally differences in interpreting complex sources. An example of the latter is 0903+169, where we give a bend of 30$^o$ compared with their value of 58$^o$ (which they flag as uncertain). In this instance there are several compact knots near the 20cm map centre, and the bend angle measured depends on which is assumed to be the core. We believe we have chosen correctly by adopting the one which is brightest in our 6cm map. Some of their bend measures depend on the optical position when the radio core was not visible: our measures are based on a homogenous set of maps obtained by ourselves, and only radio source features are used. Thus, our database is more complete and uniform. Similarly, the flux measures are all measured from our maps, and we regard them as more complete and homogenous than those quoted by Rector et al. In some cases, there are significant differences between their values and ours, which look like cases where the core or lobe fluxes were difficult to measure or may have suffered from poor dynamic range. We also note that the linear sizes used by Rector et al, correspond to H$_0$=50, and not 100 as stated in their paper. Thus, the sizes we quote (for H$_0$=100, q$_0$=0.5) are about a factor two smaller than theirs. Finally, we note that they have chosen to split high and low B$_{gq}$ at 500, which corresponds to a certain Abell richness, while we have chosen 350 on the basis of the numbers in our sample. Again, the difference is not very signficant, and a matter of choice.  The overall conclusion we reach is similar to that of Rector et al: that there are no strong correlations with B$_{gq}$. The significance level of the differences we have highlighted between high and low B$_{gq}$ sources are at the 2 to 3$\\sigma$ level, either by K-S tests between the two groups, or by the correlation with B$_{gq}$ as a parameter.  The lack of strong correspondence with B$_{gq}$ suggests that this quantity does not measure the IGM well. There have been recent reports of correspondence between radio and hot gas (X-ray) structure, in low redshift clusters such as Virgo and Perseus (e.g. Brinkman 1996 and Owen 1996). If QSOs are at the centre of their clusters, they may have little movement through the cluster, and any IGM movement may be infall (i.e. not transverse to the radio structure). It would be of interest to look for radio structure correlations with the QSO position within the cluster, and the velocity structure of the clusters. Ellingson, Green and Yee (1991) have published velocity studies of a few of the QSO clusters in the sample, but there are too few measurements to make any comparisons. They do note that  the cluster velocity dispersions are fairly low, and that they see no significant differences in the relative velocities of companion galaxies between radio-loud and radio-quiet QSOs.  We have counts of sources in the QSO cluster fields, which are lower than those found in very rich clusters at comparable redshifts. We have not yet been able to determine whether these extra sources are associated with cluster galaxies. If they are cluster members, the mean difference in redshift corresponds to a difference of a factor close to 2 in limiting luminosity. If they are background sources, there may be implications of gravitational lensing. We are pursuing these studies with optical investigations.  Finally, we note that our new data add to and reinforce our results from earlier studies of QSO radio structure in the distribution of bend angle and core flux with apparent size. Thus, QSOs in cluster environments do not appear to have large differences from the behaviour of QSOs selected without knowledge of their environment.  We thank Erica Ellingson and Travis Rector for useful discussions. The new radio maps were made in collaboration with Erica Ellingson and will be published separately.   \\newpage"
+    },
+    "9602/astro-ph9602014_arXiv.txt": {
+        "abstract": "We present a new method to compute stellar ages in Globular Clusters (GC) that is  ten times more precise than the traditional isochrone fitting procedure. The method relies on accurate stellar evolutionary tracks and on photometry for GCs complete down to the main sequence, and it is based on counting number of stars in two different regions of the CMD: the red giant branch and the main-sequence. We have applied this method to the globular cluster M68 and found an age of 16.4$\\pm$0.2 Gyr for $(m-M)_V=15.3$. This new method reduces the error associated to the uncertainty in the distance modulus by a factor of two, the error due to the choice of the value for the mixing length parameter to almost zero and the error due to the colour-$T_{\\rm eff}$ transformation to zero. ",
+        "introduction": "Globular clusters are among the oldest objects in the Universe and also the tracers of the collapse of the Galaxy, and they are very important cosmological probes for the age of the Universe. The determination of the age of GCs is still an open problem. GCs ages have been recently reviewed by Chaboyer (1995). When all possible random and systematic errors are taken into account, an error bar of 5 Gyr is associated to any age determination using the main sequence turnoff isochrone fitting method (MSTO). An alternative method has been proposed in order to cure some of the problems of the MSTO procedure (Jimenez et al. 1996), giving lower ages than the MSTO.  The main problem in the MSTO comes from the fact that the isochrone has to fit the position of the main sequence turnoff that is not a point on the CMD, but rather an extended region. The same is true for the alternative method developed by Iben \\& Renzini (1984).  In this paper we present a new method to determine ages of GCs that does not rely on fitting any particular morphological feature in the CMD diagram and that allows us to reach a precision of 0.2 Gyr for a given distance modulus. The method is based on a careful computation of stellar evolutionary tracks and on counting stars in two different regions of the CMD: the red giant branch (RGB) and the main sequence (MS), down to a magnitude where the sample is complete.  The paper is organized as follows. In section 2 we describe the theoretical stellar evolution models used in this work. In section 3 we present the method and compute the age for M68. Section 4 contains a comparison with other methods. We finish with the summary and conclusions. ",
+        "conclusions": "We have developed a new method to determine the ages of GCs. Using theoretical evolutionary tracks we have predicted the relative number of stars in the main sequence and in the RGB, as a function of age and distance modulus.  The dependence of the age on the distance modulus is twice smaller than what is found using the traditional isochrone fitting method, but the accuracy of the age determination for a given distance modulus is ten times higher, because the present method is based just on counting stars in different luminosity bins and therefore does not have troubles with fitting the morphology of the MSTO (mixing length parameter and colour calibration).  In Table 2 we show a comparison of the errors involved in computing GCs ages using the MSTO and our method.   We have applied this method to the old halo GC M68 and found an age of $16.4 \\pm 0.2$ Gyr, if the distance modulus $(m-M)_{\\rm V}=15.3$ determined by Jimenez et al. (1996), is used.  This value is in good agreement with previous age determinations found by Chaboyer, Sarajedini, Demarque (1992), using isochrone fitting, and Jimenez et al. (1996), using the HB morphology technique."
+    },
+    "9602/gr-qc9602029_arXiv.txt": {
+        "abstract": "Using the Hubble parameter as new `inverse time' coordinate ($H$-formalism), a new method of reconstructing the inflaton potential is developed also using older results which, in principle, is applicable to any order of the slow-roll approximation. In first and second order, we need three observational data as inputs: the scalar spectral index $n_s$ and the amplitudes of the scalar and the tensor spectrum. We find constraints between the values of $n_s$ and the corresponding values for the wavelength $\\lambda $. By imposing a dependence $\\lambda (n_s)$, we were able to reconstruct and visualize inflationary potentials which are compatible with recent COBE and other astrophysical observations. From the reconstructed potentials, it becomes clear that one cannot find only one special value of the scalar spectral index $n_s$. ",
+        "introduction": "Because the scalar field is not observable, we also specify the dependence of the potential on the wavelength $\\lambda$. From (\\ref{lamH}) together with (\\ref{eps}) and (\\ref{gh}), we get \\be \\frac {\\lambda }{\\lambda_0} \\simeq \\left ( 1 - \\frac {\\Delta }{A H^2} \\right )^{1/(2\\Delta )}  \\label{wave1} \\quad , \\qquad H^2 \\simeq \\frac {\\Delta }{A} \\left ( 1 - \\left [ \\frac {\\lambda }{\\lambda_0} \\right ]^{2\\Delta } \\right )^{-1} \\; , \\label{wave2} \\ee where $\\lambda_0$ is an integration constant. The value of the inflationary potential (\\ref{xxx}) is determined by \\be U(\\lambda , A, n_s) \\simeq \\frac {\\Delta }{\\kappa A (1-[\\lambda/\\lambda_0]^{2\\Delta })} \\left ( \\Delta + 3 - \\frac {\\Delta }{1-[\\lambda/\\lambda_0]^{2\\Delta }} \\right )  \\; . \\label{un3} \\ee  For $n_s=1$ and $A>0$, we obtain the solution \\ben H^2 & \\simeq & - \\frac {1}{2A\\ln [\\lambda/\\lambda_0] } \\; , \\label{wave3} \\\\ U   & \\simeq & - \\frac {1}{2 \\kappa A \\ln [\\lambda/\\lambda_0] } \\left ( 3 + \\frac {1}{2\\ln [\\lambda/\\lambda_0] } \\right ) \\; . \\label{un4} \\een From (\\ref{un}), (\\ref{un1}), (\\ref{un2}), (\\ref{un3}), and (\\ref{un4}) we recognize that $A$ merely scales the potential. From the exact equation (\\ref{lamHgen}) we find $\\lambda /\\lambda_0 = \\exp (-1/(2AH^2))/H$ for $n_s=1$.  From (\\ref{wave2}) or (\\ref{wave3}), respectively, and the scalar field solutions (\\ref{pp}) and (\\ref{pp2}), we can deduce the following reality conditions: \\ben 0 < n_s < 1 \\; , \\quad A \\neq 0 \\quad & \\Longleftrightarrow & \\quad [\\lambda/\\lambda_0]^{2\\Delta } > 1  \\label{cond1} \\\\ 1 < n_s < \\infty \\; , \\quad A \\neq 0   \\quad & \\Longleftrightarrow & \\quad 0 < [\\lambda/\\lambda_0]^{2\\Delta } < 1  \\label{cond2} \\\\ n_s = 1 \\; , \\quad A>0  \\quad & \\Longleftrightarrow & \\quad 0 < [\\lambda/\\lambda_0] < 1  \\label{cond3} \\een These conditions determine a possible dependence $\\lambda =\\lambda (n_s)$ which may be seen in future observations. The result (\\ref{cond1})-(\\ref{cond3}) for such a functional dependence is clarified by three examples. Let us suppose, the behavior is $\\lambda/\\lambda_0 (n_s) := n_s^2 \\; ;$ then the potential $U(\\phi )$ has the shape shown in Fig.~\\ref{fig.3}. For the functional dependence $\\lambda/\\lambda_0 (n_s) := 1/n_s^2 \\; ,$ the potential of Fig.~\\ref{fig.4} results. For $\\lambda/\\lambda_0 (n_s) := 1/\\sqrt{n_s}$, we find Fig.~\\ref{fig.5}.   \\section {\\bf Second-order approximation}   Our $H$-formalism allows, in principle, to determine $g(H)$ to arbitrary high order in the Taylor expansion. In second order, we obtain a nonlinear second order equation for the graceful exit function.  The second-order result of the scalar spectral index is \\be 1-n_s = 4 \\epsilon - 2 \\eta + 8 (1+C) \\epsilon^2 - (6+10 C) \\epsilon \\eta + 2 C \\epsilon \\xi  \\; , \\label{diff} \\ee where $C=-2+\\ln (2) +\\gamma \\sim -0.73$ and $\\gamma \\sim 0.577$ is the Euler constant \\cite{stelyt}. In the $H$-formalism, Eq.~(\\ref{diff}) converts into the nonlinear second order equation \\be \\Delta = 2 \\frac {g}{y} - g' - 4 (1+C) \\left ( \\frac {g}{y} \\right  )^2 + (3+4C)\\frac {g}{y} g' - 2 C g g''  \\label{g2y} \\ee for the graceful exit function $g$, where $y:=H^2$ and $'=d/d(H^2)$. In terms of $\\epsilon = -g/y$, cf.~(\\ref{eps}), we can rewrite this condition as \\be 2 C \\epsilon\\, {\\buildrel{\\bullet\\bullet}\\over\\epsilon} - (2C+3)\\epsilon\\, {\\buildrel\\bullet\\over\\epsilon} - {\\buildrel\\bullet\\over\\epsilon} + \\epsilon^2  + \\epsilon + \\Delta =0  \\; , \\label{diff2} \\ee where $\\bullet \\hat= d/d\\ln y$. Equation (\\ref{g2y}) has the exact $\\Delta$--dependent solution \\be g = \\left ({1\\over2} \\pm \\sqrt{{1\\over4} - \\Delta}\\right) y \\; . \\label{g1} \\ee The solution with the plus sign can be ruled out because we require $g<0$, while the solution with the minus sign possesses an inflationary part if and only if $n_s<1$. The potential for this $g$ was already constructed \\cite{schunck}, it belongs to the class of power-law models: \\be U(\\phi ) = \\frac {3-A}{\\kappa } C_3{}^2 \\exp (\\pm \\sqrt{2 \\kappa A}\\; \\phi ) \\; . \\ee Note, however, that in second order the integration constant $A={1\\over2} - \\sqrt{{1\\over4} - \\Delta}$ is now fixed by the observational data for the scalar spectral index $n_s$. Because of the nonlinearity of (\\ref{g2y}), further solutions exist. In order to obtain more inflaton potentials $U(\\phi)$ in second order, we rewrite (\\ref{phiH}) into the form \\be \\frac {d \\phi}{dy} = \\mp \\sqrt{\\frac {2}{\\kappa }} \\frac {1}{\\sqrt{-4 g y}} \\; . \\label{4gy} \\ee This equation together with (\\ref{g2y}) combines to a coupled system which can be solved numerically, for example by using MATHEMATICA. The potential $U(\\phi )$ is given by the parametric solution $\\{ \\phi (H), U(H)=(3H^2+g)/\\kappa \\}$; see Fig.~\\ref{fig.6}. ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602146_arXiv.txt": {
+        "abstract": "We discuss the likely sources of turbulence in the ISM and explicitly calculate the detailed grand source function for the conventional sources of turbulence from supernovae, superbubbles, stellar winds and HII regions. We find that the turbulent pumping due to the grand source function is broad band, consequently expanding the inertial range of the cascade. We investigate the general properties of the turbulent spectrum using a simple approach based on a spectral transfer equation derived from the hydrodynamic Kovasznay approximation. There are two major findings in this paper: I. The turbulent pressure calculated from the grand source function is given by $p_{turb} \\sim 10-100 p_{thermal}$.  II. With the scale dependent energy dissipation from a turbulent cascade the multi-phase medium concept can be generalized to a more natural continuum description where density and temperature are functions of scale. Approximate power-law behavior is seen over a large dynamic range.  We discuss some of the implications of models for the ISM, particularly addressing the energy balance of the warm neutral phase. However, this paper is not concerned with the detailed structure of the ISM of the Galaxy. Rather, we attempt here to broaden our approach to the global turbulent structure of the ISM and also to move from the current many-phase description to a more natural scale-dependent continuum of phases. ",
+        "introduction": "Interstellar turbulence in the multi-phase interstellar medium (ISM) has an inertial range stretching over at least 9 orders of magnitude. The substantial viscous and magnetic dissipation at large wavenumbers requires significant energy sources to maintain the spectrum. The evidence for a turbulent spectrum in the multi-phase ISM comes mainly from pulsar observations (Armstrong \\etal 1981, Rickett 1990, Narayan 1992, Armstrong \\etal 1995) which are consistent with a cascade from scales of order $\\sim 100$~pc to scales down to $10^{11}$~cm. Furthermore, the power spectrum appears to be consistent with a Kolmogorov law as reviewed recently by Sridhar and Goldreich (1994). Larson (1981) has also noted a turbulent spectrum in diffuse neutral clouds and molecular clouds. Turbulent pressure plays a crucial role in the various phases of the ISM. This conclusion follows essentially from the work of Kulkarni \\& Fich (1985) for the cold neutral HI component, and of Reynolds (1985, 1995) for the warm ionized medium. McKee (1990) argued that turbulent pressure has to represent the dominant pressure term in the ISM. Similar conclusions, based on the analysis of the vertical structure of the HI in the Galaxy, were reached independently by Lockman \\& Gehman (1991) and Ferrara (1993); Boulares \\& Cox (1990) also recognized the importance of turbulent pressure. There is general agreement on the fact that, following the seminal suggestion by Spitzer (1978), turbulent motion must be associated with energy deposition fromhigh mass stars in the form of winds, HII regions and supernova explosions.  There are a number of outstanding problems with the studies of turbulence in the ISM. It is well known that simple estimates of turbulent dissipation are too high by three orders of magnitude or so (Scheuer 1968, Cesarsky 1980, Spangler 1991, Ferriere \\etal 1988, Falgarone \\& Puget 1995). The dissipation problem is intimately related to a Reynolds number that is too low. Using typical numbers from Spangler we find, in general, a Reynolds number of $\\R\\sim 400$. Furthermore, the associated dynamic range of the cascade is only $10^2-10^5$ which is inconsistent with the observations.  The most recent studies of Shridhar \\& Goldreich (1994), Goldreich \\& Sridhar (1995) indicate a solution to the previous problems if the spectrum is generated by a magnetized four wave cascade in the wavenumber perpendicular to the field (see also Montgomery \\& Mattheus 1995; Oughton, Priest \\& Mattheus 1994; Hossain \\etal 1995).   Let us now outline the structure of the paper and indicate where to find the principal results. In Section 2 we discuss the general properties of turbulent cascades and their usefulness in understanding the ISM. The mathematical formalism for calculating the turbulent spectrum for a given broad-band source function is presented. We also derive a formula for the turbulent pressure as a function of the grand source function. We derive in Section 3 the grand source function for turbulence in the ISM resulting from supernovae, superbubbles, stellar winds, and HII regions. The properties of the individual source functions are given in Table I. Figure 1 shows the individual contributions to the spectrum and also the total grand source function. The dynamical equations are then solved numerically with the grand source function as input and the results are given in Figure~2. In Section 4 we generalize our analysis to structures called clouds and using the scale-dependent power input from turbulence we are able to calculate the multi-phase properties of an ensemble of turbulently heated clouds. We give a characteristic p-T diagram in Figure 3. In Figures 4 and 5 we indicate how the density and temperature of clouds depends on scale. Section 5 discusses the results of these, albeit idealized, multi-phase calculations and we indicate how they may explain some of the more commonly observed phenomena in the ISM. ",
+        "conclusions": "We have discussed the physics and mathematical description of the turbulent energy cascade in the ISM (Section 2). The grand source function for turbulent energy sources has been derived including the contribution of conventional sources as supernovae, superbubbles, stellar winds, and HII regions (Section 3). The results are shown in Figure 1 and the properties of the individual source functions are given in Table I. Supernovae dominate over most of the spectral range although superbubbles can dominate on large scales. We find that the turbulent pumping due to the grand source function is broad band, consequently expanding the inertial range of the cascade. The dynamical equations are solved and the results presented in Figure 2. The turbulent pressure calculated from the grand source function (Section 2.2) is given by $p_{turb} \\sim 10-100 p_{thermal}$. With the scale dependent energy dissipation from a turbulent cascade the multi-phase medium concept becomes a continuum description where density and temperature are functions of scale (Section 4). Approximate fractal behavior is seen over a large dynamic range. The thermal properties of the ISM on the large scales are similar to the conventional models. On smaller scales turbulent heating can dominate and a multi-phase structure can develop for the hydrodynamic cascade discussed here. A typical p-T diagram is presented in Figure 3. The derived temperature and density of clouds in this multi-phase structure are functions of the scale of the observing measurement and might be regarded as an approximate fractal. Typical results exhibiting this behaviour are given in Figures 4 and 5.  We have ignored various potentially important feedback effects in this paper, regarding them as important to incorporate in the next stage of calculational complexity. For example, when we use the grand source function in the calculation of the phase diagram there are temperature dependences hidden in the $R_{0i}$'s that may create potential feedback effects on the phase diagram structure shown in Figs. 3, 4, and 5. Such feedback effects will be discussed in a subsequent paper.   This is a steady state calculation with no consideration given to the initial conditions and how the system can evolve to the steady state. In fact, it may be difficult for the system to reach this state.  If the temperature is less than the temperature associated with the maximum of the cooling curve that a substantial fraction of the gas must go over, then the hot phase will never be formed. If for other physical reasons the hot phase is formed then the equilibrium state can be maintained. The two-equilibrium temperatures are actually two foci in the phase-plane $T - {\\dot T}$ of the time-dependent equation. The essential point is that the regions of attraction of the two foci are such that the system will remain attracted to the lower equilibrium temperature unless it is put over the hump of the cooling curve and can be attracted to the higher equilibrium temperature. The condition for this is exactly as described above.  A further complication arises when a magnetized turbulent cascade is utilized. This is a necessary ingredient since the ISM is certainly magnetized. The spectrum found is similar to a Kolmogorov spectrum but is highly anisotropic and is a Kolmogorov spectrum in the wave-number perpendicular to the magnetic field i.e. $E_M(k_{\\perp})\\sim k_{\\perp}^{-5/3}$ (Shridhar \\& Goldreich 1994, Goldreich \\& Shridhar 1995,Oughton, Priest and Mattheus 1994, Montgomery \\& Mattheus 1995). As discussed in detail by Ferriere \\etal (1988), Spangler (1991), and Goldreich \\& Shridhar (1995), Alfv\\'en waves in the cascade will propagate without damping in the very highly ionized ISM but will dissipate rapidly on the edges of clouds consistent with the qualitative discussion in this paper. However, the details of the dissipative processes need further study and clarification since there seems to be a significant discrepancy between the observed inertial range of the turbulence of up to 10 orders of magnitude and the rather high dissipation rates inferred for the standard ISM.  The source of the turbulent pressure that holds up the ISM against its own weight is clearly evident in our analysis with the principal being energy input form massive stars via supernova remnants and superbubbles. The spectrum of the predicted velocity has been presented and it may be possible to constrain our results by observations of large scale turbulence via measurements of the large scale velocity fields (HI, H$\\alpha$, IR fine cooling fine structure lines, recombination lines). As we shall discuss in a later paper, many other heat sources come into play in the ISM (Wolfire \\etal 1995) including direct shock heating from supernovae and superbubbles, photoelectric heating from grains, soft X-ray heating etc. (Wolfire \\etal 1995). Many of these heat sources are scale-dependent. Therefore, the range of validity of turbulent heating effects for our Galaxy may be limited to the cooler material.  The scale-dependent phase continuum model of the ISM presented here may be a beginning of a more useful description of the ISM than the current many-phase models particularly for the understanding and integration of the many detailed observations of the ISM."
+    },
+    "9602/gr-qc9602037_arXiv.txt": {
+        "abstract": "}[1]{\\vspace{19mm}  \\noindent{\\small{\\em Abstract.} #1}\\vspace{2mm}  } \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\begin{document} \\begin{flushright} Z\\\"urich University Preprint\\\\ ZU-TH 4/96 \\end{flushright} \\title{Asymptotic Behavior of the Einstein-\\kern -2.18ptYang-Mills-Dilaton\\\\[2mm] System for a Closed Friedmann-Lemaitre Universe$\\footnote{This work was supported by the Swiss National Science Foundation.}$} \\authors{Reto Eggenschwiler, David Scialom and Norbert Straumann} \\address{Institute of Theoretical Physics, University of Z\\\"urich, Winterthurerstrasse 190,\\\\CH-8057 Z\\\"urich, Switzerland} \\abstract{We study the coupled Einstein-Yang-Mills-Dilaton (EYMD) equations for a Fried\\-mann-Le\\-mai\\-tre universe with constant curvature $k=1$. Our detailed analysis is restricted to the case where the dilaton potential and the cosmological constant vanish. Also assuming a static gauge field, we present analytical and numerical results on the behavior of solutions of the EYMD equations. For different values of the dilaton coupling constant we analyze the phase portrait for the time evolution of the dilaton field and give the behavior of the scale factor. It turns out that there are no inflationary stages in this model.} ",
+        "introduction": "For a long time the matter used in most cosmological models was described in terms of ideal or viscous hydrodynamics. This is certainly adequate at least back to the time of nucleosynthesis, because---to a first approximation---we are then allowed to regard the matter as a mixture of pressure-free matter (baryons, WIMPS, etc...) and radiation (photons, neutrinos). At very high energies the fluid description has to be replaced by field theoretical models (and, when approaching the Planck scale, by yet unknown theories of matter). \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent With the success of gauge field theories, and especially in connection with various proposed inflationary scenarios \\cite{Linde}, field theoretical models became very popular in cosmology. A description of topological defects, which could be created in cosmological phase transition, requires also field theoretical models which allow for spontaneous symmetry breaking. \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent Inflationary stages appear more or less naturally in many models, but most of those which lead to successful inflation are very phenomenological in nature and are not based directly on ``fundamental'' field theories. Since inflation is a very attractive general idea and for other reasons as well, it is important to study the evolution of a large variety of cosmological models based on interesting field theoretical matter descriptions. In addition to Yang-Mills fields, the existence of a dilaton field is theoretically well motivated on the basis of superstring theory \\cite{Green} and Kaluza-Klein theories \\cite{Abbot}. \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent In the present work we study the coupled Einstein-Yang-Mills-Dilaton (EYMD) equations for a Friedmann-Lemaitre (Robertson-Walker) space-time. The basic equations for this model lead to a dynamical system of six first-order nonlinear ordinary differential equations, which we discuss with the modern techniques of the theory of dynamical systems. In view of the rather high dimension of phase space this goal is not easy to achieve. In a first step we therefore analyze in detail only a restricted class of solutions. We do so by setting the cosmological constant and the dilaton potential equal to zero, and also require that the Yang-Mills field is static, which can be done consistently. In this case the phase space becomes three-dimensional. We hope that our analysis of this reduced system is at least methodically of some interest. An extension of the discussion to the full system is planned. \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent In a first step we compactify the phase space such that the dynamical system can be extended to the enlarged space and determine the critical points. With one exception, all critical points are either located at infinity of the original phase space or in the unphysical region. We analyze the nature of these singular points and the qualitative behavior of the flow in their vicinity. In general the equilibrium points are hyperbolic. However, for special values of the dilaton coupling constant there exists a critical point in the asymptotic physical region which becomes non-hyperbolic. For this case we use the theory of normal forms and apply the Poincar\\'e-Dulac theorem to get rid of the nonresonant nonlinear terms up to third order. Analytical formula which are valid near the critical points are given in all cases. \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent The asymptotic behavior provides one with the initial conditions for a numerical study of the phase portrait. Special attention is given to the value of the dilaton coupling constant which is obtained from superstring theory. Unfortunately, and not unexpectedly, we do not find any inflationary stages. \\hfil\\vadjust{\\vskip\\parskip}\\break\\indent The paper is organized as follows. In section~2 we derive the basic equations. Their analysis for the above-mentioned assumptions is performed in section~3. We analytically determine the asymptotic behavior of the solutions near the singular points. Section~4 is devoted to numerical results. The behavior of the scale factor as well as the phase portrait of the dilaton field are given. A short summary concludes the paper. ",
+        "conclusions": "For the special case of a static gauge field and $V=\\Lambda =0$ we have performed a complete analysis of the coupled EYMD equations for a Friedmann-Lemaitre universe with constant curvature $k=1$ (closed model). For this situation, we did {\\it not} recover any inflationary stages. It might therefore be interesting to extend our discussion to the more general case of a dynamical gauge field, perhaps with a non-vanishing cosmological constant $\\Lambda $ and/or a dilaton potential $V$. This is, however, a considerably more involved task, because we encounter then six coupled first-order nonlinear differential equations.  The analysis of the coupled EYMD equations for a Friedmann-Lemaitre universe with constant curvature $k=0$ (flat model) and $k=-1$ (open model) are also of interest, since related considerations~\\cite{David,Bel} show that inflationary stages occur for trajectories which come close enough to one of the separatrices found for $k=0$. \\vskip 2\\baselineskip \\noindent{\\bf Acknowledgments} \\vskip \\baselineskip \\noindent We are grateful to Othmar Brodbeck, George Lavrelashvili and Mikhail Volkov for very useful and clarifying discussions, and to Marcus Heusler for careful reading of the manuscript."
+    },
+    "9602/astro-ph9602136_arXiv.txt": {
+        "abstract": "We show that the short spin-up time observed for the Vela pulsar during the 1988 ``Christmas'' glitch implies that  the coupling time of the pulsar core to its crust is less than $\\sim$ 10 seconds. Ekman pumping cannot explain the fast core-crust coupling and a more  effective  coupling is necessary. The internal magnetic field of the Vela pulsar can provide the necessary coupling if the field threads the core with a magnitude  that exceeds $ 10^{13}$ Gauss for a normal interior and $ 10^{11}$ Gauss for a superconducting interior. These lower bounds favor the hypothesis that the interior of neutron stars contains superfluid neutrons and protons and challenge the notion that pulsar magnetic fields decay over million  year time scales or that magnetic flux is expelled from the core as the star slows. ",
+        "introduction": "Observations of the Vela pulsar during the December 24, 1988 ``Christmas'' glitch provide a  glimpse of the coupling  between the solid crust and the fluid interior of a neutron star. The phase residuals for this event  are shown in Figure 1 (\\cite{MHMK}). The nearly linear decay in the  residuals immediately after the glitch shows that the angular velocity of the surface  first changed  abruptly - within the $\\sim 120$ s time resolution of the observations -  and then varied much more slowly. As we show below, this abrupt change strongly constrains the internal structure of the Vela pulsar.  A neutron star is composed of a solid crust containing at most a few percent of the star's moment of inertia, a liquid core, and a superfluid neutron liquid that coexists with the crystal  lattice of the inner part of the crust. If both the neutrons and protons in the core are superfluid, they  would be strongly coupled by  Fermi liquid effects (\\cite{ALS};  \\cite{AS88}). This quantum fluid together with the electrons, which are needed to maintain charge neutrality, act as a single classical liquid  whose viscosity is supplied mainly by the electrons. In this Letter, we study the coupling of this core liquid  to the star's crust.  In current models of pulsar glitches (\\cite{PA}, \\cite{LE96}), the observed spin up is due to the transfer of angular momentum from the neutron superfluid in the inner-crust to the solid crust. In these ``crust-initiated\" models, subsequent exchange of angular momentum between the crust and the core  brings these two components into rotational equilibrium. In alternative ``core-initiated\" models, the initial spin-up occurs in the  core rather than the crust (\\cite{S89}) and the crust then catches up  to the core's angular velocity.  Models in which glitches are almost entirely a crustal phenomenon, with only weak coupling to the core, are ruled out by timing data from accreting neutron stars. This type of ``crust-only\" mechanism would require that the crust of the Vela pulsar remains decoupled from the core for months, the characteristic post-glitch relaxation time scale. However,  analyses of timing noise in accretion-powered pulsars (Boynton \\& Deeter 1979; Boynton 1981; Boynton \\etal\\ 1984)  indicate that $\\gap14$\\% of the stellar core couples to the crust  in much less than a month, ruling  out the  possibility of  ``crust-only\" glitch mechanisms.  For either the crust-initiated or core-initiated models, the  Vela Christmas glitch strongly constrains the core-crust coupling time-scale. A linear coupling model for crust-interior interaction can illustrate the types of constraints different glitch models  yield.  If $I_C$ and $\\Omega_C(t)$ are the moment of inertia and  angular velocity of the solid crust,  $I_I$ and $\\Omega_I(t)$ are the moment of inertia and angular velocity  of the liquid interior (which we take to behave as a solid body for this schematic example), and $t_{CI}$ is the  coupling time scale between the crust and the interior, we can write: \\begin{equation} I_C\\dot \\Omega_C(t) = - \\dot J_{CI}(t) + \\dot J_{C, {\\rm S}}(t) \\ , \\end{equation} \\begin{equation} I_I\\dot \\Omega_I (t) =   \\dot J_{CI}(t) + \\dot J_{I,{\\rm S}}(t) \\ , \\label{eq2}\\end{equation}  where \\begin{equation} \\dot J_{CI}(t) \\equiv I_r {\\Omega_C(t) -\\Omega_I(t) \\over t_{CI}}, \\end{equation} and $$I_r = {I_C I_I\\over (I_C+I_I)}$$ is the reduced moment of inertia. For a crust-initiated glitch the source term  is $\\dot J_{C,{\\rm S}} (t) = \\tau_{C}$,  for $ 0<t<t_{su}^C$ (and $\\dot J_{I,{\\rm S}} (t) =0$); whereas for a core-initiated glitch  $\\dot J_{I,{\\rm S}}(t) = \\tau_{I}$, for $ 0<t<t_{su}^I$ (and $\\dot J_{C,{\\rm S}} (t) =0$), where the torques $\\tau_{C}$ and $\\tau_{I}$ are taken to be constants.   We solve the linear coupling model for both a core- and crust-initiated glitch and obtain the phase residuals: \\begin{equation} \\Delta \\phi(t) = \\int_0^t [\\Omega_C(0) - \\Omega_C(t') ] dt' . \\end{equation}  For  $t \\ge t_{su}^C$ and $t \\gg t_{CI}$, the phase residuals for a crust-initiated glitch are given by, \\begin{equation} \\Delta\\phi_C (t) =-\\frac{\\tau_{C}\\, t_{su}^C}{I_C+I_I}\\,\\left(t-\\frac{t_{su}^C}{2}+\\frac{I_I} {I_C}\\, t_{CI} \\right) \\ , \\end{equation}   while, for  a core-initiated glitch and $t \\ge t_{su}^I$, \\begin{equation} \\Delta\\phi_I (t) =-\\frac{ \\tau_{I} \\, t_{su}^I}{I_C+I_I}\\,\\left(t-\\frac{t_{su}^I}{2}-t_{CI} \\right)  \\ . \\end{equation} By fitting the data from the Christmas glitch to the above equations, taking into account the uncertainty in when the glitch started, we  constrain the coupling and spin-up times. The short-term behavior of the glitch model enables us to place upper limits on these time scales by requiring that the calculated phase residuals agree with the data to within the observational uncertainty ($\\sim 0.2$ milliperiods). Acceptable models lie between the curves shown in Figure 1, where the dashed curves  show the phase residuals for two crust-initiated glitch models and the solid curves show results for two core-initiated models.  We find that for a crust-initiated glitch the   core-crust coupling time scale must satisfy \\begin{equation} t_{CI} \\lap 300 \\, {\\rm s} \\ {I_C \\over I_I} \\ , \\end{equation}  and the spin-up time  must be  $ t_{su}^C \\lap 1200$ s [this constraint on the spin-up time is in agreement with the theoretical models of Link and Epstein (1996)].  Since $I_C / I_I \\sim 0.03$, we have \\begin{equation} t_{CI} \\lap 10 \\, {\\rm s} \\ . \\end{equation}  The upper bound on the coupling time scale is much shorter than the experimental resolution of the spin-rate changes due to  the small moment of inertia of the crust.  Crust-initiated models have shown excellent agreement with the observed behavior of glitching pulsars (\\cite{LE96}) and will be the main focus of the discussion below. For completeness, we point out that for core-initiated glitch models the requirement for the core-crust coupling time scale is less severe, $t_{CI} \\lap 440$~s, and the spin-up time  must satisfy $t_{su}^I\\lap 1200$~s. ",
+        "conclusions": "We have argued that fluid dynamic processes such as Ekman pumping do not provide the adequate core-crust coupling that is implied by the Christmas glitch.  The needed coupling can be provided by a magnetic field that threads both the core and crust of the star. If both the neutrons and protons in the core are normal, the required average radial flux density in the core is $\\gap  10^{13}$ Gauss, about three times greater than the surface field deduced from the spin down rate.  Such a strong internal magnetic field argues against the core containing only normal neutrons and protons. However, if the core did possess fields of this strength,    the surface field might actually {\\it grow} as the interior field diffuses out through the crust.  It is likely that  both the neutrons and protons are superfluid and the threading field need only have a mean flux density of $\\gap 10^{11}$ Gauss.  This constraint allows some of the surface magnetic flux to be confined to the outer crust of the star from which it could  decay in $10^6 -10^7$ years. However, since the magnetic diffusion time for fields that thread the core is $\\gap 10^{10}$ years (\\cite{UVR}), the decay of the crustal magnetic field would  still leave a field of $\\gap 10^{11}$ Gauss at the surface.  It has been suggested that the magnetic flux tubes in the superfluid core may be swept out by the expanding neutron vortex lattice as the rotation of the star decreases (\\cite {DCC}, \\cite{SBMT}). Our constraints on the present internal magnetic field of  the Vela pulsar indicate that the flux tubes and vortex lines manage to pass through each other. Since the present internal magnetic field is $\\gap 10^{11}$ G,  the initial field would have to be  $B_{\\rm initial}  \\gap 9 \\times 10^{12}(P_{\\rm initial}/1\\; {\\rm ms})^{-1}$ Gauss for the sweeping interpretation to be correct. If the initial spin period of the Vela pulsar was $\\sim 1$ ms, the flux sweeping hypothesis would require that the initial internal magnetic field  be considerably stronger than the present surface field.    Future observations may  provide deeper insights into  the nature of neutron star interiors. Monitoring of the Vela pulsar  at higher time resolution could set more stringent constraints on the internal magnetic field. In addition,  XTE observations of accreting neutron stars can complement glitch observations in  the study of neutron star interiors.   We would like to thank G. Baym, B. Link, and J. Sauls for helpful discussions. M.A. and A.O. were supported in part by the DOE at Chicago, by NASA grant NAG 5-2788 at Fermi National Laboratory, and through a collaborative research grant from IGPP/LANL. This work was carried out under the auspices of the U.S. Department of Energy.     \\def\\nature{{\\rm Nature}} \\def\\nucphys{{\\rm Nuc. Phys.}} \\def\\nucphysa{{\\rm Nuc. Phys. A}} \\def\\physletb{{\\rm Phys. Lett. B}} \\def\\physrevc{{\\rm Phys. Rev. C}} \\def\\physrevd{{\\rm Phys. Rev. D}} \\def\\sovphysjetp{{\\rm Soviet~Phys.~JETP}} \\def\\ptpl{{\\rm Progr.Theor.Phys.Lett}} \\def\\ptps{{\\rm Prog.Theor.Phys.Suppl.}} \\def\\ptp{{\\rm Prog. Theor. Phys.}}"
+    },
+    "9602/astro-ph9602122_arXiv.txt": {
+        "abstract": "Our search for high-redshift Type Ia supernovae discovered, in its first years, a sample of seven supernovae.  Using a ``batch'' search strategy, almost all were discovered before maximum light and were observed over the peak of their light curves. The spectra and light curves indicate that almost all were Type Ia supernovae at redshifts $z = 0.35$ -- 0.5.  These high-redshift supernovae can provide a distance indicator and ``standard clock''  to study the cosmological parameters $q_0$, $\\Lambda$, $\\Omega_0$, and $H_0$.  This presentation and the following presentations of Kim {\\it et al.} (1996), Goldhaber {\\it et al.} (1996),  and Pain {\\it et al.} (1996) will discuss observation strategies and rates, analysis and calibration issues, the sources of measurement uncertainty, and the cosmological implications, including bounds on $q_0$, of these first high-redshift supernovae from our ongoing search. ",
+        "introduction": "Since the mid 1980's, soon after the identification of the supernova subclassifications Ia and Ib were recognized, Type Ia supernovae (SNe Ia) have appeared likely to be homogeneous enough that they could be used for cosmological measurements.  At the time, it appeared that they could be used to determine $H_0$, if their absolute magnitude at peak could be measured, i.e., if some SN Ia's distance could be calibrated.  They could also be used to determine $q_0$ from the {\\em apparent} magnitudes and redshifts of nearby and high redshift supernovae, if high redshift SNe Ia could be found.  Many of the presentations at this meeting have demonstrated significant advances in recent years in  our understanding of the SN Ia homogeneity (and/or luminosity calibration), the distances to SNe Ia, and measurements of $H_0$.  We will here discuss the search for high redshift supernovae and the measurement of $q_0$. ",
+        "conclusions": "Given the error bars, our current measurements of $q_0$ do not yet clearly distinguish between an empty $q_0=0$ and closed $q_0 > 0.5$ universe. The data do, however, indicate that a decelerating $q_0 \\ge 0$ is a better fit than an accelerating $q_0 < 0$ universe.  This is an important conclusion since it limits the possibility that $q_0 = \\Omega_0 /2 - \\Omega_\\Lambda$ is dominated by the cosmological constant $\\Lambda$ (where $\\Omega_\\Lambda$ is the normalized cosmological constant  $\\Lambda (3 H_0)^ {-2}$). In an accelerating universe, high values for the Hubble constant do not conflict with the ages of the oldest stars, because the universe was expanding more slowly in the past. However, in a ``flat'' universe, in which $\\Omega_0 + \\Omega_\\Lambda = 1$, a  cosmological constant $ \\Omega_\\Lambda \\ge 0.5$ would give an acceleration $q_0 \\le -0.25$, a poor fit to our data.  Note that extinction of the distant supernovae would make this evidence of deceleration even stronger, as would the possibility of inhomogeneous matter distribution in the universe as described by Kantowski, Vaughan, \\& Branch \\nocite{ka:clump}(1995), since both of these would lead to underestimates of $q_0$ if not taken into account.  (Our future analyses will bound or measure extinction using multiple colors to differentiate reddening from intrinsic color differences within the SN Ia family.  Riess {\\em et al.}\\nocite{ri:aigua} (1996) discussed this approach at this conference.)  It will be important to compare the $q_0$ value from our most distant supernovae to the value from our closest high-redshift supernovae to look for Malmquist bias. An estimate of the size of this bias using our studies of our detection efficiency as a function of magnitude suggests that a 0.2 mag intrinsic dispersion in calibrated SN Ia magnitudes would lead to an overestimate of  0.1 in $q_0$, if not accounted for (Pain {\\em et al.} 1996, and see also Schmidt {\\em et al.} 1996).  We will complete the observations and analyses of the next 11 supernovae soon, and be able to greatly strengthen the certainty of these conclusions.   The new data set is expected to have significantly smaller error bars, particularly after calibration of the width-brightness relation, and it should be possible to distinguish the empty, $q_0 = 0$, and closed, $q_0 > 0.5$, universes. The redshift of these supernovae are plotted in the inset of Figure 5.  Note that the highest redshift supernova, at $z \\sim 0.65$, would be over 0.35 magnitudes brighter in a closed universe than in an empty universe.  For the future, it should be possible to measure both $\\Omega_0$ {\\em and} $\\Lambda$, even if the universe is not assumed to be flat.  The apparent magnitude of a ``calibrated candle'' depends on both  $\\Omega_0$ and $\\Lambda$, but with different powers of $z$.  The discovery of still higher redshift supernovae, at $z \\sim 1$, can therefore locate our universe in the $\\Omega_0$ vs. $\\Lambda$ parameter plane, as shown in Figure 6.  The feasibility of this measurement is discussed in Goobar and Perlmutter \\nocite{go:lambda}(1995).   \\begin{figure} \\psfig{figure=aigualambda.eps,height=3in} \\caption[figurecaption]{ Theoretical bands of allowed parameter space in the $\\Omega_{\\Lambda}$ versus $\\Omega_0$ plane, for the hypothetical example of two SNe Ia discovered at $z=0.5$ and at $z =1$ in a universe with  $\\Omega_{\\Lambda} = 0$ and  $\\Omega_0$ =1 (from Goobar and Perlmutter 1995).  The width of the bands is due to uncertainty in the apparent magnitude measurement of the supernovae. The intersection region would provide a measurement of both $\\Lambda$ and $\\Omega_0$.  The shaded corners are ruled out by other observations, as indicated.} \\label{lambda} \\end{figure}   \\vspace{0.22in}  We have shown here that the scheduled discovery and follow up of batches of pre-maximum high-redshift supernovae can be accomplished routinely, that the comparison with nearby supernovae can be accomplished with a generalized K-correction, and that a width-brightness calibration can be applied to standardize the magnitudes. Ideally this will now become a standard method in the field, and SNe Ia beyond $z = 0.35$ will become a well-studied distance indicator---as they already are for $z<0.1$---useful for measuring the cosmological parameters.  The prospects for this look good:  already several other supernova groups have now started high-redshift searches.  Schmidt {\\em et al} \\nocite{sc:aigua}(1996) and Della Valle {\\em et al} \\nocite{de:aigua}(1996), at this meeting discussed two of these.  The many presentations at this exciting meeting have discussed recent advances in our empirical and theoretical understanding of Type Ia supernovae. The resulting supernova-based measurement of the cosmological parameters is an impressive consequence of these efforts of the entire supernova research community.  \\vspace{.25in}  The observations described in this paper were primarily obtained as visiting/guest astronomers at the Isaac Newton and William Herschel Telescopes, operated by the Royal Greenwich Observatory at the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias; the Kitt Peak National Observatory 4-meter and 2.1-meter telescopes and Cerro Tololo Interamerican Observatory 4-meter telescope, both operated by the National Optical Astronomy Observatory under contract to the National Science Foundation; the Keck Ten-meter Telescope; and the Siding Springs 2.3-meter Telescope of the Australian National University.  We thank the staff of these observatories for their excellent support.  Other observers contributed to this data as well; in particular, we thank Marc Postman, Tod Lauer, William Oegerle, and John Hoessel for their more extended participation in the observing. This work was supported in part by the Physics Division, E. O. Lawrence Berkeley National Laboratory of the U.~S. Department of Energy under Contract No. DE-AC03-76SF000098, and by the National Science Foundation's Center for Particle Astrophysics, University of California, Berkeley under grant No. ADT-88909616."
+    },
+    "9602/astro-ph9602064_arXiv.txt": {
+        "abstract": "We present the $R$-band luminosity function for a sample of 18678 galaxies, with average redshift $z = 0.1$, from the Las Campanas Redshift Survey (LCRS). The luminosity function may be fit by a Schechter function with $M^* = -20.29 \\pm 0.02 + 5 \\log h$, $\\alpha = -0.70 \\pm 0.05$, and $\\phi^* = 0.019 \\pm 0.001 \\ h^3$~Mpc$^{-3}$, for absolute magnitudes $-23.0 \\leq M - 5 \\log h \\leq -17.5$ and $h = H_0 / (100$~km~s$^{-1}$~Mpc$^{-1})$. Over the same absolute magnitude range, the mean galaxy density is $0.029 \\pm 0.002 \\ h^3$~Mpc$^{-3}$ for a volume extending to $cz = 60000$~km~s$^{-1}$.  We compare our luminosity function to that from other redshift surveys; in particular our luminosity function normalization is consistent with that of the Stromlo-APM survey, and is therefore a factor of two below the normalization implied by the $b_J \\approx 20$ bright galaxy counts. Our normalization thus indicates that much more evolution is needed to match the faint galaxy count data, compared to minimal evolution models which normalize at $b_J \\approx 20$. Also, we show that our faint-end slope $\\alpha = -0.7$, though ``shallower'' than typical previous values $\\alpha = -1$, results primarily from fitting the detailed shape of the LCRS luminosity function, rather than from any absence of intrinsically faint galaxies from our survey.  Finally, we find that the faint end of the luminosity function is dominated by galaxies with emission lines. Using [OII] 3727 equivalent width $W_{\\lambda} = 5$~\\AA \\ as the dividing line, we find significant differences in the luminosity functions of emission and non-emission galaxies, particularly in their $\\alpha$ values; emission galaxies have Schechter parameters $M^* = -20.03 \\pm 0.03 + 5 \\log h$ and $\\alpha = -0.9 \\pm 0.1$, while non-emission galaxies are described by $M^* = -20.22 \\pm 0.02 + 5 \\log h$ and $\\alpha = -0.3 \\pm 0.1$. The average [OII] 3727 equivalent widths do not change significantly with redshift, consistent with a star formation rate which stays constant over the depths sampled by the LCRS. This result holds for galaxies of different luminosities, and over the respective redshift ranges that these galaxies may be observed, in particular up to about $z = 0.2$ for galaxies brighter than $M^*$. ",
+        "introduction": "The galaxy luminosity function is one of the most basic and fundamentally important quantities in observational cosmology. For example, it is essential for the analysis of large-scale structure statistics from redshift surveys, and knowledge of the local luminosity function is necessary for the proper interpretation of the faint galaxy count data and galaxy evolution at high redshift (e.g., \\cite{koo92}).  Though much work has gone into determination of the local $z = 0$ optical luminosity function (see \\cite{fel77}; \\cite{2:KOSII}, 1983; \\cite{eep}; \\cite{bin88}; and recent work by \\cite{lov92}; \\cite{marz94}), there is still some uncertainty in its faint end as well as in its normalization (e.g., \\cite{coll95}). In this paper we investigate the local luminosity function by presenting results for a sample of 18678 Las Campanas Redshift Survey (LCRS) galaxies, selected from CCD photometry in a ``hybrid'' red Kron-Cousins $R$-band (see \\S~\\ref{lumdata} below). This is the largest galaxy sample for which the luminosity function and mean density have been computed. The combination of the mean survey depth of $z = 0.1$ and the large survey volume ($1.0 \\times 10^7 \\ h^{-3}$~Mpc$^3$ out to $cz = 60000$ km~s$^{-1}$) imply that our results are not biased by local inhomogeneities. In addition, the spectra of LCRS galaxies allow us to study the luminosity function for samples divided by [OII] 3727 emission strength, a rough indicator of the star-formation rate (\\cite{ken92}).  We describe relevant aspects of the LCRS construction in \\S~\\ref{lumdata}. We then detail our methods for measuring the galaxy luminosity function and mean density in \\S~\\ref{methods}. Our results are presented in \\S~\\ref{lfrholcrs}, and we compare them to those of other redshift surveys in \\S~\\ref{comparisons}. Some implications of our results for faint galaxy counts and galaxy evolution are discussed in \\S~\\ref{counts}. We summarize our conclusions in \\S~\\ref{conclusions}. ",
+        "conclusions": "\\label{conclusions}  We have used inhomogeneity-independent maximum-likelihood methods to calculate the luminosity function of LCRS galaxies in the (hybrid) Kron-Cousins $R$ absolute magnitude range $-23.0 \\leq M - 5 \\log h \\leq -17.5$, where we find that it may be reasonably described by a Schechter function with $M^* = -20.29 \\pm 0.02 + 5 \\log h$, $\\alpha = -0.70 \\pm 0.05$, and $\\phi^* = 0.019 \\pm 0.001 \\ h^3$~Mpc$^{-3}$. There appears to be an excess relative to the Schechter function at the faint end $M > -17.5$. The mean galaxy density for the range $-23.0 \\leq M \\leq -17.5$ is $0.029 \\pm 0.002$ galaxies $h^3$~Mpc$^{-3}$ out to a redshift of $60000$ km~s$^{-1}$, and the run of cumulative density with redshift suggests that the density has converged within the LCRS survey volume. The measured (isophotal) galaxy luminosity density is $\\rho_L = (1.4~\\pm~0.1) \\times 10^8 \\ L_{\\sun} \\ h$~Mpc$^{-3}$, resulting in $\\Omega \\approx 0.2$, for a typical virial mass-to-light ratio of $300 \\ h \\ (M/L)_{\\sun}$ and a rough isophotal-to-total light correction. We have also divided our sample by the cut [OII] 3727 equivalent width $W_{\\lambda} \\geq 5$~\\AA \\ , and showed that the emission and non-emission galaxies so defined have significantly different luminosity functions, with emission galaxies dominating the faint end ($\\alpha = -0.9 \\pm 0.1$), non-emission galaxies dominating the bright end ($\\alpha = -0.3 \\pm 0.1$), and about equal numbers near the full-sample $M^*$.  We then compared our results to those from the optically-selected Stromlo-APM, CfA, and SSRS2 redshift surveys. Despite the fact that the Stromlo-APM survey has an $\\alpha \\approx -1$, and the LCRS has $\\alpha = -0.7$, the shapes of two luminosity functions are remarkably similar, though they do differ in specific details. Moreover, our luminosity function normalization is consistent with that of Stromlo-APM; the CfA normalization is higher but this may be biased by density inhomogeneities in its smaller and more local survey volume. Finally, our emission galaxy results confirm the picture seen in these previous surveys that the faint end of the luminosity function is dominated by late-type, blue galaxies which are presumably also gas-rich and strong in [OII] 3727 emission.  Finally, two implications of our results may be made with respect to the faint galaxy counts and to galaxy evolution. First, our normalization of the luminosity function is a factor of two below the ``high'' normalization often adopted from the $b_J \\approx 20$ APM galaxy counts in calculations used to match fainter galaxy count data. If our normalization is not an underestimate (e.g., due to surface brightness effects), then it implies that more evolution is required in the models, compared to using the high APM count normalization. Second, for the redshifts sampled in the Las Campanas survey, up to about $z = 0.2$ for galaxies brighter than $M^*$, our results do not indicate a significant change of average [OII] 3727 equivalent width with redshift. The star-formation rate, as measured by [OII] 3727, has apparently not evolved rapidly over the volume sampled by the LCRS, unlike the large increases seen in the numbers of blue star-forming galaxies in smaller, deeper surveys at $z \\gtrsim 0.3$.  In this paper we have used the Las Campanas Redshift Survey to provide an accurate determination of the local luminosity function for the largest galaxy sample thus analyzed. Our red $R$-band results are also complementary to previous results in blue photometric bands. In addition, we have begun to explore the luminosity function for galaxy samples divided by emission properties, and in future analyses we will examine galaxy samples defined by $b_J - R$ colors (which are available for our Southern galaxies from APM observations), and if possible, also by basic morphological information to be determined from our drift-scan images. The Las Campanas Redshift Survey should thereby provide a useful and large database of local galaxy properties, including luminosity, emission, color, and possibly morphological information, upon which to anchor models of faint galaxy counts and galaxy evolution at higher redshifts."
+    },
+    "9602/astro-ph9602103_arXiv.txt": {
+        "abstract": "The correlation function $\\xi(r)$ of matter in the non-linear regime is assumed to be determined by the density profiles $\\rho(r)$ and the mass distribution $n(M)$ of virialized halos. The Press--Schechter approach is used to compute $n(M)$, and the stable clustering hypothesis is used to determine the density profiles of these Press--Schechter halos. Thus, the shape and amplitude of $\\xi(r)$ on small scales is related to the initial power spectrum of density fluctuations.  The case of clustering from scale-free initial conditions is treated in detail.  If $n$ is the slope of the initial power spectrum of density fluctuations, then stable clustering requires that $\\xi(r)\\propto r^{-\\gamma}$, where $\\gamma$ is a known function of $n$.  If halo--halo correlations can be neglected, then $\\rho(r)\\propto r^{-\\epsilon}$, where $\\epsilon = (\\gamma+3)/2 = 3(4+n)/(5+n)$.  For all values of $n$ of current interest, this slope is steeper than the value $3(3+n)/(4+n)$ that was obtained by Hoffman \\& Shaham in their treatment of the shapes of the outer regions of collapsed halos. Our main result is a prediction for the amplitude of the non-linear correlation function. The predicted amplitude and its dependence on $n$ are in good quantitative agreement with $N$-body simulations of self-similar clustering.  If stable clustering is a good approximation only inside the half-mass radii of Press--Schechter halos, then the density contrast required for the onset of stable clustering can be estimated. This density contrast is in the range $\\sim 300-600$ and increases with the initial slope $n$, in agreement with estimates from $N$-body simulations. ",
+        "introduction": "Large-scale structure in the universe is thought to arise through the gravitational clustering of matter. In the nonlinear regime most of the dark matter is in bound halos which have separated out from the expanding background. Therefore the auto-correlation function, $\\xi(r)$, of the dark matter is closely related to the shapes of halos. In this paper we explore the consequences of this connection for the shape and amplitude of $\\xi(r)$ on small scales.  There are two approximations that are commonly used to describe the evolution of gravitational clustering in the non-linear regime.  The first is the assumption that, in this regime, the clustering is statistically stable.  By this, one usually means that non-linear, virialized objects no longer participate in the expansion of the background Universe; they maintain their shapes in physical coordinates, so they shrink in comoving coordinates (e.g. Peebles 1965, 1980).  The second approximation was first formulated by Press \\& Schechter (1974).  They assumed that, on average, structures collapse spherically and then virialize, with initially more dense regions collapsing first, and less dense regions collapsing later.  Further, they assumed that when they virialize, all halos have the same density (about 200 times the background density at the time of virialization) whatever their mass.  By applying these assumptions to an initially Gaussian density field, they derived the distribution of halos as a function of mass.  In the Press--Schechter approach, the clustering is hierarchical, and virialized halos at a given epoch are progenitors of the halos that virialize at a later epoch. It is not obvious that these two idealizations, stable clustering, and virialization at a fixed multiple of the background density at the time of virialization, are compatible.  For example, consider a Press--Schechter halo with uniform density (it has a `tophat' density profile).  Assume that at the time it virializes, say $t_1$, it has mass $M_1$ and density $178\\rho_{\\rm b}(t_1)$.  This sets its radius $R_1$.  In the Press--Schechter approach, at some later time $t_2>t_1$, it will have merged with other halos into an object of greater mass $M_2>M_1$.  This more massive object will have density $178\\,\\rho_{\\rm b}(t_2)<178\\,\\rho_{\\rm b}(t_1)$.  On the other hand, if stable clustering is correct, then the core $M_1$ within $R_1$ will still have density $178\\,\\rho_{\\rm b}(t_1)$ so that the shell of mass $(M_2-M_1)$ around the core must be less dense than the core itself.  Thus, if it had a tophat density profile at time $t_1$, then the halo will not have a tophat density profile at the epoch $t_2$.  In other words, if Press--Schechter halos are to evolve consistently with the stable clustering assumption, then they cannot have tophat density profiles. This toy example illustrates that by requiring consistency between Press--Schechter and the stability approximation we can constrain the allowed density profiles of halos.  In particular, it is known that the stable clustering approximation allows one to constrain the shape of $\\xi(r)$ in the non-linear regime (e.g., Peebles 1974a; section 26 in Peebles 1980; section 5.4 in Padmanabhan 1993). Section 2 uses this stable clustering shape of $\\xi(r)$ to constrain the density profiles of Press--Schechter halos. The formalism of McClelland and Silk (1977) is then used to compute the amplitude of $\\xi(r)$ in the nonlinear regime. The predicted amplitude is compared with the results of N-body simulations. The density profile inferred in Section~2 is different from the value that is obtained in other descriptions of spherical collapse, such as the Hoffman-Shaham model.  Section~3 discusses the reasons for the difference.  Section~4 studies the density contrast required for the onset of stable clustering.  It argues that only the cores of Press--Schechter halos are likely to have evolved consistently with both Press--Schechter and stable clustering.  This has the consequence that the density required for the onset of stable clustering should be an increasing function of $n$, where $n$ is the slope of the power spectrum of initial fluctuations.  Section~5 discusses the case of non-power law density profiles and summarizes the results. ",
+        "conclusions": "We have shown that requiring consistency between the stable clustering hypothesis and the Press--Schechter approach suggests that the density profiles of Press--Schechter halos should be power laws with slope $\\epsilon = 3(4+n)/(5+n)$, where $n$ is the slope of the initial fluctuation spectrum (equation~\\ref{sc}).  This value differs from that expected from secondary infall studies of the collapse of peaks associated with initially Gaussian fields (Hoffman \\& Shaham 1985). The reasons for this difference were discussed in Section~3.  The Press--Schechter model, with the stable clustering density profile, was used to compute the amplitude of the correlation function in the highly non-linear regime (Section 2).  This amplitude was shown to be in good quantitative agreement with that measured in $N$-body simulations.  Moreover, the amplitude was shown to be a function of the slope $n$ of the initial power spectrum.  The $n$-dependence resembles that which was obtained recently by Jain \\etal (1995) and Padmanabhan \\etal (1995).  Our calculation of $\\xi$ (equation~\\ref{mcsilk}) rests on the assumption that, in this (small separation) regime, the correlation function is determined primarily by the density profiles of virialized halos.  That is, the crucial assumption is that, for these small separations, most pairs are most likely to be drawn from the same halo, so that correlations between different halos have a negligible effect on the shape of $\\xi$.  While this assumption makes intuitive sense, we have yet to show that it is accurate.  We can do this as follows.  Equation~(\\ref{mclps}) shows that the correlation function can be written as a sum over halos of different masses.  To demonstrate that our approach is at least self-consistent, we must show that the correlation function is, indeed, determined primarily by halos having diameters that are larger than the separation scale $r$.  So, consider the limit $r\\to 0$ of vanishingly small separations.  Then the amplitude of the correlation function is given by equation~(\\ref{xips}).  Of course, in this limit, all halos are larger then the separation scale.  However, as a crude approximation that should also give some indication of the result on slightly larger scales, we will compute the fractional contribution to $\\xi$ from `small' halos, which we will define to be all those with $M/M_* \\le 0.2$.  Equation~(\\ref{xips}) shows that this contribution is simply an incomplete Gamma function.  Thus, the small mass contribution is \\begin{displaymath} \\gamma\\Bigl[(11+n)/(10+2n),0.2^{n+3/3}\\Bigr]\\,\\Big/\\, \\Gamma\\Bigl[(11+n)/(10+2n)\\Bigr] , \\end{displaymath} which is 11, 15, 20, and 24 per cent for $n=1,0,-1$ and $-2$ respectively. For non-zero values of $r$, the corresponding integrals can be done numerically.  They also show that $\\xi$ is determined primarily by pairs within larger halos, so that our neglect of halo--halo correlations is justified.  The other crucial ingredient in our calculation is knowledge of the density profiles of virialized halos.  We assumed that these profiles were power laws, and determined the slopes of the power laws by requiring that they yield a correlation function with the shape required by stable clustering.  However, recent work suggests that, while the profile shape may be described by an average power law slope that is consistent with equation~(\\ref{sc}), the density profiles of halos that form in $N$-body simulations of gravitational clustering from scale free initial conditions are not simple power laws (Navarro, Frenk \\& White 1995; Cole \\& Lacey 1995; Tormen 1995).  Rather, the density profiles seem to be well fitted by a function of the form $\\rho \\propto 1/[r(r+b)^2]$, where the core radius $b$ depends both on the mass $M$ and the shape of the initial power spectrum.  \\begin{figure} \\centering \\mbox{\\psfig{figure=rhos.ps,height=2.5in,bbllx=54pt,bblly=70pt,bburx=563pt,bbury=449pt}} \\caption{ The density profile $\\rho(s) \\propto 1/[s(s+a)^2]$, where $s=r/r_{\\rm vir}$, $a = b/r_{\\rm vir}$, and $b$ is the core radius, for a range of values of $a$.  When $s \\gsim 0.1$, the power law profile that corresponds to $n=-2$ (solid bold curve) describes the $a=0.25$ profile quite well.  The dashed bold curve shows the power law corresponding to the $n=1$ case; when $s \\gsim 0.1$ it fits the $a=0.1$ profile reasonably well.} \\label{nfw} \\end{figure}  We expect our formalism to work for these non-power-law profiles also because, except for the inner regions $r\\sim 0.1\\,r_{\\rm vir}$, where $r_{\\rm vir}$ is the virial radius, they are reasonably well described by the power-law profiles we have been considering.  To illustrate this, Fig.~\\ref{nfw} shows the non-power law profile $\\rho(s) \\propto 1/[s(s+a)^2]$, where $s=r/r_{\\rm vir}$ and the core radius is given by $a = b/r_{\\rm vir}$. The plot shows this profile for a range of values of the scaled core radius $a$. These values were chosen because $N$-body simulations show that $a\\sim 0.25$ describes the $n=-2$ simulations quite well, whereas $a$ is smaller ($\\sim 0.15$) when $n=0$ (Cole \\& Lacey 1995).  For comparison, the thicker curves show the corresponding stable clustering power law profiles for the two extreme cases studied here, $n=-2$ and $n=1$.  Notice that, in the outer regions where $s \\gsim 0.1$, the $n=-2$ (solid bold) curve describes the $a=0.25$ profile quite well, whereas the $n=1$ (dashed bold) curve fits the $a=0.1$ profile.  Now, the correlation function in the non-linear regime is essentially an integral over the density profile (to give $\\lambda_M$), followed by an integral over the mass function. For the choices of $n$ and $a$ shown in Fig.~1, the corresponding $\\lambda_M(r)$ curves for the power law and the non-power law cases are similar, on scales larger than about $0.1\\,r_{\\rm vir}$.  This is hardly surprising, since Fig.~\\ref{nfw} shows that the Navarro \\etal and power law density profiles $\\rho(s)$ are similar in the outer regions where $s \\gsim 0.1$ (which contain most of the mass of a halo), and $\\lambda_M$ is simply the density profile convolved with itself.  Since, in our formalism, $\\xi$ is simply the sum of many $\\lambda_M$ curves, this suggests that $\\xi$ for these non-power law profiles should be similar to that for the power law profiles considered in Section~2, at least for separations that are on the order of about $0.1\\, r_*$ and larger.  We have verified that this is indeed the case: for scales on which $\\xi(r)\\lsim 1000$, the predicted amplitude of $\\xi$ changes by less than about $10\\%$ for $-2<n<1$, if we use the Navarro \\etal profile (with the values of $a$ shown in Fig.~1) instead of power law profiles.  On scales smaller than $\\sim 0.1\\, r_*$, it is necessary to take into account the dependence of the core radius on the mass of the halo. Simulations show that, typically, $a$ decreases as $M$ decreases. Navarro \\etal (1995) argue that the qualitative form of this relation can be understood in terms of the formation times of Press--Schechter halos.  Including this trend improves the agreement of our predicted $\\xi$ with the result from the Navarro et al profile. The asymptotic behavior of $\\xi(r)$ as $r\\to 0$ however, is more subtle, as it is very sensitive to the $M \\to 0$ behavior of the core radius $a$ and the Press-Schechter $n(M)$. Here we simply note that by requiring that $\\xi(r)$ have the stable clustering shape, we can obtain an independent constraint on the relation between $a$, $M$, and $n$ for these non-power law profiles.  Since all of the results of this paper concern the stable clustering regime, Section~4 discussed the scale on which stable clustering should become a good approximation.  It used the Lacey \\& Cole (1993) analysis of the merger histories of Press--Schechter halos to show that the onset of stable clustering occurs at higher density contrasts as the slope of the initial fluctuation power spectrum $n$ increases.  This is in qualitative agreement with $N$-body simulations (Jain 1995).  Our derivation of the amplitude of the correlation function in the stable clustering regime, using the Press--Schechter mass multiplicity function, can be extended to cosmological models in which $\\Omega\\le 1$.  This is because the effect of $\\Omega$ on the density profiles of halos has been calculated (e.g. Hoffman \\& Shaham 1985; Hoffman 1988; White \\& Zaritsky 1992), as has the effect on the evolution of the Press--Schechter mass function (Lacey \\& Cole 1993).  Furthermore, recall that both the density profiles and the Press--Schechter multiplicity function depend strongly on the shape of the initial power spectrum.  Thus, our calculation of $\\xi$ shows explicitly how the non-linear correlation function depends on the shape and amplitude of the linear power spectrum.  Therefore, if our approach is correct, then by requiring consistency between the shape of the mass multiplicity function and the shape of the non-linear correlation function, one can estimate the shape and amplitude of the initial perturbation spectrum."
+    },
+    "9602/astro-ph9602045_arXiv.txt": {
+        "abstract": "Experimental determinations of the $pp$ and $pep$ fluxes have great potentialities. We briefly review the reasons that make such measurements privileged tests of neutrino properties. We discuss the predictions for these fluxes given by four {\\em good} solutions to the solar neutrino problem: small- and large-angle MSW and Just-So oscillations into active neutrinos, and small-angle MSW oscillations into sterile neutrinos. In addition, we examine the impact of the planned Hellaz detector, which should measure separately the $\\nu_e$ and $\\nu_{\\mu}$ fluxes in the $pp$ energy window and the signal from the $pep$ neutrinos, for distinguishing among the different solutions and for determining the solar central temperature. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602051_arXiv.txt": {
+        "abstract": "We derive mass functions (MF) for halo red dwarfs (the faintest hydrogen burning stars) and then extrapolate to place limits on the total mass of halo brown dwarfs (stars not quite massive enough to burn hydrogen).  The mass functions are obtained from the luminosity function of a sample of 114 local halo stars in the USNO parallax survey (Dahn \\etal 1995).  We use stellar models of Alexander \\etal (1996) and make varying assumptions about metallicity and about possible unresolved binaries in the sample.  We find that the MF for halo red dwarfs cannot rise more quickly than $1/m^2$ as one approaches the hydrogen burning limit.  Using recent results from star formation theory, we extrapolate the MF into the brown-dwarf regime.  We see that likely extrapolations imply that the total mass of brown dwarfs in the halo is less than $\\sim 3\\%$ of the local mass density of the halo ($\\sim 0.3\\%$ for the more realistic models we consider). Our limits apply to brown dwarfs in the halo that come from the same stellar population as the red dwarfs. ",
+        "introduction": "\\setcounter{footnote}{0} \\renewcommand{\\thefootnote}{\\arabic{footnote}}  In this paper we determine a mass function for low mass stars and then extrapolate to derive an upper limit on the number of brown dwarfs. A key quantity for describing a stellar population is the mass function $\\xi(\\mass)$, the density of stars with mass between $\\mass$ and $\\mass+d\\mass$; for $\\xi(\\mass)$ we use units \\#stars/pc$^3/\\msun$.  The mass function (hereafter, MF) of halo stars can address the question of what makes up the dark matter of our Galaxy. Although it has been clear for some time that main sequence stars heavier than our sun are not sufficiently abundant to explain the mass of the Galaxy, the question has remained whether or not there can be large numbers of lower mass objects.  Very low mass stars or substellar objects have been considered the most plausible candidates for baryonic dark matter.  [White dwarfs, another possibility for baryonic dark matter, are not considered in this paper.]  Recent work has ruled out red dwarfs, stars just barely massive to burn hydrogen ($m > 0.092\\msun$), as a significant source of dark matter (Graff \\& Freese 1996; Boeshaar, Tyson \\& Bernstein 1994; Bahcall \\etal 1994).  Brown dwarfs, star-like objects that are just barely too small to burn hydrogen, are the remaining low mass candidates.  These are extremely difficult to observe optically (see \\eg Burrows and Liebert 1993 for a review) although there have been searches for them with gravitational microlensing (Alcock \\etal 1995, Ansari \\etal 1996).  As our main result, we determine a MF of halo red dwarfs, and then extrapolate into the brown dwarf regime to obtain an upper limit on the number of halo brown dwarfs.  If brown dwarfs are a primary component of halo dark matter, then we would expect a steeply rising MF as we go to decreasing mass.  Instead, we show that likely extrapolations imply that brown dwarfs are not nearly numerous enough to be a significant component of halo dark matter.  As a caveat, let us point out that our limits apply only to brown dwarfs in the halo that come from the same stellar population as the red dwarfs (see the discussion at the end of the paper).  We also rely on recent results from star formation. The goal of our paper is to show that brown dwarfs coming from known stellar populations are not present in great abundance. \\footnote{There have been several attempts to determine the mass function of local stars (Miller \\& Scalo 1979, Kroupa, Tout \\& Gilmore 1993, Henry 1995). Local stars are predominantly Population I disk stars. A population II MF for low mass stars from globular cluster NGC 1261 can be found in Zoccali ${\\it et \\, al}$ (1995). In our paper we focus instead on local field halo stars, which are not in globular clusters.}  Unfortunately, the MF cannot easily be obtained by observations, since it is difficult to measure the mass of a star.  Instead, we must convert from the measured quantity, the luminosity function $\\psi(M)$, the number of stars in a magnitude range $M \\rightarrow M+dM$ (note that $M$ refers to magnitude while $\\mass$ refers to mass).  For the luminosity function (hereafter LF), we use units \\#stars/pc$^3/M$.   To obtain a mass function for red dwarfs, we use the LF from the extensive parallax survey being carried out by the US Naval Observatory.  Preliminary results were reported in Monet ${\\it et \\, al}$ (1992) with a more complete update in Dahn ${\\it et \\, al}$ (1995).  They published an LF based on 114 stars with velocities tangential to the line of sight $V_{\\rm tan}\\ge 220$km/sec. These high velocity stars should all be halo stars (temporarily in the disk) since true disk stars have a local velocity dispersion of only $\\sim$30 km/sec. M\\'{e}ra, Chabrier \\& Schaeffer (1996) used this LF to derive an MF as discussed further below; see their Fig. 1.  In this paper we present a more complete derivation of the mass function for red dwarfs, and a different interpretation.  We take into account several potentially important effects not considered by M\\'{e}ra et al, such as unresolved binaries in the sample and the effects of different choices of metallicity. In particular, many of the `stars' in the sample are likely to be unresolved binaries, and substantial numbers of binaries may drastically alter the results. The MF derived by M\\'{e}ra et al falls into the range of possibilities bracketed by the parameter range we consider. However, our interpretation is different.  We emphasize a different slope of the MF near the hydrogen burning limit.  (By ``slope'', we mean the logarithmic slope ${\\rm d}\\log\\xi / {\\rm d}\\log\\mass$.) In addition we use theoretical results from star formation to extrapolate to lower masses and limit the abundance of halo brown dwarfs.  Because we are using parallax data, we can only discuss the local population of halo stars. ",
+        "conclusions": "Red dwarfs are a tiny fraction, ~0.3\\%, of the mass of the galactic halo (\\cite{gf96}).  In order for any appreciable fraction of the halo mass to consist of low mass star-like objects, there must be substantially more mass in brown dwarfs than in red dwarfs.  We estimated the mass in brown dwarfs by extrapolating the mass function (MF) of red dwarfs below the hydrogen burning limit.  Our extrapolation was based on current theories of star formation, which predict that the MF is either flat or concave down on a log-log plot (or in fig. 1), {\\it i.e.}, the MF rises no faster than a power-law.  We made several assumptions about the metallicity and binary composition of the stars in the sample. We found that likely extrapolations imply that brown dwarfs make up less than $\\sim3\\%$ of the local mass density of the halo.  For our most realistic models, the limit on the total mass in brown dwarfs is roughly the total mass in red dwarfs, $\\sim 0.3\\%$ of the local mass density of the halo.  To repeat an earlier caveat: our limits on the brown dwarf density assume that the brown dwarfs in the halo come from the same stellar population as the red dwarfs.  The possibility always remains that there are large numbers of brown dwarfs from an entirely different population of stars (Population III).  However, theoretical work on star formation (Adams and Fatuzzo 1996) indicates that an earlier population of stars is likely to be skewed towards a predominance of high mass stars, not brown dwarfs.  There may also be some yet undiscovered mechanism of star formation which allows the MF to rise faster than power-law. In this paper we have presented red dwarf mass functions and have shown that brown dwarfs which are members of known stellar populations are not present in the halo in great abundance."
+    },
+    "9602/astro-ph9602098_arXiv.txt": {
+        "abstract": "We have measured the radial brightness distributions in the disks of three nearby face-on spirals: M~51, NGC~3631, and NGC~4321 (M~100) in the photometric bands $B$ through $I$, with the addition of the $K$ band for M~51 only.  The measurements were made by averaging azimuthally, in three modes, the two-dimensional surface brightness over the disks in photometric images of the objects in each band: (a) over each disk as a whole, (b) over the spiral arms alone, and (c) over the interarm zones alone.  From these profiles scale-lengths were derived for comparison with schematic exponential disk models incorporating interstellar dust. These models include both absorption and scattering in their treatment of radiative transfer. The model fits show that the arms exhibit greater optical depth in dust than the interarm zones.  The average fraction of emitted stellar light in $V$ which is extinguished by dust within 3 scale-lengths of the center of each galaxy does not rise above 20\\% in any of them.  We show that this conclusion is also valid for models with similar overall quantities of dust, but where this is concentrated in lanes. These can also account for the observed scale-lengths, and their variations. ",
+        "introduction": "Photometric mapping of spiral galaxies yields both detailed morphology and a set of global physical parameters.  Azimuthally averaged profiles of the surface brightness can usually be decomposed into a central bulge and an exponential disk, whose relative importance is a function of morphological type.  In this paper we will concentrate on the properties of the disks which will be characterized by their observed exponential scale-lengths.  From the fairly extensive detailed studies which have been published to date some general conclusions have been drawn which are still subject to empirical and theoretical debate. Most spiral galaxies have disks with a broadly exponentially declining radial surface brightness distribution (e.g. Kent 1984, de Jong 1995). It seems to be emerging that the scale-lengths observed tend to decrease systematically, but not strongly, with increasing wavelength, an effect which can be attributed to extinction by moderate amounts of dust, with radial metallicity gradients as a possible contributing factor (Elmegreen \\& Elmegreen 1984; Prieto \\etal 1992; Evans 1994; Peletier \\etal 1994, Beckman \\etal 1995, de Jong 1995) The disk surface brightness extrapolated to the axis of a galaxy, $\\mu_0$, lies within a narrow range, for a given photometric band, which for the $B$ band is $\\mu_0$ = 21.6 ($\\pm$ 0.3) mag~arcsec$^{-2}$ (Freeman 1970). For bands further into the red the range is larger, and the central disk surface brightness depends more strongly on spectral type (Peletier \\etal 1994).  These studies of disks have used one-dimensional, azimuthally averaged profiles to derive the parameters cited above.  Galaxy disks are not, however, axisymmetrical; they almost invariably show structures, such as arms and bars. It is obvious from inspection of galaxy images that the arm surface brightnesses are higher at all optical and NIR wavelengths than those of the inter-arm zones of the disk, that arms often have dust lanes at their edges, and that the arms are sites of more concentrated star formation.  Because of the latter two conditions, the brightness profile of the interarm zones may well be a more faithful tracer of the underlying stellar mass distribution, and thus offer a more appropriate way to determine the properties of the exponential disk. In an azimuthally averaged profile of a complete galaxy disk the arm contribution may well be dominant, especially in the shorter wavelengths, $B$ and $V$ bands.  Dust will affect the arms and the interarm disk to different degrees at different wavelengths, and with different geometries.  For these reasons we considered it a valuable exercise to separate the arms from the inter-arm zones, and to perform separate profile analyses. In a related paper (Knapen \\& Beckman 1996), radial profiles in several gas and star formation tracers were studied in detail for NGC~4321, for arm and interarm regions separately, with special emphasis on the role of the atomic hydrogen in the massive star formation process. In the present paper we concentrate on dust, and present our results for three well-studied nearby late-type spirals: M~51 (NGC~5495), NGC~3631, and NGC~4321 (M~100).  In Section 2 we describe the data sets we used, and the methods of analysis applied, in Section 3 we present the observational results, which we compare, in Section 4, to the predictions of a set of theoretical models which incorporate absorption and scattering in their extinction parameters.  In Section 5 we draw conclusions, and deal semi-quantitatively with the consequences of departures from homogeneously exponential radial dust distributions. ",
+        "conclusions": "We have obtained azimuthally averaged radial intensity profiles in the $B$ through $I$ photometric bands, of three nearby spiral galaxies: M~51, NGC~3631, and NGC~4321; for the first of these we also obtained a profile in the $K$ band. As the arms and the interarm disk zones are manifestly different photometrically, we went on to derive separate profiles for the two types of zone in each galaxy, and measured the radial intensity scale-lengths, separately, for the arms, and the interarm zones, as well as for the galaxy disks as a whole.  The results for M~51 and NGC~3631 show scale-lengths which decline systematically, but not steeply, with increasing wavelength, and arm scale-lengths which are longer than the interarm scale-lengths at short wavelengths, but converge to the interarm values at longer wavelength.  For NGC~4321 we find  an inner zone which behaves in the same way as the disks of the other two objects, and an outer zone which contains less dust. However, our observations go out further for NGC~4321, and we might well expect similar results in the outer parts of the other two objects. We have interpreted these results using radiative transfer models which take into account scattering as well as absorption by dust. We characterise the dust distribution by an integrated optical depth, $\\tau_V$ along the central axis of the disk, as well as by a vertical scale-height, and an intrinsic radial scale-length.  Since $\\tau_V$ is a quantity extrapolated inwards to the axis, it represents an upper limit to the values found everywhere in the disk. For M~51 the best model fits to the observed scale-lengths yield values for $\\tau_V$ of close to 10 for the arms, close to 1 in the interarm disk, and a mean of 3 over the disk as a whole.  The corresponding values for NGC~3631 are 10, 4, and 6 respectively.  As well as being upper limits these values yield a misleadingly high impression of the effect of the dust extinction because we tend to think of an optical depth in terms of the most familiar scenario in our own Galaxy where the source lies entirely behind the absorbing dust cloud.  The $\\tau_V$ in the models used here is an integrated value through the center of the (extrapolated) disk, but the stars and dust are distributed throughout the disk, so the net absorption will be much lower than if all the dust were between us and the stellar component.  As an example for $\\tau_V$ = 5 (the on-axis value), and a stellar to dust scale-height ratio of 1, the dust extinction for the absorption + scattering case for the face-on system along its axis is 60\\% of the total light. For a ``screen\" model (all the dust in front of the stars) the equivalent would be 99\\%.  The corresponding figures for a stellar to dust scale-height ratio of 3 are 50\\% and (of course) 99\\% respectively.  We may also include, for comparison, models with pure absorption and no scattering.  For $\\tau_V$ = 5 and the scale-height ratio 1, extinction would take out 80\\% of the light, and for a scale-height ratio 3 this would be 67\\%.  These figures are larger than those for the models including scattering, but they should not be taken as representative of real dust. It is also instructive to make a further calculation: of the fraction of light removed by dust averaged over the disk. To be specific, we have integrated the transfer equations out to 3 scale-lengths from the axis.  Setting $\\tau_V$ = 5 and with the stellar to dust scale-height ratio unity, we find that only 15\\% of the total flux has been removed by dust, this figure is virtually unchanged if we change the scale-height ratio to 3.  It is notable that in a screen scenario the figure would be 70\\%.  From the numbers presented in the previous paragraph we can conclude that the light extinguished in M~51 and in NGC~3631 by the presence of dust in their disks will be of order 15\\% of the total starlight emitted. The qualitative conclusion is that, for these two galaxies, it is most unlikely that a major fraction of their light is being extinguished by interstellar dust.  A further consideration to be taken into account is the non-uniformity of the geometrical distribution of the dust, notably its compression into lanes.  It is clear from any image of typical disk galaxies, such as those in Plate 1, that the dust is in fact concentrated in lanes, especially prominent along the edges of the spiral arms, where the gas is being compressed by the passage of a density wave. We have run models to test for the effects of a dust lane on the scale-length of (say) an arm. To first order we find that a dust lane which occupies a fraction x of an arm, with the rest of the arm dust-free, yields the same effect on the scale-length as if the whole arm were characterized by an optical depth $\\tau_V$, provided that the on-axis extrapolated optical depth of the dust-lane is $\\tau_V$/x.  A modeled lane occupying only 30\\% of the surface of an arm will give essentially the same averaged photometric arm profile as if the arm had dust uniformly distributed across its surface, but the dust lane requires some 3 times the on-axis optical depth as the uniformly dusty arm. These considerations will be dealt with in more depth in a paper devoted purely to the modelling technique.  However we can infer here that the organization of the dust into lanes, and the details of such non-uniformities, would not alter the qualitative conclusion that there is only moderate to low net visible extinction across the disks of M~51 and NGC~3631. One should note that a somewhat similar technique to measure the amount of extinction was used by D.~Elmegreen (1980). She used UBVRI colors of dusty regions in the disk of a few face-on spiral galaxies, and compared them with nearby regions without much dust. Applying a radiative transfer code that includes absorption and scattering she found values for the visual absorption of 0.3 -- 1.2 mag in the interarm region, and 1.5 -- 4 mag in the arms. These values are consistent with what we find here. The difference here is that we also require that the radial density distribution of the dust is exponential, which means that the models are somewhat better constrained.  The results for NGC~4321 are more complicated because a more extended part of the disk has been observed. In the inner disk, we find similar results to those of the other two galaxies. The measured scale-lengths, both in arms and interarm zone, decrease with increasing wavelength, but the systematics can be well modeled using dust. The outer disk is best described as a system low in dust, and in both inner and outer disk the arms show longer intrinsic scale-lengths than the interarm zone, as might be predicted from the active star formation  evident in the arms. It would also be possible to take into account a fractional dust-lane cover in the inner zone of NGC~4321 in just the same way as that described above for M~51 and NGC~3631, and with a similar conclusion.  The overall results of the observational study presented here are thus:  \\begin{itemize} \\item It is feasible, valuable, and physically instructive, to separate the arm regime from that of the interarm disk when measuring the radial luminosity profile of a disk galaxy, and to extract separately the respective scale-lengths, whose wavelength dependence is to be analyzed. In general the arm scale-lengths are longer than those of the interarm zones.  \\item  In two of the three slightly inclined galaxies observed we find a simple coherent picture of the wavelength dependence of the scale-lengths in terms of dust extinction.  The dust optical depth is measurably greater in the arms than in the interarm disk, a not unexpected result.  Our model fits imply that the dust removes less than 20\\% of the emitted starlight in the $V$ band from those portions of the disks within two intrinsic scale-lengths from the axis (at greater radii this figure would be even lower). In the third galaxy, where we have been studying a greater radial range, the inner disk shows similar dust properties to the disks of the other two galaxies, while in the outer disk the most reasonable interpretation of the results is that there is not enough dust to have a significant effect on the scale-lengths.   \\item The results would remian essentially the same if we assumed that the dust was compressed into lanes  occupying a fraction x of the surface area of the zone under consideration, with axial optical depth greater by a factor 1/x than for the uniform case.  \\item Better fits of the scale-length vs. wavelength observations are obtained for models with stellar to dust scale-height ratio of 1 rather than 3, notably for the arms. This tentative result would be most interesting to confirm with a broader data base. \\end{itemize}  The present study should be seen as a trial in two senses.  Firstly we have measured only three objects: one Sbc galaxy, and two Sc galaxies.  There are observational reasons to believe that there may be less dust in later-type galaxies, and any work aimed at drawing statistical conclusions should have a more representative base, simply in terms of numbers of objects, but also in terms of the spread of morphological types. Secondly we have not explicitly taken into account radial variation in the intrinsic colors of the stellar populations. The fact that the observed radial color variations are small in these objects suggests strongly that such stellar effects are not important here (ruling out the fortuitous cancellation between dust reddening and stellar population blueing as a function of radius), but any tendency of the stars in the outer parts of disks to be bluer would tend to make us under-estimate the dust content.  Here again, with a wider data set, where comparisons between galaxy types can be made, we may obtain a clearer idea of how well such population effects can be quantified. In spite of these limitations, we feel that we have shown, in this article, how the scale-length test may be employed with considerable value as a dust diagnostic in disk galaxies."
+    },
+    "9602/astro-ph9602117_arXiv.txt": {
+        "abstract": "We present a set of structural parameters for the central parts of 57 early-type galaxies observed with the Planetary Camera of the Hubble Space Telescope.  These parameters are based on a new empirical law that successfully characterizes the centers of early type galaxies. This empirical law assumes that the surface brightness profile is a combination of two power laws with different slopes $\\gamma$ and $\\beta$ for the inner and outer regions.  Conventional structural parameters such as core radius and central surface brightness are replaced by break radius $r_b$, where the transition between power-law slopes takes place, and surface brightness $\\mu_b$ at that radius. An additional parameter $\\alpha$ describes the sharpness of the break. The structural parameters are derived using a $\\chi^2$ minimization process applied to the mean surface brightness profiles.  The resulting model profiles generally give very good agreement to the observed profiles out to the radius of $\\sim 10^{\\prime\\prime}$ imaged by the Planetary Camera.  Exceptions include galaxies which depart from pure power-laws at large radius, those with strong nuclear components, and galaxies partly obscured by dust.  The uncertainties in the derived parameters are estimated using Monte-Carlo simulations which test the stability of solutions in the face of photon noise and the effects of the deconvolution process.  The covariance of the structural parameters is examined by computing contours of constant $\\chi^2$ in multi-dimensional parameter space. ",
+        "introduction": "Recent Hubble Space Telescope (HST) observations of early-type galaxies reveal that the central parts have surface-brightness distributions that are different from the extrapolation of traditional fitting formulae derived from ground-based observations (Crane {\\it et al.} 1993; Kormendy {\\it et al.} 1995; van den Bosch {\\it et al.} 1994; Forbes, Franx, \\& Illingworth 1995, Lauer {\\it et al.} 1995, hereafter Paper I).  The surface brightness profiles generally consist of two distinct regions; a steep power-law regime ($I(r) \\propto r^{-\\beta}$) at large radius, and a shallower power law ($I(r)\\propto r^{-\\gamma}$) at small radius.  Galaxies with a quite shallow inner power-law slope ($\\gamma <$ 0.3) are classified as ``core'' galaxies, while those with $\\gamma > 0.5$ are labeled power-law galaxies (Paper I). Note that $\\gamma$ is significantly different from zero in most objects studied with HST thus far, with only rare exceptions (van den Bosch {\\it et al.} 1994; Paper I; Ajhar {\\it et al.} 1995).  Conventional parameters characterizing the luminosity structure of the central regions are the central surface brightness and the core radius, defined as the radius where the surface brightness is half the central value.  These parameters have been very useful in understanding elliptical galaxies and bulges of later type galaxies, especially through the fundamental plane correlations (Kormendy 1982, 1984, 1985, 1987a, b; Lauer 1985).  However, as we do not find constant surface brightness cores in most HST observations, we set about finding an alternative parameterization (presented below) to better describe the observed surface brightness distributions.  For a more detailed history on this subject, readers are referred to Paper I.  An alternative approach to the one taken here is that of Gebhardt {\\it et al.} 1996), who carry out a non-parametric analysis of essentially the same sample considered here. The Gebhardt {\\it et al.} analysis is both more thorough and more accurate, but it largely verifies the simple fits we describe below. The major advantage of the approach taken in the present paper is that it reduces each galaxy to a set of 5 numbers. On the other hand, the non-parametric approach does not force galaxies into possibly ill-fitting boxes.    A one-parameter family of models for spherical stellar systems which yields double power-law surface brightness profiles was introduced independently by Dehnen (1993; $\\gamma$ models) and by Tremaine \\etal (1994; $\\eta$ models). In the notation of Tremaine \\etal, the density  \\begin{equation} \\rho_\\eta(r) \\equiv {\\eta \\over 4 \\pi}{1 \\over r^{3-\\eta}(1+r)^{1+\\eta}}, \\hspace{1cm}0 < \\eta \\le 3. \\end{equation}  However, the majority of the actual galaxy profiles is not well-fit by $\\eta$-models.  Although these models allow for a range of inner power-law slopes, the variations apparent in the sharpness of the transition from inner to outer power-laws cannot be reproduced, and the fixed asymptotic outer slope of $r^{-3}$ does not fit all galaxies.  To parametrize the HST-observed surface brightness profiles, we introduced in Paper I a more general empirical double power-law (the ``Nuker'' law):  \\begin{equation} I(r) = I_b ~2^{\\frac{\\beta-\\gamma}{\\alpha}} \\left( \\frac{r}{r_b}\\right) ^{-\\gamma} \\left[ 1 +  \\left( \\frac{r}{r_b} \\right) ^ \\alpha \\right] ^ {- \\frac{\\beta-\\gamma}{\\alpha}} \\end{equation}  \\noindent{}The break radius, $r_b$, is the radius at which the steep outer profile, $I(r)\\propto r^{-\\beta}$, ``breaks'' to become the inner shallow profile, $I(r) \\propto r^{-\\gamma}$.  More precisely, $r_b$ is the radius at which the absolute logarithmic curvature in the surface brightness profile, $|d^2\\log I(r) /d(\\log r)^2|$, is a maximum. $I_b$ is the vertical scaling parameter and is simply the surface brightness at $r_b$. As shown in Paper I, the parameter $\\alpha$ is necessary to account for the variations in the sharpness of the break.   The Nuker law contains many simpler fitting formulae as special cases;  \\begin{itemize}  \\item The Hubble-Reynolds law (Reynolds 1913; Hubble 1930), $$I(r)={I_0a^2\\over (r+a)^2},$$ corresponds to $\\alpha=1$, $\\beta=2$, $\\gamma=0$.  \\item The modified Hubble or analytic King law (Binney \\& Tremaine 1987), $$I(r)={I_0r_c^2\\over r_c^2+r^2},$$ corresponds to $\\alpha=2$, $\\beta=2$, $\\gamma=0$.  \\item de Vaucouleurs' $r^{1/4}$ law (de Vaucouleurs 1948), $$I(r)=I_0\\exp(-kr^{1/m}), \\qquad m=4,$$ and its generalization to arbitrary $m$ (Sersic 1968), corresponds to the parameters $\\alpha={1\\over m}$, $\\gamma=0$, in the limit $r_b\\to\\infty$, $\\beta={1\\over m}kr_b^{1/m}\\to\\infty$.  \\item  Ferrarese et al. (1994) fit HST photometry of early-type galaxies to the law, $$I(r)=2^{1/2}I_c\\left(r\\over r_c\\right)^{-\\beta_1}\\left[1+\\left(r\\over r_c\\right)^{2(\\beta_2-\\beta_1)}\\right]^{-1/2},$$ which corresponds to $\\alpha=2(\\beta_2-\\beta_1)$, $\\beta=\\beta_2$, $\\gamma=\\beta_1$.  \\end{itemize}  Galaxies in which the three-dimensional luminosity density varies smoothly near the center are said to have ``analytic cores''.  An analytic core requires that the surface brightness near the center has the form $I(r)=I_0+ I_1 r^2 + \\hbox{O} (r^4)$.  Thus only galaxies with $\\alpha=2$, $\\gamma=0$ have analytic cores.  The Nuker law also provides excellent fits to the family of $\\eta$-models (Figure 1).  The behavior of the five Nuker parameters for each of the $\\eta$-models is illustrated in Figure 2 in order to show the parameter space occupied by $\\eta$-models.  As will be seen below, real galaxies cover a much wider range in these parameters.  The Nuker law is convenient not only for describing the observed surface brightness distribution but also for computing quantities such as the luminosity density and the local and line-of-sight velocity dispersion distributions.  A FORTRAN code that carries out these tasks is available on request from S. Tremaine.  In this paper, we fit the Nuker law and derive the associated structural parameters for 57 early-type galaxies observed with the HST. With the exception of NGC 4881, all objects were observed prior to the HST refurbishment with the Planetary Camera ($0.\\prime\\prime043$ per pixel) using filter F555W (corresponding roughly to Johnson $V$). The sample includes early-type galaxies observed as part of our Guest Observer (GO) programs (\\#2600, \\#3912, PI Faber), as part of several WFPC Instrument Definition Team programs (\\#1105, \\#3229, \\#3286, \\#3639, \\& \\#5233, PI Westphal), and includes F555W PC images of early-type galaxies publically available in the HST archive by June 23, 1994. The latter were taken as part of GO programs \\#1038 (PI Ford) and \\#2607 (PI Jaffe).  While small amounts of dust could generally be avoided in determining surface brightness profiles, galaxies with severe nuclear dust obscuration were rejected. Two or more exposures were used for each galaxy to permit cosmic ray rejection, and the peak counts per pixel in the coadded frames were typically $\\sim$2000 before deconvolution. Further details of data acquisition and reduction of the GO \\#2600/3912 sample are given in Paper I. These reduction procedures have been uniformly applied to the additional galaxies in the present sample. The only exception is NGC 4881, which was observed after the refurbishment of HST and therefore did not require extensive deconvolution.  The parameters determined in this paper are used in the analysis and interpretation of core properties presented in Faber {\\it et al.}  (1996, Paper III). ",
+        "conclusions": "In this paper, we derive a set of structural parameters for 57 early-type galaxies observed with HST.  These parameters are based on the empirical ``Nuker'' law, which successfully describes the central surface-brightness distributions within the radius range from $\\sim$0.$^{\\prime\\prime}$1 to $\\sim$10$^{\\prime\\prime}$ for most galaxies in the sample. These parameters have been used in Paper I and Faber \\etal (1996) to identify distinct classes of objects, characterized by significant differences in one or more of the Nuker parameters.  While these parameters can in principle also be used to evaluate luminosity densities and (with the availability of suitable spectral information) the run of velocity dispersion with radius, we believe this would be better left to non-parametric methods (Gebhardt \\etal 1996), which are more faithful to small but significant undulations in the observed profiles."
+    },
+    "9602/astro-ph9602007_arXiv.txt": {
+        "abstract": "We investigate whether a simple cosmology can fit GRB results in both time dilation and Log N - Log P simultaneously.  Simplifying assumptions include: all GRBs are spectrally identical to BATSE trigger 143, $\\Omega=1$ universe, and no luminosity and number density evolution.  Observational data used includes: the BATSE 3B peak brightness distribution (64-ms time scale), the Pioneer Venus Orbiter (PVO) brightness distribution, and the Norris et al. time dilation results for peak aligned profiles presented at this meeting. We find acceptable cosmological fits to the brightness distributions when placing BATSE trigger 143 at a redshift of 0.15 $\\pm$ 0.10. This translates into a $(1 + z_{dim}) / (1 + z_{bright})$ factor of about 1.50 $\\pm$ 0.50 between selected brightness extremes of the Norris et al. sample. Norris et al. estimate, however, that $(1 + z_{dim}) / (1 + z_{bright})$ $\\approx$ 2.0 $\\pm$ 0.5 when considering duration tests. The difference is marginal and could be accounted for by evolution. We therefore find that evolution of GRBs is preferred but not demanded. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602069_arXiv.txt": {
+        "abstract": "As part of a large study to map the distribution of star formation across galactic disks, we have obtained deep H$\\alpha$ images of the nearby Sculptor Group spirals NGC~247 and NGC~7793.  These images are of sufficiently high quality that they allow identification and analysis of diffuse H$\\alpha$ emission at surface brightness levels ranging from those of extremely low density HII regions to those of the local Galactic disk diffuse emission. This paper presents a study of the large scale distribution and global energetics of diffuse ionized gas (DIG) in these galaxies and investigates the association between DIG and discrete HII regions.   Our results support the hypothesis that the DIG is photoionized by Lyc photons which leak out of traditional HII regions, and suggest that the local HI column density plays a role in regulating the amount of leakage which can occur.   This interpretation has profound implications for the derivation of star-formation rates based on H$\\alpha$ emission--line fluxes since HII region counts alone will lead to significant underestimates of the true rate.  The contribution of the diffuse H$\\alpha$ component to the total H$\\alpha$ emission, ie. the diffuse fraction, in these galaxies is found to be similar to values found in other disk galaxies with differing Hubble types and star formation rates.  The constancy of the diffuse fraction is rather unexpected and implies that the overall fraction of photons which can leak out of HII regions and ionize the ISM over large scales is relatively invariant from one galaxy to another. ",
+        "introduction": "Understanding the relationship between the star formation process and the interstellar medium is a key step towards unravelling the complex history of galaxy evolution.  A longstanding problem has been the nature and importance of the feedback process by which massive stars deposit energy into the interstellar medium via photoionization, stellar winds and supernovae.  This feedback mechanism affects the physical and dynamical state of the interstellar medium and hence influences the rate and distribution of subsequent star formation.   OB stars clearly provide the largest source of Lyman continuum (Lyc) photons in typical spiral galaxies (Abbott 1982); a major issue concerns whether the bulk of these photons are deposited over localized regions, such as the Str\\\"omgren spheres which define traditional  HII regions,  or whether a significant fraction can escape from the regions of recent star formation where they were created and can ionize the interstellar medium over much larger scales.    If large--scale ionization due to Lyc photon leakage from HII regions were a common feature of spiral galaxies, then the resulting diffuse ionized gas (DIG) component would provide an important tracer of the feedback process due to star formation, as well the structure of the interstellar medium.  It would be of crucial importance to include this component in calculations of radial and global star formation rates derived on the basis of H$\\alpha$ emission line fluxes.  Furthermore, the existence of such a DIG component would have direct bearing on the energy balance of the interstellar medium, regardless of the source of photons that maintain it in an ionized state. \\par Indeed, the existence of widespread diffuse ionized gas lying outside the boundaries of traditional HII regions has been known for over 20 years (Reynolds {\\it et al.} 1971;  Monnet 1971), but the source of this ionization remains the subject of much debate. Pulsar dispersion measures and optical line emission have been used to trace the Galactic DIG component (the `Reynolds Layer') and have established it to be an extremely important component of the interstellar medium, occupying more than 20\\% of the interstellar volume, contributing at least $\\onethird$ of the total HI column at the solar circle and constituting $\\sim$~90\\% of the ionized hydrogen mass in the Galaxy (see review by Reynolds 1991 and references therein).  The location of the Sun close to the plane of the disk of the Milky Way means that the vertical structure of the ionized gas is more easily studied than is its radial distribution; it has been established that the Galactic DIG is distributed in a thick disk with a scale--height locally of $\\sim$~1~kpc.  Much of the DIG (sometimes referred to as the Warm Ionized Medium; WIM) is observed to be in the form of discrete structures, such as loops, filaments and shells (which are not obviously associated with discrete HII regions) while the remainder is in the form of an apparently unstructured, diffuse background (Reynolds 1993).    The power required to maintain this low-density ($n_e \\sim$~0.1~cm$^{-3}$), warm ($<$~10000~K) ionized component is large -- 1.0 $\\times$~10$^{-4}$~erg~s$^{-1}$~cm$^{-2}$ of disk  -- and can be met easily only by ionizing photons from OB stars, although the kinetic energy from supernovae could possibly suffice (Reynolds 1984; Kulkarni and Heiles 1988). \\par The identification of OB stars as the source of the ionization of the Galactic DIG would require that Lyman continuum photons be able to travel several hundred parsecs both perpendicular to the plane, to account for the derived vertical scale--height of a kiloparsec, and parallel to the plane, to account for the DIG seen at large radial distances from OB stars (see Reynolds 1990).  These requirements are even greater in irregular galaxies, where ionized gas has been observed at distances of $\\gtrsim$~1~kpc from bright OB associations (Hunter and Gallager 1990; Hunter and Gallagher 1992).  Recent calculations suggest that it is indeed possible for a large fraction of ionizing photons to penetrate sufficiently large distances from their point of origin, depending on cloud properties and distributions, as well as on their radial gas distributions (Dove and Shull 1994; Miller and Cox 1993).  The observed optical emission line ratios of the Galactic DIG can be explained by a rather dilute radiation field and hence argue further for photoionization by distant OB stars (D\\\"omgorgen and Mathis 1994).  Alternative sources for the ionizing radiation that maintains the DIG which have been proposed include shocks, decaying massive neutrinos (Sciama 1990), turbulent mixing layers (Slavin, Shull and Begelman 1993) and Galactic microflares (Raymond 1992).   In addition, significant contributions to the ionizing flux may come from evolved stellar objects such as planetary nebula nuclei, hot white dwarfs and blue horizontal branch stars.  \\par The large scale radial distribution of the DIG across galactic disks can provide stringent constraints on the source of its ionization.  For example, if the amount of Lyc photons produced in star--forming regions were the only factor responsible, then the DIG distribution and intensity should be expected to correlate, over both small and large scales,  with that of discrete HII regions.  Furthermore,  DIG properties should also be expected to vary systematically with the morphological type of the host galaxy, reflecting the underlying variation in star formation and structure within the ISM (Kennicutt, Edgar and Hodge 1989).  Study of such correlations requires observations of external galaxies at moderate inclinations. To date, few quantitative studies of the DIG in external galaxies have been carried out, and the emphasis has largely been on edge--on systems (eg.  Rand, Kulkarni and Hester 1990; Dettmar 1992; Veilleux {\\it et al.} 1995).  Walterbos and Braun (1994; hereafter WB94) studied the DIG component in selected areas of the nearby spiral M31, but their small field of view restricted them to local analysis and they could not assess the overall distribution and energetics of the ionized gas they detected. \\par The first quantitative study of the large--scale distribution and global energetics of diffuse ionized gas in moderately face--on external disk galaxies is presented in this paper.  Our data consist of large field of view, deep H$\\alpha$ images of Sculptor Group late--type spirals NGC~247 and NGC~7793 which were obtained as part of a large study to map the distribution of star formation across galactic disks. NGC~7793 has been previously noted for having a significant DIG component (Monnet 1971) but this earlier work was limited to photographic data, with surface brightness and spatial resolution limits significantly lower than the CCD data presented here.  Basic properties of these galaxies are listed in Table 1.   In the present paper, we derive the radial and azimuthal distributions of the DIG and investigate the association between DIG and discrete HII regions.  We also study the global and radial variation of the diffuse fraction and the Lyc power requirements, which allow tight constraints to be placed on the origin and nature of DIG component in spiral galaxies.  Our results support the hypothesis that the DIG is photoionized by Lyc photons which leak out of discrete HII regions, and suggest that the local HI column density plays a role in regulating the amount of leakage which can occur.   Given this hypothesis, we discuss the great importance of taking account of the DIG component when deriving star formation rates based on H$\\alpha$ emission line fluxes. ",
+        "conclusions": "\\subsection{Small and Large Scale Distribution of the DIG}  The azimuthal and radial intensity distributions of H$\\alpha$ emission in these galaxies indicate that the DIG is highly correlated with bright HII regions over a variety of scales.  On scales of a few hundred parsecs, the bright DIG is highly concentrated around luminous HII regions, with the mean surface brightness decreasing as the distance to the nearest HII region increases.  This close small-scale association between DIG and HII regions was also observed by WB94 in selected fields in M31 and lends strong support for a model in which the Lyc photons which ionize the DIG leak out of individual HII regions.  In this picture, the bulk of the H$\\alpha$ recombinations would be expected to take place in the immediate vicinity of the leaking HII region, where the Lyc photon flux is higher.   At faint levels, the DIG is ubiquitous, though generally follows the overall pattern of star formation e.g. being enhanced in spiral arm structures if present, as in NGC~247. \\par Over larger scales, the DIG has spatial and intensity distributions which closely follow those of the discrete HII regions in each galaxy.  A key observation is that the DIG is detected over the entire extent of the bright star--forming disks in both galaxies.   The faint, outer disk DIG we detect has surface brightnesses comparable to the large--scale flat--fielding errors, hence we are unable to address the important issue of whether the region of DIG emission actually extends beyond the edge of the HII region distribution.  Deeper CCD data we have just obtained for these galaxies will allow a more detailed future study of the outer disk DIG emission. \\par The large scale radial distributions reveal that the DIG intensity falls off with the mean surface brightness of HII regions and, hence, the mean level of star formation.   Hester and Kulkarni (1990) and Veilleux {\\it et al.} (1995) have also noted a large scale radial variation of DIG intensity, in studies of M33 and NGC~3079 respectively, but they did not make a comparison with the distribution of discrete HII regions in these galaxies.  The similarity between the DIG and the discrete HII region surface brightness profiles observed further supports a leakage origin for the ionizing photons which create the DIG.  Moreover, the correlation between the mean surface brightness of discrete HII regions and the DIG suggests that large--scale radial photon transport, on scales comparable to the widths of our annular bins, 250~pc, does not involve a significant number of ionizing photons. \\par The small and large scale distributions of the DIG are easily understood in terms of a model where Lyc photon leakage from star--forming regions provides the dominant source of the ionizing photons.  The available evidence indicates that the star formation rate per unit area (traced by the surface brightness of discrete HII regions) is correlated with the DIG surface brightness, in the sense that a higher mean star formation rate per unit area leads to a higher number of H$\\alpha$ recombinations occurring outside of discrete HII regions.  The leakage model is not the only model which can explain our observations of the large scale distribution of the DIG,  however, and any method of producing ionizing photons which is tied to Population I objects is equally plausible.  A model in which  a significant fraction of the OB star population resides outside the boundaries of discrete HII regions may also suffice (eg. Patel and Wilson 1995), but it remains unclear whether ionization by {\\it in situ} OB stars can explain other DIG properties, such as the observed line ratios. Decaying neutrinos might be expected to produce an ionized gas component with a radial profile similar to the density profile of the dark matter halo, ie.  volume density proportional to the inverse square of the radius (Sciama \\&  Salucci 1990).  While such a distribution is not necessarily inconsistent with the profiles  derived here, it is difficult to understand  why the DIG  emission would be enhanced around individual HII regions, unless the gas becomes optically thin in these regions as well.  \\subsection{The Diffuse Fraction}  The global diffuse fractions of the galaxies presented here, together with those found in other galaxies, show a striking similarity to each other.  If the amount of Lyc photons produced in HII regions were the only governing factor in determining the properties of the DIG, then galaxies with low global star formation rates, such as the SMC, would be expected have different DIG properties than more actively star-- forming galaxies, such as NGC~7793 presented here.   The constancy  of the diffuse fraction among galaxies of differing Hubble types and star formation rates suggests a conspiracy between the level of star formation and some additional factor, with the result that the fraction of photons which escape from discrete HII regions is more or less constant from galaxy to galaxy. \\par A similar conclusion can also be reached from inspection of the radial variation of the diffuse fractions in NGC~247 and NGC~7793.  Since the surface brightness of H$\\alpha$ emission from HII regions, hence Lyc production, decreases with increasing galactocentric radius, one might expect less diffuse emission to be present in the outermost parts. Instead, the diffuse fraction in NGC~247 is constant across the disk, while the diffuse fraction in NGC~7793 increases with galactocentric radius.  WB94 also found a relatively constant diffuse fraction in each of the fields they surveyed in M31.  This behavior suggests that the additional factor which regulates Lyc photon leakage in galaxies is one which has some radial dependence within disks, and one which varies from galaxy to galaxy. \\par This additional factor may be the HI column density; the (qualitative) agreement between the observed radial profiles of the diffuse fraction and the predictions of the Dove and Shull models lend support for this idea.  Dust may also play an important role in regulating the fraction of Lyc photons which can leak over large distances.  NGC~7793 is known to have a moderately steep metallicity gradient (Kennicutt, Zaritsky and Huchra 1994), implying more dust per unit gas mass (and per unit star formation rate) in the central regions than in the outer parts. This could result in enhanced absorption of Lyc photons in the central regions, when compared to the outer disk, and provide at least a partial explanation for the trend observed in the diffuse fraction. Furthermore, a higher dust fraction in the central regions would also mean that less of the DIG produced here would actually be observed.  A detailed study of the metallicity and extinction in HII regions across the disks of NGC~7793 and NGC~247 is needed to address this issue further.  \\subsection{Lyc Power Requirements}  The minimum Lyc power requirements we have derived for the DIG in NGC~247 and NGC~7793 are huge and pose the greatest challenge to all models for the origin of the DIG.   The recent evolutionary synthesis models of Leitherer and Heckman (1995) can be used to place constraints on several possible sources  of the ionizing photons.  These models may be used to estimate the global star formation rates based on the observed H$\\alpha$ luminosities.  We make the assumption that the DIG is powered by star formation, and include its contribution in the total H$\\alpha$ luminosities of the galaxies.   Independent estimates of the SFRs, based on the far--infrared luminosities, may then be used as a check on how realistic an assumption this is. \\par The [NII] and extinction--corrected total H$\\alpha$ (DIG plus HII regions) luminosities were measured to  be 4.1$\\times$10$^{40}$~erg~s$^{-1}$ for NGC~7793 and 1.5$\\times$10$^{40}$~erg~s$^{-1}$ for NGC~247.  Formula (9) of Leitherer and Heckman (1995) then translates these values into Lyc photon production rates of $3.03 \\times 10^{52}$~s$^{-1}$ and $1.10 \\times 10^{52}$~s$^{-1}$.  These photon production rates correspond to global star formation rates (M$>$ 1M$_{\\sun}$) of 0.096 M$_{\\sun}$ yr$^{-1}$ for NGC~7793 and 0.035 M$_{\\sun}$ yr$^{-1}$ for NGC~247, assuming a Salpeter IMF with M$_{upper}$ = 100 M$_{\\sun}$ and solar metallicity (Leitherer and Heckman 1995).    It should be noted that metallicity only influences the derived quantities to a small extent. Adopting  M$_{lower}$=0.1M$_{\\sun}$ would increase these star formation rates by a factor of 2.5.  These values are all tabulated in Table 2. \\par The integrated infrared luminosities of the galaxies provide an independent estimate of the global star formation rates.   Using the measured IRAS fluxes (Rice {\\it et al.} 1988) and assuming the distances in Table 1, the far infrared luminosities were calculated to be 1.9 $\\times$ 10$^{42}$ erg~s$^{-1}$ for NGC~7793 and 4.5 $\\times$ 10$^{41}$ erg~s$^{-1}$ for NGC~247. From the derivations of Hunter {\\it et al.} (1989) (their equation 14), these luminosities can be converted into star formation rates of 0.24~E$^{-1}$ M$_{\\sun}$ yr$^{-1}$ for NGC~7793 and 0.06~E$^{-1}$ M$_{\\sun}$ yr$^{-1}$ for NGC~247, where E is a constant of order unity measuring  the coupling efficiency between the total power radiated by dust grains and that included in the IRAS--based measurement of L$_{IR}$.    While these calculations are based on different stellar lifetimes than those used in the Leitherer and Heckman models and have several uncertainties, such as the value of the lowest main--sequence stellar mass that significantly contributes to L$_{IR}$, they agree with the star formation rates calculated on the basis of the total H$\\alpha$ luminosities to within a factor of 2.  Had we been incorrect to include the H$\\alpha$ luminosity of the DIG in our estimate of the star formation rate, we could have expected much larger discrepancies between the rates calculated based on these different methods.  Thus, it seems quite reasonable that the DIG is indeed powered by Lyc photons from OB stars. \\par We can use the set of Leitherer and Heckman (1995) models which are based on a constant SFR to estimate the mechanical energy that is injected into the ISM  from Type~II supernovae and stellar winds. These models give a mechanical luminosity of $7.6 \\times 10^{40}$~erg~s$^{-1}$ for NGC~7793 and $2.7 \\times 10^{40}$~erg~s$^{-1}$ for NGC~247, based on the assumptions mentioned above and using the star formation rates calculated from the H$\\alpha$ luminosities (note that these mechanical outputs would be increased by a factor of two if we assumed a Miller--Scalo IMF).  These numbers can be compared with the minimum Lyc power required to ionize the DIG, derived in Section 4.4; they fall short by a factor of 2.5 in NGC~7793 and by a factor of 4.4 in NGC~247.   If we compare them to the more realistic power requirements, calculated on the basis of the {\\it corrected} DIG H$\\alpha$ luminosities,  then they fall short by factors of 3.6 in NGC~7793 and 4.5 in NGC~247.  Type~I supernovae will provide an additional  contribution to the mechanical luminosity, but given the estimates of the relative rates of Type~I and Type~II supernovae in typical disk galaxies (van den Bergh and Tammann 1992), it is unlikely that this will change the available energy by more than a factor of two.  We can thus conclude that  supernovae and stellar winds are unable to produce a substantial fraction of the DIG ionization in these galaxies, even assuming that they can transfer their energy to the ISM with 100\\% efficiency.  This result contrasts with that found for the solar neighborhood  where the energy from supernovae can almost account for the observed level of DIG ionization (Reynolds 1984). \\par Despite the fact that OB stars are the only known source of ionizing photons which can easily meet the overall power requirements for the DIG,  several unresolved issues remain.  As pointed out by Reynolds (1995), observations of HI clouds suggest that there are 3--4 clouds with N$_{HI}$ $\\sim$~3$\\times$ 10$^{19}$ cm$^{-2}$ along every 300~pc line of sight near the midplane of the Galactic disk, with the number possibly increasing dramatically towards lower column densities (Kulkarni and Heiles 1987; Dickey and Garwood 1989).  Such clouds are expected to present a significant opacity to Lyc photons and should severely limit the distance which photons can travel radially within disks similar to that of our Galaxy.  In addition, the recent non-detection of the HeI $\\lambda$5876 recombination line in the direction of relatively high Galactic DIG surface brightness implies that the local DIG ionization is due to a significantly softer spectrum than that expected from the bulk of the O~star population in the solar neighborhood (Reynolds and Tufte 1995).  This result implies that either stars of spectral type O8 and later provide most of the ionizing photons for the local DIG, or that some mechanism is in place which effectively softens the radiation from early type O stars. \\par Late-type O and early B stars account for 24\\% of the total Lyc photons produced in the solar neighborhood (Vacca, Garmany and Shull 1995), although this number may be an underestimate (e.g. Cassinelli {\\it et al.} 1995). Assuming this value, and using the Leitherer and Heckman models to predict the total output in Lyc photons for the inferred star formation rates, we can estimate the amount of ionizing luminosity produced by such stars in NGC~7793 and NGC~247.  The Leitherer and Heckman models predict total ionizing outputs of $9.6~\\times~10^{41}$~erg~s$^{-1}$ in NGC~7793 and $3.4~\\times~10^{41}$~erg~s$^{-1}$ in NGC~247.   The  power from soft Lyc photons produced directly by late-type O and early B stars is then estimated to be $2.3~\\times~10^{41}$~erg~s$^{-1}$ in NGC~7793 and $8.2~\\times~10^{40}$~erg~s$^{-1}$ in NGC~247.  Comparing these numbers with Lyc power requirements in each galaxy, we find that the minimum power requirements can be met in NGC~7793 but not in NGC~247.  Both numbers fall short of satisfying the more realistic, corrected Lyc power requirements.   Furthermore, these estimates of the soft Lyc power available should be regarded as upper limits since a considerable fraction of the power is consumed in the immediate vicinity of OB stars and hence is not available to ionize ISM over large scales.  Our simple calculations indicate that such stars are not likely to produce enough direct soft Lyc radiation to match the power requirements and imply that another mechanism may indeed be needed to produce the soft radiation field which a low HeI$\\lambda$5876/H$\\alpha$ ratio requires. The absorption and re--emission of Lyc photons as they travel through the ISM could have such an effect; possible sites for this process include `chimney' walls (Norman 1991) and the extended envelopes of HII regions.  Detailed maps of the ionization structure in the HII region -- DIG transition zone would be required to  investigate these issues more thoroughly and to place tighter constraints on the actual soft Lyc power that is available. \\par Of course, as mentioned above, there are many known sources of Lyc photons, and they all must contribute at some level to produce the widespread ionization of the DIG which is observed.    We have presented a study of the large scale distribution and global energetics of widespread diffuse ionized gas lying outside the boundaries of discrete HII regions, identified from deep H$\\alpha$ images of the nearby Sculptor spirals NGC~247 and NGC~7793.   We have separated the total H$\\alpha$ emission into that from discrete HII regions and that from diffuse emission by means of an isophotal cut in surface brightness, corresponding to an emission measure of 80 pc cm$^{-6}$.  Most of the H$\\alpha$ emission lying below this value is clearly diffuse and/or filamentary. \\par Radial and azimuthal intensity distributions of the DIG reveal that it is highly correlated with bright HII regions over both small and large scales.  On the scales of a few hundred parsecs, the bright DIG is localised around individual luminous HII regions in frothy, filamentary halos.  At faint surface brightnesses, the DIG is ubiquitous, though maintains the overall pattern and extent of the bright star--forming disk.    Over larger scales, the DIG has spatial and intensity distributions which closely follow those of the discrete HII regions, and, hence, of the mean level of star formation. \\par The observed DIG H$\\alpha$ luminosities are considerable and the DIG contributes 50\\% of the measured total H$\\alpha$ luminosity in NGC~247 and 29\\% in NGC~7793.   These values are lower limits derived by assuming that the intensity of the DIG overlying discrete HII regions is zero.  More realistic values of the diffuse fractions, calculated assuming that DIG fills the entire area of the elliptical annuli used, are 53\\% in NGC~247 and 41\\% in NGC~7793.  These values are remarkably similar to the global diffuse fractions found in other actively star--forming galaxies, which span a range of morphological types and star formation rates.   This unexpected result suggests some global regulation of the fraction of photons which can leak from HII regions and are available to ionize the ISM over large scales.  The radial variation of the diffuse fraction shows an increase across the disk of NGC~7793 and is relatively constant across the disk of NGC~247, despite the declining star formation rate per unit area (ie. Lyc production rate) with increasing radius.  This behaviour suggests that the additional factor which regulates photon leakage from HII regions is one which has some radial dependence.  The qualitative agreement between our observations and the predictions of the Dove and Shull (1995) models for photoionization of the DIG suggest that this factor may be the HI column density.  We note the role that dust might play in regulating photon leakage in galaxies.  More detailed metallicity gradients are required for these galaxies in order to further investigate this issue. \\par The integrated miminum Lyc powers required to sustain the DIG in these galaxies are enormous and place tight constraints on the source of the ionizing photons.  Only massive star formation can easily satisfy the power requirements, with the mechanical luminosity from supernovae and stellar winds falling short by more than a factor of two.  This result contrasts with that found for the local Galactic DIG where the energy from supernovae can almost account for the observed level of ionization.  In view of the new constraints imposed on the origin of the DIG photons from the non--detection of the HeI$\\lambda$5876 recombination line in the Galaxy, we estimated the Lyc power supplied by stars of spectral type O8 and later.   The available power falls short of realistic estimates of the required Lyc power, implying that direct radiation from such stars alone cannot acount for the entire ionization of the DIG.  This suggests that either some mechanism is in place which softens the radiation field of early O stars, which produce the bulk of the Lyc photons, or that the photons which ionize the DIG come from a variety of sources,  all contributing at some level to produce the widespread ionization which is observed. \\par Our results strongly support an interpretation of the DIG as being due to Lyc photon leakage from discrete HII regions.   In such a scenario, the  DIG is a direct product of recent massive star formation and its contribution to the H$\\alpha$ emission line flux of the galaxy should be included in order to derive accurate star formation rates.  We have illustrated the importance of taking the DIG into account when studying the correlation between star formation rate per unit area  and HI gas surface density.  If HII regions alone are used to trace the star formation rate per unit area, then these plots show a sharp cut--off in star formation activity  at low HI column densities, commonly interpreted as being the result of a threshold for massive star formation (Kennicutt 1989).   This evidence for an abrupt cut--off is not seen in our data when the DIG component is included in determining the star formation rate per unit area."
+    },
+    "9602/astro-ph9602155_arXiv.txt": {
+        "abstract": "We have made polarimetric observations of three Type Ia supernovae (SN Ia) and two type II supernovae (SN II). No significant polarization was detected for any of the SN Ia down to the level of 0.2\\%, while polarization of order $1.0\\%$ was detected for the two SN II 1994Y and 1995H. A catalog of all the SNe with polarization data is compiled that shows a distinct trend that all the 5 SN II with sufficient polarimetric data show polarizations at about 1\\%, while none of the 9 SN Ia in the sample show intrinsic polarization. This systematic difference in polarization of supernovae, if confirmed, raises many interesting questions concerning the mechanisms leading to supernova explosions. Our observations enhance the use of SN Ia as tools for determining the distance scale through various techniques, but suggest that one must be very cautious in utilizing Type II for distance determinations. However, we caution that the link between the asphericity of a supernova and the measured ``intrinsic'' polarization is complicated by reflected light from the circumstellar material and the intervening interstellar material, the so-called light echo. This effect may contribute more substantially to SN II than to SN Ia. The tight limits on polarization of SN Ia may constrain progenitor models with extensive scattering nebulae such as symbiotic stars and other systems of extensive mass loss. ",
+        "introduction": "Polarization measurements of supernovae are very rare. So far, detailed observations with both broad band and spectral polarimetry taken at different epochs after explosion are only available for SN 1987A and SN 1993J, generally the two most well-studied supernovae. The observed polarization for SN 1987A was around 0.5\\% and evolved with time. Published data are available only in the first year after explosion, and unfortunately do not extend much beyond the photospheric phase (Jeffery 1991b and references therein). Spectropolarimetric data of SN 1987A were also obtained on days 523 and 643 after explosion, but the data are still not published (Stathakis 1995). Observations of SN 1993J yield polarizations of about 1.5\\% (Trammell, Hines, \\& Wheeler 1993; Doroshenko, Efimov, \\& Shakhovskoi 1995). The published data are even more limited, but cover epochs both before and after the supernova passed its second optical maximum. The observed polarization for these two supernovae is attributed to aspherical photospheres in their early phases of evolution (H\\\"oflich 1987; Jeffery 1991a; H\\\"oflich et al. 1995), or to asymmetric distributions of radioactive materials which may be the result of instabilities produced during the shock break out (Chugai, 1992).  Polarized light is the best, but not the only, signature of deviations from spherical symmetry in these two supernovae. Strong deviations from spherical symmetry are also recorded in spectroscopic observations of SN 1987A and SN 1993J. In SN 1987A, various observations show that the spectral lines of most chemical species have fine structures the strengths of which evolve with time. Most noticeable is the the fine structure in \\Ha (Hanuschik \\& Dachs 1987). These fine structures are modeled in terms of high velocity $^{56}$Ni clumps ejected during the explosion (Chugai 1988). Evidence of the existence of high velocity radioactive clumps is also obvious in the spectra taken at about 5 years after the explosion of SN 1987A (Wang et al 1996). The asymmetric \\Ha profile of SN 1987A five years after explosion implies that $^{44}$Ti has a similar spatial distribution as $^{56}$Ni. Extensive evidence also exists for macroscopic chemical mixing in the ejecta, supported by both observations and hydrodynamical calculations. For SN 1993J, spectral evidence for deviations from spherical symmetry is perhaps most clearly observed in the late stage nebular phase spectra. The first remarkable feature is the blue shift of the \\OID lines (Clocchiatti et al 1994; Wang \\& Hu 1994). Later evidence is given by the peculiar evolution of the \\Ha line (Filippenko et al, 1994). The early phase \\OID lines were interpreted by Wang \\& Hu (1994) as due to a highly deformed photosphere. The blue shifts, however, were also observed in the late nebular stage when the ejecta became optically thin, and this is indicative of a global geometrical distortion. A nickel sphere that is off-center with respect to the the bulk of the oxygen shell provides a reasonable fit of the late time \\OID profile (Chugai \\& Wang 1995)  A major difficulty of polarimetric studies of supernovae is to subtract the polarization produced by interstellar material. For SN 1993J, Trammell et al. (1993) accomplished this by assuming \\Ha to be intrinsically unpolarized, and that the interstellar polarization follows Serkowski's law (Serkowski, 1970; Wilking, Lebofsky, Rieke, 1982). The assumption that the spectral lines are not polarized is based on the assumption that the continuum photons are polarized by Thomson scattering above an asymmetric supernova photosphere. The spectral lines are produced by resonance scattering which tends to destroy polarized photons. This picture is quantitatively justified in H\\\"oflich et al. (1995), and has proven to be useful for SN 1993J. This technique can only be applied when spectropolarimetry is available.  The observed polarization in SN 1987A and SN 1993J implies that at least some supernova explosions are aspherical. Considering the fact that supernovae are now becoming an increasingly important tool for the determination of the extra-galactic distance scale, it is important to investigate whether polarization is a common phenomenon implying that supernovae deviate significantly from spherical symmetry. Asphericity will introduce additional uncertainties in the distances determined by using SN Ia as ``calibrated'' candles (Reiss, Press \\& Kirshner 1995; Hamuy et al. 1995), and by using the expanding photosphere method for SN II (Kirshner \\& Kwan 1975; Eastman \\& Kirshner 1989).  Only two SN Ia have been previously observed with spectropolarimetry. These are SN 1983G, reported in McCall et al. (1984), and SN 1992A reported in Spyromilio (1993). McCall et al. (1984) observed polarization of about 2\\% for SN 1983G, but no intrinsic polarization of the supernova can be identified. For SN 1992A, no noticeable polarization was detected down to the level of 0.3\\%. Theoretically, SN Ia are associated with binary stars and the ejecta should be intrinsically asymmetric at some level; however, the existing theories are not detailed enough to provide a clear picture of how the asymmetries evolve and whether they will be large enough to produce detectable degrees of polarization. On the other hand, SN Ia are bright enough close to maximum light for observational studies. A typical SN Ia in the Virgo cluster can reach 11th magnitude at maximum light. It is possible to collect enough photons so that accuracies down to 0.1\\% can be achieved in polarimetry.  There is no doubt that spectropolarimetry is the most powerful tool for extracting information on asymmetries in supernova explosions. Most supernovae, however, are quite faint and a systematic spectropolarimetric study seems impractical. In this paper, we report broad band polarimetry of the supernovae SN 1994D, SN 1994Y, SN 1994ae, SN 1995D and SN 1995H. The observations are described in \\S 2, and the results are given in \\S 3. A brief summary is given in \\S 4. ",
+        "conclusions": "\\subsection{Supernovae with polarization measurements}  We have observed five supernovae with broad band polarimetry. Among them, SN 1994Y and SN 1994H are of Type II, and SN 1994D, SN 1994ae and SN 1995D are of Type Ia. The observations show that the two SN II are polarized at a level of about 1.0\\%, while none of the Type Ia supernovae shows significant intrinsic polarization. The Type Ia supernovae were observed at different epochs after the outburst. No firm evidence of intrinsic polarization of Type Ia supernovae at any epoch can be established.  The supernovae for which we can find  polarization measurements in the literature are shown in Table 7. The table gives the SN designation (column 1) and type (column 2), the parent galaxy (column 3), the approximate date of observation past maximum light (column 4), the representative value of the observed polarization (column 5), comments on the detection of intrinsic supernova polarization (column 6), and the references (column 7). In column 6, a positive detection means that the observed polarization satisfies at least one of the criteria listed in \\S 2, while undetermined means that no firm evidence exists for any intrinsic polarization of the supernova.  The qualities of the measurements are very different for the supernovae listed in Table 7. The data on SN 1968L were reported by Serkowski (1970), but without any comments. As pointed out by Shapiro and Sutherland (1982), the extinction due to our own Galaxy alone is $A_v = 0.33$, which suggests that polarization due to the Galaxy alone could conceivably be as large as 1\\%. The measurement therefore has very little diagnostic value. The data on SN 1981B was quoted by Shapiro and Sutherland (1982) from Breger's unpublished observations. The observations were conducted only once in unfiltered light, and are apparently insufficient to establish any useful conclusion for the intrinsic polarization of the supernova. Spectropolarimetric observations of SN 1983G and SN 1983N were obtained by McCall et al. (1984) and McCall (1985). McCall et al. (1984) detected polarization of about 2\\% for SN 1983G with no clear changes of polarization parameters with wavelength. McCall et al. (1984) conclude that no intrinsic polarization is detectable. Their model analysis indicates that the apparent axis ratio of the expanding atmosphere was greater than 0.5. McCall (1985) also discussed spectropolarimetric data of  SN 1983N, the only Type Ib SN with polarimetric data. The observations were obtained close to maximum light and covered the wavelength range between 3700\\AA and 5200\\AA. Preliminary reductions indicated that the polarization spectrum dips from about 1.4\\% to about 0.8\\% at the position of the Fe II peak centered on 4600\\AA. The details of the data have not been published, but this behavior of polarization satisfies criterion $B$ in \\S 2, and therefore implies intrinsic polarization of around the same level, 1\\%.  Wolstencroft and Kemp (1972) claim that both linear and circular polarization are detected for the Type Ia supernova SN 1972E. The measured linear polarization is at a level of $0.35\\pm0.2\\%$, while the degree of circular polarization is $-0.028\\pm0.04\\%$. As argued by Shapiro and Sutherland (1982), however, the lack of time or wavelength dependence in the measurements makes a determination of the magnitude of the intrinsic part of the polarization very difficult.  The linear polarization of SN 1975N in NGC 7723 was measured by Shakhovskoi (1976). The observations were made in three colors at epochs corresponding to maximum light and a month past maximum light, respectively. No clear evidence of time evolution or wavelength dependence of polarization parameters was recorded. The mean level of polarization is about 1.6\\% and the error is typically 0.2\\%.  Because the extinction due to our Galaxy is likely to be small, Shakhovskoi (1976) argued that the polarization was intrinsic. As shown by Della Valle and Panagia (1992), however, the extinction due to the host galaxy can be as large as E$(B-V)$ = 0.36. This would easily produce polarizations of a few percent. The observations of Shakhovskoi (1976) are plotted in Fig. 3, together with a fit to the Serkowski interstellar polarization law. The parameters for the model fit are $P_{max}$ = 1.6 and $\\lambda_{max}$ = 4200\\AA. The observations were satisfactorily fitted with typical residuals of 0.25\\%. SN 1975N, like other supernovae of Type Ia, shows no convincing evidence for intrinsic polarization. Note also that the derived wavelength for maximum polarization is low compared to the average of 5400\\AA\\ in the ISM of our Galaxy (Serkowski, Mathewson \\& Ford 1975), however, it is comparable with the value of 4300\\AA\\ found by Martin \\& Shawl (1979) for M31 and Hough et al. (1987) for Centaurus A (NGC 5128). This suggests that the aligned grains which contribute to the polarization are about 20\\% smaller than those in our Galaxy, assuming their chemical characteristics are the same.  \\begin{figure}[h] \\plotone{fig3.ps} \\caption{The polarizations of SN 1975N fitted by the Serkowski law for interstellar polarization. Observations of Nov. 7, 1975 (circles), Nov. 8, 1975 (open squares), and Dec. 11, 1975 (solid squares) are plotted. The central wavelengths of the Nov. 8 and Dec. 11 data were shifted by +100\\AA\\ and -100\\AA\\ , respectively, for clarity. The success of the fit implies that no intrinsic polarizations can be identified down to the noise level of the data ($\\sim 0.2\\%$).} \\label{threebarrel} \\end{figure}  SN 1986G in Centaurus A (NGC 5128) is another SNIa with good polarimetric observations. Hough et al. (1987) report UBVRIJH broad band polarimetry of SN 1986G, and show that the observed polarimetry can be well fitted by Serkowski's law for interstellar polarization. The authors used SN 1986G to derive properties of dust distribution in the host galaxy of SN 1986G. They found no detectable intrinsic polarization for SN 1986G down to the noise level of about 0.1\\%.  Spyromilio and Baily (1992) report spectropolarimetry of the Type Ia supernova SN 1992A in the galaxy NGC 1038. The observations were obtained 2 weeks and 7 weeks after maximum light. The extinction towards the supernova is perhaps small, as no NaD lines are present at either the red shift of our Galaxy or the parent galaxy in an echelle spectrum. The mean level of polarization is around $0.3\\%$, and is comparable in strength to the noise level. No polarization is detected across any of the spectral features.  \\subsection{Are Type Ia different from Type II and Type Ib/c ?}  It is remarkable that in none of the SN Ia events have we positively identified any intrinsic polarization, while five Type II supernovae are found to be intrinsically polarized. Although the number of the supernovae observed is small and the quality of the observations varies, the present study seems to suggest that there is a significant difference in polarized light between Type Ia and Type II supernova. The SN Ia are less likely to be polarized, while all SN II show polarization around 1\\%. The only Type Ib/c supernova with polarization measurements is SN 1983N, which seems to be intrinsically polarized from preliminary analyses (McCall 1985). Fully reduced data have not been published for SN 1983N, and with a single event no statistical conclusion can be reached.  Several mechanisms can be invoked to explain polarization in a supernova atmosphere. The conventional picture involves electron scattering in an asymmetric envelope (Shapiro \\& Sutherland 1982; H\\\"oflich et al., 1995, and references therein). The mechanism producing the asymmetry is not clear, but may  perhaps be related to stellar rotation or the existence of a companion star. The asymmetries can also be produced by an asymmetric collapse or explosion, or due to a high velocity $^{56}$Ni clump (Chugai, 1992). To reproduce the observed amount of polarization for the observed Type II supernovae with an asymmetric envelope, the supernovae ejecta would have to possess asymmetries with major axis $\\sim 1.5$ times longer than the minor axis. In this picture, our measurements would suggest that the supernova envelope is more asymmetric for Type IIs than for Type Ia's. If this is the case, SN Ia might make better distance calibrators and SN II would have to be considered with considerable caution.  Scattering by dust particles in the circumstellar and interstellar material, the so-called ``light echo'' of a supernova, may also produce polarized light (Wang \\& Wheeler 1996a). In this picture, the supernova light is scattered by the circumstellar or interstellar dust, and due to the light travel time across the circumstellar or interstellar material, some photons from earlier epochs will be integrated into the observation. The scattered light will be highly polarized, and when the scattering material is asymmetrically distributed, the integrated scattered light can also be polarized. Chevalier (1986) studied this situation for a particular case with circumstellar density distribution given by $\\rho\\,=\\,Cr^{-2}(1-\\beta\\,\\cos 2\\phi)$, where $C$ is a constant, $\\beta$ is a measurement of the asymmetry and $\\phi$ is the angle from the symmetry axis. The integrated polarization of the scattered light can easily be as large as several tens of a percent. As can be derived from the equations given in Chevalier (1986), the total flux of the scattered light can be a few percent of that of the supernova near maximum light. This mechanism can therefore easily produce polarizations at a level of around 1\\%. Dust blobs ejected by the progenitor can be even more efficient in producing polarized light.  For such a dust scattering mechanism, our observations  would imply that Type II supernovae are more likely to be associated with a significant circumstellar material than Type Ia supernovae. In this case, the polarization measurements will not endanger the distance determination using the expanding photosphere method for Type II supernova (Eastman \\& Kirshner, 1989).  \\subsection{Progenitors of Supernovae}  Polarimetry can set strong constraints on possible progenitors of supernovae. The low level of the observed polarization in Type Ia, in particular, may put strong constraints on various models for Type Ia. A popular model for Type Ia is that they arise in a binary system when a red-giant companion to a previously formed white dwarf finally evolves and transfers mass to the white dwarf, leading to thermonuclear explosion of the white dwarf (Iben \\& Tutukov 1984; Paczynski 1985; Wheeler \\& Harkness 1990; Wheeler 1996). The observed low level of polarization may therefore give interesting constraints on progenitor system models with cool companions such as cataclysmic variables and symbiotic systems. Some of these systems show polarizations as high as 15\\% (e. g. Aspin 1988; Magalhaes \\& Schulte-Ladbeck 1988). Scattering by circumstellar dust particles plays an important role in producing the observed polarization. It is then interesting to ask why, if they result from systems with high polarization, SN Ia are not polarized. The low level of polarization for Type Ia may also help to clarify the link between novae and supernovae. There is both observational and theoretical evidence that novae are antithetical to SN Ia. Polarimetry of novae generally show, in contrast to SN Ia, polarization at a level around 1\\% (Bjorkman et al. 1994). Does this rule out nova-like systems as progenitors of SN Ia ? Another possibility is that SN Ia ejecta expand much faster than nova ejecta, and the asymmetries or dust particles are quickly destroyed. A closer look into these problem is necessary and is underway (Wang \\& Wheeler 1996b).  Another question is why SN II are polarized ? While models of SN Ia generally involve binary systems, most models of SN II, on the other hand, involve only massive single stars. If the observed polarization is due to an asymmetric ejecta, it then implies that the supernova atmosphere or circumstelar environment are non-spherical. What is producing the asphericity ? Can rotation alone produce asymmetries large enough to explain the asymmetry ? Does this imply that all or most of the SN II are produced by binary systems as suggested for SN 1987A (Podsiadlowski 1992) and SN 1993J (Podsiadlowski et al. 1993; Shigeyama et al. 1994; Wheeler et al. 1993; Woosley et al. 1994) as well ?  \\subsection{Perspectives of Future Supernova Polarimetry}  It is obvious that more polarimetric data are needed for a clear picture of the geometric structure of supernovae. We have suggested that there may be a significant difference in polarization between supernova of Type Ia and Type II. More observations are critical to verify this proposed dichotomy and to determine if SN Ib/c are as commonly polarized as SN II. We need to understand the origin of the polarization of SN II and why, if they occur in binary systems with substantial mass transfer, SN Ia display so little polarization.  An enlarged sample of high quality polarimetric data of supernovae is not only important for studying the supernova phenomenon, but may also be useful for the study of the physical properties of the interstellar medium. If, as suggested earlier, Type Ia supernovae are not intrinsically polarized, we should then be able to use them as good indicators of interstellar polarization along the line of sight to the supernova. This will be an important method for gaining insight into interstellar polarization in galaxies other than our own.  This research is supported in part by NASA Grant GO 5652, NAGY 2905 and NSF Grant 9218035. We are grateful to the staff at McDonald observatory for help in obtaining these data. We are indebted to B. Wills for providing the transmission curves for the Breger polarimeter filter system, and D. Wills for obtaining observations of SN 1994D and reading an initial draft of this paper. We thank the referee, M. L. McCall, for many useful comments."
+    },
+    "9602/astro-ph9602013_arXiv.txt": {
+        "abstract": "We present a study of the surface density profiles of the clusters of galaxies from the Palomar Distant Cluster Survey (Postman et al. 1996).  The survey contains a total of 79 clusters of galaxies, covering the estimated redshift range of $0.2 \\simless z \\simless 1.2$.  We have analyzed the richest clusters in this sample and find that the typical Palomar cluster has a surface density profile of $r^{-1.4}$ ($r \\ge 0.10~h^{-1}~{\\rm Mpc}$) and a core radius of $0.05~h^{-1}~{\\rm Mpc}$. There may be an indication that the slope of the surface density profile steepens with increasing redshift, though the observational uncertainty is at present too large to be conclusive. Our cluster population is inconsistent at a 99.9\\% confidence level with a population of azimuthally symmetric clusters. ",
+        "introduction": "Clusters of galaxies provide a powerful probe of the nature of galaxy formation and the origin of large scale structure in the universe. Having a well-studied sample of intermediate and high-redshift clusters is vital to understanding the evolution of galaxies and large scale structure. Because the diversity of cluster properties and the effects of evolution are significant at high redshift (Gunn \\& Dressler 1988; Bower et al. 1994; Castander et al. 1994; Postman et al. 1996), it is essential to sample many distant clusters with different properties. Large scale simulations have shown that cluster properties such as profile shape and substructure can constrain the nature of large scale structure formation theories and the mass density of the universe (Evrard et al. 1993; Crone, Evrard \\& Richstone 1995; Jing et al. 1995; Tyson \\& Fischer 1995).  Detailed studies of these observational parameters have already been made in some individual, intermediate redshift clusters.  Using weak lensing to map the mass distribution in several clusters, Smail et al. (1995) have studied two X-ray luminous clusters 1455+22 ($z = 0.26$) and 0016+16 ($z = 0.55$), Tyson \\& Fischer (1995) have studied Abell 1689 ($z = 0.18$), and Squires \\etal (1995) have studied Abell 2218 ($z = 0.175$).  All four clusters show moderate to extreme degrees of structure.  However, the distribution of structure in mass, light, and X-rays are all well correlated at radii larger than $100~h^{-1}~{\\rm kpc}$. The projected mass density (and light) profile of A1689 can be well approximated by a power-law of $r^{-1.4 \\pm 0.2}$, while the projected total mass, gas, and light surface densities of A2218 are consistent with an isothermal sphere ($r^{-1}$), though perhaps slightly steeper. Smail et al. (1995) found that, while the galaxies are very good tracers of the mass, they are less concentrated, the respective core radii for 1455+22 being $r_{c}^{mass} = 50^{+40}_{-25}~h^{-1}~{\\rm kpc}$ and $r_{c}^{gal} = 90^{+35}_{-25}~h^{-1}~{\\rm kpc}$.  A multi-color photometric study of three clusters, A3284 ($z = 0.15$), A3305 ($z = 0.16$), and A1942 ($z = 0.23$), also yield consistent core radii of $r_{c}^{gal} \\approx 120,~100,~{\\rm and}~120~h^{-1}~{\\rm kpc}$, respectively (Molinari et al.\\ 1994).  In this paper, we expand on these detailed analyses of individual intermediate redshift clusters by studying a large sample of intermediate and high redshift clusters from the Palomar Distant Cluster Survey (hereafter PDCS; Postman et al. 1996).  We examine both the individual and global properties of the cluster profiles, specifically the cluster profile slope, core radius, and degree of asymmetry.  Because the PDCS uses a completely objective and automated algorithm to detect clusters, it provides us with the largest, statistically complete sample of distant clusters.  In addition, the selection biases due to the detection technique can be well quantified through simple Monte-Carlo simulations.  In \\S 2 of this paper we briefly describe the PDCS and the cluster sample used in this analysis.  The parameters of the individual cluster profiles are discussed in \\S 3. Composite cluster profiles as a function of redshift are presented in \\S 4. Cluster morphology is discussed in \\S 5.  We present a preliminary comparison of our results to large-scale cosmological simulations in \\S 6 and summarize the results in \\S 7. ",
+        "conclusions": "We have examined the galaxy distributions of the richest clusters of galaxies from the Palomar Distant Cluster Survey.  Through an analysis of individual cluster profiles and composite profiles as a function of redshift, we find that the typical Palomar cluster has a profile of $r^{-1.4}$ at radii greater than $0.10~h^{-1}~{\\rm Mpc}$ and a core radius of $0.05~h^{-1}~{\\rm Mpc}$. The distribution of core radii in the Palomar clusters indicates that up to 30\\% of the clusters have very small or essentially no core radii. Using simple Monte-Carlo simulations, we have shown that the median cluster profile parameters appear to be an accurate representation of the actual cluster population; the measurement errors associated with these parameters are $\\sim 10\\%$ in the surface density slope $\\alpha$ and $\\sim 50\\%$ in the cluster core radius $r_{c}$. The average profile parameters of the PDCS clusters are consistent with measures of nearby clusters, as well as other intermediate redshift clusters (Molinari et al.\\ 1994; Tyson \\& Fischer 1995; Squires et al.\\ 1995; Smail et al.\\ 1995). There may be an indication that the mean cluster profile steepens with increasing redshift, though the observational uncertainty is presently too large to be conclusive.  In addition, we find that a large fraction of the PDCS clusters have a significant degree of asymmetry.  The PDCS cluster population is inconsistent with a population of circularly symmetric clusters of galaxies at a greater than 99.9\\% confidence level.  \\vskip 0.7cm  We thank the anonymous referee for comments which greatly improved this paper. We gratefully acknowledge the useful discussions with Neta Bahcall, Jim Gunn, Robert Lupton, David Spergel, and Michael Strauss. We also thank Sangeeta Malhotra for her library of SM macros, Bob Rutledge for his incredible fitting program BFIT, and Ue-Li Pen and Guohong Xu for preliminary results of their cosmological simulations. This work is supported in part by NASA contract NGT-51295 (LML).  \\vfill \\eject"
+    },
+    "9602/astro-ph9602032_arXiv.txt": {
+        "abstract": "In this paper we address the question whether age and metallicity effects can be disentangled with the aid of the  broad-band colours and spectral indices from  absorption feature strengths, so that the age of elliptical galaxies can be inferred. The observational data under examination are the  indices \\Hbeta\\ and \\MgFe,  and the velocity dispersion $\\Sigma$ for the sample of galaxies of Gonzales (1993), supplemented by the ultra-violet data, i.e. the colour (1550-V), of Burstein et al. (1988). The analysis is performed with the aid of chemo-spectro-photometric models of elliptical galaxies with infall of primordial gas (aimed at simulating the collapse phase of galaxy formation) and the occurrence of galactic winds. The galaxy models are from Tantalo et al. (1995). The study consists of four  parts. In the first one, the aims are outlined and the key data are presented. In the second part, we summarize the main properties of the infall models that are relevant to our purposes. In the third part we present the  detailed calculations of the spectral indices for single stellar populations and model galaxies. To this aim, we use the analytical relations of Worthey et al. (1994) who  give index strengths as a function of stellar parameters. In the last part, we examine the age-metallicity problem. In contrast with previous interpretations of the \\Hbeta\\ and \\MgFe\\ data  as a sort of age sequence (Gonzales 1993), we find that the situation is more complicate when the space of the four variables \\Hbeta,  \\MgFe, (1550-V), and $\\Sigma$ is examined. Galaxies in the \\Hbeta\\ and \\MgFe\\ plane do not follow a pure sequence either of age or metallicity. The observed (1550-V) colours are not compatible with young ages. Basically, all the galaxies in the sample are old objects (say as old as 13$\\div$15 Gyr) but have suffered from  different histories  of star formation. Specifically,  it seems that some galaxies  have exhausted the star forming activity at very early epochs with no significant later episodes. Others have continued to form stars for long periods of time. This is perhaps sustained by the analysis of the gradients in the \\Hbeta\\ and \\MgFe\\ indices across the galaxies. There are galaxies with no age difference among the various regions.  There are other galaxies in  which large gradients in the mean age of the star forming activity between the central and the peripheral regions seem to exist. The nucleus turns out to be younger and more metal-rich than the outer regions. Finally, there are galaxies in which the nucleus is older but less metal-rich than the external regions. All this perhaps hints not only different histories of star formation  but also different mechanisms of galaxy formation difficult to pin down at the present time. From the analysis of the  \\Hbeta,  \\MgFe, (1550-V), and $\\Sigma$ space, and of the age and metallicity gradients in single galaxies, the suggestion is advanced that the overall duration of the star forming activity is inversely proportional to the velocity dispersion $\\Sigma$ (and perhaps  galactic mass).  \\keywords {Galaxies: ages - Galaxies: photometry - Ultraviolet: galaxies } ",
+        "introduction": "In this section we summarize current information on basic data for elliptical galaxies that are either constraints or targets of our models.  \\subsection{ Hints on metallicities and ages }  Direct determinations of metallicities in elliptical galaxies are still rather limited. Most of the information comes from metal line-strength indices and their radial variations (Carollo et al. 1993, Carollo \\& Danziger 1994,  Davies et al. 1993)  or colour gradients (Schombert et al. 1993). The evidence  arises that the metallicities for elliptical galaxies are solar or larger. Furthermore, for bright elliptical galaxies there are also indications that the $\\alpha$-elements (O, Mg, Si, etc..) are enhanced with respect to Fe. The average [Mg/Fe], in particular,  exceeds that of the most metal-rich stars in the solar vicinity by about 0.2-0.3 dex, and the ratio [Mg/Fe] is expected to increase with the galactic mass up to the this value. Due to their proximity, the metallicity  of stars in the Galactic Bulge  is currently measured by spectroscopic methods (cf. Rich 1990,  McWilliam \\& Rich 1994). The mean metallicity is slightly lower than solar and long tails towards the metal-poor (${\\rm [Fe/H]=-1}$) and metal-rich end (${\\rm [Fe/H]=1}$) are also present. High metallicities seem also to be indicated by the color-magnitude diagrams (CMD) of stellar populations  in the bulges of nearby Galaxies (M31 and M32). In the case of M32 (Freedman 1989, 1992), these stars are most likely in the RGB and AGB phases and metallicities as high as [M/H]=0.1 are possible (cf. Bica et al. 1990, 1991). Concerning the age, from the magnitude of the brightest AGB stars in M32, Freedman (1989, 1992) concludes that  a fraction of these could be as young as 5 Gyr.   \\begin{figure} \\psfig{file=fig1_gonz_data.ps,height=8.5truecm,width=8.5truecm} \\caption {The fundamental plane \\Hbeta\\ versus \\MgFe\\ for the Re/8 data of Gonzales (1993) together with the error bars for each galaxy } \\label{gonz_data} \\end{figure}   \\subsection {Line strength indices \\Hbeta\\ and \\MgFe\\ }  Table~\\ref{tab_data} contains  the key data relevant to our purposes, namely the fully corrected indices \\Hbeta\\ and \\MgFe\\, and the velocity dispersion $\\Sigma$ (km sec$^{-1}$) of Gonzales (1993) for two different galaxy coverage, i.e. the  central region within Re/8, where Re is the effective radius, and the wider region within Re/2. Column (1) is the galaxy identification NGC number, column (2) the logarithm of the \\Hbeta\\ index, column (3) is the logarithm of the \\MgFe\\ index, column (4) is the logarithm of the velocity dispersion $\\Sigma$ in km sec$^{-1}$ for the Re/8-data. Columns (5), (6) and (7) show the same but for the Re/2-data. The content of the remaining columns will be described below.  The fundamental plane \\Hbeta\\ versus \\MgFe\\ for the Re/8 data is shown in Fig.~\\ref{gonz_data} together with the error bars for each galaxy. It is worth noticing that passing from the Re/8 to Re/2 data, there is a mean decrease in $\\log\\MgFe$ amounting 0.032,  and a mean increase in $\\log\\Hbeta$ amounting to 0.003.       \\begin{table*} \\caption{Basic observational data used in this study: fully corrected indices \\Hbeta\\ and \\MgFe, and velocity dispersion $\\Sigma$ (km sec$^{-1}$) for the Re/8 and Re/2 data of Gonzales (1993);  \\UVex\\  colours of Burstein et al. (1988) for the central regions; magnitudes ${\\rm M_V}$ and ${\\rm M_V}$,  and colours (B--V) and (V--K) of Bender et al. (1992, 1993), Bower et al. (1992), Marsiaj (1992) and Schweizer \\& Seitzer (1992) as explained in the text. } \\littleskip \\scriptsize \\begin{center} \\begin{tabular*}{180mm}{ccc ccc ccc ccc ccc cc} \\hline \\hline NGC &\\Hbeta&\\MgFe&$\\Sigma$&\\Hbeta&\\MgFe&$\\Sigma$&\\UVex&Note&-${\\rm M_B}$&S& -${\\rm M_B}$& B-V& B-V& V-K& V-K&-$M_V$\\\\ \\hline (1) &(2) &(3) &(4) &(5) &(6)&(7)  &(8)& (9) &(10)&(11)&(12)&(13)&(14)& (15)&(16)&(17)\\\\ \\hline &Re/8&Re/8&Re/8&R2/2&Re/2&Re/2&    &   &     &   &     &    &    &    &    &     \\\\ \\hline 221&0.36&0.45&1.86&0.33&0.45&1.81&4.50&   &     &   &15.70&0.84&    &    &    &16.64\\\\ 224&0.22&0.59&2.19&0.23&0.58&2.19&3.51&   &     &   &     &    &0.87&    &    &     \\\\ 315&0.24&0.57&2.51&0.25&0.54&2.49&    &   &     &   &23.61&1.00&    &    &    &24.61\\\\ 507&0.24&0.55&2.42&0.31&0.53&2.42&    &   &     &   &     &    &    &    &    &     \\\\ 547&0.20&0.58&2.37&0.15&0.55&2.34&    &   &     &   &     &    &    &    &    &     \\\\ 584&0.32&0.55&2.29&0.31&0.52&2.23&3.93&   &     &   &21.72&0.93&0.94&    &3.18&22.66\\\\ 636&0.28&0.55&2.20&0.27&0.52&2.17&    &   &     &   &20.65&0.92&0.92&    &    &21.57\\\\ 720&0.25&0.59&2.38&0.36&0.58&2.25&    &   &     &   &21.60&0.99&0.97&    &    &22.59\\\\ 821&0.22&0.56&2.28&0.26&0.53&2.25&    &   &     &   &21.61&0.98&0.98&    &    &22.59\\\\ 1453&0.20&0.58&2.46&0.23&0.55&2.42&    & SS&22.26&1.5&     &    &1.02&    &    &     \\\\ 1600&0.19&0.60&2.50&0.24&0.60&2.49&    &   &     &   &23.17&0.97&0.98&    &3.34&24.14\\\\ 1700&0.32&0.55&2.36&0.32&0.52&2.35&    & SS&22.50&3.7&22.28&0.92&0.91&    &3.18&23.20\\\\ 2300&0.23&0.58&2.40&0.21&0.56&2.35&    & SS&21.45&2.8&21.56&1.04&1.02&    &3.42&22.60\\\\ 2778&0.25&0.56&2.19&0.19&0.53&2.11&    &   &     &   &18.14&0.91&    &    &    &19.05\\\\ 3377&0.32&0.51&2.03&0.33&0.45&1.95&    & SS&19.06&1.5&19.49&0.90&0.90&    &2.95&20.39\\\\ 3379&0.21&0.57&2.31&0.20&0.54&2.27&3.86& SS&20.44&0. &20.17&0.97&0.96&    &3.26&21.13\\\\ 3608&0.23&0.57&2.25&0.24&0.53&2.21&    & SS&19.87&0. &19.87&0.97&0.94&    &3.17&20.84\\\\ 3818&0.23&0.58&2.24&0.26&0.52&2.18&    & SS&19.61&1.3&19.49&0.92&0.92&    &    &20.41\\\\ 4261&0.13&0.59&2.46&0.11&0.56&2.43&    & SS&21.78&1.0&21.74&0.99&0.99&    &3.29&22.73\\\\ 4278&0.19&0.56&2.37&0.22&0.53&2.31&2.88& SS&19.93&1.5&19.79&0.96&    &    &3.21&20.75\\\\ 4374&0.18&0.56&2.45&0.19&0.54&2.43&3.55&SSV&21.80&2.3&21.85&0.98&0.95&3.28&3.29&22.73\\\\ 4472&0.21&0.57&2.45&0.22&0.56&2.41&3.42& V &     &   &22.21&0.98&0.96&3.32&3.36&23.10\\\\ 4478&0.26&0.55&2.11&0.24&0.52&2.13&    & V &     &   &19.44&0.89&    &3.13&3.24&20.33\\\\ 4489&0.38&0.47&1.67&0.36&0.42&1.61&    &   &     &   &     &    &    &    &    &     \\\\ 4552&0.17&0.59&2.40&0.18&0.56&2.37&2.35& V &     &   &     &    &0.96&3.33&3.31&21.68\\\\ 4649&0.15&0.60&2.49&0.14&0.57&2.45&2.24&   &     &   &21.81&1.01&0.99&    &3.36&22.82\\\\ 4697&0.24&0.53&2.21&0.22&0.47&2.19&3.41&SSV&21.76&0. &21.51&0.95&0.94&    &3.22&22.46\\\\ 5638&0.22&0.56&2.19&0.23&0.52&2.14&    &   &     &   &     &    &    &    &    &     \\\\ 5812&0.23&0.58&2.30&0.23&0.56&2.26&    &   &     &   &     &    &    &    &    &     \\\\ 5813&0.15&0.55&2.31&0.09&0.53&2.32&    &   &     &   &     &    &    &    &3.29&     \\\\ 5831&0.30&0.56&2.20&0.31&0.57&2.18&    & SS&20.28&2.6&20.22&0.95&    &    &3.42&21.85\\\\ 5846&0.16&0.58&2.35&0.10&0.54&2.32&3.11&   &     &   &21.85&0.99&    &    &3.42&22.84\\\\ 6127&0.18&0.58&2.38&0.18&0.55&2.34&    &   &     &   &     &    &    &    &    &     \\\\ 6702&0.39&0.53&2.24&0.40&0.51&2.22&    &   &     &   &     &    &    &    &3.26&     \\\\ 6703&0.27&0.55&2.26&0.26&0.51&2.22&    &   &     &   &     &    &    &    &    &     \\\\ 7052&0.17&0.58&2.44&0.25&0.55&2.42&    &   &     &   &     &    &    &    &    &     \\\\ 7454&0.33&0.45&2.03&0.32&0.40&2.04&    &   &     &   &     &    &    &    &    &     \\\\ 7562&0.23&0.51&2.39&0.24&0.54&2.38&    &   &     &   &     &    &    &    &3.30&     \\\\ 7619&0.13&0.59&2.48&0.17&0.56&2.44&    & SS&22.23&0. &22.35&1.01&1.02&    &3.47&23.36\\\\ 7626&0.16&0.58&2.40&0.16&0.54&2.38&    & SS&22.28&2.6&22.36&1.00&1.00&    &3.39&23.36\\\\ 7785&0.21&0.56&2.38&0.18&0.54&2.34&    &   &     &   &22.22&0.95&    &    &    &     \\\\ \\hline \\hline \\end{tabular*} \\end{center} \\label{tab_data} \\normalsize \\end{table*}   \\subsection{  UV excess and  \\UVex\\ colour }  Much of the available information is from the study of Burstein et al. (1988). The main points are the following: \\littleskip \\noindent (1) All studied elliptical galaxies have detectable flux short-ward of $\\sim 2000$ \\AA, with  large galaxy to galaxy differences in the level of the UV flux.  The intensity of the UV emission is measured by the colour (1550-V). \\littleskip \\noindent (2) The colour (1550-V) correlates with the index Mg2, the velocity dispersion $\\Sigma$ and the luminosity (mass) of the galaxy. The few galaxies (e.g. NGC~205) in which active star  formation is seen do not obey these relations. \\littleskip \\noindent (3) An important constraint is posed by the HUT observations by Ferguson et al. (1991) and Ferguson \\& Davidsen (1993)  of the UV excess in the bulge of M31. In this case the UV emission  shows a drop-off short-ward of about 1000 \\AA\\ whose interpretation requires that  the temperature of the emitting source must be about 25,000 K.  Only a small percentage of the $912 \\leq \\lambda \\leq 1200 $ \\AA\\ flux can be coming from stars hotter than 30,000 K and cooler than 20,000 K. \\littleskip  The \\UVex\\ colours of Burstein et al. (1988) are listed in column (8) of Table~\\ref{tab_data}.   \\subsection{The CMR  of early-type galaxies}  Colours and magnitudes of elliptical galaxies define a mean CMR, according to which the colours get redder at decreasing absolute magnitude and hence increasing mass of the galaxies.  The slope of the CMR is commonly understood as indicative of a mass-metallicity sequence (Dressler 1984, Vader 1986) resulting from the effect of galactic winds  (Larson 1974; Saito 1979a,b; Bressan et al. 1994). In brief, independently of their total mass,  galaxies evolving as either {\\it closed or open boxes} (cf. Tinsley 1980) should reach metallicities governed by the degree of gas gas consumption. Elliptical galaxies, having exhausted their gas,  should possess almost identical (high) metallicities. Therefore, in order to have a mass-metallicity sequence, a   mechanism  halting star formation and thus leaving different metallicities must be invoked: galactic winds are the right agent (cf.  Arimoto \\& Yoshii 1986, 1987; Matteucci \\& Tornamb\\'e 1987; Yoshii \\& Arimoto 1987; Bressan et al. 1994; Tantalo et al. 1995).  The tightness of the CMR seems to depend on the galaxy environment and could reflect the process of galaxy formation.  Looking at the CMR of Bower et al. (1992a,b) for galaxies in the largest known nearby clusters Virgo and Coma, the colour dispersion is very small (uncertainty in absolute magnitudes is less of a problem here), with typical rms scatter of  0.05 mag of which 0.03 mag can be accounted for by observational errors. Bower et al. (1992a,b) using the rate of colour change ${\\partial \\UV / \\partial t}$ for SSPs of Bruzual (1983) argued that for the Hubble time $t_H=15$ Gyr and coordination in the galaxy formation process (their parameter $\\beta=1$)  the galaxy ages are confined in the range $13.5 \\simeq t_F \\simeq 15.0 $ Gyr. It is an easy matter to check that the Bower et al. (1992a,b) conclusion neither depends on the particular rate of colour change they have adopted nor the use of single SSPs instead of a manyfold of SSPs with different chemical composition weighted on the past rate of star formation (thus closer to the complexity of a real galaxy). This can be checked using the SSPs of Bressan et al. (1994) and Tantalo et al. (1995) or the models presented below.  In contrast,  examining the colour dispersion (relatively  small also in this case) of the CMR for nearby galaxies in small groups and field, Schweizer \\& Seitzer (1992) found it compatible with recent mergers of spiral galaxies spread over several billion years. Also in this case the CMR implies a mass-metallicity sequence.  Since the  interpretation of the CMR bears very much on the kind of process of galaxy formation and evolution, its use as a constrain on the galactic models has different implications. If the CMR is indicating nearly coeval, old galaxies ranked along a mass-metallicity sequence,  passive evolution after the initial star forming activity and onset of galactic winds is appropriate. This facilitates the use of the CMR as a constrain on galactic models (cf. Bressan et al. 1994, Tantalo et al. 1995). In contrast, if the CMR is compatible with mergers spreading over long periods of time, first the evolutionary history of an elliptical galaxy is more complicated, and the use of the CMR as a constraint is less secure.  In order to cast light on this topic we try a closer inspection of the data and derive the CMR for the Gonzales (1993) sample to be compared with that of Bower et al. (1992a,b) and Schweizer \\& Seitzer (1992). Cross-checking of the three lists reveals that,  while the Gonzales (1993) and the Schweizer \\& Seitzer (1992) samples have many objects in common (17 galaxies), there is only one galaxy in common to all the three groups (NGC~4374),  and four in common to Gonzales (1993) and Bower et al. (1992a,b), namely the  galaxy above plus NGC~4472, NGC~4478, and NGC~4552. The galaxy  NGC~4374 is particularly interesting because Schweizer \\& Seitzer (1992) have estimated for it a merger age of about 7 Gyr.  In addition to this, we tried to assign magnitudes and colours to as many of the Gonzales (1993) galaxies as possible making use  of magnitudes and colours taken from Bender et al. (1992, 1993) and  Marsiaj (1992). All the data are presented in Table~\\ref{tab_data}. Column (9) indicates whether the galaxy belong to the Schweizer \\& Seitzer (1992) list (SS) and/or the Virgo cluster (V or SSV); columns (10) and (11) are the ${\\rm M_B}$ magnitude and the fine structure parameter S, respectively,  of Schweizer \\& Seitzer (1992). The parameter S is a rough measure of dynamical youth or rejuvenation (0 no trace of fine structure, 10 strong fine structure). Column (12) and (13) are the magnitude ${\\rm M_B}$ and colour (B-V), respectively, taken from Bender et al. (1992, 1993). In a few cases these magnitudes are somewhat different from those of Schweizer \\& Seitzer (1992). Columns (14) is the (B-V) colour listed by Marsiaj (1992). In general, it agrees with that of Bender et al. (1992, 1993). Columns (15) and (16) are the (V-K) colours of the Bower et al. (1992a,b) and Marsiaj (1992) lists, respectively. Also in this case,  for the few galaxies in common  the agreement is remarkable. This means that Marsiaj's (1992) (V-K) colours assigned to the remaining galaxies (when possible) can be considered as reasonably good. Finally, column (17) lists the magnitude ${\\rm M_V}$ obtained from columns (12) and (13).  The resulting \\VK\\ versus ${\\rm M_V}$ CMR is shown in panel (a) of Fig.~\\ref{cmr}, whereas that of Bower et al. (1992a,b) is displayed in panel (b). In this latter, we adopt for Virgo the distance modulus \\DMo=31.54 (Branch \\& Tammann 1992) and apply to Coma the shift $\\delta\\DMo=3.65$ (cf. Bower et al. 1992a,b). In Fig.~\\ref{cmr} we also plot the loci of constant ages (17, 10, 8 and 5 Gyr) for the models presented in Sect. 3. The comparison of the two CMRs yields the following results: \\littleskip \\noindent (1) The slopes of the two CMRs agree each other and with the theoretical prediction. This means that both groups of galaxies reflect a similar mass-metallicity sequence. \\littleskip \\noindent (2) Indeed the  Bower et al. (1992a,b) CMR is less scattered than the one of Schweizer \\& Seitzer (1992), thus explaining the different interpretations. \\littleskip \\noindent 3) NGC~4374 and NGC~4697, for which  rather recent merger and accompanying stellar activities have  been proposed by Schweizer \\& Seitzer (1992), are compatible with the Bower et al. (1992a,b) conclusion at the same time. \\littleskip \\noindent 4) Finally, as galactic models are imposed to match the slope and the mean (V--K) colours of the CMR, the results do not depend on the particular CMR in usage.  Therefore, owing to its better definition, we prefer to adopt  the Bower et al. (1992) CMR to infer the slope of the underlying mass-metallicity sequence.     \\begin{figure} \\psfig{file=fig2_cmr_new.ps,height=9.0truecm,width=8.5truecm} \\caption { {\\bf Panel a}: the CMR for the sample by Schweizer \\& Seitzer for which we have obtained the magnitude $M_{V}$. The position of NGC~4374 (filled star) and NGC~4697 (open square) is indicated. {\\bf Panel b}: the CMR for galaxies in Virgo (circles) and Coma (filled circles) reproduced from Bower et al.(1992).  The galaxy NGC~4374 is shown by the open star. The typical error bars of the data are also displayed.  In addition to this, in each panel, are drawn the line of constant age for  the infall models characterized by the parameters:  $\\tau=0.10$ Gyr, $k=1$, $\\zeta=0.50$, and  $\\nu$ variable with \\ML\\ (cf. Table 2). Each line corresponds to different values of the age: (solid: 17 Gyr), (dotted: 10 Gyr), (short-dashed: 8 Gyr), and (long-dashed: 5 Gyr). See the text for more details  } \\label{cmr} \\end{figure} ",
+        "conclusions": "In this paper we have addressed the question whether age and metallicity effects on the stellar content of elliptical galaxies can be disentangled by means of broad-band colors and line strength indices derived from detailed chemo-spectro-photometric models. The main goal of our study is to cast light on a number of interwoven questions: are  elliptical galaxies made only of old stars generated in a dominant initial episode of star formation? is there any evidence of subsequent star forming  episodes ? If so, has this activity an internal origin or is it induced by the merger mechanism ?  The analysis of the observational data in the space of the parameters \\Hbeta, \\MgFe, \\logS, and \\UVex\\ and of the gradients in\\Hbeta\\ and \\MgFe\\ within individual galaxies allow us to draw a number of conclusions.  \\begin{itemize} \\item{The  distribution of galaxies in the \\Hbeta\\ - \\MgFe\\ plane does  not correspond to a mere  sequence of metallicity with bluer galaxies significantly more metal-poor than the red ones. Metallicity effects are expected to broaden the distribution mainly along the   \\MgFe\\ axis.}  \\item{Equally, the  distribution of galaxies in the  \\Hbeta\\ - \\MgFe\\ plane does not constitute a mere sequence of age with  bluer galaxies significantly younger than the red ones. Although some scatter in the age is possible, it is not likely to be as large as indicated by the formal match of data with isochrones in this diagram. M32 in particular according to its \\Hbeta\\ should have an age of about 3 Gyr, in disagreement with its \\UVex=4.5 for which much older ages are needed according to the present theory of population synthesis.}  \\item{The observed distribution of galaxies in the  \\Hbeta\\ - \\MgFe\\ plane does not agree with the CMR-strip on the notion that they evolve as passive objects in which star formation took place in an initial episode. Most likely, the history of star formation was more complicated than this simple scheme. }  \\item{The observed \\UVex\\ colours are not compatible with ages younger that about 6 Gyr. UV excesses stronger than \\UVex=5.5 at younger ages are possible only if the age is younger that $1\\div 2$ Gyr. However, in such a case the remaining UBV colours would be too blue compared with the observational data. Most likely, all galaxies in the sample are globally older than at least 6 Gyr. }  \\item{As far as ages are concerned, our analysis is still unable to give a definitive answer. Besides the group of galaxies like M32 and NGC~584 whose properties lead to contrasting results depending on the particular relationship among the four parameters \\Hbeta, \\MgFe, \\UVex\\ and $\\Sigma$ under examination, the remaining galaxies appear to span a wide range of ages going from  8 to 15 Gyr. However, this large scatter can be partly due to observational uncertainties and partly  to a real age spread. It is worth pointing out that for the very small subgroup of galaxies with all the four parameters simultaneously  available, the objects with an old age ($13\\div 15$ Gyr) like NGC~4649 show an  age spread  amounting only  to about 3 Gyr, which is   consistent with the small scatter in the CMR of cluster galaxies.}    \\item{The present analysis cannot yet cast light on the dominant mechanism by which  elliptical galaxies are formed. However, the following considerations can be made. If elliptical galaxies are all old, coeval  and evolving in isolation, the difficulties encountered with the interpretation of the \\Hbeta, \\MgFe, \\UVex, and \\logS\\ data can be removed by invoking later mild episodes of star formation. The major problem with this suggestion  is that in the case of evolution in isolation, some mechanism synchronizing the bursts of star formation should be required in order to dislocate galaxies from the CMR-strip to their observed location in the \\Hbeta\\ - \\MgFe\\ plane. Indeed, even if the statistics is limited, no galaxies are seen along the CMR-strip but for those at the red end. A simple way out could be that the star forming activity is limited to the nuclear region, thus affecting the \\Hbeta\\  index without changing significantly the global value of the broad-band colours (cf. the case of NGC~4374 and NGC~4697 for which little age scatter is indicated by the CMR, whereas a significant age scatter is suggested by the indices). No such difficulty would exist if already formed ellipticals suffer in more recent times from interaction with other small size galaxies (for instance the capture of small gas rich systems) thus inducing some mild star formation activity with some chemical enrichment. Equally, if  elliptical galaxies are the result of classical mergers of spiral galaxies. In such a case a spread in \\Hbeta mimicking an age sequence for the resulting composite population would naturally follow. The accompanying chemical enrichment can be easily adjusted to  match the small observational scatter along the \\MgFe\\ axis. }   \\item{Although based on a handful of objects, there is a interesting scenario emerging from the analysis of the four dimensional space \\Hbeta,  \\MgFe, \\UVex, and \\logS. In brief, galaxies of high velocity dispersion tend to be confined at the red end of the \\Hbeta\\ - \\MgFe\\ distribution.  Their properties are consistent with being old objects (some scatter in the age is perhaps possible) that are able to exhaust the star forming process at very early epochs with no need for later episodes of stellar activity. In contrast, galaxies of lower velocity dispersion have more contrasting properties. It seems as if in these systems star formation, following the initial activity most likely at epochs as old as in other galaxies, later went through a series of episodes taking place in different epochs that vary from galaxy to galaxy. Equivalently, it could be that star formation continued perhaps at minimal levels over significantly longer periods of time. This would mimic a sort of age  sequence. }  \\item{The analysis of the gradients in \\Hbeta\\ and \\MgFe\\ within individual galaxies has revealed that for the majority of these the nucleus is younger and more metal-rich than the peripheral regions, thus resembling a out-inward mechanism of galaxy formation, and that the global process of star formation may last very long, with duration varying from galaxy to galaxy. It seems that the duration of the star formation activity is inversely proportional to the galactic velocity dispersion and perhaps mass. There are two less numerous subgroups for which the nucleus is older and more metal-rich or older and less-metal rich than the periphery of the galaxy. The interpretation of these structure is not straightforward and it is left to future investigation.}  \\item{The fact that most galaxies are found in the group characterized by a younger and more metal-rich nucleus calls to mind the fact that recent bursts of star formation in the merger scenario are expected to peak in the centers of elliptical galaxies (Alfensleben \\& Gerhard 1994). However, an internal cause related to the building up process of galaxy formation cannot be excluded as perhaps suggest by the correlation between velocity dispersion and $\\Delta t_{NW}$. }  \\item{The idea that the overall duration of the star forming activity is inversely proportional to the velocity dispersion and perhaps  mass of galaxies does not contradict the current information of abundance ratios in elliptical galaxies  (Carollo et al. 1993, Carollo \\& Danziger 1994, Matteucci 1994). Indeed it would be consistent with the expected trend for the ratio [Mg/Fe] inferred from the narrow band indices (Faber et al. 1992; Worthey et al. 1992; Davies et al. 1993). From the comparison of the observed indices with model indices, assuming solar partition of elemental abundances, it is concluded that the average [Mg/Fe] in giant elliptical galaxies exceeds that of the most metal rich stars in the solar vicinity by about 0.2-0.3 dex. Furthermore, the ratio [Mg/Fe] is expected to increase with the galactic mass up to the this value. According to Matteucci (1994) a possible explanation of this trend is that the efficiency of star formation is an increasing function of galactic mass. As a consequence of this, massive ellipticals should have formed on shorter time scales than smaller ellipticals. In brief, the constant super-solar value for giant ellipticals hints that only an unique source of nucleosynthesis is contributing to chemical enrichment. This is just the case of massive stars exploding as type II supernovae and thus releasing elemental species (Mg and Fe in particular) in fixed proportions. This implies a short time scale of star formation, i.e. of about 0.1 Gyr. The trend decreasing with galactic mass means that another source of Fe intervenes altering the [Mg/Fe] ratio. It is worth recalling that Mg is produced by massive stars. Therefore, it is the contribution to Fe that must change. The goal is achieved invoking type I supernovae generated by mass accreting white dwarfs in binary systems. The minimum time scale for their formation is longer than 0.1 Gyr. This implies a longer duration of the star forming period. In this context, we would like to recall that in order to match the CMR with our infall  models we had to increase the efficiency parameter $\\nu$ by a factor of 12 as the galactic mass increases from 0.01 to 3 \\M12. Furthermore, the duration of the star forming activity determined by the onset of galactic winds was found to decrease with the galactic mass (see $t_{gw}$ in Table~\\ref{tab_mass}). The trend is small but in the right direction. } \\end{itemize}  \\oneskip"
+    },
+    "9602/astro-ph9602026_arXiv.txt": {
+        "abstract": "We present results from high resolution (R$\\simeq 28000$) spectra of six high-redshift QSOs taken at the ESO NTT telescope that allow the detailed study of the Lyman-$\\alpha$ population in the redshift interval $z=2.8-4.1$.  The typical Doppler parameters found for the Lyman-$\\alpha$ lines lie in the interval $b = 20 \\div 30$ km s$^{-1}$, corresponding to temperatures $T > 24000$ K, with a fraction of the order 15\\% in the range $ 10 \\le b \\le 20$ km s$^{-1}$.  These values are still consistent with models of low density, highly ionized clouds.  The observed redshift and column density distributions obtained from these spectra and from the observations of 4 additional QSOs taken in the literature allow an accurate estimate of the proximity effect from a relatively large Lyman-$\\alpha$ sample (more than 1100 lines with $\\log N_{HI}\\geq 13.3$) in the redshift interval $z=1.7-4.1$.  A Maximum Likelihood analysis has been applied to estimate {\\it simultaneously} the best fit parameters of the Lyman-$\\alpha$ statistics {\\it and} of the UV background.  After correcting for the blanketing of weak lines, we confirm that the column density distribution is best represented by a double power-law with a break at $\\log N_{HI}\\simeq 14$, with a slope $\\beta_s = 1.8$ for higher column densities and a flatter slope $\\beta_f = 1.4$ below the break.  A value $J_{LL} = 5\\pm 1 \\times 10^{-22}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$ is derived for the UV background in the redshift interval $z=1.7-4.1$, consistent with the predicted QSO contribution. No evidence is found for redshift evolution of the UVB in the same redshift interval.  The comoving volume density distributions of protogalactic damped systems, Lyman Limit systems and Lyman-$\\alpha$ clouds with $\\log N_{HI}\\magcir 14$ and radii $R\\simeq 200$ kpc are found to be similar, suggesting a possible common association with galaxies. ",
+        "introduction": "The statistics of the Lyman-$\\alpha$ absorptions, seen in QSO spectra shortward of the Lyman-$\\alpha$ emission, provide unique information about the physical and cosmological properties of the neutral gas phase of the baryonic component of the Universe.  The high sensitivity of the observations of the Lyman-$\\alpha$ lines allows to investigate very different structures, ranging from fluctuations of the diffuse intergalactic medium to the interstellar medium in protogalactic disks.  This scientific driver coupled with the advances in instrumental and detector capabilities at 4-m class telescope has led in the last few years to a flourishing of investigations on Lyman-$\\alpha$ forest (e.g. Pettini et al. 1990, Carswell et al. 1991, Rauch et al. 1993, Giallongo et al. 1993, Fan \\& Tytler 1994, Cristiani et al. 1995). The new data have sufficient resolution (R$>$20000) and signal-to-noise ratio to use line-fitting procedures to obtain reliable estimate of the L-$\\alpha$ clouds parameters like redshift, column density $N_{HI}$ and Doppler width $b$ and a more detailed physical picture of the absorbing gas has emerged.  The average Doppler parameter is found to be $b\\mincir 30$ km s$^{-1}$, corresponding to an average temperature $T\\sim 5\\times 10^4$ K. This is typical for clouds in photoionization equilibrium with the general UV ionizing background.  However, the presence of a non negligible fraction of narrow lines with $b=10-20$ km s$^{-1}$ has been recognized in the highest resolution spectra, although the assessment of the real fraction relies on the S/N ratio of the data and on the intrinsic blending of the lines, which is strong at very high redshifts. It is difficult to reconcile, using standard photoionization models, temperatures $T\\sim 15000-20000$ K with the large cloud sizes $R\\sim 50-150h^{-1}$ kpc inferred from observations of QSO pairs (Smette et al. 1992, 1995; Bechtold et al.  1994; Dinshaw et al. 1995). This has led to the introduction of particular cooling mechanisms (Giallongo \\& Petitjean 1994).  The Lyman-$\\alpha$ lines show an appreciable redshift evolution in the redshift range $z=1.7-4$ and small clustering, limited at $\\Delta v<300$ km s$^{-1}$ for lines with $\\log N_{HI}\\magcir 13.8$ (Cristiani et al. 1995).  Attempts have been done to explain these cosmological properties in the framework of the CDM cosmology (e.g. Miralda-Escud\\'e \\& Rees 1994, Cen et al. 1994).  The introduction of the HIRES echelle spectrograph at the Keck 10m telescope has given new impulse to the observations.  Recent spectra with HIRES have revealed the presence of CIV absorption associated with Lyman-$\\alpha$ lines with $\\log N_{HI} \\magcir 14$ (Cowie et al. 1995, Tytler \\& Fan 1995). The derived abundances, although dependent on the ionization state of the Lyman-$\\alpha$ clouds, seem to be similar to the ones derived for the heavy element absorptions originated in galaxy halos, suggesting a continuity in their physical properties.  The high resolution observations of the Lyman-$\\alpha$ forest provide also a powerful method to estimate the UV background at high redshifts.  It is already known from low resolution Lyman-$\\alpha$ samples that the redshift evolution of the line number density, within a single spectrum, does not follow the general cosmological trend when approaching the QSO emission redshift. This ``inverse effect'' has been interpreted and modeled as a ``proximity effect'' (Weymann, Carswell, \\& Smith 1981; Tytler 1987; Bajtlik, Duncan \\& Ostriker 1988).  It consists in a reduction of the line density in the region near the QSO emission redshift due to the increase of the ionizing flux by the QSO.  In this way, absorbers near QSOs are more highly ionized than those farther away where the general UV background is the only source of ionization.  Using a simple photoionization model Bajtlik, Duncan, \\& Ostriker 1988 showed that the predicted and observed proximity effects agreed at high redshift adopting the value $J_{LL}\\sim 10^{-21}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$ for the UV background at the Lyman limit. This estimate has been confirmed within uncertainties of $\\pm 0.5$ in $\\log J_{LL}$ by Lu, Wolfe \\& Turnshek 1991. The most comprehensive analysis of the proximity effect has been published by Bechtold 1994 on the basis of a homogeneous low resolution (R$\\sim $3500) Lyman-$\\alpha$ sample derived from the spectra of 34 QSOs in the redshift range $1.6<z<4.1$. She found a value about 3 times higher, $J_{-21}=3$ and a correlation of the proximity effect on the QSO luminosities as expected from the photoionization model. However, any analysis performed at low resolution does not allow the use of the correct column density and redshift distributions for the Lyman-$\\alpha$ lines, both affecting the estimate of the UV background from the proximity effect.  Giallongo et al. (1993) and Cristiani et al. (1995) on the basis of a few quasars observed at high resolution (R$\\magcir 25000$), suggest a rather flat power-law column density distribution for lines with $\\log N_{HI}<14$ and a value of the UV background sensibly smaller than the reference value $\\log J_{LL}=-21$ at $z\\sim 3$, even if the uncertainties were relatively large.   We discuss in this paper a larger sample including more than 1100 Lyman-$\\alpha$ lines with column densities $\\log N_{HI}\\geq 13.3$, observed at an average R$\\sim 25000$, and perform a Maximum Likelihood analysis to estimate {\\it simultaneously} the parameters of the Lyman-$\\alpha$ line distributions {\\it and} the intensity of the UV background in the redshift interval $1.7<z<4.1$. In particular, in section 2.1 our data sample is defined, in section 2.2 the $b-\\log N_{HI}<14$ distribution and its implication on the temperature of the Lyman-$\\alpha$ clouds are analysed. In section 2.3 and 2.4 we discuss the H I column density and redshift distribution respectively. Section 3 gives the details of the computation and the results on the UV background. The conclusions of our work are listed in section 4.  We adopt throughout the value H$_o = 50$ km s$^{-1}$ Mpc$^{-1}$ for the Hubble parameter and $q_o = 0.5$. ",
+        "conclusions": "A sample of absorption lines has been extracted from the spectra of 6 QSOs taken at an average resolution of 11 km s$^{-1}$ over the absorption redshift range $z=2.8-4.1$.  Merging these data with those of 4 other QSOs with observations of similar quality available in the literature has provided a relatively large Lyman-$\\alpha$ sample (more than 1100 lines), to which a Maximum Likelihood analysis has been applied to estimate simultaneously the best fit parameters of the Lyman-$\\alpha$ clouds statistics and the intensity of the UVB background in the redshift interval $z=1.7-4.1$.  The main results of our study are: \\begin{enumerate} \\item From the subsample of Lyman-$\\alpha$ lines from the spectra at higher S/N we derive, in the range $z=3-4$, a typical Doppler parameter in the range $b = 20 \\div 30$ km s$^{-1}$ and estimate a fraction $\\sim 15-18$\\% of lines with $10 \\leq b\\leq 20$ km s$^{-1}$, corresponding to average temperatures $T\\sim 15000$ K, still consistent with models of low density, highly ionized clouds . \\item We confirm the presence of a break in the column density distribution at $\\log N_{HI}=14$ with a flat slope $\\beta _f\\simeq 1.4$ for the lower column densities and a steep one, $\\beta _s \\simeq 1.8$, for the higher column densities. \\item An interesting similarity is revealed between the volume density distributions of Lyman Limit systems, Damped systems and Lyman-$\\alpha$ clouds, as derived on the basis of their estimated sizes and geometries.  This suggests a physical association between the Lyman-$\\alpha$ clouds with $\\log N_{HI} \\magcir 14$ and the halos of protogalactic systems. \\item The intensity of the UV background $J$ as derived from the proximity effect is a factor of 6 lower than previously estimated, with the same photoionization model, from data at lower resolution, for which the blanketing effects are stronger.  The value $J_{-22}\\simeq 5\\pm 1$ we have found is consistent with the one predicted for QSOs after allowing for some dust obscuration of the QSO population. \\item No evidence is found, within relatively large uncertainties, for evolution in redshift of the UVB in the range $z=1.7-4.1$. \\end{enumerate}"
+    },
+    "9602/astro-ph9602105_arXiv.txt": {
+        "abstract": "We report on the detection of dark matter in the cluster Abell~2390 using the weak gravitational distortion of background galaxies. We find that the cluster light and total mass distributions are quite similar over an angular scale of $\\simeq 7^\\prime \\;(1 \\Mpc$). The cluster galaxy and mass distributions are centered on the cluster cD galaxy and exhibit elliptical isocontours in the central $\\simeq 2^\\prime \\; (280 \\kpc)$. The major axis of the ellipticity is aligned with the direction defined by the cluster cD and a ``straight arc'' located $\\simeq 38^{\\prime\\prime}$ to the northwest. We determined the radial mass-to-light profile for this cluster and found a constant value of  $(320 \\pm 90) h\\; M_\\odot/L_{\\odot V}$, which is consistent with other published determinations. We also compared our weak lensing azimuthally averaged radial mass profile with a spherical mass model proposed by the CNOC group on the basis of their detailed dynamical study of the cluster. We find good agreement between the two profiles, although there are weak indications that the CNOC density profile may be falling more steeply for $\\theta\\geq 3^\\prime$ $(420\\kpc)$. ",
+        "introduction": "Understanding the nature of the mass distribution in clusters of galaxies is a key cosmological issue. Typically, three independent methods to probe the cluster mass distribution have been applied: 1) applications of the virial theorem to the cluster galaxy velocity distribution, 2) analyses of the X-ray emission of the diffuse, intracluster gas and 3) mass determinations based on the gravitational lens distortion of background objects by the foreground cluster.  The latter technique, the analyses of both the strong as well as the weak gravitationally induced distortions in the images of faint background galaxies, is particularly attractive as it does not require any assumptions about the geometry or the dynamical of the clusters, with the weak distortions being especially suited for mapping the cluster mass distribution out to large radii.  Over the past few years, there has been considerable activity to develop a theoretical framework for determining the mass distribution in the cluster based on a measured set of gravitationally distorted galaxy shapes (eg. \\cite{tyson90}; \\cite{ks93}; \\cite{schneider94}; \\cite{kaiser95b}; \\cite{schneider95}; \\cite{seitzc95};  \\cite{seitzs95}, \\cite{sk95}). Several groups have developed algorithms and software designed to measure and analyze very small distortions expected in the cluster outskirts (\\cite{ksb95}; \\cite{bonnet95}; \\cite{fischer96}). These techniques have been applied to the central $\\simeq 1 \\Mpc$ of several galaxy clusters (eg. \\cite{tyson90}; \\cite{fahlman94}; \\cite{bonnet94}; \\cite{tyson95}; \\cite{squires95}).  For a few clusters, the results of the weak lensing analyses has compared with the results of X-ray analyses as well as detailed spectroscopic studies. There is no {\\it apriori} reason to expect an agreement between the different analyses and in fact, such joint analyses can potentially enable to us learn much more about the relative distributions and the dynamical states of the gas, galaxies and dark matter in the clusters than is currently known.  Indeed, in the central regions of some clusters (eg. A2218 and A1689), the mass inferred {}from the presence of (strong) lensing features exceeds that derived {}from the X-ray data under the standard assumption of pressure supported hydrostatic equilibrium by a factor of $\\simeq 2$--$3$ (\\cite{meb95}).  A similar discrepancy also arises in galaxy cluster MS~1224, where the weak lensing inferred mass is $\\simeq 2$ times larger than the mass determined {}from a virial analysis of the cluster galaxies (\\cite{fahlman94}; \\cite{carlberg94}). Also, in a joint lensing/X-ray analysis of A2218  (\\cite{squires95}), the lensing mass profile was found to be systematically a factor of $\\simeq 2$ larger than the X-ray mass profile, although the two are formally in agreement within the 95\\% confidence level. A recent analysis based on ASCA data confirmed this discrepancy (\\cite{loewenstein96}). Conversely, a similar analysis of the cluster A2163 yielded good agreement between the weak lensing and X-ray mass estimates (\\cite{squires96}).  In this paper, we discuss the dark matter and galaxy distributions in the cluster of galaxies Abell~2390. Abell 2390 is an Abell richness class 1 cluster (\\cite{abell89}) at redshift $z = 0.2279$ with a measured line of sight velocity dispersion $\\sigma = 1093$~km/s (\\cite{carlberg95}). Abell~2390 has an X-ray luminosity of $L_x = 9.77 \\times 10^{44}$~ergs/s (H~=~75~km/s/Mpc) in the 2-6~keV energy range (\\cite{kowalski84}) and $L_x = 4.67 \\times 10^{44}$~ergs/s in the 0.7-3.5~keV range (\\cite{ulmer86}). A2390 has been the focus of several studies (eg. \\cite{pello91}, \\cite{kassiola92}, \\cite{narasimha93}) seeking to account for the observed lensing features such as the ``straight arc'' located approximately 38$^{\\prime\\prime}$ away {}from the central galaxy (\\cite{pello91}).  In this analysis, the total mass distribution in A2390 is determined using the weak gravitational distortion of faint background images induced by the cluster. The data reduction and analysis procedure used here is very similar to that we employed in our analysis of A2218 (\\cite{squires95}) and A2163 (\\cite{squires96}) and we refer to those papers for further detail. Here we will concentrate mainly on the results. We also compare our results for the morphology of the  mass distribution and the estimate for the  total mass with published results {}from galaxy spectroscopy and X-ray analyses. ",
+        "conclusions": "We have mapped the light and mass distributions in the cluster of galaxies A2390 over a scale of $\\simeq 250^{\\prime\\prime}$ {}from the cluster center. The projected light distribution was determined using a color selected galaxy catalogue. The total radial light profile was calculated using all of galaxies detected, which we expect should include some contributions {}from galaxies not associated with the cluster. To map the mass distribution, we measured the shapes of faint galaxies which have been gravitationally distorted, correcting for systematic effects in the shape determinations such as psf anisotropy and seeing.  The projected light and mass distributions are displayed in Figure \\ref{fig:a2390_lightandmass}. Both the light and mass distributions are centered on the cD galaxy. The light and mass isocontours are elliptical in the central region, with the major axis running {}from the southeast to northwest. There is a smaller, secondary concentration in both the cluster galaxy and light distributions located approximately $4^\\prime$ to the northwest of the cluster center. The mass map gives the impression that the mass distribution is extended in the direction of the secondary concentration seen in the cluster light distribution but there is no corresponding peak in the mass map.  As we noted previously, this is not surprising given the number of galaxies used in constructing the mass map.  Several authors (eg. Pello \\etal 1991; Kassiola \\etal 1992) have attempted to account for the lensing features seen in A2390 by suggesting that the cluster has bi-modal mass distribution, with a small secondary mass concentration located near the observed arc position.  This model is further supported by the cluster X-ray emission, which shows elliptical isocontours centered on the cluster cD galaxy as well as a significant secondary peak located $40^{\\prime\\prime}$ to the northwest of the cluster center (\\cite{pierre95}). The weak lensing mass map presented here is not inconsistent with the notion of a bi-modal mass distribution. As we have noted previously, the small-scale results {}from the weak lensing analysis are subject to greater uncertainties arising {}from the combined effects of nonlinearity in the lensing, contamination by cluster galaxies, and the statistical uncertainties due to the small number of galaxies used in the mass reconstructions.  We calculated the mass-to-light ratio as a function of radius (Figure \\ref{fig:a2390_MassvsR}) and found a constant value of  $M/L_V = (320 \\pm 90) h  M_\\odot/L_{\\odot V}$, although we have noted that the mass estimates may be biased downwards in the innermost $\\simeq 1^\\prime$ due to contamination by faint cluster galaxies. In the course of their study of several intermediate redshift cluster, the CNOC group derived a mass-to-light ratio, K-corrected to redshift zero, of $M/L_r = (331 \\pm 54) h M_\\odot/L_{\\odot r}$ for the cluster. We convert the Gunn r luminosity into V-band luminosity by applying the transformation V-r=0.16 (which is typical for a nearby early type galaxy) and find that the CNOC result corresponds to $M/L_V = (370 \\pm 60) h  M_\\odot/L_{\\odot V}$, which is in excellent agreement with out determination.  We furthermore compared the radial mass profile determined {}from this analysis with the CNOC model. We projected the CNOC 3D mass model and calculated the 2D radial mass profile, correcting for mass in the control annulus to facilitate direct comparison with the lensing results (Figure \\ref{fig:a2390_MassvsR}). All of the mass estimates based on the lensing analysis agree within the 95\\% CL with the CNOC model (although we note that the CNOC model assumes spherical symmetry and there is some indications that this might be invalid on the scales probed here). At the very largest radius, there is a slight indication that the lensing inferred mass might be falling slower than the CNOC model, although the values are consistent within $2 \\sigma$. Larger scale observation of this cluster will prove useful to determine if this trend continues at large radii.  As with many of the weak lensing studies performed to date, this analysis was confined to a relatively small field of view. This has the immediate consequence that, since the mass determinations made by this technique are differential, the masses quoted form a {\\em lower} bound on the total mass at any radius. The future prospects for this method are exciting as the mass distributions are mapped out to large distances in the cluster halos. 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. With the types of observations possible with these instruments, the weak lensing technique offers a unique opportunity to probe  many of the outstanding puzzles regarding the dark matter content and distribution in the universe.  \\clearpage"
+    },
+    "9602/astro-ph9602057_arXiv.txt": {
+        "abstract": "The globular cluster M79 was observed with the Hopkins Ultraviolet Telescope (HUT) during the Astro-1 space shuttle mission in 1990 December. The cluster's far-UV integrated spectrum shows strong absorption in the Lyman lines of atomic hydrogen. We seek to use this spectrum, together with optical photometry, to constrain the stellar mass distribution along its zero-age horizontal branch (ZAHB). We find that a Gaussian distribution of ZAHB masses, with a mean of $0.59 M_{\\sun}$ and standard deviation $0.05 M_{\\sun}$, is able to reproduce the cluster's $(B,V)$ color-magnitude diagram when subsequent stellar evolution is taken into account, but cannot reproduce the cluster's far-UV spectrum.  Model stellar spectra fit directly to the HUT data indicate a surprising distribution of atmospheric parameters, with surface gravities (and thus implied masses) significantly lower than are predicted by canonical HB evolutionary models. This result is consistent with the findings of Moehler et al. [A\\&A, 294, 65 (1995)] for individual HB stars in M15. Further progress in understanding the mass distribution of the HB must await resolution of the inconsistencies between the derived stellar atmospheric parameters and the predictions of HB evolutionary models.  Improved stellar spectral models, with higher spectral resolution and non-solar abundance ratios, may prove useful in this endeavor. ",
+        "introduction": "M79 (NGC 1904) is a centrally condensed, intermediate-metallicity ([Fe/H] = $-1.69$) globular cluster \\shortcite{DK86,Djorgovski93} with an extremely blue horizontal branch \\shortcite{SH77}, extending at its high-temperature end as faint as the main-sequence turnoff \\shortcite{Ferraro92}. The cluster's integrated UV spectrum is nearly flat from 1500 to 3300 \\AA\\ \\shortcite{vAlbada81}, and observations with the Ultraviolet Imaging Telescope (UIT; \\shortciteNP{UIT_M79}) indicate that most of the far-UV flux emanates from individual blue HB (BHB) and post-HB stars.  In our current understanding \\shortcite{Renzini81p319}, stars with zero-age-main-sequence mass $M_{ZAMS} \\sim M_{\\sun}$ lose about $0.2 M_{\\sun}$ while ascending the red-giant branch (RGB) and eventually produce zero-age horizontal branch (ZAHB) stars with core masses $M_C \\sim 0.5 M_{\\sun}$ and envelope masses $M_{env} \\sim 0.1 M_{\\sun}$. A star's location on the ZAHB is determined by the mass of the hydrogen-rich envelope remaining above the helium-rich core, which for typical globular cluster ages is a function only of the abundances of helium and heavier elements $(Y,Z)$. Stellar evolutionary models imply a spread in $M_{env}$, with the bluest HB stars having the lowest masses \\shortcite{Dorman93}.  The mass distributions used to derive gross cluster properties from the HB population are most often based on the function introduced by \\shortciteN{Rood73}. This function consists of a Gaussian multiplied by polynomial truncation terms designed to ensure that the simulation contains physically possible objects, \\ie, stars with $M_C \\le M \\le M_{\\rm RG}$. Otherwise, the assumption is not based on any physical understanding of the cause of the HB mass spread \\shortcite{Rood90p11}. Monte-Carlo simulations based on such syntheses have been quite successful in reproducing observed HB morphologies (\\shortciteNP{Rood73}; \\shortciteNP{LDZ90}, 1994; \\shortciteNP{Cat93}); however, the polynomial truncation terms give distributions that are skewed away from the theoretical limits of the mass range, a situation that may not be desirable for the extremely blue HB of M79.  In this paper we analyze the far-UV spectrum of M79 obtained with the Hopkins Ultraviolet Telescope (HUT). Our object is to derive constraints on the ZAHB mass distribution from our far-UV spectrum and from optical photometry. \\shortciteN{Ferraro92} note that the distribution of stars in the color-magnitude diagram (CMD) of M79 seems quite different from a Gaussian. We investigate whether a Gaussian distribution of masses along the ZAHB can reproduce this asymmetry by combining the HB evolutionary models of \\shortciteN{Dorman93} with the synthetic stellar flux distributions of \\shortciteN{Kurucz92} to fit the cluster's optical CMD \\shortcite{Ferraro92}. We find that a Gaussian distribution of ZAHB masses with a mean of $0.59 M_{\\sun}$ and standard deviation $0.05 M_{\\sun}$ can indeed reproduce the cluster's CMD. We use this model to compute synthetic far-UV spectra using a Monte Carlo algorithm, compare them with the HUT spectrum, and find that it cannot reproduce the observed far-UV flux distribution. Using individual Kurucz model flux distributions, we derive surface gravities for BHB stars that are significantly lower than theoretical predictions.  This result is consistent with that of \\shortciteN{MHd95} for individual HB stars in M15. We conclude that more precise constraints on the HB mass distribution must await a resolution of this discrepancy. ",
+        "conclusions": "We have shown that the apparently asymmetric distribution of HB stars in M79 observed by \\shortciteN{Ferraro92} can be reproduced with a ZAHB mass-distribution function in the form of a Gaussian with polynomial truncation terms if stellar evolution is properly taken into account.  The function that best fits the HB of M79 has a mean mass of $0.59 M_{\\sun}$ and a standard deviation of $0.05 M_{\\sun}$. This study represents one of the first detailed investigations into the underlying stellar mass distribution on the HB.  The synthetic HB distribution that best fits the optical CMD is, however, unable to reproduce the HUT far-UV spectrum of the cluster. \\shortciteN{Kurucz92} models fit directly to the HUT spectrum have values of $\\log g$---and thus stellar masses---significantly lower than are predicted by canonical HB evolutionary theory. This result is consistent with the recent report by \\shortciteN{MHd95} that \\shortciteN{Kurucz92} stellar atmosphere models yield surface gravities for BHB stars systematically lower than are predicted by the evolutionary tracks of \\shortciteN{DLV91}. \\shortciteANP{Kurucz92} models with $\\mh = -1.0$ provide significantly better fits to the cluster spectrum than do those with $\\mh = -1.5$. These discrepancies may reflect non-solar abundances ratios in the HB stars of M79 and/or problems in the treatment of hydrogen absorption by the Kurucz model atmospheres. The deficiencies are made obvious here by the high signal-to-noise ratio of the HUT spectra and the reliability of their calibration.  Stars with 27~000 K $\\lesssim T_{eff} \\lesssim$ 32~000 K are required to reproduce the observed flux shortward of Ly$\\gamma$. Thus, at least a few such objects must reside within about 1\\arcmin\\ of the cluster center, and should be detectable using WFPC2 on \\hst. We see no evidence, however, of a large population of previously unobserved extreme HB stars."
+    },
+    "9602/astro-ph9602130_arXiv.txt": {
+        "abstract": "Is pulsar make up of strange matter? The magnetic field decay of a pulsar may be able to give us an answer. Since Cooper pairing of quarks occurs inside a sufficiently cold strange star, the strange stellar core is superconducting. In order to compensate the effect of rotation, different superconducting species inside a rotating strange star try to set up different values of London fields. Thus, we have a frustrated system. Using Ginzburg-Landau formalism, I solved the problem of rotating a superconducting strange star: Instead of setting up a global London field, vortex bundles carrying localized magnetic fields are formed. Moreover, the number density of vortex bundles is directly proportional to the angular speed of the star. Since it is energetically favorable for the vortex bundles to pin to magnetic flux tubes, the rotational dynamics and magnetic evolution of a strange star are coupled together, leading to the magnetic flux expulsion as the star slows down. I investigate this effect numerically and find that the characteristic field decay time is much less than 20~Myr in all reasonable parameter region. On the other hand, the characteristic magnetic field decay time for pulsars is $\\geq 20$~Myr. Thus, my finding cast doubt on the hypothesis that pulsars are strange stars. ",
+        "introduction": "} What is the most stable form of baryonic matter at high density? This question is of both observational and theoretical interests. With the discovery of pulsars and their identification with the neutron stars in the late 60s, many people thought that neutron star matter is the most stable form of cold condensed matter at high density. This believe was later challenged by Witten. He proposed that at sufficiently high density, deconfinement of nucleons occurs. Moreover, some of the d and u quarks on of the Fermi surface is converted to strange quarks via weak interaction. Thus, the most stable form of matter at ultra-high density is a degenerate mixture of u, d, and s quarks together with a small amount of electrons so as to maintain overall charge neutrality (\\cite{Witten}). This kind of material state is called strange matter. His idea was examined in detail later by Farhi and Jaffe (1984). \\par The possibility of self gravitating strange matter, namely, a strange star, has also been investigated. The equation of state (EOS) calculations of strange star suggests that they are indeed stable objects in certain parameter range (\\cite{Alcock86}; \\cite{Haensel86}; \\cite{Benvenuto89}; see also \\cite{Proc91}; and \\cite{Strange_Pulsar_EOS}). The mass and radius of a stable strange star are found to be similar to those of a neutron star. Consequently, some authors suggested that pulsars are in fact strange stars (see for example \\cite{Strange_Pulsar_EOS}). In principle, the question of whether pulsars are neutron or strange stars can be answered by comparing the core density of a neutron star to the strange matter transition density. Unfortunately, all the nuclear physics calculations to date do not yield a definitive answer. \\par There are a number of difficulties in explaining some of the observational properties of pulsars using the strange star model. The first, and perhaps the most serious difficulty, is the possibility of a strange star glitch. Early EOS calculations suggested that strange stars have very thin nuclear matter crust. The ratio of inertial moment of the nuclear matter crust to that of the whole strange star is of the order of $10^{-5}$. Moreover, the density of the nuclear matter crust is not high enough for neutron drip (\\cite{Alcock86}; \\cite{Alcock91}). In contrast, pulsar glitch observations tell us that the ratio of inertial moment of pinned neutron superfluid crust to that of the whole star is about $10^{-2}$ (\\cite{Glitch_Comment}; \\cite{Link}; \\cite{VCM}). Nevertheless, by taking into account the existence of strange matter bound states, Benvenuto et al. (1991ab) showed that strange star may support a ``strange matter crust'' thick enough to account for the observed pulsar glitches. Besides, a more recent calculation suggested that a strange star may be able to retain a reasonable size of nuclear matter crust by accretion as well (\\cite{Benvenuto94}). \\par The second difficulty is the possibility of type~I X-ray bursts on a strange star surface. It is commonly believed that type~I X-ray burst involves sudden thermo-nuclear runaway at the surface of an accreting neutron star (see for example \\cite{X_Burst}). The same bursting behavior cannot occur on a bare strange star surface because nuclear fuel will dissociate into constituent quarks immediately (\\cite{Jones86}). The possibility of X-ray burst,therefore, requires the existence of a nuclear matter crust over a strange star. EOS calculations showed that the nuclear matter crust is separated from the interior strange matter by a electrostatic gap of a few fermi thick, thereby preventing the nuclear matter from converting into strange matter (\\cite{Alcock86}; \\cite{Curst_SP}). A strange pulsar magnetosphere can also be formed, giving rise to the observed radio pulsation (\\cite{Alcock86}; \\cite{Strange_Pulsar_EOS}). \\par The final difficulty is the magnetic field decay time. Assuming a {\\em normal} strange matter core, Jones (1988) pointed that magnetic field inside a strange star will not decay in Hubble's time. Moreover, there is no obvious mechanism for the star to retain a small residual field after say $10^9$ years. Therefore, it is inconsistent with both the pulsar magnetic field decay hypothesis, and the spun-up formation scenario of milli-second pulsars (\\cite{Jones88}). Nevertheless, the statistical evidence for pulsar magnetic field decay over a period of some $10^9$~yr is still inconclusive. Also, milli-second pulsars may rather form by accretion induced collapse of white dwarfs (see for example \\cite{Field_Decay} and \\cite{Decay4} for discussions on neutron star magnetic field decay). Therefore, the objection of Jones may not be completely well founded. Moreover, as I shall discuss in \\S\\ref{S:SM}, the core of a strange star is likely to be superconducting. In this case, the magnetic field decay time estimated by Jones (1988) has to be modified. The question of magnetic field decay have been brought up in the review paper by Bailin and Love (1984). They briefly mentioned that the low electron density in the strange stellar core may lead to rapid flux expulsion. However, they did not provide a detail calculation. A major objective of this paper, therefore, is to perform such a calculation, taking into account the coupling between magnetic evolution and rotational dynamics of the star in the presence of collective effects such as clumping of flux tubes. \\par At this moment, one has no doubt that the standard neutron star model for pulsars is well tested by numerous observational data. However, we still have to keep a close look at the alternative hypothesis, namely, the strange pulsar model (\\cite{Strange_Pulsar_EOS}). In particular, we like to explore various physical properties of a neutron star and a strange star, and to see if there are further ways to test if which of the two models is a better candidate in explaining pulsar observations. \\par As I shall discuss in \\S\\ref{S:SM}, the core of a strange star is likely to be superconducting (\\cite[1984]{Bailin82}). However, this cast a problem on the rotation of a strange star. In order to rotate an object with one superconducting species, a uniform magnetic field \\begin{equation} {\\bf B} = \\frac{2m^* c}{q^*}\\,{\\bf\\Omega} \\mbox{~,} \\label{E:London} \\end{equation} called the London field, is set up in the object (\\cite{London_Theo}), where $m^*$ and $q^*$ are the effective mass and charge of the Cooper pair, and $\\Omega$ is the rotational angular velocity of the object. But for a strange star, u, d, and s quarks are all superconducting. These superconducting species require different values of $B$ to set up a rotation. Thus, we have a {\\em frustrated} system. The situation is further complicated by the fact that all three superconducting quark flavors interact strongly with each other. In \\S\\ref{S:Rotate}, I tackle this rotation problem using Ginzburg-Landau formalism. Instead of setting up a uniform London field, I find that vortices and localized magnetic fields are created when the superconducting quarks rotate together. Then in \\S\\ref{S:Ap}, I point out that the vortices will inter-pin with the magnetic flux tubes, which alter the magnetic field evolution and rotation of a strange star dramatically. In particular, it suggests that the (superconducting) strange pulsar hypothesis is inconsistent with the observed magnetic field decay time in pulsars. Finally, a conclusion is drawn in \\S\\ref{S:Conclude}. ",
+        "conclusions": "} In summary, I consider the rotation of a multi-superconducting species object. The rotation of such an object is made possible by the formation of vortices similar to that of superfluid helium. This finding implies the existence of vortex bundles in the core of a rotating superconducting strange star. Because it is energetically favorable for the vortex bundles to pin to magnetic flux tubes, the rotational dynamics and the magnetic field evolution of a strange star are coupled. \\par Because the nuclear crust of a strange star cannot sustain a strong magnetic stress, the magnetic field evolution of the star is dictated by dynamics of the flux tubes in the core. The core magnetic field evolution due to inter-pinning of magnetic flux tubes and vortex bundles, and clumping of flux tubes is computed numerically. I find that in all reasonable parameter range, the characteristic decay time of the magnetic field is $\\leq 20$~Myr. This finding does not agree with the hypothesis that pulsars are superconducting strange stars because current pulsar data strongly suggest that pulsar magnetic fields do not decay in 20~Myr. \\par A number of interesting questions remains. The phase diagram of interacting multiple component superconducting species in rotation remains unclear. It is also interesting to study the effects of alignment torque on magnetic and rotational history of a strange star, and the possible collapse of the nuclear matter crust due to magnetic stress. I plan to report them in my future works."
+    },
+    "9602/astro-ph9602076_arXiv.txt": {
+        "abstract": "Analytical expressions for the rates of formation and destruction of gravitationally bound systems are derived assuming that they are originated from primordial random-Gaussian density fluctuations.  The resulting formulae reproduce the time derivative of the Press-Schechter mass function in a certain limit.  Combining a theoretical model for the evolution of structures with the formation rate, we can make various cosmological predictions which are to be compared with observations.  As an example to elucidate such applicability, we evaluate the contribution of clusters of galaxies to the cosmic X-ray background radiation. With the {\\it COBE} normalization, we find that the significant fraction of the observed soft X-ray background is accounted for by clusters of galaxies in a cold dark matter universe with $\\Omega_0 \\sim 0.2$, $\\lambda_0 =1- \\Omega_0$ and $h \\sim 0.8$. ",
+        "introduction": "Formation, merging, and destruction history of dark matter haloes of astrophysical objects is of fundamental importance in describing the subsequent formation and evolution of luminous objects such as quasars, galaxies, and clusters of galaxies. In addition, early formation of structures may affect the growth of objects through the reionization of the entire universe (Sasaki, Takahara, \\& Suto 1993; Sasaki \\& Takahara 1994; Fukugita \\& Kawasaki 1994).  Key quantities in this context are the rates of formation and destruction of gravitationally bound systems. They are defined as the comoving number density of bound systems of a given mass that are formed or destroyed in unit time at a given epoch.  A number of authors have attempted to compute this quantity either numerically or analytically. A straightforward approach is to use a state-of-art numerical simulations (e.g., Evrard 1990; Cen \\& Ostriker 1992; Bryan et al. 1994; Suginohara 1994). Although the dynamical range and small-scale resolution of available codes are rapidly increasing, it is not clear at present to what extent they unambiguously describe small-scale dynamics. Besides it is not easy to grasp the underlying key physics from the results of numerical simulations. It is therefore important to develop an analytical model on the basis of simple physics.  The alternative approach in such a direction is closely related to the theory of Press \\& Schechter (1974; hereafter PS). PS derived a mass function, the number density as a function of mass and epoch, of virialized systems.  The PS mass function, however, is not directly related to the rate of formation mentioned above.  In fact, the time derivative of the PS mass function can become either positive or negative. Thus it can not be simply interpreted as the rate of formation, but rather corresponds to the balance between the rates of formation and destruction (Cavaliere, Colafrancesco \\& Scaramella 1991; Blain \\& Longair 1993a,b).  Although several authors have adopted the time derivative of the PS mass function to approximate the formation rate of bound systems (e.g., Efstathiou \\& Rees 1988; Fukugita \\& Kawasaki 1994), its validity should be examined carefully. Blain \\& Longair (1993a,b) and Sasaki (1994) correctly realized the above problem, and attempted to obtain the rates of formation and destruction separately, assuming some empirical forms for the rates. In this paper, we derive a more general formula for each rate on the basis of merger probabilities discussed by Bower (1991) and Lacey \\& Cole (1993; hereafter LC) (See also Kauffmann \\& White 1993).  Our formalism is applicable to a number of problems regarding the hierarchical clustering in the universe. In what follows we further apply the resultant formulae to the formation of clusters of galaxies and calculate their contribution to the observed cosmic X-ray background (XRB). The cluster contribution is in general supposed to be very small in the high energy band above a few keV, while it is likely to dominate in the soft energy band (Silk \\& Tarter 1973; Schaeffer \\& Silk 1988). It is, however, still uncertain to what extent the observed XRB can be contributed by emissions from clusters of galaxies. Several authors have employed the PS mass function to evaluate the cluster contribution (Evrard \\& Henry 1991; Blanchard et al. 1992; Burg, Cavaliere \\& Menci 1993), and thus formation epochs and the subsequent evolution of virialized clusters are not well specified in their treatment.  Here we present an approach which explicitly takes account of the formation epoch of clusters and their luminosity evolution.  Our predictions are made specifically in cold dark matter (CDM) universes with/without the cosmological constant whose fluctuation spectra are properly normalized by the {\\it COBE} data (Sugiyama 1995).  The results are also compared with those of the hydrodynamical simulation and the {\\it ASCA} observations.  The plan of this paper is as follows.  In Section 2, we outline the PS formalism and discuss how the time derivative of its mass function is related to the formation and destruction rates of bound systems. Then using the merger probabilities of virialized haloes deduced by Bower (1991) and LC, we derive expressions for the rates of formation and destruction. These expressions are further employed to describe the distribution of formation epoch of bound systems.  In Section 3, we apply our formulae and calculate the contribution of clusters of galaxies to the XRB.  Finally Section 4 is devoted to our conclusions and further discussion. ",
+        "conclusions": "We have extended the theory of Press \\& Schechter (1974) and the formulation of Bower (1991) and Lacey \\& Cole (1993) to derive expressions for the rates of formation and destruction of gravitationally bound systems.  Although the formal expressions of these quantities diverge, they properly reproduce the time derivative of the PS mass function. The divergence is not unphysical but can be interpreted as being originated from the somewhat improper definition of formation and destruction of bound objects.  We have proposed a phenomenological prescription to remove such divergences by specifying a realistic mass range in which the identity of an object is to be maintained. The resulting expressions are practically insensitive to the choice of the threshold masses. On the basis of this prescription, we have obtained a distribution function of formation epochs of bound systems.  This quantity is of great cosmological importance because it provides a useful theoretical tool in describing the subsequent evolution of the systems.  In order to exhibit the applicability of the above formalism, we have evaluated the contribution of clusters of galaxies to the XRB, as well as the X-ray cluster temperature and luminosity functions. Changes from the previous PS formalism are obvious especially when the cluster evolution is strong and the X-ray properties depend largely on its formation epoch. The results of our analytical procedure, although fairly phenomenological in modelling clusters' X-ray gas, are reasonably supported by those of the three-dimensional hydrodynamical simulation by Kang et al. (1994).  Thus these two approaches are complementary in investigating the evolution history of virialized structures.  We have shown that the significant fraction of the observed soft XRB can be accounted for by clusters of galaxies in a spatially flat CDM universe with $\\Omega_0 \\sim 0.2$ and $\\lambda_0=1-\\Omega_0$ with the {\\it COBE} DMR two year normalization. The $\\Omega_0=1$ CDM model tends to overproduce the observed X-ray intensity.  On the other hand, the cluster contribution becomes negligibly small in an open CDM universe with $\\Omega_0 \\sim 0.2$, in which case observed XRB should be explained entirely by other classes of sources such as quasars. Although our overall conclusion in the $\\Omega_0=1$ case agrees with Burg et al. (1993), we derived new constraints on the other models from the XRB in a consistent and clear manner.  Note that the model with $\\Omega_0 \\sim 0.2$, $\\lambda_0=1-\\Omega_0$ and $h\\sim 0.8$ is a specific example of the most successful cosmological models explaining the cosmic age, two-point correlation functions of galaxies and clusters, and small-scale velocity dispersions (Efstathiou et al. 1990; Suginohara \\& Suto 1991; Suto 1993; Watanabe, Matsubara \\& Suto 1994; Bahcall 1994). In this context it is interesting to note that the same model can simultaneously account for a major fraction of the observed XRB in the soft energy band.  Furthermore, the microwave background distortion and the resulting anisotropy due to the distant clusters in CDM models are below the current limit (Makino \\& Suto 1993).  Obviously there remain several areas for further theoretical work. With regard to the XRB predictions, the greatest uncertainty probably lies in the evolution of intracluster medium. Although we have assumed that it basically traces the growth of the gravitational potential, this assumption is still an open question and different evolutionary models would alter the predicted luminosity functions and the XRB spectra. The contamination of the observed spectra by the Galactic emission and absorption is also uncertain. Moreover, we have omitted the line emissions which may affect the soft band intensity. Therefore we did not aim to rigorously reproduce the XRB spectra from theory, but rather to examine the possibility that clusters of galaxies play a crucial role in the origin of the XRB.  As to the formalism, on the other hand, our method contains a formal divergence for which somewhat artificial prescription has been proposed. In addition, our argument itself has further room for improvement; for instance, if $M_1$ in equation (\\ref{formdef}) does not correspond to the largest parent of the final object $M$, the rate of formation thereby defined possibly overcounts the true rate.  This problem is in fact ascribed to a fundamental limitation of the current formalism which does not take into account the history of every single object involved in a merger process.  Clearly the validity of our treatment needs to be examined through a direct comparison with other classes of approach such as N-body simulations. Other ways of formulation are also desirable. In this context we note that LC proposed a differential distribution function of halo formation times based on what they called `the halo counting argument' (see their \\S 2.5.2). As notified by LC, however, their definition leads to a slight self-inconsistency of predicting a negative probability density depending on the slope of the fluctuation spectrum.  As a matter of fact, we compared their result with ours and basically confirmed that both agree well in the mass range of astrophysical interest.  For practical purposes, therefore, our present formulae are expected to be applicable to a number of further problems in cosmology such as cluster luminosity functions in various cosmological models, formation of high redshift objects and reionization history of the universe. These issues, together with a quantitative comparison with LC, will be described elsewhere (Kitayama \\& Suto, in preparation).  \\vspace{1cm} \\noindent {\\large\\bf ACKNOWLEDGEMENTS}\\\\[4mm] \\noindent We thank Shin Sasaki, Katsuhiko Sato, Tatsushi Suginohara, Hajime Susa, and Takahiro Tanaka for helpful discussions and comments. We are grateful to Renyue Cen, Keith Gendreau and Naoshi Sugiyama for providing the relevant results of their current work which are adopted in the present paper. TK acknowledges the fellowship from the Nippon Telephone and Telegram Co. .  This research is supported in part by the Grants-in-Aid by the Ministry of Education, Science and Culture of Japan (05640312, 06233209, 07740183, 07CE2002)."
+    },
+    "9602/astro-ph9602018_arXiv.txt": {
+        "abstract": "We review recent work aimed at showing how global topological defects influence the shape of the angular power spectrum of the cosmic microwave background radiation on small scales. While Sachs--Wolfe fluctuations give the dominant contribution on angular scales larger than about a few degrees, on intermediate scales, $0.1^\\circ \\mathrel{\\raise.3ex\\hbox{$<$\\kern-.75em\\lower1ex\\hbox{$\\sim$}}} \\theta \\mathrel{\\raise.3ex\\hbox{$<$\\kern-.75em\\lower1ex\\hbox{$\\sim$}}} 2^\\circ$, the main r\\^ole is played by coherent oscillations in the baryon radiation plasma before recombination. In standard cosmological models these oscillations lead to the  `Doppler peaks' in the angular power spectrum. Inflation--based cold dark matter models predict the location of the first peak to be at $\\ell \\sim 220 / \\sqrt{\\Omega_0}$, with a height which is a few times the level of anisotropies at large scales. Here we focus on perturbations induced by global textures. We find  that  the height of the first peak is reduced and is shifted to $\\ell\\sim 350$. ",
+        "introduction": "The importance of the observations of anisotropies on intermediate and small angular scales cannot be over-emphasized. Presently a number of experiments scrutinize different regions of the sky trying to uncover the real amplitude and structure of the relic radiation. The CMB anisotropies are a powerful tool to discriminate among the different models of structure formation  by  purely linear analysis. These anisotropies are parameterized in terms of $C_\\ell$'s, defined as the coefficients in the expansion of the angular correlation function \\begin{equation} \\label{twotwo} \\langle C_2(\\vartheta) \\rangle = {1\\over 4\\pi}\\sum_\\ell(2\\ell+1)C_\\ell P_\\ell(\\cos\\vartheta). \\end{equation} For Harrison--Zel'dovich spectra of perturbations $\\ell (\\ell + 1) C_\\ell$ is constant on large angular scales, say $\\ell \\mathrel{\\raise.3ex\\hbox{$<$\\kern-.75em\\lower1ex\\hbox{$\\sim$}}} 50$. Both inflation and topological defect models lead to approximately scale invariant spectra on large scales. The next important step towards discriminating between competing models is the measurement of degree scale anisotropies. Disregarding Silk damping, gauge invariant linear perturbation analysis leads to (Durrer 1994) \\begin{equation} {\\delta T\\over T} = \\left[ - {1\\over 4}D_g^{(r)} + V_j \\gamma^j -\\Psi+\\Phi\\right]_i^f + \\int_i^f (\\Psi' - \\Phi' ) d\\tau  ~, \\label{dT} \\end{equation} where $\\Phi$ and $\\Psi$ are quantities describing the perturbations in the geometry and $\\bf V$ denotes the peculiar velocity of the radiation fluid with respect to the overall Friedmann expansion. In Eq.(\\ref{dT}), $D_g^{(r)} = \\delta\\rho_r/\\rho_r$ specifies the intrinsic density fluctuation in the radiation fluid; below we use $D_r = (\\delta\\rho_r +\\delta\\rho_b)/(\\rho_r+\\rho_b)$, about 5\\% smaller than $D_g^{(r)}$ for scales inside the horizon, which are the ones we are interested in.  The Sachs--Wolfe (SW) contribution in the above equation has been calculated for both inflation and defect models, yielding mainly a normalization of the spectra. We present below a computation for the Doppler contribution from global topological defects; in particular we perform our analysis for $\\pi_3$--defects, textures (Turok 1989), in a universe dominated by cold dark matter (CDM). We believe that the main ideas of our analysis remain valid for all global defects. The Doppler contribution to the CMB anisotropies is approximately given by\\footnote{We are neglecting here the SW effect, which decays on subhorizon scales like $\\ell^{-2}$. We also neglect the contribution of the neutrino fluctuations. These approximations lead to an error of less than about 30\\% in the amplitude of the first Doppler peak; see discussion in (Durrer, Gangui \\& Sakellariadou 1995).} \\begin{equation} \\left[ {\\delta T \\over T}({\\vec x},\\hat\\gamma ) \\right]^{Doppler} \\approx {1\\over 4}D_r({\\vec x}_{rec},\\eta_{rec}) - {\\vec V}({\\vec x}_{rec},\\eta_{rec})\\cdot {\\hat\\gamma} ~, \\label{Doppler} \\end{equation} where ${\\vec x}_{rec} = {\\vec x} + {\\hat\\gamma} \\eta_0$. In the previous formula ${\\hat\\gamma}$ denotes a direction in the sky and $\\eta $ is the conformal time, with $\\eta_0 $ and $\\eta_{rec} $ the present and recombination times, respectively. ",
+        "conclusions": "\\label{sec-disctextu}  As we see from the figure, the angular power spectrum $\\ell (\\ell + 1) C_{\\ell}$  yields the Doppler peaks. For $\\ell <1000$, we find three peaks located at $\\ell=365$, $\\ell=720$ and $\\ell = 950$. Silk damping, which we have not taken into account here, will decrease the relative amplitude of the third peak with respect to the second one; however it will not affect substantially the height of the first peak. The integrated SW effect, which also has been neglected, will shift the position of the first peak to somewhat larger scales, lowering $\\ell_{peak}$ by (5 -- 10)\\% and possibly increasing its amplitude slightly (by less than 30\\%) (Durrer, Gangui \\& Sakellariadou 1995).  Our second important result regards the amplitude of the first Doppler peak, for which we find $\\ell(\\ell+1)C_\\ell\\left|_{{~}_{\\!\\! \\ell\\sim 365 }}=5\\epsilon^2~.\\right.$ It is interesting to notice that the position of the first peak is displaced by  $\\Delta \\ell\\sim 150$  towards smaller angular scales than in standard inflationary models (Steinhardt 1995). This is due to the difference in the growth of super--horizon perturbations (Sugiyama 1996), which is $D_r \\propto \\eta^{3/2}$ in our case, and $D_r \\propto \\eta^2$ for inflationary models.  Let us now compare our value for the Doppler peak with the level of the SW plateau. Unfortunately the numerical value for the SW amplitude is uncertain within a factor of about 2, which leads to a factor 4 uncertainty in the SW contribution to the power spectrum: the results are $\\ell(\\ell+1)C_\\ell\\left|_{{~}_{\\!\\! SW}} \\sim 2 \\epsilon^2 \\right.$ (Bennett \\& Rhie 1993; Pen et al. 1994) and $\\ell(\\ell+1)C_\\ell\\left|_{{~}_{\\!\\! SW}} \\sim 8 \\epsilon^2 \\right.$ (Durrer \\& Zhou 1995). According to the first two groups, the Doppler peak is a factor of $\\sim 3.4$ times higher than the SW plateau, whereas it is only about 1.5 times higher if the second result is assumed. (We allow for about 30\\% of the SW amplitude to be added in phase to the Doppler amplitude of $\\sim 2.24\\epsilon$.) Improved numerical simulations or analytical approximations are needed to resolve this discrepancy. However, it is apparent that the Doppler contribution from textures is somewhat smaller than for generic inflationary models.  We believe that our results, about the position and amplitude of  the first Doppler peak, are basically valid for all global defects. This depends crucially on the $ 1/\\sqrt{\\eta}$ behavior of $(\\phi')^2$ on large scales (cf. Eq.~(\\ref{fit})). The displacement of the peaks towards smaller scales is reminiscent of open models, where the first Doppler peak is located at $\\ell_{peak} \\simeq 220 / \\sqrt{\\Omega_0}$. A value $\\Omega_0 \\simeq 0.3$ would produce $\\ell_{peak} \\simeq 365$ for the first peak, as we actually find. However, the amplitude of the peaks within open inflationary models is generically larger than the one we find for global defects.   The problem treated in this paper has also been studied by (Crittenden \\& Turok 1995).  They obtained the same position of the first Doppler peak but with a somewhat higher amplitude."
+    },
+    "9602/gr-qc9602025_arXiv.txt": {
+        "abstract": "In the context of the open inflationary universe, we calculate the amplitude of quantum fluctuations which deform the bubble shape. These give rise to scalar field fluctuations in the open Friedman-Robertson-Walker universe which is contained inside the bubble. One can transform to a new gauge in which matter looks perfectly smooth, and then the perturbations behave as tensor modes (gravitational waves of very long wavelength). For $(1-\\Omega)<<1$, where $\\Omega$ is the density parameter, the microwave temperature anisotropies produced by these modes are of order $\\delta T/T\\sim H(R_0\\mu l)^{-1/2} (1-\\Omega)^{l/2}$. Here, $H$ is the expansion rate during inflation, $R_0$ is the intrinsic radius of the bubble at the time of nucleation, $\\mu$ is the bubble wall tension and $l$ labels the different multipoles ($l>1$). The gravitational backreaction of the bubble has been ignored. In this approximation, $G\\mu R_0<<1$, and the new effect can be much larger than the one due to ordinary gravitational waves generated during inflation (unless, of course, $\\Omega$ gets too close to one, in which case the new effect disappears). ",
+        "introduction": "The possibility of an open inflationary universe , in which the cosmological density parameter $\\Omega$ is less than 1, has been intensively studied in recent years \\cite{martin, linde,yama,lyth,tama,sasaki,selftama}. According to this model, the universe is initially in a de Sitter phase, driven by the potential energy of a scalar field trapped in a false vacuum $\\sigma_f$ (see Fig. 1). A bubble of the new vacuum $\\sigma_t$ nucleates and its interior undergoes a second period of inflation. The homogeneity of our universe is then attributed to the O(3,1) symmetry of the bubble: the interior of the light cone from the nucleation event is isometric to an open Friedman-Robertson-Walker (FRW) universe \\cite{coleman}. The second period of inflation has to be suficiently long to generate the observed entropy, but if it is too long then $\\Omega$ is driven exponentially close to 1. As a result, to obtain $\\Omega<1$ the parameters of the scalar field potential have to be fine tunned to some extent \\cite{martin,yama}. Nevertheless, if observations ultimately determine that $\\Omega$ is smaller than one, then open inflation may be regarded as a  `natural' scenario \\cite{linde}.  The cosmic microwave background anisotropies produced by a nearly massless scalar field in open inflation were analyzed in \\cite{yama}. It was shown that a `super-curvature' mode \\cite{lyth}, which is not normalizable on the open FRW space-like sections, would give a significant contribution to the low multipoles if $\\Omega<.1$ (see also \\cite{ sasaki,ratra}). A more complete study of cosmological perturbations in open inflation requires the quantization of fields in the presence of a bubble. Quantum field theory in a bubble background was pioneered in \\cite{tanmayalex} and further developped in \\cite{tama,sasaki,selftama}. As noted in \\cite{linde}, large contributions to microwave perturbations may result from the quantum fluctuations of the bubble wall itself \\cite{jaumealex}. The purpose of this paper is to calculate the amplitude of such fluctuations and their effect on the microwave background.  In Section 2 we briefly describe the bubble geometry. In Section 3 we calculate the amplitude of wall fluctuations for bubbles nucleated during inflation. This extends previous work for bubbles in flat space \\cite{jaumealex}(see also \\cite{selftama}). In Section 4 we show that the effect of wall perturbations can be described in terms of long wavelength tensor modes (analogous to gravitational waves), and we evaluate their impact on the microwave sky. Finally, in Section 5 we summarize our conclusions and compare them with recent related work \\cite{linde,juan}. ",
+        "conclusions": "We have calculated the amplitude of small fluctuations in the shape of bubbles nucleated during inflation. In the case of thin bubbles, the result takes the  simple form (\\ref{demol}). Since the de Sitter modes $Y_{-3lm}$ grow like the intrinsic radius of the bubble $r=R_0 \\cosh \\rho$ ($\\rho R_0$ is the proper time coordinate on the bubble wall), the relative amplitude of proper local radial wall displacements is given by $$ {\\delta r\\over r}\\sim (R_0^3\\mu)^{-1/2}, $$ where $R_0$ is the radius of the bubble at the time of nucleation, given by (\\ref{r0}), and $\\mu$ is the wall tension. This coincides by order of magnitude with the estimate of Ref.\\cite{linde} in the case when the size of the bubbles at the time of nucleation is much smaller than the de Sitter horizon $H^{-1}$.  The propagation of wall perturbations to the interior of the light cone gives rise to shear deformations of the reheating surface, and consequently, of the surface of last scattering. The simplest way to study the cosmological evolution of such perturbations is to use a synchronous coordinate system in which matter looks perfectly smooth. Somewhat surprisingly, the metric fluctuations in this gauge are transverse and traceless, just like usual gravitational waves. This transmutation of scalar into tensor-like modes is only possible because of the peculiar (supercurvature) eigenvalue of the Laplacian for the scalar modes corresponding to wall fluctuations, $k^2=-3$ \\cite{sasaki}.   For $(\\Omega-1)<<1$, the anisotropies of the microwave sky produced by these waves are given in (\\ref{deltatemp}). The dominant effect is in the quadrupole, with $$ {\\delta T\\over T}\\sim {H\\over\\sqrt{R_0\\mu}}(1-\\Omega). $$ Within the limits of validity of our approximation [see (\\ref{weak})], and unless $\\Omega$ is too close to 1, this can be much larger than the distortion produced by usual gravitational waves produced during inflation, which is of order $G^{1/2}H$. The case with strongly gravitating domain walls may yield a different result \\cite{juan}, and is currently under investigation.  Finally, we would like to compare our results with those of Ref. \\cite{juan}. There it is claimed that the $k^2=-3$ modes produce no distortions of the microwave background. Here we have calculated the amplitude of such perturbations and their nonvanishing effect. In Ref. \\cite{juan} the amplitude of a $k=0$ homogeneous fluctuation is estimated by considering the small change $\\Delta S$ in the instanton action when the radius of the bubble is changed by an amount $\\delta a$. This is used to estimate the `typical deviation' $\\delta a$ as the one that corresponds to $\\Delta S \\sim 1$. However, the physical meaning of that prescription is unclear, because the radius of the bubble cannot have a homogeneous fluctuation. In flat space, this would violate energy conservation, and a similar argument can be applied in curved space. The only homogeneous radial fluctuations allowed in the thin wall limit are time translations of the nucleation point."
+    },
+    "9602/astro-ph9602124_arXiv.txt": {
+        "abstract": "This work is based on the first results from a systematic search for high redshift Type Ia supernovae. Using filters in the $R$-band we discovered seven such SNe, with redshift $z=0.3 - 0.5$, before or at maximum light.  Type Ia SNe are known to be a homogeneous group of SNe, to first order, with very similar light curves, spectra and peak luminosities.  In this talk we report that the light curves we observe are all broadened (time dilated) as expected from the expanding universe hypothesis. Small variations from the expected $1+z $ broadening of the light curve widths can be attributed to a width-brightness correlation that has been observed for nearby SNe ($z<0.1 $). We show in this talk the first clear observation of the cosmological time dilation for macroscopic objects. ",
+        "introduction": "In an ongoing systematic search, we have recently discovered and studied seven supernovae (SNe) at redshifts between $z = 0.35$ and 0.46 as discussed in Perlmutter's talk at this conference. For details on search technique, light curves, and spectra, see Perlmutter et al. 1996. ",
+        "conclusions": "There is one  source of concern  that must still be addressed: how certain are the Type Ia identifications? We expect to find primarily Type Ia since these are the brightest SNe by typically 2 magnitudes. For three of our SNe we have  spectral identification and their light curves are consistent with the other four. One of the remaining four SNe was discovered in an eliptical galaxy, and therefore is highly likely to be a Type Ia.  In sum, the current set of data are very likely to all be SNe Ia."
+    },
+    "9602/astro-ph9602108_arXiv.txt": {
+        "abstract": "We show that the simplest assumptions for the dynamics of particle production allow us to understand the fluxes of hadrons and photons at mountain altitudes as well as the structure of individual events. The analysis requires a heavy nuclear component of primary cosmic rays above the ``knee\" in the spectrum with average mass number $<A> = 7.3 \\pm 0.9$. ",
+        "introduction": "The energy spectrum and chemical composition of primary cosmic rays have been determined from direct observation above the earth's atmosphere using balloons and spacecraft~\\cite{shibata}. The technique is limited by the small size of the detectors and the short exposure times. As a result of the steep energy spectrum no direct observations are available above a primary energy of roughly $10^{14}$~eV. In the interesting region of the ``knee\" in the spectrum and above information on composition has to be inferred from indirect measurements of air showers at sea level or mountain altitudes~\\cite{gaisser}, by using  large area detectors for long periods of time.  In this paper we infer the composition of the cosmic rays from measurements at  mountain altitudes of the hadronic and electromagnetic component of the air  cascades initiated at the top of the atmosphere. A connection between the nature of the primary particle and air shower observations requires the detailed understanding of particle interactions at very high energies and forward scattering angles where no information is available from accelerator-based experiments. The basic problem is that one is faced with the  impossibility of deducing two unknowns, the composition and the dynamics of  particle interactions, from a single measurement. We nevertheless pursue this  challenge because we are confident that we understand particle interactions  with sufficient accuracy to meaningfully approach this problem. We have indeed  formulated a model which is based on the most straightforward assumptions and which respects the spirit of quantum chromodynamics~\\cite{costa}. More importantly, it describes in quantitative detail single events, i.e. shower cores in their early stage of development, observed in emulsion chamber experiments at mountain altitudes~\\cite{lattes}. Here we will show that this model describes the observed hadronic and electromagnetic spectrum at mountain  altitudes, provided the mass number of the primary cosmic rays is $<A> = 7.3 \\pm  0.9$. This is consistent with the result obtained by other indirect means.   Our model of particle production~\\cite{costa} is guided by the features of QCD-inspired models: approximate Feynman scaling in the fragmentation region and an inelasticity slowly varying with energy. The  rapidity density of secondary charged pions is parametrized as  \\begin{equation} {dN\\over dy} = x {dN\\over dx} = a {(1-x)^n \\over{x}} \\;, \\label{rapdens} \\end{equation} where $y$ is the rapidity of the secondaries and the Feynman variable $x$ is given by the ratio of the energy $E$ of the secondary particle to the incident energy $E_0$. With $a = 0.12$ and $n = 2.6$ ($n = 3$ is expected on the basis  of  counting rules) the overall features of the hadronic component of single  events detected in emulsion chambers were quantitatively reproduced. For  illustration, we present in Fig.~\\ref{fig:families} the hadronic integral spectrum of two events detected by the Brazil-Japan Collaboration at  Mt. Chacaltaya, Bolivia (atmospheric depth $540$ g/cm$^2$)~\\cite{chinellato,bjc-1}, which are successfully described by our  model~\\cite{fluctuations}. In the present paper we use the same model to calculate both the hadronic and electromagnetic integral spectra of atmospheric showers detected in large emulsion chamber experiments. The calculation is performed by solving the cosmic-ray diffusion equations using the rapidity distribution for particle production given by Eq.~(\\ref{rapdens}). Starting with the measured all-particle primary spectrum at the top of the atmosphere, we propagate the  particle showers down to the mountain altitude detection levels of the various  experiments and investigate our results as a  function of the assumed average  composition of the primary cosmic radiation. ",
+        "conclusions": "We conclude that with the simplest assumptions for the production of secondaries based on approximate scaling in the fragmentation region, it is possible to explain a broad set of experimental data on very high energy cosmic rays in the atmosphere, namely the lateral spread and the integral spectra of superfamilies (as in Ref.~\\cite{costa} and Fig.~\\ref{fig:families}), and the energy spectra of hadronic and electromagnetic showers detected in large emulsion chamber experiments (as in Figs.~\\ref{fig:chacaltaya} and~\\ref{fig:fujikanb}). This scenario requires a  primary composition with average mass number $7.3 \\pm 0.9$. We investigated the sensitivity of this quantity to different parametrizations of the all-particle  spectrum ~\\cite{jacee-2,biermann} and the best fit invariably yields $<A> \\simeq  7$. Our  result is also consistent with underground muon measurements~\\cite{macro} which yield an $A$-value of $ 10 \\pm 4$  in the 1 to 1,000 TeV energy range.   It has been noted elsewhere~\\cite{yuda-halzen} that it is difficult to establish whether one must adopt a heavy primary composition along with a model of particle production based on scaling or, alternatively, a proton dominant composition along with a strong violation of Feynman scaling. It should be noted however that in our analysis the particle interaction model was determined on the basis of an independent study of single events initiated by protons deep in the atmosphere. The a posteriori analysis of the hadron and photon spectra at mountain altitude presented here, required the introduction of heavy primaries."
+    },
+    "9602/gr-qc9602009_arXiv.txt": {
+        "abstract": "\\baselineskip=1.5em We study the existence and stability of spherical membranes in curved spacetimes. For Dirac membranes in the Schwarzschild--de Sitter background we find that there exists an equilibrium solution. By fine--tuning the dimensionless parameter $\\Lambda M^2,$ the static membrane can be at any position outside the black hole event horizon, even at the stretched horizon, but the solution is unstable. We show that modes having $l=0$ (and for $\\Lambda M^2<16/243$ also $l=1$) are responsible for the instability. We also find that spherical higher order membranes (membranes with extrinsic curvature corrections), contrary to what happens in flat Minkowski space, {\\it do} have equilibrium solutions in a general curved background and, in particular, also in the ``plain'' Schwarzschild geometry (while Dirac membranes do not have equilibrium solutions there). These solutions, however, are also unstable. We shall discuss a way of by--passing these instability problems, and we also relate our results to the recent ideas of representing the black hole event horizon as a relativistic bosonic membrane. ",
+        "introduction": "\\setcounter{equation}{0}  The idea of representing physical objects by extended ones such as membranes, can be traced back to the pioneering work of Dirac\\ \\cite{D62}. In this model for the electron, tension forces tending to collapse the membrane were balanced by repulsive electric forces. Later, canonical\\ \\cite{CT76,AS90} and semiclassical\\ \\cite{DIPSS88} quantization of membranes in a flat background was considered. On the other hand, starting with the works of `t Hooft\\ \\cite{tH85}, a series of papers\\cite{STU93,M94,L295} considered the quantum degrees of freedom of a black hole, localized around its event horizon. It is also worth noting the classical approach to black holes by the, so called, ``membrane paradigm'', summarized in Ref.\\ \\cite{TPM86}.  In Ref.\\ \\cite{L95} it was considered the approach of effectively representing the quantum degrees of freedom of a black hole by a relativistic membrane. The total Lagrangian of the system was composed of the classical curved geometry plus a relativistic bosonic membrane. The membrane, located at the ``stretched horizon'', was decomposed into a spherically symmetric term plus small fluctuations around it. The first order fluctuations have then been quantized and the energy spectrum was obtained. From this, the entropy of the membrane at the temperature of the black hole (given by the background geometry) was computed. It was however left open the question of the existence and stability of the zeroth order solution representing a spherical membrane at rest very close to the black hole event horizon. In the present paper we address the above question.  This paper is organized as follows: In Sec.\\ II we study the Dirac membrane in its simplest form. No coupling to the electromagnetic field is considered, instead the membrane is let to freely move in a curved background. Since our final aim is to study black hole properties, and for the sake of simplicity, we consider spherically symmetric backgrounds. We derive the condition for the existence of spherical membranes to be in equilibrium at a given radius $r_m$. We also find the stability condition and apply it to a membrane in the Schwarzschild--de Sitter space. Sec.\\ III deals with a covariant description of small fluctuations around the equilibrium position $r_m$. We find that the fluctuations are governed by a Klein--Gordon--like equation in the 2+1 dimensional world--volume swept by the membrane. We identify modes with angular momentum label $l=0,1$ as responsible for the instabilities in the Schwarzschild--de Sitter background. In Sec.\\ IV we study a more general membrane described by a Lagrangian with up to quadratic terms in the extrinsic curvature. This introduces two more arbitrary constants $A$ and $B$, in addition to the membrane tension $T$. This freedom allows us, in the plain Schwarzschild background,  to choose $r_m$ as close to the event horizon as we want (for instance at the stretched horizon). We study the radial small fluctuation equation (fourth order differential equation) and find the equilibrium and stability conditions. In Sec.\\ V we discuss the way of extending the validity of the perturbative approach around a maximum of the potential, instead of a minimum, when we take the comoving system of reference to describe the zeroth order motion. In Appendix\\ A we give the analytic expressions, corresponding to the plots of Fig.\\ \\ref{fig1}, for the cosmological and event horizon, and $r_m$ in the Schwarzschild--de Sitter spacetime. ",
+        "conclusions": "\\setcounter{equation}{0}  In this section we will show that it is possible to avoid the instabilities by considering a dynamical (contracting) membrane instead of a static equilibrium membrane. We will perform the computations for a Dirac membrane in the Schwarzschild--de Sitter background, i.e. the model considered in Secs.\\ II--III. However, the results of Sec.\\ IV, concerning the higher order membrane (especially the existence of unstable static equilibrium solutions), strongly indicate that the following arguments can also be carried through for the higher order membrane in the plain Schwarzschild background, that is to say, the presence of the cosmological constant does not seem to be essential.  Let us return to the potential (\\ref{1.5}) in the case of the Schwarzschild--de Sitter spacetime. It is apparent from  Fig.\\ \\ref{fig2} that there are the following qualitatively different kinds of ``orbits'' for spherical membranes  If $\\tilde E^2>\\tilde V^2(r_m)$ membranes are always ``over'' the potential barrier, and either expand for ever or contract towards the black hole, ending trapped by the singularity at $r=0$.  If $\\tilde E^2<\\tilde V^2(r_m)$ membranes might bounce if they are expanding and are ``inside'' the potential well, or if they are contracting but are ``outside'' the potential well.  We are particularly interested in the intermediate case, i.e. $\\tilde E^2=\\tilde V^2(r_m),$ of which we have analyzed the stability properties of the static equilibrium configuration $r=r_m$ in Sec.\\ II. However, for $\\tilde E^2=\\tilde V^2(r_m), $ there is also a contracting solution as well as an expanding solution. For the static solution we have: \\begin{equation} \\tilde E^2=(4\\pi\\tilde{T})^2 r^4_m \\left( 1-\\frac{2M}{r_m}-\\frac{\\Lambda}{3}r^2_m\\right), \\label{V.1}\\end{equation} \\begin{equation} 2-\\frac{3M}{r_m}-\\Lambda r^2_m=0. \\label{V.2}\\end{equation} Now consider a contracting solution with \"energy\" given by (\\ref{V.1}). We are interested in solutions contracting from a finite coordinate distance $r_0$ (where $r_m<r_0<r_{CH}$) towards $r_m.$ The {\\it proper} time to reach the summit of the potential barrier ($r=r_m$) happens to be logarithmically divergent, and can be derived as follows from Eq.\\ (\\ref{1.5}) \\begin{equation} \\int_{\\tau_0}^{\\tau_m}d\\tau=-\\int_{r_0}^{r_m}{dr\\over \\sqrt{\\tilde E^2-\\tilde V^2(r)}}= -{1\\over (4\\pi\\tilde{T})M}\\int_{x_0}^{x_m}{dx\\over (x-x_m)\\sqrt{f(x)}}~, \\label{V.3}\\end{equation} where \\begin{equation} f(x)={1\\over3}\\left\\{\\tilde\\Lambda x^4+2\\tilde\\Lambda x^3x_m+ 3(\\tilde\\Lambda x_m^2-1)x^2+2(2\\tilde\\Lambda x_m^3-3x_m+3)x+ 5\\tilde\\Lambda x^4_m-9x^2_m+12x_m\\right\\}, \\end{equation} and we have taken the dimensionless variables $x=r/M$ and $\\tilde\\Lambda=\\Lambda M^2$. One can thus show that the proper time distance becomes infinite. In fact, \\begin{equation} \\tau_m-\\tau_0=-{1\\over(4\\pi\\tilde T)M}{1\\over\\sqrt{f(x_m)}} \\ln(x-x_m)\\biggr\\vert_{x_0}^{x_m}+... \\label{V.5}\\end{equation} where \\begin{equation} f(x_m)=x_m\\left(6-6x_m+5x_m^3\\tilde\\Lambda\\right)~, \\end{equation} and the ellipsis stands for finite terms as $x\\to x_m$.  We have stressed that this infinity occurs in the proper time description. This is clearly different from the infinite ``Schwarzschild'' time it takes a particle to reach the event horizon (while it only takes a finite proper time, even to reach the singularity). In fact, the above logarithmic divergence arises at a coordinate radius $r_m>r_{EH}$. This is a particular property of the Schwarzschild--de Sitter spacetime background, that generates a maximum in the effective potential of the membrane, and this does {\\it not} occur for the Dirac membrane in the plain Schwarzschild geometry.  On the other hand, the above discussion applies quite generally to any potential barrier possessing a maximum. The infinite time to reach the top of the potential is well--known in classical mechanics, for instance, an ordinary pendulum experiences the same process when going towards the vertical (unstable) position with precisely the energy corresponding to being in the static, unstable equilibrium position. The results of Sec.\\ IV show that the dynamics of the higher order membrane in the plain Schwarzschild background is also governed by some kind of finite potential barrier, thus also in that case we can expect to find the above mentioned type of solution.  We now consider fluctuations around the contracting Dirac membrane in the Schwarzschild--de Sitter background, using equation (\\ref{2.2}). In general, in the Schwarzschild--de Sitter spacetime \\begin{equation} {\\cal V}(r)=\\Lambda+\\frac{6\\tilde{E}^2}{(4\\pi\\tilde{T})^2 r^6}, \\end{equation} so that for the contracting membrane with $\\tilde{E}$ given by Eq.\\ (\\ref{V.1}) \\begin{equation} {\\cal V}(r)=\\Lambda+\\frac{6r^4_m a(r_m)}{r^6},\\;\\;\\;\\;\\;\\;\\;r_m<r<r_0 \\end{equation} where $a(r)$ is given by (\\ref{II.14}) and $r_m$ is the solution of\\ (\\ref{V.2}). The \"frequency\" in equation (\\ref{2.2}) now also depends on $r$ \\begin{equation} \\omega_l^2(r)=r^4 (4\\pi\\tilde{T})^2\\left( \\frac{l(l+1)}{r^2}-\\Lambda-\\frac{6r^4_m a(r_m)}{r^6}\\right). \\label{V.9}\\end{equation} By careful (and partly numerical) study of the formulas for $r_m$ and $r_{CH}$ (See Appendix\\ A), one finds that: \\begin{eqnarray} &\\omega_0^2(r)<0&\\nonumber\\\\ &\\omega_1^2(r)>0&\\;\\;\\;\\;\\;\\Leftrightarrow \\;\\;\\;\\;\\tilde{\\Lambda}>0.07737...\\\\ &\\omega_l^2(r)>0&\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; \\mbox{for all}\\;\\;l\\geq 2\\nonumber \\end{eqnarray} and these results hold for all $r_m\\leq r\\leq r_0.$ Notice that the condition for $\\omega_1^2$ being positive is somewhat more restrictive than what was obtained at the end of Sec.\\ III. The reason is that we now consider $\\omega_1^2$ for arbitrary $r,$ not only for $r=r_m.$  We are particularly interested in the case when the membrane approaches a position close to the Schwarzschild event horizon. This is the case when $\\Lambda M^2\\equiv\\tilde{\\Lambda}$ is close to the maximal value $\\tilde{\\Lambda}_{\\mbox{max}}=1/9,$ c.f. the discussion at the end of Sec.\\ III. From the above results follow that for such contracting membranes, {\\it all} $\\;l\\geq 1$--modes are stable everywhere (that is, for all $r$) during the contraction. At first sight it seems, on the other hand, that the zero-mode is always unstable. However, for a contracting spherical membrane, the zero-mode does not represent a real physical fluctuation. The zero-mode merely represents a time-translation, and therefore should be eliminated from the spectrum. We thus conclude that {\\it all} physical fluctuations around the contracting membrane are stable.  To actually compute the classical and quantum spectrum of physical fluctuations ($l\\geq 1),$ we must solve the equation \\begin{equation} \\ddot{\\phi}_l+\\omega_l^2(r) {\\phi}_l=0, \\label{V.11}\\end{equation} where $\\omega_l^2(r)$ is given by (\\ref{V.9}), and $r=r(\\tau)$ is obtained by inverting (\\ref{V.3}). Eq.\\ (\\ref{V.3}) can be solved explicitly in terms of Jacobi Elliptic Functions \\cite{AS65}, and Eq.\\ (\\ref{V.11}) is then of (generalized) Lam\\'{e}-type. However, we shall not attempt here to obtain the exact solution for $\\phi(\\tau)$ in closed form. Instead we will use the result that the membrane takes infinite proper time to reach the summit of the potential, c.f. Eq.\\ (\\ref{V.5}). Physically this means that after some finite time, the membrane is effectively always very close to the top of the barrier. We can thus approximate $\\omega_l^2(r)$ by a constant which is essentially $\\omega_l^2(r_m),$ as given by Eqs.\\ (\\ref{III.16}) and (\\ref{III.17}). We are then left with an ordinary harmonic oscillator equation, which can be easily solved. In fact, in Ref.\\ \\cite{L95} it was assumed that this equilibrium point existed and from the quantization of the harmonic oscillations a discrete spectrum of energies for the membrane obtained. The entropy associated to this discrete energy levels, in the thermal bath of the background black hole, can be computed and one obtains a leading term proportional to the two--surface of the membrane plus a non--leading logarithmic term.  In conclusion, we have learned that in a curved background, as opposite to what happens in flat space, spherical membranes can exist in equilibrium. Higher order membranes provide enough arbitrariness to locate the membrane as close to the event horizon as we want, even in the plain Schwarzschild background. The issue of the stability of spherical membranes is delicate. We identified modes $l=0$ and sometimes also $l=1$ as responsible for the instabilities. We discussed, however, a certain range of validity of the perturbative description and eventual quantization of the fluctuations. Further study of the non--radial fluctuations of higher order membranes is needed, as well as of the situation in more general curved backgrounds (for example, to include rotation) to see if a ``true'' stability mechanism can be found."
+    },
+    "9602/astro-ph9602087_arXiv.txt": {
+        "abstract": "The prediction of cross sections for nuclei far off stability is crucial in the field of nuclear astrophysics. In recent calculations the nuclear level density -- as an important ingredient to the statistical model (Hauser-Feshbach) -- has shown the highest uncertainties. We present a global parametrization of nuclear level densities based on the back-shifted Fermi-Gas formalism. Employment of an energy-dependent level density parameter $a$ and microscopic corrections from a recent FRDM mass formula by M\\\"oller et al.\\ leads to a highly improved fit of level densities at the neutron-separation energy in the mass range $20\\le A \\le 245$. The importance of using proper microscopic corrections from mass formulae is emphasized. The resulting level description is well suited for astrophysical applications.  The level density can also provide clues to the applicability of the statistical model which is only correct for a high density of excited states. Using the above description one can derive a ``map'' for the applicability of the model for reactions of stable and unstable nuclei with neutral and charged particles. ",
+        "introduction": "Explosive nuclear burning in astrophysical environments produces unstable nuclei, which again can be targets for subsequent reactions. In addition, it involves a very large number of stable nuclei, which are not fully explored by experiments. Thus, it is  necessary to be able to predict reaction cross sections and thermonuclear rates with the aid of theoretical models. Explosive burning in supernovae involves in general intermediate mass and heavy nuclei. Due to a large nucleon number they have intrinsically a high density of excited states. A high level density in the compound nucleus at the appropriate excitation energy allows to make use of the statistical model approach for compound nuclear reactions (e.g. Hauser \\& Feshbach 1952, Mahaux \\& Weidenm\\\"uller 1979), which averages over resonances.   As the capture of an alpha particle leads usually to larger Q-values than neutron or proton captures, the compound nucleus is created at a higher excitation energy and especially in the case of alpha-captures it is often even possible to apply the Hauser-Feshbach formalism for light nuclei. Another advantage of alpha-captures is that the capture Q-values vary little with the N/Z-ratio of a nucleus, for nuclei with Z$\\le$50. For Z$>$50, entering the regime of natural alpha-decay, very small alpha-capture Q-values can be encountered for proton-rich nuclei. Such nuclei on the other hand do not play a significant role in astrophysical environments, maybe with exception of the p-process. This means that in the case of alpha-captures the requirement of large level densities at the bombarding energy is equally well fulfilled at stability as for unstable nuclei. Opposite to the behavior for alpha-induced reactions, the reaction Q-values for proton or neutron captures vary strongly with the N/Z-ratio, leading eventually to vanishing Q-values at the proton or neutron drip line. For small Q-values the compound nucleus is created at low excitation energies and also for intermediate nuclei the level density can be quite small. Therefore, it is not advisable to apply the statistical model approach close to the proton or neutron drip lines for intermediate nuclei. For neutron captures close to the neutron drip line in r-process applications it might be still permissable for heavy and often deformed nuclei, which have a high level density already at very low excitation energies.  In astrophysical applications usually different points are emphasized than for investigations in pure nuclear physics. Many of the latter in this well established field were focused on specific reactions, where all or most \"ingredients\", like optical potentials for particle and alpha transmission coefficients, level densities, resonance energies and widths of giant resonances to be implementated in predicting E1 and M1 gamma-transitions, were deduced from experiments. This of course, as long as the statistical model prerequisits are met, will produce highly accurate cross sections. For the majority of nuclei in astrophysical applications such information is not available. The real challenge is thus not the well established statistical model, but rather to provide all these necessary ingredients in as reliable a way as possible, also for nuclei where none of such informations are available. In addition, these approaches should be on a similar level as e.g. mass models, where the investigation of hundreds or thousands of nuclei is possible with manageable computational effort, which is not always the case for fully microscopic calculations.  The major uncertainty in all existing calculations stems from the prediction of nuclear level densities (Truran et al.\\ 1966; Holmes et al.\\ 1976; Cowan, Thielemann \\& Truran 1991, and extended references therein), which in earlier calculations gave uncertainties even beyond a factor of 10 at the neutron separation energy (Gilbert \\& Cameron 1965), about a factor of 8 (Woosley et al. 1978), and a factor of 5 even in the most recent calculations (e.g. Thielemann, Arnould \\& Truran 1987; see also Fig.\\ 3.16 in Cowan, Thielemann \\& Truran 1991).  \\section {Thermonuclear Rates from Statistical Model Calculations}   A high level density in the compound nucleus permits to use averaged transmission coefficients $T$, which do not reflect a resonance behavior, but rather describe absorption via an imaginary part in the (optical) nucleon-nucleus potential (for details see Mahaux \\& Weidenm\\\"uller 1979). This leads to the well known expression \\begin{eqnarray} \\sigma^{\\mu \\nu}_{i} (j,o;E_{ij})& = & {{\\pi \\hbar^2 /(2 \\mu_{ij} E_{ij})} \\over {(2J^\\mu_i+1)(2J_j+1)}} \\nonumber \\\\ & & \\times \\sum_{J,\\pi} (2J+1){{T^\\mu_j (E,J,\\pi ,E^\\mu_i,J^\\mu_i, \\pi^\\mu_i) T^\\nu_o (E,J,\\pi,E^\\nu_m,J^\\nu_m,\\pi^\\nu_m)} \\over {T_{tot} (E,J,\\pi)}} \\label{hf} \\end{eqnarray} for the reaction $i^\\mu (j,o) m^\\nu$ from the target state $i^{\\mu}$ to the exited state $m^{\\nu}$ of the final nucleus, with center of mass energy E$_{ij}$ and reduced mass $\\mu _{ij}$. $J$ denotes the spin, $E$ the excitation energy, and $\\pi$ the parity of excited states. When these properties are used  without subscripts they describe the compound nucleus, subscripts refer to states of the participating nuclei in the reaction $i^\\mu (j,o) m^\\nu$ and superscripts indicate the specific excited states. Experiments measure $\\sum_{\\nu} \\sigma_{i} ^{0\\nu} (j,o;E_{ij})$, summed over all excited states of the final nucleus, with the target in the ground state. Target states $\\mu$ in an astrophysical plasma are thermally populated and the astrophysical cross section $\\sigma^*_{i}(j,o)$ is given by \\begin{equation} \\label{astro} \\sigma^*_{i} (j,o;E_{ij}) = {\\sum_\\mu (2J^\\mu_i+1) \\exp(-E^\\mu_i /kT) \\sum_\\nu \\sigma^{\\mu \\nu}_{i}(j,o;E_{ij}) \\over \\sum_\\mu (2J^\\mu_i+1) \\exp(-E^\\mu_i/kT)}\\quad. \\end{equation} The summation over $\\nu$ replaces $T_o^{\\nu}(E,J,\\pi)$ in Eq.~\\ref{hf} by the total transmission coefficient \\begin{eqnarray} T_o (E,J,\\pi) & = &\\sum^{\\nu_m}_{\\nu =0} T^\\nu_o(E,J,\\pi,E^\\nu_m,J^\\nu_m, \\pi^\\nu_m) \\nonumber \\\\ & &+ \\int\\limits_{E^{\\nu_m}_m}^{E-S_{m,o}} \\sum_{J_m,\\pi_m} T_o(E,J,\\pi,E_m,J_m,\\pi_m)\\rho(E_m,J_m,\\pi_m) dE_m \\quad. \\label{tot} \\end{eqnarray} Here $S_{m,o}$ is the channel separation energy, and the summation over excited states above the highest experimentally known state $\\nu_m$ is changed to an integration over the level density $\\rho$. The summation over target states $\\mu$ in Eq.~\\ref{astro} has to be generalized accordingly.  In addition to the ingredients required for Eq.~\\ref{hf}, like the transmission coefficients for particles and photons, the width fluctuation corrections $W(j,o,J,\\pi)$ have to be employed. The important ingredients of statistical model calculations as indicated in Eqs.~\\ref{hf} through \\ref{tot} are the particle and $\\gamma$-transmission coefficients $T$ and the level density of excited states $\\rho$. Therefore, the reliability of such calculations is determined by the accuracy with which these components can be evaluated. ",
+        "conclusions": "\\noindent We were able to improve considerably the prediction of nuclear level densities by employing an energy-dependent description for the level density parameter $a$ and by correctly including microscopic corrections. All nuclei can now be described with a single parameter set consisting of just three parameters. The globally averaged deviation of prediction from experiment of about 1.5 translates into somewhat smaller deviations of the final cross sections due to the dominance of transitions to states with low excitation energies. This will also make it worthwile to recalculate the cross sections and thermonuclear rates for many astrophysically important reactions in the intermediate and heavy mass region.  We also presented a ``map'' as a guide for the application of the statistical model for neutron-, proton- and $\\alpha$-induced reactions. The above plots can give hints on when it is safe to use the statistical model approach and which nuclei have to be treated with special attention at a given temperature. Thus, information on which nuclei might be of special interest for an experimental investigation may also be extracted. It should be noted that we used very conservative assumptions in deriving the above criteria for the applicability of the statistical model."
+    },
+    "9602/astro-ph9602020_arXiv.txt": {
+        "abstract": "We discuss a new test of inflation from the harmonic pattern of peaks in the cosmic microwave background radiation angular power spectrum. By characterizing the features of alternate models and revealing signatures essentially unique to inflation, we show that inflation could be validated by the next generation of experiments. ",
+        "introduction": " ",
+        "conclusions": ""
+    },
+    "9602/astro-ph9602147_arXiv.txt": {
+        "abstract": "We study the photometric properties and content of the distant clusters of galaxies Cl~1613+3104 ($z = 0.415$) and Cl~1600+4109 ($z = 0.540$). The former is a rich, concentrated cluster which shows a strong evidence for segregation in luminosity and color. Its population of galaxies separates into two different kinds of objects: a red population compatible with the colors expected for E/S0 galaxies located preferentially in the central part of the cluster, and a sparse blue population explained in part by the presence of normal S/Im galaxies, which is much less concentrated. In the case of Cl~1600+4109, there is a lack of E/S0 galaxies and the population is mainly made up of S/Im galaxies, without any evidence of segregation in magnitude or in color. Both clusters exhibit an important fraction of blue objects, increasing with magnitude. This result is due in part to the presence of normal S/Im galaxies, and also to extremely blue objects probably undergoing an episode of intense star-formation. ",
+        "introduction": "Photometry of rich clusters of galaxies at high redshift remains one of the most important tools to understand the evolution of galaxy populations in dense environments. The discovery of a large fraction of blue galaxies in rich clusters has been understood as evidence for a strong, recent evolution of galaxies in clusters (Butcher \\& Oemler 1984, hereafter BO; Sharples et al.\\ 1985). The interpretation of this blue excess is a difficult challenge because the galaxy content varies from cluster to cluster. Dressler et al.\\ (1985) suggested that it can arise from several reasons: different initial conditions, a difference in the time scales for cluster formation and/or environmental influences. The presence of luminosity and color segregation in nearby clusters (Capelato et al.\\ 1980) as well as in high redshift clusters (Mellier et al.\\ 1988) is often interpreted in terms of dynamic friction and environmental influence (galaxy-galaxy and galaxy-ICM interactions). Nevertheless, it is not still clear to what extent these properties are innate or the result of later evolutionary processes. Dressler \\& Gunn (1992) published a paper on the photometry and spectroscopy of seven clusters in the redshift range $0.35 \\leq z \\leq 0.55$, including a thorough discussion on the statistical properties of the sample. Such studies are crucial to an understanding of the evolutionary processes that take place in clusters. Recently, Dressler et al.\\ (1994a,b) and Coutch et al.\\ (1994) studied the morphology of the blue galaxies in clusters at $0.3\\ltsm z \\ltsm 0.4$, using HST images. They found that the blue objects are predominantly disk-dominated or irregular galaxies, involved in merger or interaction processes. Thus, it is important to enlarge the sample of well known high and medium redshift clusters, and photometry allows to study a larger sample of galaxies than spectroscopy.  Cl~1613+3104 ($\\alpha_{1950}=16^{\\rm h} 13^{\\rm m} 49\\fs2$, $\\delta_{1950}=31\\degr 4\\arcmin$ $57\\farcs0$; $l=50\\fdg58$, $b=+45\\fdg56$) is a mid-redshift rich cluster, which exhibits X-ray (Henry et al.\\ 1982) and radio emission (Jaffe 1982). Its redshift ($z=0.415$) was obtained by Sandage et al.\\ (1976) from the  spectrum of the first-ranked galaxy. V\\'\\i l\\-chez--G\\'o\\-mez et al.\\ (1994) reported the existence of a diffuse light extending up to 265 h$^{-1}$ kpc from its center, associated with a stellar component immersed in the gravitational potential of the cluster. Cl~1600+4109 ($\\alpha_{1950}=16^{\\rm h}~0^{\\rm m}~23\\fs0$, $\\delta_{1950}=41\\degr~9\\arcmin~39\\farcs0$; $l=65\\fdg21$, $b=+48\\fdg72$) is another mid-redshift cluster ($z=0.540$; Henry et al.\\ 1982) included in the sample of Gunn et al.\\ (1986), whose population of galaxies has not been studied to date. It is not a strong X-ray emitter: its upper limit for the X-ray emission is $1.6\\times10^{44}$~erg~s$^{-1}$ in the 0.5--4.5~keV range (Henry et al.\\ 1982). The database used here comes from the photometric survey by Pe\\-ll\\'{o} \\& V\\'{\\i}l\\-chez--G\\'{o}\\-mez (1996), hereafter referred as Paper I.  The layout of this paper is the following. A summary of the observational procedure is presented in Sect.\\ 2. Section 3 reviews the methods used in the analysis of the photometric catalogs and the results obtained are given in Sects.\\ 4 and 5. In Sect.\\ 6 we discuss these results and we present the conclusions of this work. We assume $H_{0}=50~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$ and $q_{0}=0.1$ in a standard Friedmann cosmology. ",
+        "conclusions": "Cl~1613+3104 is a rich, concentrated cluster, dominated by a compact group of luminous galaxies within a radius of 121 h$^{-1}$ kpc around the center. The cluster shows strong evidence for segregation in luminosity and color, and both effects seem to be independent. Red and bright objects tend to concentrate around the center of the cluster. The distribution of galaxies in colors shows two different populations of objects. The red population is compatible with the colors expected for E/S0 type galaxies at the cluster redshift. These objects dominate the population of cluster galaxies at low magnitudes and they are located preferentially in the central part of the cluster. The blue population is explained in part by the presence of normal star-forming S/Im galaxies, but a population of galaxies bluer than normal Im galaxies at the cluster redshift is also detected. Moreover, the 2D distribution for the red population is clearly more concentrated around the center of the cluster than the blue one.  The detection level in Cl~1600+4109 is lower than in Cl~1613+3104, and most cluster galaxies are found beyond the completeness limits in magnitude. Nevertheless, its concentration and richness parameters indicate that we are probably dealing with a rich, concentrated cluster. It is dominated by a bright galaxy in its central region, with colors fully compatible with those expected for E galaxies at this redshift. The fraction of E/S0 galaxies is low in this cluster (less than $15 \\%$ within the completeness limit in r), and the population of galaxies is dominated by S/Im galaxies. There is no evidence of segregation in magnitude or in color.  In both clusters, the luminosity function in B is compatible with the standard values derived by Colless (1989) for a sample of rich low-redshift clusters. The absolute magnitudes derived for their brightest galaxies, after applying the standard k+e correction, are similar. The main galaxy in Cl~1600+4109 is slightly brighter than that of Cl~1613+3104 (${\\rm M}_{\\rm r} = -23.40$ compared to ${\\rm M}_{\\rm r} = -23.01$, within the isophote corresponding to $1\\sigma$ of the background sky noise). Comparison of the two histograms in M$_{\\rm B}$ for Cl 1600+4109 and Cl~1613+3104, after correction through the comparison field (Figs.\\ 1 and 6), shows that the completeness magnitude in Cl~1613+3104 is 1.2 mag fainter than in Cl~1600+4109 when a k correction for E/S0 galaxies is used, and this value reduces to 0.9 mag when a standard k+e correction is applied. Nevertheless, most galaxies with magnitudes around B$_{\\rm com}$ in both clusters are S/Im galaxies. When the k correction used is that of an Sb galaxy, the completeness magnitude in Cl~1613+3104 is 0.6 mag fainter than in Cl~1600+4109, and the difference reduces to 0.3 mag after a k+e correction. So the absolute detection level in Cl~1600+4109 is poorer than in Cl~1613+3104. The main reason for this is probably the difference in the seeing conditions between both clusters. Cl~1613+3104 was observed with an averaged seeing of $\\sim 1\\farcs3$, whereas the seeing was $\\sim 2\\farcs1$ for Cl~1600+4109, but the isophotes used to integrate fluxes were 0.5 mag fainter in the later (see Paper I). A straightforward simulation of photometry has been done, taking the best r image of Cl~1613+3104 (seeing $= 1\\farcs1$) and debasing it to an equivalent seeing of $\\sim 2\\farcs5$. As expected, a general tendency appears in the isophotal magnitudes: objects in the bad-seeing image tend to be measured as fainter than in the good-seeing one, and this effect increases with the magnitude. The brightest objects (${\\rm r} \\ltsm 20$), such as the main cluster-galaxies, are almost unaffected, the difference being lower than 0.1 mag. This difference is about 0.2 mag for objects with ${\\rm r} \\simeq 23$, and it increases to 0.4 mag for the faintest objects in the sample, up to ${\\rm r} = 25$. Using a fainter isophote to integrate magnitudes, as it was done for Cl~1600+4109, strongly reduces the effect. When the isophotal magnitude is the same in both the good and the bad seeing images, the difference in magnitude for the faintest objects can be as large as 0.7 to 1 mag.  \\begin{table} \\caption[]{Expected colors for a single burst of star formation as a function of its age, as seen at the redshifts of both clusters} \\begin{flushleft} \\begin{tabular}{lrrrl} \\hline\\noalign{\\smallskip} Age       &{\\rm B} $-$ {\\rm r} & {\\rm g} $-$ {\\rm r}&  {\\rm B} $-$ {\\rm r} & {\\rm g} $-$ {\\rm r} \\\\ $\\left[ {\\rm log(yr)} \\right]$     & \\multispan2 [{\\it z}~=~0.415] & \\multispan2 [{\\it z}~=~0.540]  \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 5   &         $-$0.10     &       0.04   &           $-$0.06    & 0.08 \\\\ 6   &         $-$0.13     &       0.02   &           $-$0.09    & 0.06 \\\\ 7   &          0.32       &       0.40   &            0.27      & 0.36 \\\\ 8   &          0.85       &       0.83   &            0.87      & 0.86 \\\\ 9   &          1.44       &       1.30   &            1.51      & 1.36 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} \\end{table}  In Cl~1613+3104, the population of galaxies is dominated by the E/S0 type, at least up to the completeness limit in r. In the case of Cl~1600+4109, there is a lack of E/S0 galaxies (they are less than $15 \\%$ of the sample) and the population is mainly composed of S/Im galaxies up the completeness limit in r. The extremely blue population dominates the sample at faint magnitudes (${\\rm r} > 24$) in both clusters. The general trend is that the fraction of blue galaxies in excess with respect to an empty field increases with magnitude. Nevertheless, getting the absolute distribution in morphological types is only possible when the sample of galaxies is limited within the magnitude of completeness in {\\sl all} the filters involved. Otherwise, a bias is introduced, as can be seen in Figs.\\ 2 and 7. The result is that the fraction of blue objects detected in both clusters should be considered as an upper bound because the detection of red objects is not complete. In Cl~1613+3104, the color-color diagrams are only complete for objects bluer than B $-$ r = 0.25 and g $-$ r = 0.5 in the interval $22 \\leq {\\rm r} \\leq 24$ and B $-$ r = $-$0.5 and g $-$ r = $-$0.25 when ${\\rm r} \\geq 24$. In Cl~1600+4109, this condition is fulfilled for objects bluer than B $-$ r = 0.75 and g $-$ r = 0.75 if $22 \\leq {\\rm r} \\leq 24$, and B $-$ r = $-$0.25 and g $-$ r = $-$0.25 when ${\\rm r} \\geq 24$. However, even if the absolute ratio of blue to red objects is biased, the both cluster fields show clearly an excess of extremely blue objects in comparison with an empty field, and the number of such objects increases with the magnitude. Table 3 gives the location in the color-color plane of a single burst of star formation, as seen at the redshifts of both clusters, as a function of its age. It is assumed that the light coming from the stellar mass created by the burst is responsible for the bulk of the energy in the spectra. When we compare these values with the color-color distributions shown in Figs.\\ 3 and 8, it appears that most extremely blue objects in both clusters are compatible with the colors expected for such burst models younger than a few $10^{7}$~yr. High quality images of these clusters could confirm the hypothesis that these galaxies are going through mergers or tidal interactions. These observations would reinforce the idea that triggers of starburst processes were more frequent in the past than in the present epoch."
+    },
+    "9602/astro-ph9602129_arXiv.txt": {
+        "abstract": "We have monitored a sample of 27 nearby globular clusters in the hard X-ray band ($20-120$~keV) for $\\sim$1400 days using the BATSE instrument on board the Compton Gamma-Ray Observatory.  Globular clusters may contain a large number of compact objects (e.g., pulsars or X-ray binaries containing neutron stars) which can produce hard X-ray emission. Our search provides a sensitive ($\\sim$50 mCrab) monitor for hard X-ray transient events on time scales of $\\gtrsim1$~day and a means for observing persistent hard X-ray emission. We have discovered no transient events from any of the clusters and no persistent emission. Our observations include a sensitive search of four nearby clusters containing dim X-ray sources: 47~Tucanae, NGC~5139, NGC~6397, and NGC~6752. The non-detection in these clusters implies a lower limit for the recurrence time of transients of 2 to 6~years for events with luminosities $\\gtrsim10^{36}$~erg~s$^{-1}$ ($20-120$~keV) and $\\sim20$~years if the sources in these clusters are taken collectively. This suggests that the dim X-ray sources in these clusters are not transients similar to Aql~X-1. We also place upper limits on the persistent emission in the range $(2-10)\\times10^{34}$~erg~s$^{-1}$ ($2\\sigma$, $20-120$~keV) for these four clusters. For 47~Tuc the upper limit is more sensitive than previous measurements by a factor of 3.  We find a model dependent upper limit of 19 isolated millisecond pulsars (MSPs) producing gamma-rays in 47~Tuc, compared to the 11 observed radio MSPs in this cluster. ",
+        "introduction": "Globular clusters (GCs) are remarkable stellar systems where a variety of compact objects may form and evolve. Numerous millisecond pulsars (MSPs) inhabit GCs as revealed by radio observations (Lyne 1996), e.g. 11 MSPs in 47~Tuc (Robinson et al. 1995). In the X-ray band, the number of bright sources ($1-10$~keV) per unit stellar mass is two orders of magnitude larger for GCs than for the remainder of the galaxy (Clark 1975). Persistent X-ray sources in GCs divide into two luminosity groups: one with low luminosities ($L_{x}\\lesssim 10^{34.5}$~erg~s$^{-1}$, $0.5-4.5$~keV), and the other with high luminosities ($L_{x}\\gtrsim10^{36}$~erg~s$^{-1}$) (Hertz and Grindlay 1983, Hertz and Wood 1985, Verbunt et al. 1995). The nature of the low luminosity, dim X-ray sources (DXSs) has been controversial (for reviews see Bailyn 1995, Verbunt et al. 1994, Grindlay 1993b, Hut et al. 1992). Suggestions for their origin include cataclysmic variables, with mass accreting onto a white dwarf (Grindlay 1993a), low mass X-ray binaries (LMXBs) with accretion at quiescent levels onto a neutron star (NS) or black hole (BH) (Verbunt, Van Paradijs and Elson 1984), plerionic MSP binaries (Tavani 1991) or chance superpositions of foreground/background objects (Margon and Bolte 1987).   We have employed the BATSE instrument on board the Compton Gamma-Ray Observatory to search for transient and persistent hard X-ray emission ($\\gtrsim20$~keV). Hard X-rays may be produced by GC compact objects in a number of ways. In a binary, hard X-ray emission can be powered by mass accretion onto the BH/NS primary, or in pulsars, from shock interactions in the relativistic pulsar winds in the binary. Hard X-ray emission may also be produced by magnetospheric emission of isolated MSPs. Transient emission has been observed previously from GCs. NGC~6440, a GC containing DXSs, has exhibited a high luminosity transient episode (Forman, Jones and Tananbaum 1976). M15 may also have shown a transient event (Pye and McHardy 1983). In the hard X-ray band, SIGMA has observed transient emission from the globular cluster Terzan~2. Recent ROSAT HRI observations indicate that the SIGMA source is most likely associated with the X-ray burster XB~1724-30 (Mereghetti et al. 1995). SLX~1732-304 in the GC Terzan~1 has also been detected as a hard X-ray source (Churazov et al. 1994). Our search does not include Terzan~2, Terzan~1 or NGC~6440 due to problems with nearby interfering sources. We have been able to observe M15 however.   In Section 2, we describe hard X-ray emission mechanisms in detail. Section 3 outlines the BATSE earth occultation flux measurement technique and the parameters of our search. In Sections 4 and 5 we present the results of our search, providing upper limits on transient events and constraints on the recurrence times of transients from DXSs. Section 6 contains upper limits on persistent emission and constraints on the number of isolated and interacting MSPs. Section 7 is a summary and conclusion. ",
+        "conclusions": "We have monitored 27 galactic globular clusters (GCs) with BATSE during a period of approximately four years each. We have detected no distinct hard X-ray outburst episodes and no persistent emission. The lack of detected events has several interesting implications.  For the GCs with dim X-ray sources (DXSs), we find a lower limit for the outburst recurrence time from DXSs of $\\tau_{r}\\sim2-6$~years. This limit excludes the existence of `Aql~X-1-like' objects (i.e., persistent X-ray sources subject to major X-ray outbursts with a time scale of $\\sim 1$ year), since $\\tau_{r}$ is comparable to or greater than this outburst recurrence time scale. This suggests that the DXSs in these clusters are not LMXBs similar to Aql~X-1.  The lack of strong outburst events in our search means we also have no evidence of accreting BHs in GCs. If active BH binaries exist in GCs, they would be clearly detectable with hard X-ray luminosities of order of $10^{36}-10^{37}$~erg~s$^{-1}$.  We have calculated upper limits on persistent hard X-ray emission from GCs. The limiting luminosity is $\\sim(2-10)\\times10^{34}$~erg~s$^{-1}$ for the closest clusters. The limit for 47~Tuc is somewhat more sensitive than previous measurements. A model-dependent but reasonable estimate of the expected hard X-ray magnetospheric emission from isolated MSPs (Chen 1991) implies that the number of isolated MSPs emitting hard X-rays in 47~Tuc is less than 19. Taking the observed 11 MSPs in 47~Tuc, and a beaming factor of 2 for short period pulsars, this upper limit is comparable to actual number of MSPs in 47~Tuc. Our results are therefore a constrain the magnetospheric model of MSP hard X-ray emission.  Our upper limits on persistent hard X-ray emission also provide a mild constraint on the number of interacting pulsars in binaries. The observed efficiency (a few percent) of conversion of pulsar spindown luminosity into hard X-ray emission in the case of the periastron passage of the Be star/pulsar system PSR~B1259-63 (Tavani et al. 1996, Grove et al. 1995) together with the observed MSP spindown average luminosity of $10^{34}$~erg~s$^{-1}$ implies a limit to the GC population interacting pulsar binaries. We find limits of $14-66$ interacting MSPs in the clusters studied."
+    }
+}
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